Datasets:

charlieoneill

/

jsalt-astroph-dataset

like 1

1012

1012.2604_arXiv.txt

{NA61/SHINE is a fixed-target experiment to study hadron production in hadron-nucleus and nucleus-nucleus collisions at the CERN SPS. Due to the very good acceptance and particle identification in forward direction, NA61/SHINE is well suited for measuring particle production to improve the reliability of air shower simulations. Data with proton and pion beams have been taken in 2007 and 2009. First analysis results for the pion yield in proton-carbon interactions at 31 GeV will be shown and compared to predictions from models used in air shower simulations.} \FullConference{35th International Conference of High Energy Physics - ICHEP2010,\\ July 22-28, 2010\\ Paris France} \newcommand{\pip}{\ensuremath{\pi^+}} \newcommand{\pim}{\ensuremath{\pi^-}} \newcommand{\pth}{p-$\theta~$} \newcommand{\FlukaLong}{{\scshape Fluka2008}\xspace} \newcommand{\Fluka}{{\scshape Fluka}\xspace} \newcommand{\UrqmdLong}{{\scshape Urqmd1.3.1}\xspace} \newcommand{\Urqmd}{{\scshape Urqmd}\xspace} \newcommand{\GheishaLong}{{\scshape Gheisha2002}\xspace} \newcommand{\GheishaOld}{{\scshape Gheisha600}\xspace} \newcommand{\Gheisha}{{\scshape Gheisha}\xspace} \newcommand{\Corsika}{{\scshape Corsika}\xspace} \begin{document}

Ultra-high energy cosmic rays initiate extensive air showers (EAS) when they collide with the nuclei of the atmosphere. The interpretation of EAS data as for instance recorded by the Pierre Auger Observatory~\cite{Auger} or the KASCADE air shower array~\cite{KASCADE} relies to a large extent on the under- \begin{wrapfigure}[17]{r}[0pt]{6cm} \includegraphics[width=\linewidth]{talkICRCAuger.eps} \caption[particle energies]{ Particle types and energies involved in the last interaction leading to muon detected at ground level ($E_0=10^{19}$~eV, detection distance $\approx$~1~km)~\cite{Maris:2009x1}.} \label{grandmothers} \end{wrapfigure} standing of these air showers and specifically on the correct modeling of hadron-air interactions that occur during the shower development. The relevant particle energies span a wide range from primary energies of $\gtrsim 10^{20}$~eV down to energies of~$10^{9}$~eV. The mesons that decay to muons at ground level typically originate from low energy interactions in the late stages of an air shower. Depending on the primary energy and detection distance, the corresponding interaction energies are between 10 and 1000~GeV (cf.~Fig~\ref{grandmothers}). As it has been noted in e.g.\ \cite{Heck:2003br, Drescher:2003gh, Meurer:2005dt, Maris:2009uc}, the modeling of low energy interactions contribute at least 10\% to the overall uncertainty of the predicted muon number at ground. Since these energies are within the reach of fixed target experiments, precise measurements of hadronic particle production at, for instance, the Super Proton Synchrotron at CERN can help to diminish the uncertainties of air shower simulations. An example of current difficulties to describe air shower measurements at ultra-high energies is the excess of the number of ground level muons wrt.\ to air shower simulations~\cite{qgsjet01,qgsjetII,sibyll2.1} as reported by the Pierre Auger Observatory~\cite{augerMuon}. \begin{figure}[b!] \includegraphics[width=.545\linewidth]{baryonFraction.eps} \includegraphics[width=.455\linewidth]{baryPropa.eps} \caption{ Left panel: Energy fraction of produced baryons and anti-baryons in $\pi$-air collision as a function of pion momentum. Right panel: Expected number of charged tracks as a function of secondary momentum for the NA61 $\pi^-$-C data set at 350 GeV/c. } \label{baryonPrediction} \end{figure} A solution to this inconsistency was proposed in~\cite{epos}, where it was pointed out that an increased production of baryons and anti-baryons in hadron-air collisions, would lead to an increase in the number of muons at ground level. As can be seen in Fig.~\ref{baryonPrediction}, an enhanced (anti-)baryon production, as e.g.\ currently implemented in the {\scshape Epos}-model, should be easily distinguishable from previous model assumptions at fixed target energies.\\ Unfortunately, there exist no comprehensive and precise particle production measurements for the most numerous projectile in air showers, the $\pi$-meson. Therefore, new data with pion beams at 158 and 350 GeV/c on a thin carbon target (as a proxy for nitrogen) has been recently collected by the NA61 experiment at the SPS.

1012.2604

1012

1012.2574_arXiv.txt

We study the effect of the relative velocity of dark matter and baryonic fluids after the epoch of recombination on the evolution of the first bound objects in the early universe. Recent work has shown that, although relative motion of the two fluids is formally a second order effect in density, it has a dramatic impact on the formation and distribution of the first cosmic structures. Focusing on the gas content, we analyze the effect of relative velocity on the properties of halos over a wide range of halo masses and redshifts. We calculate accurately the linear evolution of the baryon and dark matter fluctuations, and quantify the resulting effect on halos based on an analytical formalism that has been carefully checked with simulations in the case with no relative velocity. We estimate the effect on the abundance of and gas fraction in early halos. We find that the relative velocity effect causes several changes: (i) the characteristic mass that divides gas-rich and gas-poor halos is increased by roughly an order of magnitude, from $2\times 10^4\,M_\odot$ to about $2\times 10^5\,M_\odot$; (ii) this characteristic mass has a large scatter (full width at half maximum is $\sim 1.5 \times 10^5 M_{\odot}$ at $z=20$); (iii) the fraction of baryons in star-less gas minihalos is suppressed by a factor of 4 at $z=20$; (iv) the fraction of baryons in halos that can cool and form stars is suppressed by a factor of 1.5 at $z=20$; and (v) there are enhanced spatial variations of these various fractions.

One of the most important questions in astrophysics today is understanding the formation and evolution of the first bound structures. Significant theoretical and observational efforts are devoted to understanding the properties of the first galaxies and minihalos, at what redshifts they form and how they influence the epoch of reionization. Observations, most notably of the cosmic microwave background (CMB), have established the basic parameters for the initial conditions for structure formation \citep{WMAP1}, thus providing a foundation for theoretical work on the first structures. Advances in computation have made it possible to simulate the formation of the first stars \citep{abel, OShea05, Yoshida08}. Meanwhile, several efforts are underway to probe the structure of the intergalactic medium (IGM) during the reionization epoch using the 21 cm line of hydrogen, and second-generation experiments may be able to explore the early stages of reionization. To answer these and many other questions it is imperative to know the correct initial conditions that led to the formation of the first bound objects and to account for all subtle effects that influence evolution of the density perturbations after recombination. The critical role of initial conditions has been discussed by \citet{NYB}, who showed that three commonly used setups lead to significantly different abundances and properties of the first star-forming gas clouds as well as first gas-rich minihalos. There are two major classes of early-type objects that must be analyzed. The first class consists of large halos in which the gas can cool and form stars; these are the presumed sites of the first dwarf galaxies, which represent the first source of metals in the Universe, and provide ultraviolet photons that begin the decoupling of the hydrogen spin temperature from the CMB \citep{MMR} and eventually start the epoch of reionization. The second class consists of smaller halos (``minihalos'') that are too small for molecular cooling, but still affect the epoch of reionization by acting as a sink for ionizing photons \citep{Haiman01, BL02, Iliev05, Ciardi05} and may generate a 21 cm signal from collisional excitation of H{\sc\,i} \citep[e.g.][]{Iliev03, FO06}. It is important to understand both the abundance and distribution of halos, as well as the precise boundaries separating halos that undergo cooling and star formation, those that collect baryons in their potential wells but do not cool, and the lightest halos that exist only as dark matter structures and do not collect gas. An important effect that was previously overlooked is that of the relative velocity of dark matter and baryonic fluids \citep{TH10}. This effect leads to power suppression on scales that correspond to the first bound halos between $10^4 \ M_{\odot}$ and $10^8 \ M_{\odot}$ and delays the formation of the first objects. More importantly this effect introduces scale-dependent bias and stochasticity, leading to significant qualitative changes in the distribution of the first objects. The relative velocity effect is especially important on the small scales where the first stars and galaxies form. Introduction of this effect dramatically changes the gas distribution inside the first halos and changes the characteristic mass of gas-rich objects. \citet{Dalal10} recently calculated analytically the effect on the gas content of halos and found a large effect on the fluctuations of the Lyman-$\alpha$ background at high redshifts. Their analysis, however, was based on a very simplified model of which halos can form stars and in what abundance. In this paper we carry out a detailed analytical study of the distribution of gas and stars in the first halos. The rest of the paper is organized as follows. Section~\ref{S:IC} reviews the relative velocity effect (Sec.~\ref{sec1}) and improves the analysis of \citet{TH10} to account for spatial variation of the sound speed (Sec.~\ref{sec2}). Section~\ref{S:FH} then investigates the early halos and their gas content, focusing on computation of the filtering mass (Sec.~\ref{sec3}) and then examining the fraction of baryons in minihalos and in larger halos that can cool, including an analysis of spatial variations in the baryon budget (Sec.~\ref{sec4} and ~\ref{sec5}). We summarize our results in Sec.~\ref{sec6} and compare them to other recent work. The numerical results and plots shown in this paper assume a cosmology with present-day baryon density $\Omega_{\rm b,0}=0.044$, CDM density $\Omega_{\rm c,0}=0.226$, dark energy density $\Omega_{\rm \Lambda,0}=0.73$, Hubble constant $H_0=71$ km$\,$s$^{-1}\,$Mpc$^{-1}$, and adiabatic primordial perturbations of variance $\Delta^2_\zeta(k_\star)=2.42\times 10^{-9}$ at $k_\star = 0.002$ Mpc$^{-1}$ and with slope $n_s=0.96$.

\label{sec6} We have shown that the relative velocity of baryons and dark matter has a significant impact on the properties of the first bound objects and has to be considered in detailed studies of the epoch of reionization and especially earlier epochs. The supersonic motion of the baryonic fluid relative to the underlying potential wells created by the dark matter causes advection of small-scale perturbations by large-scale velocity flows, leading to a significant suppression of gas accretion during halo formation and dramatically increasing the characteristic mass of gas-rich objects at high redshifts ($z > 10$). In particular, instead of this characteristic filtering mass being close to the Jeans mass of $2\times 10^4 M_{\odot}$ at $z=20$, it varies among various regions from this value up to $\sim 10^6 M_{\odot}$, with a $1\sigma$ value (and global average) around $M_F = 2\times 10^5 M_{\odot}$, i.e., an order of magnitude higher than without the relative velocity effect. The relative velocity effect also modifies the star formation history, delaying star formation and causing significant spatial fluctuations. However, since the minimum mass for H$_2$ cooling ($\approx 6\times 10^5 M_{\odot}$ at $z=20$) is somewhat higher than the average $M_F$, the suppression effect of $\vbc$ is limited to about a factor of 1.6 at $z=20$ (added on top of the spatially-uniform factor of 1.2 from the still-depressed baryon perturbations on large scales), compared to a much larger effect (a factor of 3.3) on the gas fraction in star-less gas minihalos. The importance of the relative velocity grows steadily with redshift, so that at $z=40$ the suppression factors due to $\vbc$ increase to 2.5 for galaxies (on top of a pre-existing factor of 1.5) and 21 for minihalos. In our detailed treatment, we included the spatial variation of the baryonic sound speed, the suppression of baryonic perturbations on large scales, and the effect of the relative velocity, through the modified power spectrum, both on the halo mass function and the internal gas fractions in halos. In order to gauge the induced spatial variability, we further calculated the full probability distribution functions of the characteristic mass and of gas fractions inside of the first collapsed halos. These results are important for understanding of the relative velocity effect on large scales, and we plan to study them further. Our results significantly extend the work done recently by \citet{Dalal10}. For example, we find a suppression factor due to $\vbc$ at $z=20$ of 1.6 and 3.3, for star-forming halos and minihalos, respectively. In their approach \citet{Dalal10} did not separate these two categories, and found a factor of 2.5 suppression in the collapsed fraction, which under their approximation can be interpreted as a suppression of star formation. In our work we removed this and many other approximations used in \citet{Dalal10}. Comparing to our work, we expect that their calculation of Lyman-$\alpha$ flux fluctuations is qualitatively correct but may be somewhat overestimated and requires a more detailed analysis. As we were finishing this paper, two simulation papers appeared on the effect of $\vbc$ at high redshift \citep{Maio,Stacy}. While both found a small suppression of star formation, their results appear at first glance to show a smaller suppression effect than we predict. This difference is not surprising if we note that these simulation papers focused on star-forming halos around $z\approx 15$, while the largest effects that we find occur for star-less minihalos at higher redshifts. At $z>20$, \citet{Maio} find tens of percents difference in the gas fractions, although the statistical errors are large. \citet{Stacy} find a delay in gas collapse by $\Delta a/a=0.14$ for $v_{\rm bc}/\sigma_{\rm vbc}=1$. We also note that the choice of initial conditions should be carefully considered: standard initial condition codes do not properly treat the separate baryonic and dark matter perturbations or the gas temperature perturbations, leading to a filtering mass that is too high by a factor of $\sim 2$ at $v_{\rm bc}=0$ \citep{NYB}; as such they may underestimate the effect of relative velocities. The simulations by \citet{Maio} and \citet{Stacy} clearly represent a very important step and will serve as a good foundation for simulations with larger boxes and improved initial conditions. Eventually we hope that simulations including $v_{\rm bc}$ will advance to the point where improved fitting functions for the local halo mass function and gas mass fraction become available. We note that various feedback effects may reduce or mask some of the effect of the relative velocity. For galaxies, local feedback from star formation may effectively raise the minimum halo mass for star formation (except for the very first generation of stars). The possibilities include supernova feedback as well as radiative feedback acting via photoheating and photoevaporation or suppression of H$_2$ formation, although ``positive'' feedback due to X-ray ionization enhancing H$_2$ formation has also been suggested \citep{hrl96, hrl97}. For minihalos, astrophysical heating, e.g., from an early X-ray background, may heat the gas and raise the filtering mass above the value due to $\vbc$. There are many unknowns, but these various effects could begin to be significant by $z \sim 20$, and very likely by the time of significant cosmic reionization. Still, the relative velocity between baryons and dark matter is the main determinant of the gas content of halos at the highest redshifts.

1012.2574

1012

1012.0571_arXiv.txt

One of the first stages of planet formation is the growth of small planetesimals and their accumulation into large planetesimals and planetary embryos. This early stage occurs much before the dispersal of most of the gas from the protoplanetary disk. Due to their different aerodynamic properties, planetesimals of different sizes/shapes experience different drag forces from the gas at these stage. Such differential forces produce a wind-shearing effect between close by, different size planetesimals. For any two planetesimals, a wind-shearing radius can be considered, at which the differential acceleration due to the wind becomes greater than the mutual gravitational pull between the planetesimals. We find that the wind-shearing radius could be much smaller than the \emph{gravitational} shearing radius by the Sun (the Hill radius), i.e. during the gas-phase of the disk wind-shearing could play a more important role than tidal perturbations by the Sun. Here we study the wind-shearing radii for planetesimal pairs of different sizes and compare it with gravitational shearing (drag force vs. gravitational tidal forces). We then discuss the role of wind-shearing for the stability and survival of binary planetesimals, and provide stability criteria for binary planetesimals embedded in a gaseous disk.

1012.0571

1012

1012.5953_arXiv.txt

A three dimensional (3D) tomographic reconstruction of the local differential emission measure (LDEM) of the global solar corona during the whole heliosphere interval (WHI, Carrington rotation CR-2068) is presented, based on STEREO/EUVI images. We determine the 3D distribution of the electron density, mean temperature, and temperature spread, in the range of heliocentric heights 1.03 to 1.23\,\rsun. The reconstruction is complemented with a potential field source surface (PFSS) magnetic-field model. The streamer core, streamer legs, and subpolar regions are analyzed and compared to a similar analysis previously performed for CR-2077, very near the absolute minimum of the Solar Cycle 23. In each region, the typical values of density and temperature are similar in both periods. The WHI corona exhibits a streamer structure of relatively smaller volume and latitudinal extension than during CR-2077, with a global closed-to-open density contrast about 6\% lower, and a somewhat more complex morphology. The average basal electron density is found to be about $2.23$ and $1.08\times 10^8\,{\rm cm^{-3}}$, in the streamer core and subpolar regions, respectively. The electron temperature is quite uniform over the analyzed height range, with average values of about 1.13 and 0.93 MK, in the streamer core and subpolar regions, respectively. {Within the streamer closed region, both periods show higher temperatures at mid-latitudes and lower temperatures near the equator.} Both periods show $\beta>1$ in the streamer core and $\beta<1$ in the surrounding open regions, with CR-2077 exhibiting a stronger contrast. Hydrostatic fits to the electron density are performed, and the scale height is compared to the LDEM mean electron temperature. Within the streamer core, the results are consistent with an isothermal hydrostatic plasma regime, with the temperatures of ions and electrons differing by up to about 10\%. In the subpolar open regions, the results are consistent with departures from thermal equilibrium with $T_{\rm ions}>\Te$ (and values of $T_{\rm ions}/\Te$ up to about 1.5), and/or the presence of wave pressure mechanisms linear in the density.

Advancing our understanding of the processes that heat and accelerate the coronal plasma now requires empirical knowledge of its three-dimensional (3D) structure. Coronal images are two-dimensional projections of the 3D structure, and a number of methods have been used to recover the 3D information from the 2D images. Available techniques include a variety of approaches, with diverse aims, strengths, and limitations. Towards this end, solar rotational tomography (SRT) constitutes a powerful empirical technique. Since the original work by Altschuler and Perry (1972), SRT has been developed and applied to polarized white-light image time series, allowing for the reconstruction of the 3D structure of the coronal electron number density. A modern implementation of SRT can be found in Frazin (2000) and Frazin and Janzen (2002), and a comprehensive review of its development in Frazin and Kamalabadi (2005). One of the primary goals of NASA's dual-spacecraft \textit{Solar Terrestrial Relations Observatory} (STEREO) mission is precisely to determine the 3D structure of the corona (Kaiser \emph{et al.}, 2008). The \textit{Extreme UltraViolet Imager} (EUVI) on the STEREO mission returns high-resolution ($1.6''$) narrow-band images centered over Fe emission lines at 171, 195, 284 \AA, and the He\,{\sc ii} 304 \AA\ line (Howard \emph{et al.}, 2008). In this context, we have developed a novel technique, named differential emission measure tomography (DEMT). The technique was theoretically proposed by Frazin \etal{} (2005), and fully developed and applied to STEREO/EUVI data by Frazin \etal{} (2009; henceforth FVK09). DEMT takes advantage of the solar rotation to provide the multiple views required for tomography, as well as of the dual view angles provided by the STEREO spacecraft, the use of which allows for a reduced data-gathering time. Based on the input of EUV-image time series, DEMT produces maps of the 3D EUV emissivity, and a 3D DEM analysis free of 2D projection effects. As explained in FVK09, the first three moments of this local DEM (or LDEM) analysis give 3D maps of the electron density, the mean electron temperature, and the electron temperature spread. A major advantage of DEMT is that it obviates the need for \textit{ad-hoc} modeling of specific structures of interest. Its main (current) limitation is the assumption of a static corona during the data-gathering process, implying that the reconstructions are reliable only in coronal regions populated by structures that are stable throughout their disk transit in the images. In contrast to other approaches, DEMT does not require background subtraction, and is global (\textit{i.e.} it considers the entire corona), but it does not resolve individual loops. In V\'asquez \etal{} (2009), we published the first empirically derived 3D density and temperature structure of coronal-filament cavities, structures that are particularly interesting to study as filament eruptions are the progenitors of about 2/3 of all CMEs (Gibson \etal, 2006). In V\'asquez \etal{} (2010, VFM10 hereafter), we presented the first EUVI/STEREO DEMT analysis of the global corona, specifically for the period CR-2077, belonging to the Solar Cycle 23 extended solar-activity minimum period. In the present work, we develop a similar DEMT analysis for EUVI/STEREO data corresponding to the Whole Heliosphere Interval (WHI) period CR-2068 (20 March 2008, 01:14 UT through 16 April 08:05 UT). We also show a potential field source surface (PFSS) magnetic-field model based on the \textit{Michelson Doppler Imager} (MDI/SOHO) magnetograms of the same period. The \emph{Comparative Solar Minima} (CSM) working group (WG), sponsored by Division II (Sun and Heliosphere) of the International Astronomical Union (IAU), focuses on the research of the coupled Sun--Earth system during solar minimum periods. It seeks to characterize the system at its most basic, ``ground state" and aims to understand the degree and nature of variations within and between solar minima. In this context, we present here the results of the DEMT+PFSSM analysis of the WHI period. We discuss the implication of our model for the thermodynamical structure of the equatorial streamer belt, as well as for the surrounding magnetically open regions, both at the latitudes of the so-called ``streamer legs'', and at the higher subpolar latitudes. To address the central interests of the IAU/CSM WG, we compare our results both with our similar analysis of CR-2077 (VFM10), as well as with results from studies of the Whole Sun Month period (WSM, CR-1913, 22 August through 18 September 1996), belonging to the previous solar-cycle minimum.

\begin{enumerate} \item In the magnetically closed regions of the streamer core, our results are consistent with a plasma regime of hydrostatic isothermal equilibrium, allowing for either over-density or $\Te > \Tions$, with a temperature difference of up to about 10\%. This is seen both in the streamer core central latitudes, as well as closer to the boundary with the surrounding open magnetic structures. \item On the open field lines, we found considerably larger values of $b\Tfit-\left<\Tm\right>$, the difference between the scale-height temperature and the LDEM electron temperature, than those in the streamer core. Nevertheless, the hydrostatic fits have small residuals, indicating one or both of these scenarios: \begin{enumerate} \item Pressure mechanisms other than thermal are acting, and these mechanisms are approximately linear in the density. \item Isothermality among species is not met. In this case, our analysis indicates that $T_{\rm ions}>\Te$, and that the temperature ratio can typically reach values up to order $T_{\rm ions}/\Te\approx1.5$ at subpolar latitudes. \end{enumerate} \end{enumerate} Ion excess temperature ($\Tions>\Te$) in quiet- and active-Sun regions have been found observationally by several authors (Seely \etal, 1997; Tu \etal, 1998). Landi (2007) analyzed quiet region off-disk spectra observed by SUMER, for three different periods of the last solar cycle (1996, 1999, 2000). Landi estimated the possible temperature ranges of several ions, between the limb and 1.25 \rsun. In all cases, the author estimated electron temperatures in the range 1.25 to 1.35 MK, and a systematic trend for $\Tions>\Te$, with indication of a larger average ions excess temperature for increasing activity of the Sun. Typical temperature ratios [$\Tions/\Te$] were found in the range one to two, with peak values of up to order three, which is consistent with our findings. In their stereoscopically reconstructed CH plumes study, Feng \etal{} (2009) reported density scale height temperatures to be about 70\% larger than their electron temperature values (derived from SUMER observations), assigning the discrepancy to ion excess temperatures. To further explore our conclusions, we plan to develop an extensive analysis of other solar regions, refining the comparisons by developing non-potential field extrapolation for selected regions. Another important improvement (currently under development) is the incorporation of the PSF deconvolution of the images used for tomography, allowing us to extend DEMT analysis to coronal holes. The global character of DEMT analysis make it suitable to serve as observational constraint to global coronal models. As an example, we have recently used the CR-2077 DEMT results to provide electron density and temperature basal boundary conditions for a two-temperature MHD model of the solar wind (van der Holst, 2010). The usefulness of the DEMT results in general, and as constraint for models in particular, will improve much by the inclusion of the PSF deconvolution. Immediate future work includes the application of the DEMT technique to use the six Fe bands of the SDO/AIA instrument. Taking advantage of the increased number of bands and temperature coverage provided by SDO/AIA, we aim to refine the LDEM determination. Finally, we also plan to implement time-dependent DEMT through the application of the Kalman-filtering method (Frazin \etal, 2005; Butala \etal, 2008).

1012.5953

1012

1012.0747_arXiv.txt

{ After the discovery of V391 Peg b, the first planet detected around a post Red Giant phase star (Silvotti et al. 2007), the EXOTIME (EXOplanet search with the TIming MEthod) project is focused on the search for new planets with similar characteristics. The aim of the project is to organize a global observing network to collect as much data as possible for a sample of five subdwarf B (sdB) stars and share them in order to obtain a more precise analysis. These evolved pulsators may have extremely regular oscillation periods. This feature makes these stars suitable to search for planetary companions with the timing method as in the case of pulsars. In this contribution we present the project and some preliminary results for the star PG 1325+101 (QQ Vir) after the first two years of activity. }

The EXOTIME project (EXOplanet search with the TIming MEthod) is a coordinated observing program aimed at the search for substellar companions around pulsating subdwarf B stars and the derivation of evolutionary timescales. The required observations span over a very long time (of the order of years), so this method is quite expensive in terms of observing time, but with the advantage to be sensitive to wide orbits and relatively low masses (down to $\sim $ M$_{Jup}$). Using the timing method we can measure both the variations of the pulsation periods and the phase variations; the latter potentially allow us to detect the presence of substellar companions as in the case of V391 Peg b (Silvotti et al. 2007), the first planet discovered with this technique around a pulsating star. The discovery of V391 Peg b has raised the interest to investigate evolved planetary systems beyond the main sequence and beyond the red giant branch. The orbital distance of this planet, lower than 2 AU, suggests that this planet may have ``survived'' to the RG expansion of the parent star. Recently, two substellar companions have been detected orbiting the sdB eclipsing binary HW Vir (Lee et al. 2009), suggesting that substellar objects might be a relatively common phenomenon around sdB stars. Further goals of EXOTIME are the characterization of the targets using asteroseismic methods and the measurement of the secular variation of the oscillation period (the so called P-dot, $\dot{P}$), giving a further example of the synergy between asteroseismology and the search for exoplanets, already known in the case of solar-like stars. Measuring $\dot{P}$ allows a precise determination of the evolutionary status of a star and can help the identification of the pulsation modes. As a final goal, EXOTIME wishes to improve our understanding of the formation and evolution of the sdB stars. The formation processes of sdB stars still represent an unclear topic in stellar evolution. Different scenarios have been proposed, both for single stars and binaries. The presence of a secondary body such as a planet, in particular for single stars, has been suggested to play a role in this process, enhancing the mass loss near the RGB tip (Soker 1998).

1012.0747

1012

1012.2756_arXiv.txt

Recent observations of near supernova show that the acceleration expansion of Universe decreases. This phenomenon is called the transient acceleration. In the second part of work we consider the 3-component Universe composed of a scalar field, interacting with the dark matter on the agegraphic dark energy background. We show that the transient acceleration appears in frame of such a model. The obtained results agree with the latest cosmological observations, namely, the 557 SNIa sample (Union2) was released by the Supernova Cosmology Project (SCP) Collaboration.

At the begining of 21 century, the standard cosmological model (SCM) has become the dominant model of the universe, replacing the hot model of the universe (Big Bang). SCM is based on two important observational results: accelerated expansion of the universe and the Euclidean spatial geometry. In addition, it is assumed that the early universe is adequately described by the theory of inflation. SCM fixes a number of parameters of the universe and, in particular, its energy structure. According to the SCM the Universe is currently dominated by dark energy (in the form of a cosmological constant $\Lambda$), required to explain the accelerated expansion and dark (non-baryonic) matter (DM). Existence of DM gives a possibility to solve a number of contradictions in the Big Bang model (non-decreasing behavior of rotation curves, the structure of galactic halo, a chronology of the structures formation, etc.). Attributing the acceleration of the universe expansion exclusively to the negative pressure generated by the cosmological constant drastically reduced the possibilities of the scale factor dynamics and condemned the universe to eternal accelerated expansion. In fact, the dynamics of the scale factor in SCM is described by Friedmann equations \begin{equation} \left( {\frac{{\dot a}} {a}} \right)^2 = H_0^2 \left[ {\Omega _{m0} \left( {\frac{{a_0 }} {a}} \right)^3 + \Omega _\Lambda } \right] \end{equation} The solution of this equation reads \begin{eqnarray} a(t) = a_0 \left( {\frac{{\Omega _{m0} }} {{\Omega _{\Lambda 0} }}} \right)^{1/3} \left[ {sh\left( {\frac{3} {2}\sqrt {\Omega _{\Lambda 0} } H_0 t} \right)} \right]^{2/3} , \\ a\left( {t \ll H_0^{ - 1} } \right) \propto t^{2/3} ;\;a\left( {t \gg H_0^{-1}} \right) \propto e^{H_0 t} . \end{eqnarray} We see that the asymptotic behavior of the solution described the era of matter domination in the early Universe $t\ll H_0^{-1}$ and dark energy domination for the later evolution $t\gg H_0^{-1}$. We now find the dependence of the deceleration parameter $q\equiv -a\ddot{a}/\dot{a}^2 = -\ddot{a}/\left( aH^2\right)$ on the redshift $z$ for a universe filled with arbitrary components of the state equation $p_i = w_i\rho_i$. In this case, \begin{equation} q = \frac{3} {2}\frac{{\sum\limits_i \Omega _i^{(0)} \left( {1 + w_i } \right)(1 + z)^{3\left( {1 + w_i } \right)} }} {{\sum\limits_i \Omega _i^{(0)} (1 + z)^{3\left( {1 + w_i } \right)} }} - 1 \end{equation} For the SCM, this expression takes the form (see Figure 1) \begin{equation} q = \frac{1} {2}\frac{{\Omega _M^{(0)} (1 + z)^3 - 2\Omega _\Lambda ^{(0)} }} {{\Omega _M^{(0)} (1 + z)^3 + \Omega _\Lambda ^{(0)} }} \end{equation} In particular, at the present time \begin{equation} q_0 = \frac{{1 - 3\Omega _{\Lambda 0} }}{2}\simeq - 0.6. \end{equation} A characteristic feature of the dependence $q(z)$ - a monotonous tendency to a limiting value of $q(z) = -1$ for $z \to -1$. Physically, this means that once the dark energy became dominant component (at $z \sim 1$), the universe in SCM is doomed to eternal accelerated expansion. Questioned the adequacy of this result is expressed repeatedly \cite{Andreas}, \cite{Barrow}.

The original agegraphic dark energy model was proposed in~\cite{0707.4049} based on the K\'{a}rolyh\'{a}zy uncertainty relation, which arises from quantum mechanics together with general relativity. The interacting agegraphic dark energy model has certain advantages compared to the original agegraphic or holographic dark energy model. Many studies show that this model gives an opportunity to explain the accelerated expansion of the Universe without a cosmological constant or some form of the scalar field. All the three models give dynamics of the Universe which are virtually indistinguishable from SCM, but without most of its problems, such as the cosmological constant, fine tuning and coincidence problems. Some authors have recently suggested that the cosmic acceleration have already peaked and that we are currently observing its slowing down \cite{Barrow,Starobinsky,Lima}. Under a kinematic analysis of the most recent SNe Ia compilations, the paper \cite{Lima,LiWuYu,LiWuYu1} shows the existence of a considerable probability in the relevant parameter space that the Universe is already in a decelerating expansion regime. One of the deficiencies of original ADE model is the inability to explain the phenomenon of transient acceleration. Density of holographic dark energy is determined by the surface terms in action, while volume terms are usually ignored. We take into account both surface and volume terms, where the latter correspond to (described by) homogeneous scalar field with exponential potential $V(\varphi)$. We consider a model of Universe consisting of dark matter interacting with a scalar field on the agegraphic background. It is shown that this model can explain the transient acceleration. This model also is in accordance with the observational data.

1012.2756

1012

1012.0615_arXiv.txt

Mid-infrared data, including {\it Spitzer} warm-IRAC [3.6] and [4.5] photometry, is critical for understanding the cold population of brown dwarfs now being found, objects which have more in common with planets than stars. As effective temperature ($T_{\rm eff}$) drops from 800~K to 400~K, the fraction of flux emitted beyond 3~$\mu$m increases rapidly, from about 40\% to$ >$75\%. This rapid increase makes a color like $H$-[4.5] a very sensitive temperature indicator, and it can be combined with a gravity- and metallicity-sensitive color like $H-K$ to constrain all three of these fundamental properties, which in turn gives us mass and age for these slowly cooling objects. Determination of mid-infrared color trends also allows better exploitation of the WISE mission by the community. We use new {\it Spitzer} Cycle 6 IRAC photometry, together with published data, to present trends of color with type for L0 to T10 dwarfs. We also use the atmospheric and evolutionary models of Saumon \& Marley to investigate the masses and ages of 13 very late-type T dwarfs, which have $H$-[4.5]$>$3.2 and $T_{\rm eff} \approx$500~K to 750~K. Note: This is an updated version of \citet{legg10a}; a photometry compilation is available at {\it www.gemini.edu/staff/sleggett}.

The last decade has seen a remarkable increase in our knowledge of the bottom of the stellar main-sequence and of the low-mass stellar and sub-stellar (brown dwarf) population of the solar neighborhood. Two classes have been added to the spectral type sequence following M --- L, and T; T dwarfs with effective temperatures ($T_{\rm eff}$) as low as 500~K are now known and we are truly finding objects that provide the link between the low-mass stars and the giant planets. As $T_{\rm eff}$ decreases, brown dwarfs emit significant flux at mid-infrared wavelengths \citep[e.g.][]{burr03,legg09}. At $T_{\rm eff}=$1000~K 30\% of the total flux is emitted at wavelengths longer than 3~$\mu$m, while at 600~K 60\% of the flux is emitted in this region (according to our models). \begin{figure}[!hb] \plotfiddle{Leggett_S_fig1.eps}{9cm}{270}{50}{50}{-200}{290} \caption{Observed (black lines) and modelled (red lines) SEDs of the 750~K T8 dwarf 2MASS 0415-09 \citep[lower spectrum,][]{saum07}, and the 600~K T9 dwarf ULAS 0034-00 \citep[upper spectrum,][]{legg09}. Principal absorption features are identified; NH$_3$ is also likely for the T9 dwarf near 1.0, 1.2, 1.3, 1.5 and 1.8~$\mu$m. The MKO-system $YJHK$, IRAC and WISE filter bandpasses are indicated. } \end{figure} Figure 1 shows spectral energy distributions (SEDs) for a 750~K T8 dwarf and a 600~K T9 dwarf; the flux emerges from windows between strong bands of, primarily, CH$_4$ and H$_2$O absorption. As the temperature drops from 750~K to 600~K, the ratio of the mid- to the near-infrared flux increases dramatically, and more flux emerges through the windows centered near 5~$\mu$m and 10~$\mu$m. Filter passbands are indicated for the near-infrared $YJHK$ MKO-system \citep{toku02}, as well as the {\it Spitzer} IRAC bands \citep{fazi04} and the three shortest-wavelength WISE bands \citep{liu08}. The IRAC and WISE filters sample regions of both high and low flux, thus both cameras are sensitive to cold brown dwarfs, which can be identified by extreme colors in their respective photometric systems. We used {\it Spitzer} Cycle 6 time to obtain IRAC photometry of late-type T dwarfs. Here we combine these data with published photometry to examine trends in colors with spectral type, which will be useful for the design and use of ongoing and planned infrared surveys. We also examine in detail the colors of the $500 \leq T_{\rm eff}$ K $\leq 800$ dwarfs for correlations with the photospheric parameters effective temperature, gravity, and metallicity. We find that various colors do provide indicators of temperature and gravity, which in turn constrains mass and age when combined with evolutionary models.

The IRAC data implies that the majority of the UKIDSS 500~K to 600~K dwarfs are young and low mass, a result not currently understood in terms of selection effects or the mass function. Photometric data at wavelengths longer than 3~$\mu$m are both important and useful for the latest-type T dwarfs with $500 \leq T_{\rm eff}$ K $\leq 800$, and will be even more so for the cooler objects expected to be found by WISE and other sky surveys. It is only possible to get such data from the ground for very bright objects and mid-infrared space missions are crucial for continued progress in this field.

1012.0615

1012

1012.0938_arXiv.txt

Spectra of Galactic Cosmic Rays (GCRs) measured at the Earth are the combination of several processes: sources production and acceleration, propagation in the interstellar medium and propagation in the heliosphere. Inside the solar cavity the intensity of GCRs is reduced due to the solar modulation, the interaction which they have with the interplanetary medium. We realized a 2D stochastic simulation of solar modulation to reproduce CR spectra at the Earth, and evaluated the importance in our results of the Local Interstellar Spectrum (LIS) model and its agreement with data at high energy. We show a good agreement between our model and the data taken by AMS-01 and BESS experiments during periods with different solar activity conditions. Furthermore we made a prediction for the differential intensity which will be measured by AMS-02 experiment.

Models of the Heliosphere need to be very accurate in order to reproduce the complex structure of the solar cavity and its effect on the Cosmic Rays propagation. In particular simulations of CR intensity have to be compared with experimental values, which have been measured in the last years by several space experiments (e.g. AMS-01, BESS) and will be measured even more accurately in the near future (e.g. AMS-02 on the ISS). We have developed a two dimensional (radius and helio-colatitude) model of GCR propagation in the Heliosphere\cite{six,VO2010}\! that is a function of measured values of the Solar Wind velocity in the ecliptic plane ($V_0$) and of the neutral sheet tilt angle ($\alpha$), as well as of stimated values of the diffusion parameter ($K_{0}$). This model is including curvature, gradient and current sheet drifts, which are depending on the charge sign of particles and magnetic field polarity\cite{art_midrift}\!. Particles modulation strictly depends on the Local Interstellar Spectrum (LIS), up to now not measured and supposed to be constant with time outside of the heliosphere. In the model we include also the effects of the evolving solar activity conditions experienced by a CR particle inside the heliosphere, due to the time spent by the solar wind to reach the outer border of the solar cavity. Here we deal with the differential intensity observed at 1 AU: we do not take into account the effects of the Earth magnetosphere\cite{mi_jgr}\!.

\label{sec:conc} We built a 2D stochastic Monte Carlo code for particles propagation inside the heliosphere. Our model takes into account drift effects and shows a good agreement with measured values, in periods with positive as well as with negative polarity. Proton spectra, as predicted by the model, are decreasing with increasing tilt angle and solar wind velocity. We use as input LIS the model published by Burger \& Potgieter\cite{10}\!, corrected in order to fit the high energy measured intensity. Recent measurements have pointed out the needs to reach a high level of accuracy in the modulation of the differential intensity, in relation to the charge sign of the particles and the solar field polarity\cite{mi_ap}\!. This aspect will be even more crucial in the next generation of experiments.

1012.0938

1012

1012.4807_arXiv.txt

We use archival data of NASA's Chandra X-ray telescope to compile an X-ray light curve of all four images of the quadruply lensed quasar Q2237+0305 ($z_{Q}$=1.695) from January 2006 to January 2007. We fit simulated point spread functions to the four individual quasar images using Cash's C-statistic to account for the Poisson nature of the X-ray signal. The quasar images display strong flux variations up to a factor of $\sim 4$ within one month. We can disentangle the intrinsic quasar variability from flux variations due to gravitational microlensing by looking at the flux ratios of the individual quasar images. Doing this, we find evidence for microlensing in image A. In particular, the time-sequence of the flux ratio A/B in the X-ray regime correlates with the corresponding sequence in the optical monitoring by OGLE in the V-band. The amplitudes in the X-ray light curve are larger. For the most prominent peak, the increase of the X-ray ratio A/B is larger by a factor $\sim 1.6$ compared to the signal in the optical. In agreement with theory and other observations of multiply imaged quasars, this suggests that the X-ray emission region of this quasar is significantly smaller than the optical emission region.

The quasar Q2237+0305 was discovered in 1984 \citep{Huchra} during a spectroscopic survey of nearby galaxies. The spectrum of the nucleus of the barred spiral galaxy 2237+0305 ($z_{G}$=0.0394) was found to be superimposed by a quasar component at a redshift of $z_{Q}$ = 1.695.\\ The first high resolution observations of the system showed three images of the quasar \citep{tyson}. This number was soon corrected by \citet{Yee} who first observed all four known point like quasar images around the core of the galaxy. The images are arranged in a nearly symmetric way, hence the name `The Einstein Cross'.\\ Spectroscopy proved that they are images of a single quasar \citep{lensmodel_schneider, 1989A&A...208L..15A}. The images are separated by up to 1.8\arcsec\; from each other, and are labeled A through D.\\ After the prediction by \citet{chang} and initial theoretical studies (e.g., \citet{kayser, schneiderweiss, pacz}), microlensing was first detected in Q2237+0305 \citep{irwin, wampaczschneid}. In fact, this detection of `quasar microlensing' was the first evidence for microlensing generally, including stellar microlensing in the galaxy which was not discovered until 1993 \citep{aubourg, alcock, udal93}. Today `The Einstein Cross' is one of the best studied multiply lensed quasars and there have been a lot of monitoring programs in the optical \citep{corrigan, monitoring_pen, ostensen, OGLE1, OGLE2, monitoring_APO, monitoring_alcade, OGLE3} where much microlensing activity has been observed. The quasar has also been detected in the UV \citep{Blanton}, the NIR \citep{AgolNIR}, the MIR \citep{AgolMIR} and the radio regime \citep{Falco}. However, there is no published light curve of Q2237+0305 in the X-ray regime yet.\\ The X-ray emission of Q2237+0305 was first detected with \textit{ROSAT/HRI} observations in 1997 \citep{joachimROSAT}. Since then there have been other X-ray observations of Q2237+0305 like a single spectroscopic observation with \textit{XMM-Newton} from 2002 (data set ID: 0110960101; PI: Watson) \citep{fedorova} and two \textit{Chandra} observations \citep{photonindex}. In this paper we study ten archival \textit{Chandra} observations ranging from January 2006 until January 2007 and compile the first X-ray light curve of Q2237+0305.

In this paper we have analysed archival \textit{Chandra} data of the gravitationally lensed quasar Q2237+0305. We compiled an X-ray light curve for all four images. The data set comprises ten epochs ranging from January 2006 to January 2007. Because of the blended nature of the four images (see Figure \ref{6831_raw_labled}) it was necessary to simultaneously fit appropriate PSFs to all four images in order to obtain proper photometry. For this, we used a simulated PSF which accounts for the optical properties of the \textit{Chandra} observatory and the spectrum of the quasar X-ray emission. The fitting was accomplished by a two-dimensional fitting algorithm and a grid search by minimising Cash's C-statistic. Finally we analysed the light curve and found evidence for microlensing variations in quasar image A (see Figure \ref{ratiosixplot}). The X-ray microlensing signal in image A coincides with the signal in the optical OGLE light curve. Assuming that this parallel behaviour is caused by the same process, i.e. one source that is microlensed, the amplitude of the microlensing signal is a direct measure for the source size in the particular wavelength regime \citep{wampacz,livingreview, Kochanek, timo_paper, Eigenbrod}. As the microlensing signal in the X-ray regime is much more prominent than in the optical, this suggests that the X-ray emission region is much smaller than the optical emission region of the quasar \citep{pooley2, dai}. In a future paper (Zimmer et al. 2010, \textit{in prep.}) we will make use of this effect to measure the size of the X-ray and the optical emission region in Q2237+0305. While the time-delays in Q2237+0305 are negligible the intrinsic variations have to be considered in the analysis. Figure \ref{cashlightcurve} clearly shows how strong these variations are in the X-ray regime. We do not find such strong fluctuations in the optical. This indicates that the mechanism powering the quasar is variable and leads to brightness variations by a factor of $\sim$ 4 on time-scales of less than 30 days.\\

1012.4807

1012

1012.1613_arXiv.txt

There exists overwhelming evidence, most recently from the Wilkinson Microwave Anisotropy Probe (WMAP) \cite{Komatsu:2008hk}, that non-baryonic cold dark matter comprises around 23 percent of the Universe's energy density. Identifying this dark matter, presumably an elementary particle, is one of the foremost contemporary challenges in particle physics and cosmology. The goals for successful identification of dark matter are: (1) Detection of the relic dark matter particle, and measurement of its mass and distribution directly. (2) Production of the dark matter particle at the LHC and future linear colliders, and measurement of its properties. (3) Testing the consistency between these measurements, namely in astrophysics and particle physics, and reproduction of the relic abundance of the particle from the measured properties in order to confirm that the dark matter particle (possibly more than one species of particles) really makes up 23 percent of the Universe's energy density. % One of the most compelling features of low scale supersymmetry (SUSY), supplemented with $R$-parity conservation, is that it can provide an attractive cold dark matter candidate with the correct relic abundance, provided the lightest neutralino $\tilde\chi_1$ is also the lightest SUSY particle (LSP) \cite{DMreview}. If the LSP neutralino is bino dominated (in an admixture of bino, wino, and Higgsinos), it often leads to an over-abundance of dark matter, unless (co)annihilation processes reduce the relic density to levels compatible with WMAP. Many solutions have been proposed to accomplish this \cite{DMreview,silkphysrep}. One attractive scenario for realizing the correct relic abundance is to consider an appropriate bino-Higgsino mixture in the composition of the LSP \cite{focus, hyper, funnel}. In this mixed bino-Higgsino LSP (called bino-Higgsino dark matter) scenario, two neutralinos and one chargino have masses that are close to the LSP mass, such that (co)annihilation processes among them can reproduce the desired relic density. The spin-independent (SI) cross section on nuclei in this scenario is enhanced, which is an advantage from the point of view of direct detection experiments searching for the LSP. Indeed, the recent candidate events reported by CDMSII \cite{Ahmed:2009zw} and EDELWEISS-II \cite{Edelweiss} would suggest that the SI cross section is $O(10^{-8})\,{\rm pb} $. This is of the right order of magnitude for the bino-Higgsino dark matter scenario \cite{CDMSlightWIMP,CDMSotherSUSY}. As the bounds on the cross section get lowered by the ongoing and planned measurements by XENON100 \cite{Aprile:2009zzc}, SuperCDMS \cite{Schnee:2005pj}, and XMASS \cite{Abe:2009zz}, the WMAP compatible bino-Higgsino mixing solutions will be among the first ones to be tested. Moreover, it is known that a significant Higgsino component in the LSP neutralino also gives a large spin-dependent (SD) cross section, which would make the indirect detection of this dark matter through self-annihilation into neutrinos and other particles more feasible. Thus, the bino-Higgsino dark matter solution will be tested by the IceCube/Deep Core neutrino observatory~\cite{Hultqvist:2010xy}. It is important to observe both SI and SD cross sections and to see their correlation \cite{Cohen:2010gj} in order to adequately test the bino-Higgsino dark matter scenario. The SI cross section is enhanced if the mass $m_A$ of the CP-odd Higgs boson is small and $\tan\beta$ (ratio of Higgs vacuum expectation values (VEVs)) is large. With $m_A$ small, pair annihilation processes are enhanced, and a reduced bino-Higgsino mixing can give rise to the desired WMAP relic density. We refer to this case as bino-Higgsino-like dark matter, if we need to distinguish among the WMAP solutions. On the other hand, when we specify a mixed bino-Higgsino LSP solution where (co)annihilation processes via scalars are negligible, we refer to it as well-tempered bino-Higgsino dark matter \cite{Masiero:2004ft}. {}From the particle physics point of view, the rare decay $B_s\to\mu^+\mu^-$ is one of the most interesting processes in the region of large $\tan\beta$ and small $m_A$ \cite{Choudhury:1998ze}. The Tevatron will provide a bound ($\sim 2\times10^{-8}$) on this branching ratio in run II \cite{CDF}, and LHCb, within a few years, will probe the standard model prediction $(3-4)\times 10^{-9}$. (The exclusion limit from 1 fb$^{-1}$ of data expected by the end of 2011 will be $\sim 6 \times 10^{-9}$) \cite{Lenzi:2007nq}. It is thus important to investigate the regions of parameter space that provide larger SI cross sections (small $m_A$ and/or small Higgsino mass $\mu$), and to explore their predictions. At the LHC, the neutralino LSP is created from cascade decays of squarks and gluinos, and manifests itself as missing energy. As mentioned above, to identify dark matter, one major goal is to reproduce the LSP relic density from the collider measurements \cite{Feng:2001ce}. However, the inverse problem at the LHC \cite{ArkaniHamed:2005px} is not so easy in general, since it is hard to measure the mass spectrum and the couplings directly. Several techniques have been developed, and several reliable relic density simulations have been explored for various WMAP solutions \cite{Polesello:2004qy,Arnowitt:2006jq,Nath:2010zj,Dutta:2010uk}. The assumptions of universality and/or unification of the SUSY breaking mass parameters are crucial simplifications for the collider measurements of the relic density. For the bino-Higgsino dark matter, universality of the sfermion masses is less important since, by definition, the coannihilation processes via sfermions are negligible as far as the relic density is concerned. The gaugino mass spectrum (which should also be addressed at the LHC \cite{Altunkaynak:2009tg}) will be more important in restricting the relic density with LHC measurements. In this paper we investigate bino-Higgsino dark matter and its implications for direct and indirect detection, and for the LHC measurements. We will study both the well-tempered bino-Higgsino dark matter and bino-Higgsino-like dark matter with smaller $m_A$. In the study of the well-tempered mixing solution, it is assumed that the sfermions are sufficiently heavy, without specifying the SUSY breaking scenario or any underlying theory. This is done in order to make the bino-Higgsino dark matter relic density, and the SI and SD cross sections insensitive to these masses. The bino-Higgsino mixing needed to satisfy the desired WMAP relic density depends on wino and bino mass ratio, and thus the cross sections implicitly depend on this ratio. For simplicity, we assume gaugino unification at the grand unification scale, $M_{\rm GUT}$, in order to investigate the cross sections. We also study the possibility of non-universal gauginos, which also can be tested experimentally within our framework. Our results should be applicable to any model where the well-tempered bino-Higgsino dark matter solution can be realized. On the other hand, in order to exhibit our study of the correlation of cross sections and Br($B_s \to \mu^+\mu^-$) for bino-Higgsino-like dark matter, we employ non-universal Higgs mass boundary conditions, where $m_A$ and $\mu$ are free low energy parameters. In this case, for a given LSP mass and $m_A$, the proper WMAP relic density constrains the Higgsino mass $\mu$. As a result, the chargino contribution to Br($B_s \to \mu^+\mu^-$) is predictable for a given stop mass, if gaugino mass unification is assumed. In our presentation we first study the constraints and implications from SI and SD cross sections for the bino-Higgsino(-like) dark matter solution. If the SI cross section is large ($\sigma_{\rm SI} \gtrsim 10^{-8}$ pb), the bino-Higgsino mixing is large and/or $m_A$ is small. The SD cross section is restricted, for given $m_A$ if the bino-Higgsino mixing is determined by the WMAP observation. If the CP-odd Higgs mass $m_A$ is small, the amount of bino-Higgsino mixing required to satisfy the WMAP relic density is not very large. The SD cross section, accordingly, is then also not very large. Hence it is worth making clear the conditions under which we can observe the SD cross section by indirect detection, as well as the corresponding prediction for Br$(B_s\to\mu^+\mu^-)$, while satisfying the other experimental constraints. When $m_A$ is large, the bino-Higgsino mixing needs to be well-tempered and the SD cross section must be large. We also investigate the bound on SD cross section for smaller neutralino masses, $\lesssim 100$ GeV, since it is already bounded by the recent CDMSII / XENON100 data. We then proceed to study the implication from LHC measurements. If the bino-Higgsino mixing is well-tempered, three of the mass eigenvalues of the neutralino mass matrix can be degenerate to within $O(M_Z)$, depending on the neutralino mass parameters. In such a case, the dilepton invariant mass distribution from the heavier neutralino decay with missing energy will give us important information on the neutralino mass parameters. Due to the large SI cross section, the mass of the bino-Higgsino dark matter particle is expected to be measured from the distribution of the recoil energy of the heavy nuclei in the direct detection experiments, and there arises a possibility to extract the parameters for reproducing the LSP relic density and the gaugino mass spectrum. This article is organized as follows. In Section 2, the SI and SD cross sections of neutralino-nucleon interactions are briefly studied. In Section 3, the correlation between the bino-Higgsino dark matter solution and SI cross section is described. Within the bino-Higgsino dark matter scenario, the interplay between the SD cross section and Br$(B_s\rightarrow\mu^+\mu^-)$ is presented in Section 4. We discuss in Section 5 several possible signatures of this scenario at the LHC, and in Section 6 we summarize our results.

We have investigated the direct and indirect detection of the bino-Higgsino dark matter scenario, and explored its implications for the LHC. We first presented the prediction of SD cross section and the branching ratio of $B_s \to \mu^+\mu^-$, assuming that the WMAP relic density constraint is satisfied. Because the relic density restricts the bino-Higgsino mixing, the SD cross section is predicted for given $m_A$. For the WMAP compatible solution, we have two regions for bino-Higgsino dark matter : (1) $m_A < 2 m_{\tilde \chi_1}$, (2) $m_A > 2 m_{\tilde \chi_1}$. Since pair annihilation channels can open up in region (1), the bino-Higgsino mixing here should be smaller than in region (2). In region (1), the SD cross section is therefore smaller, and the neutralino should be sufficiently heavy for the SD cross section to be observed. We find that the SD cross section can be observed indirectly by the neutrino flux from the sun if $m_{\tilde \chi_1} \gtrsim 300$ GeV. In region~(1), the branching ratio for $B_s\to\mu^+\mu^-$ can be enhanced, and Br$(b\to s\gamma)$ constraint gives a lower bound in flavor universal models which can be tested at LHCb. The CDMSII bound can exclude a large branching ratio of $B_s\to\mu^+\mu^-$. In region~(2), on the other hand, the bino-Higgsino mixing is well-tempered and the SD cross section is large enough to be observed. For a neutralino mass less than about 100 GeV, the CDMSII bound constrains the SD cross section even in region~(2). A SD cross section just below the maximal value, for given neutralino mass, is excluded by the CDMSII experiment. This exclusion depends on the strange sea-quark content $f_s$ in the nucleon (multiplied by the strange mass). If the maximal SD cross section for $m_{\tilde\chi_1} \simeq 80-100$ GeV is observed, a smaller value ($f_s\sim0.03$) consistent with the recent results from lattice calculations will be preferred. We next studied the LHC phenomenology of bino-Higgsino dark matter. Because the mass differences of the neutralinos in the case of well-tempered bino-Higgsino dark matter are small, they can be measured by the dilepton invariant mass distribution. From the neutralino mass differences, we may be able to infer whether gaugino masses are unified or not. For this, it turns out that $\tan\beta$ is an important parameter. If we find that $\tan\beta$ is large, say from an observation such as Br$(B_s\to\mu^+\mu^-)$, the gaugino mass ratio at the weak scale can be obtained from the mass differences. The shape of the dilepton invariant mass distribution depends on the relative signatures of the neutralino mass eigenvalues. This distribution will be a powerful tool in providing important information about neutralino masses and the relative signatures of the gaugino masses. If gaugino mass unification is assumed and two of the mass differences of the neutralinos are measured, $\tan\beta$ can be determined, and the bino-Higgsino dark matter relic abundance is then reproduced. The relic density thus deduced from collider measurements provides a strong hint for identifying the nature of dark matter if it coincides with the WMAP data. If the two do not coincide, we cannot decide whether gaugino unification is not satisfied or the neutralino LSP alone does not saturate the WMAP measured relic abundance. A model-independent measurement of $\tan\beta$ provides a strong hint to solve this dilemma. In general, it is hard to measure $\tan\beta$ model-independently at the LHC, but it is possible at a future linear collider. The polarization of $\tau$ lepton may give us a hint of the size of $\tan\beta$ if sleptons are light enough in the bino-Higgsino dark matter scenario \cite{Nojiri:1994it}. We have in this paper assumed that all the sfermions are heavy in order to make their mass parameters insensitive to our discussion, but this assumption can be relaxed. The large bino-Higgsino mixing can provide various features for collider phenomenology, such as $\tau$ polarization, if on-shell sleptons appear in the cascade decays. From a theoretical point of view, a confirmation of the bino-Higgsino dark matter scenario can provide important impetus to investigations of SUSY breaking. A bino-Higgsino dark matter needs a relatively small Higgsino mass $\mu$. In fact, small $\mu$ is preferable if it is a parameter independent of the SUSY breaking scale, while $\mu$ can be large among the electroweak symmetry breaking vacua if it depends on a single SUSY breaking scale parameter \cite{Giudice:2006sn}. Therefore, testing the bino-Higgsino dark matter scenario can serve as an important avenue for distinguishing among the various models of SUSY breaking. In conclusion, the lightest neutralino with an appropriate composition of bino and Higgsino components is a compelling dark matter candidate. This will soon be tested by the ongoing and planned direct detection experiments, and indirectly at the IceCube neutrino telescope through pair annihilation. A mixed bino-Higgsino dark matter particle can also lead to characteristic signals at the LHC as we have discussed.

1012.1613

1012

1012.1339_arXiv.txt

The non-isotropic nature of the neutrino emission from a supernova (SN) core might potentially affect the flavor evolution of the neutrino ensemble, via neutrino-neutrino interactions in the deepest SN regions. We investigate the dependence of these ``multi-angle effects" on the original SN neutrino fluxes in a three-flavor framework. We show that the pattern of the spectral crossings (energies where $F_{\nu_e}= F_{\nu_x}$, and $F_{\overline\nu_e}= F_{\overline\nu_x}$) is crucial in determining the impact of multi-angle effects on the flavor evolution. For neutrino spectra presenting only a single-crossing, synchronization of different angular modes prevails over multi-angle effects, producing the known ``quasi single-angle'' evolution. Conversely, in the presence of spectra with multiple crossing energies, synchronization is not stable. In this situation, multi-angle effects would produce a sizable delay in the onset of the flavor conversions, as recently observed. We show that, due to the only partial adiabaticity of the evolution at large radii, the multi-angle suppression can be so strong to dramatically affect the final oscillated neutrino spectra. In particular three-flavor effects, associated with the solar parameters, could be washed-out in multi-angle simulations.

\label{intro} The characterization of the flavor conversions for neutrinos emitted by a stellar collapse is a field of intense activity. In particular, the flavor transformation probabilities in supernovae (SNe) not only depend on the matter background~\cite{Matt,Dighe:1999bi}, but also on the neutrino fluxes themselves: neutrino-neutrino interactions provide a nonlinear term in the equations of motion~\cite{Pantaleone:1992eq, Sigl:1992fn,McKellar:1992ja} that causes collective flavor transformations~\cite{Qian:1994wh,Samuel:1993uw, Kostelecky:1993dm, Kostelecky:1995dt, Samuel:1996ri, Pastor:2001iu, Wong:2002fa, Abazajian:2002qx, Pastor:2002we, Sawyer:2004ai, Sawyer:2005jk}. Only recently~\cite{Duan:2005cp,Duan:2006an,Hannestad:2006nj} it has been fully appreciated that in the SN context these collective effects give rise to qualitatively new phenomena (see, e.g.,~\cite{Duan:2010bg} for a recent review). The main consequence of this unusual type of flavor transitions is an exchange of the spectrum of the electron species $\nu_e$ (${\bar\nu}_e$) with the non-electron ones $\nu_x$ (${\bar\nu}_x$) in certain energy intervals. These flavor exchanges are called ``swaps'' marked by the ``splits'', which are the boundary features at the edges of each swap interval~\cite{Duan:2006an,Duan:2010bg, Fogli:2007bk,Fogli:2008pt, Raffelt:2007cb, Raffelt:2007xt, Duan:2007bt, Duan:2008za, Gava:2008rp, Gava:2009pj,Dasgupta:2009mg,Friedland:2010sc,Dasgupta:2010cd}. The location and the number of these splits, as well as their dependence on the neutrino mass hierarchy, is crucially dependent on the flux ordering among different neutrino species~\cite{Fogli:2009rd,Choubey:2010up}. In this context, one of the main complication in the simulation of the flavor evolution is that the flux of neutrinos emitted from a supernova core is far from isotropic. The current-current nature of the weak-interaction Hamiltonian implies that the interaction energy between neutrinos of momenta ${\bf p}$ and ${\bf q}$ is proportional to $(1-{\bf v}_{\bf p} \cdot {\bf v}_{\bf q})$, where ${\bf v}_{\bf p}$ is the neutrino velocity~\cite{Qian:1994wh,Pantaleone:1992xh}. In a non-isotropic medium this velocity-dependent term would not average to zero, producing a different refractive index for neutrinos propagating on different trajectories. This is the origin of the so-called ``multi-angle effects''~\cite{Duan:2006an}, which hinder the maintenance of the coherent oscillation behavior for different neutrino modes~\cite{Raffelt:2007yz,Fogli:2007bk,EstebanPretel:2007ec,Sawyer:2008zs}. In~\cite{Raffelt:2007yz} it has been shown that in a dense neutrino gas initially composed of only $\nu_e$ and $\overline\nu_e$ with equal fluxes, multi-angles effects would rapidly lead to flavor decoherence, resulting in flux equilibration among electron and non-electron (anti)neutrino species. On the other hand, in the presence of relevant flavor asymmetries between $\nu_e$ and ${\overline \nu}_e$ multi-angle effects can be suppressed. In particular, during the early SN accretion phase, one expects as neutrino flux ordering $\Phi_{\nu_e}\gg \Phi_{\overline\nu_e}\gg \Phi_{\nu_x}= \Phi_{\overline\nu_x}$ ~\cite{Raffelt:2003en, Fischer:2009af, Huedepohl:2009wh}, defined in terms of the total neutrino number fluxes $\Phi_\nu$ for the different flavors. This case would practically correspond to neutrino spectra with a \emph{single crossing} point at $E \to \infty$ (where $F_{\nu_e}= F_{\nu_x}$, and $F_{\overline\nu_e}= F_{\overline\nu_x}$) since $F_{\nu_e}(E) > F_{\nu_x}(E)$ for all the relevant energies (and analogously for $\bar\nu$). For spectra with a single crossing, it has been shown in~\cite{EstebanPretel:2007ec} that the presence of significant flavor asymmetries between $\nu_e$ and $\overline\nu_e$ fluxes guarantees the synchronization of different angular modes at low-radii ($r\lesssim 100$~km), so that essentially nothing happens close to the neutrinosphere because the in-medium mixing is extremely small. Therefore, the possible onset of multi-angle effects is delayed after the synchronization phase. Then, the flavor evolution is adiabatic to produce spectral splits and swaps but not enough to allow the multi-angle instability to grow and produce significant decoherence effects. Therefore, the resultant neutrino flavor conversions would be described by an effective ``quasi single-angle'' behavior. In this case, the self-induced spectral swaps and splits would be only marginal smeared by multi-angle effects, as explicitly shown in~\cite{Duan:2006an,Fogli:2007bk,Fogli:2008pt}. This reassuring result has been taken as granted in most of the further studies that typically adopted the averaged single-angle approximation. However, this nice picture does not represent the end of the story for multi-angle effects.% \footnote{ We mention that recently multi-angle effects have been included also in the study of the flavor evolution of the $\nu_e$ neutronization burst in O-Ne-Mg supernovae. We address the interested reader to~\cite{Cherry:2010yc}.} A different result has been recently shown in~\cite{Duan:2010bf}, where multi-angle effects are explored, assuming a flux ordering of the type $\Phi_{\nu_x} \gtrsim \Phi_{\nu_e} \gtrsim \Phi_{\overline\nu_e}$, possible during the SN cooling phase, where one expects a moderate flavor hierarchy among different species and a ``cross-over'' among non-electron and electron species is possible~\cite{Raffelt:2003en, Fischer:2009af, Huedepohl:2009wh}. This case would correspond to neutrino spectra with \emph{multiple crossing} points, i.e. with $F_{\nu_e}(E) > F_{\nu_x}(E)$ at lower energies, and $F_{\nu_e}(E) < F_{\nu_x}(E)$ at higher energies (and analogously for $\bar\nu$). For such spectra, it has been shown in~\cite{Raffelt:2008hr} that the synchronization is not a stable solution for a neutrino gas in presence of a large neutrino density. Therefore, collective flavor conversions would be possible at low-radii in the single-angle scheme~\cite{Raffelt:2008hr}, in a region where one would have naively expected synchronization. However, it has been shown in~\cite{Duan:2010bf} that the presence of a large dispersion in the neutrino-neutrino refractive index, induced by multi-angle effects, seems to block the development of these collective flavor conversions close to the neutrinosphere. This recent result extends the finding obtained with a toy model in~\cite{Raffelt:2008hr}. The delay of the self-induced flavor conversions for this case is also visible in the multi-angle simulations in~\cite{animations}. So, it is apparent that multi-angle effects are relevant not only for fluxes with small flavor asymmetries, where they trigger a quick flavor \emph{decoherence}, but also in cases of spectra with multiple crossing points, where multi-angle effects can \emph{suppress} flavor conversions at low-radii. Triggered by the contrasting impact of the multi-angle effects for fluxes typical of the accretion and cooling phase, we take a closer look at the dependence of these effects on the neutrino flux ordering. The plan of our work is as follows. In Sec.~II we introduce our supernova flux models, and describe the equations for the flavor conversions in the multi-angle and single-angle case. In Sec.~III we show and explain our numerical results for the single-angle and multi-angle flavor evolution for some representative choices of SN neutrino fluxes. In particular, we select three cases corresponding respectively to $a)$ single-crossed neutrino spectrum with $\Phi_{\nu_e}\gg \Phi_{\overline\nu_e}\gg \Phi_{\nu_x}$, producing a ``quasi single-angle'' flavor evolution, $b)$ multiple-crossed spectrum with $\Phi_{\nu_x} \gtrsim \Phi_{\nu_e} \gtrsim \Phi_{\overline\nu_e}$, where single-angle and multi-angle evolutions give significantly different final neutrino spectra, $c)$ small flavor asymmetries, i.e. $\Phi_{\nu_e} \approx \Phi_{\overline\nu_e}\approx\Phi_{\nu_x}$, where the multi-angle suppression is small, and multi-angle decoherence produces a partial flavor equilibration among the different species. Finally, in Sec.~IV we draw inferences from our results and summarize. Technical aspects are discussed in the Appendix.

We have performed an exploration on the dependence of multi-angle effects in self-induced supernova neutrino oscillations on the original neutrino fluxes. Most of the previous studies~\cite{Duan:2006an,Fogli:2007bk,Fogli:2008pt,EstebanPretel:2007ec} focused on neutrino fluxes typical of the accretion phase, with a pronounced $\nu_e$ excess, \emph{de facto} behaving like spectra with a single crossing. In this situation, the synchronization of different angular modes at low-radii prevails over multi-angle effects. Then, when flavor conversions start, these are adiabatic to produce the spectral swaps and splits, but not enough to allow multi-angle decoherence to emerge. The result is the known ``quasi single-angle'' evolution. However, one has to be cautious in generalizing this reassuring result. In this context, we have shown that multi-angle effects can produce significant deviations in the flavor evolution with respect to the three-flavor single-angle case, for neutrino fluxes with a moderate flavor hierarchy and a $\nu_x$ excess, as possbile during the supernova cooling phase. In this situation, the presence of multiple crossing points in the original neutrino spectra destabilizes the synchronization at large neutrino densities~\cite{Raffelt:2008hr}. In absence of synchronization, in the single-angle scheme collective conversions would be possible at low-radii. However, multi-angle effects introduce a large dispersion in the neutrino-neutrino potential, that prevents any possible collective flavor conversion at low radii. As a consequence, in the multi-angle scheme there would be a significant delay of the onset of the flavor conversions, as recently observed~\cite{Duan:2010bf}. We have shown that this multi-angle delay can produce dramatic changes, not only in the deepest supernova regions, as shown in~\cite{Duan:2010bf}, but also in the final oscillated neutrino spectra. In particular, the multi-angle suppression can be so strong to allow the onset of the flavor evolution only at a large radius, when the evolution is less adiabatic. Depending on the violation of adiabaticity, in the inverted mass hierarchy there could be a suppression of the three-flavor effects, associated to the solar sector. This would dramatically change the pattern of swaps and splits in the final neutrino spectra. Finally, if the flavor asymmetries between $\nu_e$ and ${\overline \nu}_e$ are very small, the multi-angle suppression occurs only close to the neutrinosphere. In this situation, the stronger adiabaticity of the evolution allows multi-angle effects to act efficiently also after the onset of the conversions, tending to establish a three-flavor equilibration in both neutrino and antineutrino sectors. In Table~I we summarize our results on the role of multi-angle effects, $3\nu$ effects and spectral splits for different SN neutrino fluxes. In this work we have explicitly shown numerical results only for the case of neutrino inverted mass hierarchy. However, we have checked that the impact of multi-angle effects is qualitatively similar also for the normal hierarchy case. From our numerical explorations, it results that self-induced flavor transformations of supernova neutrinos during the cooling phase are a continuous source of surprises. The richness of the phenomenology, in the presence of neutrino spectra with multiple crossing points, was first realized in~\cite{Dasgupta:2009mg} with the discovery of the possibility of multiple spectral splits in both the mass hierarchies. Then, it was realized that for these neutrino fluxes, three-flavor effects can play a significant role in inverted mass hierarchy, changing the splitting pattern expected from two-flavor calculations~\cite{Dasgupta:2010cd,Friedland:2010sc}. Now, we show that also multi-angle effects are crucial in characterizing the flavor evolution in this case, and could potentially kill the three-flavor effects. In general, self-induced flavor conversions for spectra with multiple crossing points challenge most of the naive expectations on which was based the original picture of the collective supernova neutrino conversions: low-radii synchronization, subleading role of multi-angle and three-flavor effects. The discovery of these new effects adds additional layers of complications in the simulation of the flavor evolution for supernova neutrinos. In particular, our result shows that during the cooling phase three-flavor multi-angle simulations are crucial to obtain a correct result. Multi-angle effects would be taken into account to assess the impact of collective neutrino oscillations on the r-process nucleosynthesis in supernovae, as recently investigated in~\cite{Duan:2010af}. The impact of the multi-angle effects would crucially depend on different SN input: neutrino luminosities and flavor asymmetries, neutrinosphere radius, etc. Since all these quantities significantly change during the neutrino emission, one would expect time-dependent effects. At this regard, the possibility to detect signatures of these effects in the next galactic supernova neutrino burst~\cite{Choubey:2010up} would motivate further analytical and numerical investigations. \subsection*{Appendix} We discuss here a few technical aspects of the multi-angle numerical simulations, we performed on our local computer facility (with Fortran 77 codes running a Linux cluster with 48 processors per CPU with 128 Gb of shared RAM memory). Equation~(\ref{eomMA}), after discretization, provides a set of $16\times N_E \times N_u$ ordinary differential equations in $r$, where $N_E$ and $N_u$ are the number of points sampling the (anti)neutrino energy $E$ and emission angle. In particular, we find convenient to label the neutrino angular modes in terms of the variable~\cite{EstebanPretel:2007ec} \begin{equation} u = \sin^2 \theta_R \,\ , \end{equation} where $\theta_R$ is the zenith angle at the neutrino sphere $r = R$ of a given mode relative to the radial direction. With this choice, the parameter $u$ is fixed for every neutrino trajectory. We have then performed our simulations of the three-flavor neutrino evolution using a Runge-Kutta integration routine taken from the CERNLIB libraries~\cite{cernlib}. We fixed the numerical tolerance of the integrator at the level of $10^{-6}$ and increased the number of sampling points in angle and energy till we reach a stable numerical result. In this situation we estimate a numerical (fractional) accuracy of our results better than $10^{-2}$. In order to have a clear energy resolution of the spectral splits we took $N_E= 10^2$ energy points, equally distributed in the range $E \in [0.1,80]$~MeV. The number of angular modes is also a crucial choice, since it is well known that a sparse sampling in angle can lead to numerical artifacts that would destroy the collective behavior of the neutrino self-induced conversions~\cite{EstebanPretel:2007ec}. Typically, numerical stability would require $N_u = {\mathcal O}(10^3)$ angular modes. Since we can claim to have reached stable numerical simulations, we are confident in the accuracy of the results obtained in this work. Moreover, we have been able to reproduce previous results presented in literature for cases similar to the ones we are investigating (see, e.g.,~\cite{Fogli:2007bk,Duan:2010bf}).

1012.1339

1012

1012.4988_arXiv.txt

We consider a holographic cosmological model in which the infrared cutoff is fixed by the Ricci's length and dark matter and dark energy do not evolve separately but interact non-gravitationally with one another. This substantially alleviates the cosmic coincidence problem as the ratio between both components remains finite throughout the expansion. We constrain the model with observational data from supernovae, cosmic background radiation, baryon acoustic oscillations, gas mass fraction in galaxy clusters, the history of the Hubble function, and the growth function. The model shows consistency with observation.

Holographic models of late acceleration have become of fashion, among other things, because they link the dark energy density to the cosmic horizon, a global property of the universe, and have a close relationship to the spacetime foam \cite{jng1}, \cite{arzano}. For a summary motivation of holographic dark energy, see section 3 in \cite{cqg-wd}. The cosmological model we present here assumes a spatially flat homogeneous and isotropic universe dominated by dark matter (DM) and dark energy (DE) (subscripts $M$ and $X$, respectively), the latter obeying the holographic relationship \begin{equation} \rho_{X}= \frac{3 M_{P}^{2}\, c^{2}}{L^{2}}\, , \label{rhox} \end{equation} where $c^{2}$ is a dimensionless parameter that we will take as constant -though strictly speaking it may slowly vary with time \cite{jcap-nd}-, and $L = (\dot{H} +2H^{2})^{-1/2}$ denotes the Ricci's lenght. We adopt the latter because it corresponds to the size of the maximal perturbation leading to the formation of a black hole \cite{brustein}. For models that identify $L$ with Hubble's length, $H^{-1}$, see \cite{jcap-idw} and references therein. The other main assumption is that DE and DM interact non-gravitationally with each other according to \begin{eqnarray}\label{eq:EvolIn} \dot{\Omega}_{M} - \left(1-\frac{2\Omega_{X}}{c^{2}}\right)(1-\Omega_{X})\, H&=&QH \, ,\\ \label{eq:EvolIn2} \dot{\Omega}_{X}+ \left(1-\frac{2\Omega_{X}}{c^{2}}\right)(1-\Omega_{X})\, H&=&-QH \, . \end{eqnarray} Here $\Omega_{M}$ and $\Omega_{X}$ stand for the fractional densities of the components, and the interaction term \begin{equation} Q = -\frac{r_{f}}{(1+r_{f})^{2}}\left(1+r_{f}-\frac{2}{c^{2}}\right) \, , \label{eq:QChim} \end{equation} is such that ratio $r \equiv \Omega_{M}/\Omega_{X}$ evolves from a constant value at early times to a final, finite, value -denoted as $r_{f}$- at late times (see \cite{ivan-diego} for details and Fig. \ref{fig:r}). This clearly alleviates the coincidence problem (namely, ``why are the densities of matter and dark energy of the same order precisely today?" \cite{steinhardt}), something beyond the reach of the $\Lambda$CDM model. A consequence of the model is that the equation of state of dark energy $w \equiv p_{X}/\rho_{X}$ varies with expansion as shown in Fig. \ref{fig:w}. \begin{figure}[htb] \begin{center} \includegraphics[width=10cm]{rRicci} \end{center} \caption{\label{fig:r} Plot of the ratio $r$ between the energy densities vs. redshift for the best-fit model. As the inset shows, $r_{f} \equiv r(z)$ does not vanish when $z \rightarrow -1$. In this, and the next figure, the red swath indicates the region obtained by including the $1\sigma$ uncertainties of the constrained parameters used in the calculation.} \end{figure} \begin{figure}[htb] \hspace{-0.5 cm} \begin{minipage}{0.3\textwidth} \centering \includegraphics[width=8cm]{wRicci} \end{minipage} \hspace{3.5 cm} \begin{minipage}{0.3\textwidth} \centering \includegraphics[width=8cm]{wRicciobs} \end{minipage} \caption{The equation of state parameter vs. redshift up to $z = 8$ (left panel), and up to $ z = 1.2$ only (right panel) for the best fit holographic model. At high redshifts $w$ approaches the equation of state of non-relativistic matter and at low redshifts it does not depart significantly from $-1$. The observational data are taken from \cite{serra-prd}. The error bars indicate a $2\sigma$ uncertainty.} \label{fig:w} \end{figure}

The statistical analysis sketched above shows that the ratio $\chi^{2}_{total}/dof$ is lower than unity for the holographic interacting model of section \ref{introduction}, whereby it is compatible with observation. However, as Table \ref{table:chi2} shows, the $\Lambda$CDM model fits better the same sets of data. Yet, the latter cannot explain the cosmic coincidence problem while the former can. \ack{ID was funded by the ``Universidad Aut\'{o}noma de Barcelona" through a PIF fellowship. This research was partly supported by the Spanish Ministry of Science and Innovation under Grant FIS2009-13370-C02-01, and the ``Direcci\'{o} de Recerca de la Generalitat" under Grant 2009SGR-00164.}

1012.4988

1012

1012.1755_arXiv.txt

H.E.S.S is an array of atmospheric Cherenkov telescopes dedicated to GeV-TeV $\gamma$-ray astronomy. The original array has been in operation since the beginning of 2004. It is composed of four 12-meter diameter telescopes. The installation of a fifth 28-meter diameter telescope is being completed. This telescope will operate both in stereoscopic mode and in monoscopic mode \textit{i.e.} without a coincident detection on the smaller telescopes. A second-level trigger system is needed to supress spurious triggers of the 28-meter telescope when operated in monoscopic mode. This paper gives the motivation and principle of the second-level trigger. The principle of operation is illustrated by an example algorithm. The hardware implementation of the second level trigger system of H.E.S.S. phase 2 is described and its expected performances are then evaluated.

\label{sec:intro} The H.E.S.S. (High Energy Stereoscopic System) instrument is an array of four imaging atmospheric Cherenkov telescopes working in stereoscopic mode. It is located in Namibia, in the Khomas Highland and has started its operations in 2004. Each of the present ``small'' Cherenkov telescopes (SCT) has a 12-meter diameter mirror and is equipped with a 960-pixel (photomultiplier) camera at its focal plane. It detects photons in the 100 GeV-50 TeV energy range. {\hlch In addition to photons showers, the combinatorial background from diffuse sky photons and charged cosmic showers can trigger the telescopes. } Stereoscopy allows one to achieve a large rejection of the single muon triggers. These single muons come from very distant hadronic showers and dominate the single-telescope {\hlch particle} triggers \cite{Funk2005}. {\hlcg The single muon trigger rate is discussed in section \ref{sect:part_trig_rates}.} The H.E.S.S. collaboration has started to build a fifth, 28-meter-diameter large Cherenkov telescope (LCT). The LCT will be equipped with a 2048-pixel camera at its focal plane. The LCT will be sensitive to photons down to 10~GeV. In normal operation, the SCTs are triggered only in case of a coincidence with another telescope (LCT or SCT). However, the energy threshold of SCTs is too high to efficiently detect low energy ($\le 50$ GeV) $\gamma$ rays. To increase its acceptance at low photon energies, the HESS instrument will have to accept standalone LCT triggers. {\hlch Assuming similar first level trigger conditions on the SCTs and the LCT, these standalone LCT triggers would have a rate which is typically a factor of five larger than H.E.S.S stereoscopic triggers. As shown later in section \ref{sect:part_trig_rates}, these triggers are mostly background triggers.} The H.E.S.S. collaboration has decided to build a second level (L2) trigger board in the camera of the LCT to improve the rejection of accidental night-sky background triggers and single muon triggers. The L2 trigger board is programmable, which gives flexibility in the choice of the trigger algorithms. For instance, low energy selection algorithms similar to the trigger used by the MAGIC collaboration to detect the pulsed emission from the Crab pulsar \cite{MAGICCrab} can be implemented. These low energy selection algorithms allow to lower the energy threshold on the LCT. Alternatively, at constant energy threshold, the gain in bandwidth obtained by rejecting the background events can be used to transfer timing informations on fired pixels to the acquisition farm, in addition to the total charge. The timing information may be useful for analyzing {\hlcg single telescope} events, as has also been shown by the MAGIC collaboration \cite{MAGICtiming}. Topological triggers have been previously used on other Cherenkov instruments. For example, the MAGIC collaboration \cite{Bastieri2001} {\hlch uses a N-next-neighbor logic in its first level trigger and} has designed a second-level trigger which can perform a rough event analysis and {\hlcg can apply} topological cuts to the images. In the first part of this paper, the various contributions to the LCT instrument trigger rate are listed and evaluated. The next section is devoted to the L2 concept and an example L2 trigger algorithm is given. The actual L2 trigger board is described in section \ref{sect:hardware}. Finally, the on-board implementation of the L2 algorithm is discussed in section \ref{sect:firmware}.

This paper describes the design and implementation of the L2 trigger system for the second phase of the H.E.S.S. experiment. The L2 trigger will be used to reject night sky background related and isolated muon events and thus reduce the trigger rate. The principle of the trigger is to build a 2-bit (``combined'') map of the camera pixels at the time of trigger. The night sky background events can then be rejected by demanding clusters of pixels on the combined map. Further rejection of the hadronic background can be obtained by using quantities such as the center of gravity of the triggered pixels. A possible, illustrative, algorithm for the L2 trigger system has been given in section \ref{section2:algorithms}. This example algorithm shows that the required rejection of night sky background and isolated muon triggers is achievable. The hardware and software integration into the LCT camera of the previously described system based on a single Virtex-4FX12 FPGA has been achieved. The L2 system still needs to be fully integrated in the H.E.S.S. acquisition and tested with real data. This will be achieved at the beginning of the HESS-2 phase.

1012.1755

1012

1012.3750_arXiv.txt

We build a model for the density and integrated Sachs-Wolfe (ISW) profile of supervoid and supercluster structures. Our model assumes that fluctuations evolve linearly from an initial Gaussian random field. We find these assumptions capable of describing N-body simulations and simulated ISW maps remarkably well on large scales. We construct an ISW map based on locations of superstructures identified previously in the SDSS Luminous Red Galaxy sample. A matched filter analysis of the cosmic microwave background confirms a signal at the $3.2-\sigma$ confidence level and estimates the radius of the underlying structures to be $55 \pm 28 \mpc$. The amplitude of the signal, however, is $2-\sigma$ higher than $\Lambda$CDM predictions.

Since the first sky maps of the Wilkinson Microwave Anisotropy Probe (WMAP) were published, there have been claims for the existence of circular features (spots and rings) in the cosmic microwave background (CMB) \citep{Cruz2005,Granett2008}. Their origin and statistical significance are still debated. Sources, such as foreground contamination, integrated Sachs-Wolfe (\citealp{SW1967}; ISW) effect and the more exotic conformal cyclic cosmology or cosmic texture have been considered viable candidates to explain circular CMB anomalies \citep{Rudnick2007,Inoue2010,Cruz2007,bs}. A notable feature is the cold spot \citep{Cruz2005}, which has a mean temperature of $-70\mu K$ in a $5^{\circ}$-radius aperture. Additionally, $10\mu K$ hot and cold spots have been identified on $4^{\circ}$ scales associated with super structures \citep{Granett2008}. In this paper we focus on the integrated Sachs-Wolfe effect and investigate whether it is possible that large spherical fluctuations (supervoids and superclusters) in the dark matter field produce the aforementioned features in the CMB. The large-scale ISW effect is expected in a universe with accelerated cosmic expansion arising either from a dark energy term in flat cosmological models, or from spatial curvature. The effect is sensitive to both the expansion history and the rate of structure formation and provides constraints on alternative cosmological models (e.g. \citealp{Giannantonio2008b}) as well as initial non-Gaussianity (e.g. \citealp{Afshordi2008}). Cross correlating a galaxy catalog and a CMB temperature map is the standard way of studying the ISW signal and it has an extensive literature \citep{Scranton2003,Afshordi2004,Padmanabhan2005,Raccanelli2008,Giannantonio2008,Ho2008,Sawangwit2010}. Our work is similar considering what we measure is related to the correlation between the dark matter and the CMB. However, we only focus on parts of the sky where the signal is expected to be large, regions corresponding to supervoids and superclusters. Since reports \citep{Cruz2005,Granett2008} indicate that the scale of these regions is beyond nonlinear scales, our model for their average density profile and ISW imprint is derived from the statistics of the linearly evolving primordial Gaussian density field \citep{Papai2010}. This is in contrast with \cite{Inoue2010}, who assumed a top hat density profile, which is the asymptotic final state of a void with steep initial density profile \citep{Sheth2004}. We use the publicly available Hubble Volume Simulation of the Virgo Supercomputing Consortium (\citealp{Colberg2000}; HVS) to test the Gaussian model. This is the N-body simulation with the largest volume, which is relevant considering that the gravitational potential is correlated more strongly than the density. We simulate ISW maps by ray tracing through the potential and calculating the linear part of the ISW effect. These are compared to partial ISW maps generated from sets of spherical regions based on our model. After gaining some confidence by studying simulations we create an ISW map from real data. We select locations on the sky based on \cite{Granett2008}. They compiled a list of supervoids and superclusters found in the Sloan Digital Sky Survey (SDSS) Luminous Red Galaxy (LRG) sample. The list can be found in \cite{supplement}. We build an ISW map by placing the theoretical profiles to the given R.A., decl. coordinates. This map is fitted to a WMAP temperature map. The structure of the paper is the following: in Section \ref{PinS} we measure the expected density and ISW profile of spherical dark matter fluctuations in N-body simulations; in Section \ref{Match} we apply the matched filter technique to detect the signature of superstructures in the CMB; in Section \ref{discussion} we discuss our results and views of the relationship between the linear ISW effect and circular features on the CMB.

\label{discussion} This paper is the continuation of the work started in \cite{Papai2010}. Our goal is to estimate the ISW imprint of large spherical dark matter overdensities and underdensities. In Section \ref{PinS}, first, we tested a simple Gaussian model for density profiles in the HVS. We created an ISW map by ray tracing through the simulation and computing the linear ISW effect. We neglected the nonlinear part of the ISW effect, as it is small in comparison at $z=0$ \citep{Cai2010}. Another simplification was that our ISW map was Euclidean. For the purpose of testing this has no relevance. As it can be seen from Figure \ref{HVsig2}, the density profiles followed linear theory given by Equation (\ref{aveDelta}) within $1$-$\sigma$ uncertainty up to $400\mpc$, the largest scale in our study. This is not surprising, since the density profile is equivalent of the two-point function and on large scales high-order clustering is not important. When we averaged the density profiles we only used locations where the average density contrast in a certain radius was in the $2$-$\sigma$ (positive or negative) tail of the PDF. This demonstrates that not even the extreme cases are affected by nonlinearities significantly. We demonstrated that cosmic variance was a much more important factor for the ISW effect than for the density. Since cosmic variance is large for low $k$-modes and the potential is non-vanishing as $k\rightarrow 0$, these modes have relatively large effect on the ISW profile. Surely large variance modes should be removed from the CMB and the ISW maps if the ISW maps are based on the statistics of density fluctuations and not on the observed density. First we measured the ISW profiles after filtering out the low $k$-modes from the HVS density and consequently the ISW map. Filtering the two-point function in the same way gave a good match to the data. (See Figure \ref{iswprofHV}.) If the ISW map is given, as in the CMB, another route has to be taken. We filtered the complete ISW map and compared it to the theoretical profile after removing the same modes. The result is shown in Figure \ref{remove} In Section \ref{Match} we used matched filter technique to detect the ISW signal of superstructures in the CMB. We closely followed the procedure described in \cite{Granett2009}. The major difference is that our ISW map was built from theoretical assumptions about the shape of ISW profiles, while in \cite{Granett2009} the authors used an analytic curve which fitted the measured profiles the best. Because of this, the significance of our measurement can readily be interpreted. We estimated the marginalized significance of our measurement to be $\sigma = 3.24$. (See Subsection \ref{results} for details.) The interpretation of the best fitting amplitude of the ISW map is still not without difficulties. It is several times higher than the anticipated signal (Figures \ref{significance} and \ref{single}), although it appears to be $2$-$\sigma$ higher than $\Lambda$CDM predictions. Despite the difference in amplitude, the theoretical and the measured curves run parallel with each other, which supports that the signal is related to the ISW effect. We caution though that the uncertainty can still be underestimated, since it was derived from data alone. The ISW effect provides an explanation for another feature of observations. On average supervoids tend to have a hot ring around an initial cold dip and the opposite is true for superclusters \citep{Granett2009}. When we erased large scale fluctuations in the microwave background, the expected ISW profile changed sign as seen in Figure \ref{sphiswprof}. This means that in certain, not particularly unique realizations the average profiles of 50 supervoids or superstructures from a finite area of the sky can have rings in the ISW context, simply due to the cosmic variance of low $k$-modes. Overall we can say that the shape of the measured signal follows the predicted ISW profile while its amplitude exceeds expectation. This is a good reason to investigate further by studying galaxy surveys other than the SDSS. A key to a quantitative study is the well-measured galaxy density. Since the model described in this paper is a model for density fluctuations around a particular location, which is not necessarily a maximum, it is enough to know the density at this location accurately. We thank Mark Neyrinck for sharing his knowledge on void and cluster finders. The authors were supported by NASA grants NNX10AD53G and NNG06GE71G, and the Pol\'anyi program of the Hungarian National Office for the Research and Technology (NKTH). The Hubble Volume Simulations were carried out by the Virgo Supercomputing Consortium using computers based at the Computing Centre of the Max-Planck Society in Garching and at the Edinburgh parallel Computing Centre. The data are publicly available at http://www.mpa-garching.mpg.de/NumCos.

1012.3750

1012

1012.4346_arXiv.txt

X-ray observations of solar flares routinely reveal an impulsive high-energy and a gradual low-energy emission component, whose relationship is one of the key issues of solar flare study. The gradual and impulsive emission components are believed to be associated with, respectively, the thermal and nonthermal components identified in spectral fitting. In this paper, % a prominent $\sim 50$ second hard X-ray (HXR) pulse of a simple GOES class C7.5 flare on 20 February 2002 is used to study the association between high energy, non-thermal and impulsive evolution, and low energy, thermal and gradual evolution. We use regularized methods to obtain time derivatives of photon fluxes to quantify the time evolution as a function of photon energy, obtaining a break energy between impulsive and gradual behavior. These break energies are consistent with a constant value of $\sim 11$ keV in agreement with those found spectroscopically between thermal and non-thermal components, but the relative errors of the former are greater than $15\%$ and much greater than the a few percent errors found from the spectral fitting. These errors only weakly depend on assuming an underlying spectral model for the photons, pointing to the current data being inadequate to reduce the uncertainties rather than there being a problem associated with an assumed model. The time derivative method is used to test for the presence of a `pivot energy' in this flare. Although these pivot energies are marginally consistent with a constant value of $\sim 9$ keV, its values in the HXR rise phase appear to be lower than those in the decay phase. Assuming that electrons producing the high-energy component have a power law distribution and are accelerated from relatively hot regions of a background plasma responsible for the observed thermal component, a low limit is obtained for the low-energy cutoff. This limit is always lower than the break and pivot energies and locates in the tail of the Maxwellian distribution of the thermal component.

\label{intro} High-energy observations of solar flares with the Reuven Ramaty High-Energy Solar Spectroscopic Imager (RHESSI) \citep{lin2002rhessi} allows high resolution studies over a broad energy range from 3 keV soft X-rays to $\gamma$-rays up to 17 MeV. The photon flux in the energy range of $\sim 20 - 100$~keV can be reasonably well fitted with a power-law function, and its time-variability increases with the photon energy \citep{asch2005psc,mcateer2007bursty}. It is commonly assumed that this emission is produced by an electron population distinct from electrons forming a thermal background plasma, which is presumed to produce the low-energy X-ray emission \citep[e.g.][]{asch2002paa}. The impulsive high-energy emission originates predominantly from the chromospheric footpoints, while --- at least later in the flare --- the more slowly-varying low-energy emission is dominated by a hot coronal source, observed in many cases to be located near EUV flare loops \citep[e.g.][]{gallagher2002}. These observations are usually interpreted in the framework of the standard flare model where the hard X-ray (HXR) emission at the chromospheric footpoints of magnetic loops is bremsstrahlung of non-thermal high-energy electrons moving downward along flare loops from acceleration sites higher up in the corona \citep{brown1971}, with the resulting footpoint heating and evaporation leading to the hot (usually dense) coronal thermal component \citep{neupert1968, petrosian1973, fisher1989}. Note, we do not automatically adopt this assumed relationship between accelerated and heated particles. In fact, in Section~\ref{s:discussion}, we interpret the observations in a framework where the non-thermal electrons are accelerated out of a heated thermal background. The `non-thermal' electron distribution is usually assumed to have a low-energy cutoff, the presence of which ensures that the total electron number and power are finite. However, it is not clear that there is a theoretical mechanism for particle acceleration which can naturally lead to the low-energy cutoff distinguishing non-thermal from thermal particles \citep{benz1977, miller1997jgr, petrosian_liu2004}. Indeed, it has been argued by \citet{emslie2003} that the low energy cutoff may be a redundant concept. \citet{hannah2009wave} suggested that a sharp cutoff in the injected electron spectrum disappears with the inclusion of wave-particle interactions. A dip in the electron distribution obtained through the inversion of the observed photon spectrum of some flares may be associated with the low-energy cutoff. \citet{Kontar2008cutoff}, however, showed that such a feature vanishes when isotropic albedo correction is applied. The time correlation between the impulsive HXR and/or radio emission and the derivative of the gradual emissions at certain energies, the so-called `Neupert effect', \citep{neupert1968, dennis1993} carries with it the implication that most gradual emissions are a ``by-product", resulting from energy deposition by non-thermal electrons. This is also suggested by the high non-thermal electron energy content resulting from application of the standard collisional thick-target model \citep[e.g.][]{emslie2004,emslie2005}, which points to high-energy electrons forming a dominant channel in the energy conversion process. However, more quantitative examination of relevant observations show that the picture is somewhat less clear. Some flares involve heating of thermal coronal plasma in the absence of a power-law emission component \citep{battaglia2009}, many show footpoints with impulsive phase emission within the energy range usually considered as thermal \citep{mrozek2004}, and the `Neupert effect', which never represents a perfect correlation, does not hold in all flares or at all (thermal) energies \citep{mctiernan1999,veronig2002, veronig2005}. So the possibility of heating of the solar plasma as a direct part of the energy release before and/or during the acceleration is still an open issue for investigation \citep{petrosian_liu2004, liu2009elementary}. By deriving abundances of elements with low first ionization potentials, such as calcium and iron, \citet{feldman2004} found that at least the hot plasmas of some flares result from direct in situ heating of corona plasma, possibly due to a compression process. This approach may also lead to a measurement of the partition of hot flare plasmas originated from the corona and chromosphere. The general question of how the pre-flare magnetic energy is converted into radiation, plasma bulk motion, thermal and non-thermal particle energy may not have a simple answer \citep{emslie2005}. Although flares share the same kind of energy source, different flares can have quite different appearances, and possibly involve different physical processes. Nevertheless, some well-observed characteristics can still set constraints on the overall energy dissipation process. The soft-hard-soft spectral evolution of some HXR pulses is one of the most important characteristics of high-energy emissions \citep{kane1970,grigis2004spectral} and may point to a turbulent particle acceleration mechanism \citep{grigis05}. Early analyses \citep{gan1998invariable} suggested that there is a value of photon energy at which the non-thermal flux does not change, so that the power-law pivots about this location, a possible further model constraint. \citet{grigis2004spectral} showed that there is no single `pivot energy', rather there is a small range. \cite{battaglia2006} determined that in the rise phase this energy may be lower than that in the decay phase. In the context of stochastic particle acceleration from the thermal background plasma, the pivot energy should evolve with the background plasma properties \citep{petrosian_liu2004,liu2010elementary}. However, distinguishing between different models on the basis of observations remains a challenging task \citep{grigis05}. Other constraints based on the evolution of HXR light curves include the observation that sub-second HXR pulses peak earlier in high than in low energies, consistent with a time-of-flight dispersion if the electrons producing these pulses are accelerated at some distance from the location where the bremsstrahlung radiation is produced \citep{asch1996scaling}. However, this does not mean that all the energetic electrons have to be associated with these sub-second pulses. The reverse delay in the longer timescale (seconds) HXR pulses could indicate collisional escape from a coronal trap \citep{asch1996loop, asch1998deconvolution, krucker2008HXR}, but could also be a result of a more gradual acceleration process \citep[e.g.][]{bai1979}. In this paper we investigate the characteristics of flare emission across a range of photon energies, and examine the association between temporal, spatial and spectral characteristics, with particular interest in the region between thermal and non-thermal parts of the spectrum. The paper is organised as follows. We first review theoretical considerations and present a simple model for the flare with an isothermal and a power-law X-ray emission component (Section \ref{theory}). A simple RHESSI flare on 20th February 2002, with distinct gradual low-energy and impulsive high-energy emissions is analysed in detail (Section \ref{s:obs}). An overview of the flare is presented in Section \ref{s:overview}. The semi-calibrated photon flux is then used to derive the rate of change of photon fluxes at different energies during a prominent HXR pulse, and two temporal components are identified (Section \ref{s:semi_ph}). This is repeated in Section \ref{s:spectral_fit} but using a full spectral fit. The evolution of model parameters and the corresponding photon fluxes are used to check self-consistency of the model, and in Section \ref{s:pivot} we look at the pivot energy derived from the rate of change of the photon fluxes. In Section \ref{s:discussion}, we discuss the implications of these results, and conclusions are drawn in Section \ref{s:conclusions}.

\label{s:conclusions} We have developed a new method to study in detail the temporal evolution of thermal and non-thermal photon fluxes in solar flares. The application of this method to a flare on Feburary 20 2002 demonstrates that as expected, the low energy part of the spectrum evolves slowly, and the high energy part evolves rapidly, with an intermediate range between a few keV and 20 keV where the behavior is in transition. The data support the scenario in which the non-thermal component of the flare spectrum is impulsive, and the thermal component is gradual, in that the transition energies between these two behaviors are the same within errors whether examined in time or in energy. However, although in the spectral fitting exercise it is possible to make a clean separation between a non-thermal, impulsive component, and a thermal, gradual component, time evolution gives a more ambiguous picture, due to the large error bars. Imaging is also ambiguous, with no clear distinction between footpoints and loops in the energy range around 9-25~keV. Therefore we must leave open the possibility that the electrons form a continuous distribution over this range. Further studies with larger flares should help to improve the precision with which we can identify the transition between gradual and impulsive behavior. Finally, the presence of a single pivot point throughout the flare is not supported by our analysis, though a pivot `range' is. There is some evidence of a slightly higher value for this pivot range in the decay phase than in the rise phase.

1012.4346

1012

1012.0345_arXiv.txt

We present some results about a novel damping mechanism of r-mode oscillations in neutron stars due to processes that change the number of protons, neutrons and electrons. Deviations from equilibrium of the number densities of the various species lead to the appearance in the Euler equations of the system of a dissipative mechanism, the so-called rocket effect. The evolution of the r-mode oscillations of a rotating neutron star are influenced by the rocket effect and we present estimates of the corresponding damping timescales. In the description of the system we employ a two-fluid model, with one fluid consisting of all the charged components locked together by the electromagnetic interaction, while the second fluid consists of superfluid neutrons. Both components can oscillate however the rocket effect can only efficiently damp the countermoving r-mode oscillations, with the two fluids oscillating out of phase. In our analysis we include the mutual friction dissipative process between the neutron superfluid and the charged component. We neglect the interaction between the two r-mode oscillations as well as effects related with the crust of the star. Moreover, we use a simplified model of neutron star assuming a uniform mass distribution.

Fermionic superfluidity might be realized in the interior of neutron stars where the temperature is low and the typical energy scale of particles is extremely high. In particular, in the inner crust of standard neutron stars the attractive interaction between neutrons can lead to the formation of a BCS condensate. A way to detect or to discard the presence of a superfluid phase in compact stars consists in studying the evolution of the r-mode oscillations~\cite{Andersson:2000mf}. R-modes are non-radial oscillations of a fluid with the Coriolis force acting as the restoring force. They provide a severe limitation on the rotation frequency of a star through coupling to gravitational radiation. When dissipative phenomena damp these oscillations the star can rotate without losing angular momentum to gravitational radiation. If dissipative phenomena are not strong enough, the r-mode oscillations will grow exponentially fast in time and the star will keep slowing down until some dissipation mechanism is able to damp the r-mode oscillations. In this way one can put some constrains on the stellar structure ruling out phases that do not have large enough viscosity. For such studies it is necessary to consider in detail all the dissipative processes and the interactions among the various layers of a star. In Ref. \cite{Colucci:2010wy} we have studied a novel dissipative process associated with the change in the number of protons, neutrons and electrons. In real neutron stars these processes can take place in the outer core and in the inner crust of the star and are related to beta decays and interactions between the neutron fluid and the crust. Both these processes lead to the appearance of a dissipative force (the so-called rocket term) in the Euler equations of the system. This force is due to the fact that when two or more fluids move with different velocities a change of one component into the other results in a momentum transfer between the fluids. This change in momentum is not reversible, because it is always the faster moving fluid that will lose momentum. The resulting dissipative force is proportional to the mass rate change, and to the relative velocity between the fluids. The name ``rocket effect" reminds that the same phenomenon takes place in the dynamical evolution of a rocket whose mass is changing in time as it consumes its fuel. As far as we know, the dissipative force due to the rocket term has not been considered in the context of r-mode oscillations. In the hydrodynamical equations corresponding to the mass conservation laws, it is in general assumed that the neutron, proton and electron components are separately conserved quantities. Assuming that the change in the particle densities is due to out of equilibrium direct Urca processes, we have determined the typical timescale associated with the rocket effect and we have found that it is sufficiently short to damp countermoving r-mode oscillations, with the normal and superfluid component oscillating out of phase.

1012.0345

1012

1012.0035_arXiv.txt

We study a paradigmatic system with long-range interactions: the Hamiltonian Mean-Field Model (HMF). It is shown that in the thermodynamic limit this model does not relax to the usual equilibrium Maxwell-Boltzmann distribution. Instead, the final stationary state has a peculiar core-halo structure. In the thermodynamic limit, HMF is neither ergodic nor mixing. Nevertheless, we find that using dynamical properties of Hamiltonian systems, it is possible to quantitatively predict both the spin distribution and the velocity distribution functions in the final stationary state, without any adjustable parameters. We also show that HMF undergoes a non-equilibrium first-order phase transition between paramagnetic and ferromagnetic states.

1012.0035

1012

1012.1175_arXiv.txt

Cosmological runaway solutions may exhibit an exact dilatation symmetry in the asymptotic limit of infinite time. In this limit, the massless dilaton or cosmon could be accompanied by another massless scalar field - the bolon. At finite time, small time-dependent masses for both the cosmon and bolon are still present due to imperfect dilatation symmetry. For a sufficiently large mass the bolon will start oscillating and play the role of dark matter, while the cosmon is responsible for dark energy. The common origin of the mass of both fields leads to an effective interaction between dark matter and dark energy. Realistic cosmologies are possible for a simple form of the effective cosmon-bolon-potential. We find an inverse bolon mass of a size where it could reduce subgalactic structure formation.

Dilatation symmetry and its anomaly could play an important role for cosmology \cite{Wetterich:1987fm}. Models with a dilatation symmetric fixed point could provide a dynamical solution for the cosmological constant problem \cite{Wetterich:1994bg, Wetterich:1987fm, Wetterich:2002wm}. As one of the most characteristic features such models have predicted the presence of a homogeneous dark energy component \cite{Wetterich:1987fm}, long before its observational discovery. Recent investigations of higher dimensional settings have shed new light on such theories \cite{Wetterich:2008bf}. It has been shown that dimensional reduction of a dilatation symmetric quantum effective action in higher dimensions leads to four-dimensional models with a vanishing cosmological constant. If the cosmological solution approaches the dilatation symmetric fixed point in the limit $t\rightarrow \infty$, such models can naturally give rise to an asymptotically vanishing cosmological constant. In such a scenario the observed particle masses are due to the spontaneous breaking of dilatation symmetry by the cosmological or vacuum solution. At the fixed point one will therefore encounter a massless goldstone boson - the dilaton. For finite time, the dilatation symmetry is broken by anomalous terms which generate a small time-dependent mass for the dilaton. This picture can give rise to a quintessence cosmology where the dilaton rolling towards the fixed point plays the role of the scalar "cosmon"-field with a slowly decreasing mass \cite{Wetterich:1987fm, Wetterich:1994bg}. This higher dimensional dilatation symmetric setup naturally leads to scaling solutions where the cosmon mass tracks the Hubble parameter. First investigations of the manifold of extrema of a dilatation symmetric higher-dimensional quantum effective action often reveal the presence of additional massless scalar fields besides the dilaton \cite{Wetterich:2008bf}. These may correspond to a change of the characteristic length scale of internal space - the radion -, or other changes in geometry, similar to the moduli fields in string theory. In this work we will concentrate on one such field and name it the "bolon'', since it will be ultimately responsible for dark matter and therefore for the emergence of structure (i.e. "lumps" or "bola") in the cosmos. We start with a quick revision of the key concept of dilatation symmetry. Then we will investigate the coupled system of cosmon and bolon and show that it can reproduce the standard cosmological evolution at the background level for rather simple potentials. After a radiation dominated period, during which the cosmon and bolon act as early dark energy and play a subdominant role, the bolon starts to oscillate around a partial potential minimum. Its oscillation energy (potential and kinetic) is diluted as non-relativistic matter and thus it will eventually become dominant, enforcing a transition to a matter dominated period. After the transition the fluid equations for the energy density of bolon fluctuations obey the standard form for cold dark matter, with a small coupling to the cosmon. Thus the simple model with two scalars describes a cosmology with coupled dark matter and dark energy. We further proceed to slight modifications of the simplest potential for which dark energy finally dominates the energy density of the universe. This can be achieved by an effective stop in the cosmon evolution, induced either by a characteristic change in the scalar potential or a leaping kinetic term, or if the scaling behaviour of the cosmon gets terminated by a cosmic trigger event. An example of the last scenario is given by the growing neutrino quintessence model \cite{Amendola:2007yx}.

In summary, we have found rather simple models for coupled cosmon and bolon scalar fields which realize a consistent cosmology without the need of other dark matter particles. The essential ingredients are the presence of a period where the bolon energy density is much smaller than radiation, and the increase of the bolon mass relative to the Hubble parameter due to $\beta \ll \alpha$. Remarkably, the inverse bolon mass in the present cosmological period, \beq \label{bolonmass} m_\chi^{-1} = \sqrt{\frac{1}{3}} \frac{\chi_{\rm eq}}{M} \, H_{\rm eq}^{-1} \, e^{ \beta \Delta \varphi / M} \approx \left( \frac{10 \, \chi_0}{M} \right)^4 \, {\rm pc} \, , \eeq is typically found at subgalactic scales and could reduce the clustering of dark matter on these and smaller scales \cite{Sahni:1999qe, Matos:2000ss}. Here $\Delta \varphi = \varphi_{\rm today} - \varphi_{\rm eq}$ and the last factor is neglected in the weak-coupling approximation. We found that for the type of models discussed in section \ref{sec5} values of $\chi_0 / M \approx 1$ or somewhat larger are naturally realized, independently of the precise initial conditions. (A bolon mass of a galactic scale of $10$ kpc is realized for $\chi_0 / M \approx 1$.) Since a coupling of the bolon to ordinary matter has presumably at most gravitational strength, a direct detection by the searches for WIMP-like particles or axions seems excluded. On the other hand, a coupling $\beta$ between dark energy and dark matter is a generic feature of our setting. It reflects the common origin of the cosmon-bolon potential and the cosmon and bolon masses from the deviations from dilatation symmetry. Interestingly, if cosmological measurements should indicate an equation of state w$_{\rm de} < -1$ for uncoupled dark energy, the coupling $\beta$ can be used to explain such observations \cite{Das:2005yj}. Furthermore $\beta$ influences both the behaviour of the cosmological solution and the properties of dark matter on smaller length scales. A test of the interesting observational consequences can constrain $\beta$ or give hints in the direction of the coupled dark energy and dark matter of our model.

1012.1175

1012

1012.4799_arXiv.txt

Optical spectropolarimeters can be used to produce maps of the surface magnetic fields of stars and hence to determine how stellar magnetic fields vary with stellar mass, rotation rate, and evolutionary stage. In particular, we now can map the surface magnetic fields of forming solar-like stars, which are still contracting under gravity and are surrounded by a disk of gas and dust. Their large scale magnetic fields are almost dipolar on some stars, and there is evidence for many higher order multipole field components on other stars. The availability of new data has renewed interest in incorporating multipolar magnetic fields into models of stellar magnetospheres. I describe the basic properties of axial multipoles of arbitrary degree $\ell$ and derive the equation of the field lines in spherical coordinates. The spherical magnetic field components that describe the global stellar field topology are obtained analytically assuming that currents can be neglected in the region exterior to the star, and interior to some fixed spherical equipotential surface. The field components follow from the solution of Laplace's equation for the magnetostatic potential.

The solution of Laplace's equation by separation of variables for the electrostatic potential in a region external to a charge distribution is a standard topic in graduate and undergraduate courses in electromagnetism.\cite{jea27,ble76,hea95,wyl99} As an example application, the long-range interaction between molecular charge clouds can be determined via a multipole expansion of the electrostatic potential obtained by solving Laplace's equation.\cite{gra84} The solution for the magnetostatic potential in a region devoid of current sources is equivalent to the electrostatic case.\cite{gra84} In this paper I show how this approach can be adapted to construct models of the large scale magnetic fields of stars and planets. I will derive the equation for the field lines of an axial magnetic multipole of arbitrary degree $\ell$ with a source surface. The source surface is a spherical surface of radius $R_s$ used by the solar physics community to mimic the effects of the solar wind. The solar wind is continuous streams of outflowing charged particles that open the large scale magnetosphere (the region external to a star, or equivalently a planet, consisting of closed magnetic field lines) of the Sun.\cite{alt69,sch69} At $R_s$ the field is assumed to be purely radial. This model has been successfully adapted to produce models of stellar magnetospheres via field extrapolation from observationally derived magnetic surface maps.\cite{jar02,gre06,gre08} As shown in Fig.~\ref{fig_colored_dipole} multipole magnetic fields with a source surface boundary condition incorporate regions of closed field line loops, as well as regions of open field lines along which outflows (stellar winds) are launched. A basic assumption of the model is that currents can be neglected in the region external to the star, and interior to $R_s$. With this assumption and with the source surface boundary condition, the complicated problem of solving Poisson's equation to determine the magnetic field components is bypassed. Instead, the $\Bv$-field is derived from the general solution of the familiar Laplace's equation. The layout of the paper is as follows. In Sec.~\ref{intro} some recent results from the study of stellar magnetic fields are presented. I define the field components of a multipolar large scale stellar (or equivalently a planetary) magnetosphere in Sec.~\ref{comps}. In Sec.~\ref{comps_part2} I discuss how models of stellar magnetospheres can account for the distortion of the large scale magnetic field caused by stellar outflows (winds) by the use of the source surface boundary condition, and demonstrate in Sec.~\ref{comps_part3} how the field components are modified. The modified field components are derived from the solution of Laplace's equation in spherical coordinates, subject to the assumptions discussed at the beginning of Sec.~\ref{comps_part3}. In Sec.~\ref{equ} I solve the differential equation for the path of the field lines for an arbitrary multipole $l$ with a source surface. I summarize the main results in Sec.~\ref{sum}. Throughout the paper I consider only stellar/planetary magnetic fields that are in a ``potential state.'' This terminology, which is common in the solar/stellar physics literature, refers to a magnetic field in which the current density $\Jv =\mathbf{0}$ everywhere within the stellar magnetosphere, and therefore the field $\Bv$ can be written in terms of the gradient of a magnetostatic scalar potential.

\label{sum} Over the past five years the current generation of optical spectropolarimeters has allowed stellar magnetic field topologies to be probed in unprecedented detail. Large observational programs have provided maps of the magnetic fields of stars spanning a range of masses, rotation rates, and evolutionary stages.\cite{don09} In particular, magnetic maps of forming solar-like stars can now be obtained.\cite{don07} The large scale field topologies vary from simple dipoles to more complex magnetic fields consisting of several multipole components.\cite{gre10} Given the availability of the observational datasets, new models that incorporate magnetic fields are under development such as models that consider how the magnetospheres of forming solar-like stars interact with their surrounding planet-forming disks.\cite{gre08,gre10,rom10} I have discussed how the solution of Laplace's equation using separation of variables can be applied to construct models of stellar magnetospheres where the influence of outflows on the global field structure is considered. With the assumption that currents have a negligible effect on the large scale structure of stellar magnetic fields in the region between the stellar surface and the source surface, the field components have been derived by solving the Laplace boundary value problem, thus by passing the more complicated problem of solving Poisson's equation. Basic properties of multipole magnetic fields have been discussed, and from the field components, Eqs.~(\ref{Brsource}) and (\ref{Bthetasource}), an analytic equation for the field lines of an arbitrary axial multipole with a source surface has been derived. The resulting expression, Eq.~(\ref{final}), is straightforward to solve using a root finding algorithm and its use only requires knowledge of the roots of the $m=1$ $\ell$th associated Legendre function $P_{\ell1}(\cos{\theta})$. I have concentrated on deriving the equation of the field lines for a modified magnetic multipole in spherical coordinates by solving Eq.~(\ref{path}). Spherical coordinates are the most natural coordinate system to use in the description of stars and planets and their immediate environments. Other authors have derived analytic expressions for the field lines of axisymmetric multipoles without the source surface boundary condition in spherical polar coordinates.\cite{wil87,jef88,bac88} The external field line equation for the magnetic multipoles is identical to that obtained for the field lines of point electric linear multipoles. (However, the internal and contact field lines of magnetic and electric multipoles of the same order differ.\cite{gra10}) For example, the path of the electrostatic field lines for a linear quadrupole in spherical coordinates, can be found in Appendix~A of Ref.~\onlinecite{gra09}. For the electrostatic case it is much more common for the path of the field lines to be determined in Cartesian coordinates. Reference~\onlinecite{smy50} provides a Cartesian expression for the field lines of an arbitrary linear electric multipole. Several other overviews describe algorithms for simple computer programs that can be adapted to visualize magnetostatic or electrostatic fields.\cite{kri85,kir85,kir86} The magnetospheric structures considered in this paper have certain limitations. The fields are assumed to be axisymmetric. The assumption of symmetry allows the analytic equation for the field lines to be derived. In reality, stellar magnetic fields can be highly complex, with many high order field components. (See Ref.~\onlinecite{don09} for a review of recent results on stellar magnetic topologies.) A different approach is to model stellar magnetic fields numerically via extrapolation from observationally derived magnetic maps.\cite{jar02,gre06,gre08} Both the numerical approach and the analytic work in this paper assume that the stellar magnetic fields are current free. As discussed in Sec.~\ref{comps_part2}, the global field topologies obtained in potential field source surface models approximately match those obtained from more computationally complex magnetohydrodynamic models.\cite{ril06,rom10} In this paper I have assumed that the $\ell$th order multipole moment symmetry axis is aligned with the stellar rotation axis, and that both lie in the same stellar meridional plane (planes with $\phi={\rm constant}$; for example, the $x$-$z$ plane). However, Eqs~(\ref{final})--(\ref{final_source}) apply generally to tilted multipole symmetry axes (that is, large scale magnetospheres tilted arbitrarily in the polar and azimuthal directions with respect to the stellar rotation axis, assumed to be the $z$ axis) with appropriate coordinate and vector frame transformations. The equation for calculating the co-latitude of the footpoints of the largest closed field line loop within each region of closed field lines can be used as a basis for calculating the amount of unsigned open (or closed) flux relative to the total flux (open plus closed) through the stellar surface. Such detailed calculations, from which models of stellar rotational evolution can be developed (for example, Refs.~\onlinecite{iva03} and \onlinecite{mat09}), are deferred to a future paper. \appendix*

1012.4799

1012

1012.1343_arXiv.txt

Quantum fields in compact stars can be amplified due to a semiclassical instability. This generic feature of scalar fields coupled to curvature may affect the birth and the equilibrium structure of relativistic stars. We point out that the semiclassical instability has a classical counterpart, which occurs exactly in the same region of the parameter space. For negative values of the coupling parameter the instability is equivalent to the well-known ``spontaneous scalarization'' effect: the plausible end-state of the instability is a static, asymptotically flat equilibrium configuration with nonzero expectation value for the quantum fields, which is compatible with experiments in the weak-field regime and energetically favored over stellar solutions in general relativity. For positive values of the coupling parameter the new configurations are energetically disfavored, and the end-point of the instability remains an open and interesting issue. The vacuum instability may provide a natural mechanism to produce spontaneous scalarization, leading to new experimental opportunities to probe the nature of vacuum energy via astrophysical observations of compact stars.

We look for static, spherically symmetric equilibrium solutions of the field equations admitting a nonzero scalar field with metric \be ds^2=-f(r)dt^2+\left(1-2m(r)/r\right)^{-1}dr^2+r^2d\theta^2+r^2\sin^2\theta\,d\phi^2\,.\nonumber \ee Following LMV~\cite{Lima:2010na}, we study a nonminimally coupled scalar field in the presence of a perfect-fluid star in GR. We consider the action \be S=\frac{1}{16\pi G}\, \int d^4x\sqrt{-g} R + \int d^4x\sqrt{-g} \, {\cal L}_m\,, \label{action} \ee where ${\cal L}_m={\cal L}_{\rm scalar}+{\cal L}_{\rm perfect \, fluid}$. The Lagrangian for a scalar field with conformal coupling $\xi$ is given by \be {\cal L}_{\rm scalar}= -\xi\,R\,\Phi^2-g^{\mu\nu}\Phi_{,\mu}\Phi_{,\nu}-\mu^2\Phi^2\,, \ee where $\mu$ denotes the scalar field mass. Here we set $\mu=0$; we will consider the massive case in a follow-up paper. For $\xi =1/6$, $\mu=0$ the action is invariant under conformal transformations ($g_{\mu\nu}\rightarrow \Omega^2g_{\mu\nu}\,,\Phi\rightarrow\Omega^{-1}\Phi$). For $\xi=0$, $\mu=0$ one recovers the usual minimally coupled massless scalar. The above Lagrangian corresponds to a viable theory of gravity, passing all weak-field tests \cite{Bronnikov:1973fh,Capozziello:2005bu}. \begin{figure*}[htb] \begin{center} \begin{tabular}{cc} \epsfig{file=fig1a.eps,width=7.5cm,angle=0}& \epsfig{file=fig1b.eps,width=7.5cm,angle=0} \end{tabular} \caption{Left panel: Existence diagram for constant density stars, showing the values of the ratio $M/R_s$ of uniform density compact objects which support a dynamical scalar field. The dash-dotted line (red in the online version) corresponds to the analytical prediction in the Newton ian limit, Eq.~(\ref{analytics}). The vertical dashed line indicates the conformal-coupling value $\xi=1/6$. We present this plot in ``standard'' geometrical units so that the compactness of our models can be easily compared to the maximum compactness for constant-density stars in GR: $M/R_s\leq 4/9\simeq 0.444$ (Buchdahl's limit). Modulo this trivial unit conversion, the shaded regions in the left panel match {\it exactly} the shaded regions in LMV's Fig.~1 (see main text). In the upper-right region of the diagram, different critical lines correspond to ``excited'' equilibrium configurations and the integers refer to the number $N$ of nodes in $\Phi(r)$ (see the main text). Right panel: Same for stars with a polytropic EOS. The dashed line (orange online) through the negative--$\xi$ region marks the configuration with maximum mass (i.e., the radial stability limit) for each $\xi$.\label{fig:diagram}} \end{center} \end{figure*} The theory can be recast as a scalar-tensor theory of gravity in the Einstein frame~\cite{Damour:1993hw} via the transformation \be g_{\mu\nu}\to(1-8\pi\xi\Phi^2)g_{\mu\nu}\,,\quad d\Phi\to\frac{\sqrt{1-8\pi\xi(1-6\xi)\Phi^2}}{1-8\pi\xi\Phi^2}d\Phi\,.\nonumber \ee In this form, both the classical counterpart of the LMV instability and its final state have been thoroughly studied \cite{Damour:1993hw,Harada:1997mr,Novak:1997hw}. We can explicitly construct static, asymptotically flat, spherically symmetric solutions to the theory~(\ref{action}) with nonzero scalar field and verify that our solutions match those found in \cite{Damour:1993hw,Harada:1997mr,Novak:1997hw}, as follows. The equations of motion following from the Lagrangian are (hereafter we set $c=G=1$) \beq G_{\mu\nu}&=&8\pi\,T_{\mu\nu}\,,\\ \nabla_{\alpha}\nabla^{\alpha} \Phi&=&(\mu^2+\xi\,R)\Phi\,, \eeq where \begin{widetext} \be T_{\mu\nu}=\left(2\Phi_{,\mu}\Phi_{,\nu}-g_{\mu\nu}g^{\alpha\beta}\Phi_{,\alpha}\Phi_{,\beta}-\mu^2g_{\mu\nu}\Phi^2-2\xi\Phi^2_{,\mu;\nu}+2\xi\,g_{\mu\nu}\Phi^2_{,\alpha;\beta}g^{\alpha\beta}+T_{\mu\nu}^{\rm perfect\, fluid}\right)\left(1-16\pi\,\xi\Phi^2\right)^{-1}\,. \ee \end{widetext} We consider perfect-fluid, spherically symmetric stars with energy density $\rho(r)$ and pressure $P(r)$ such that $T^{\mu\nu}_{\rm perfect\,fluid}= \left(\rho+P\right)u^\mu\,u^\nu+g^{\mu\nu}P$, where the fluid four-velocity $u^\mu=(1/\sqrt{f},0,0,0)$. We specify some equation of state (EOS) $P=P(\rho)$ and we impose regularity conditions at the center of the star, i.e. \be m(0)=0\,,\quad \rho(0)=\rho_c\,,\quad \Phi(0)=\Phi_c\,,\quad \Phi'(0)=0 \,. \ee We also require continuity at the stellar radius $R_s$, defined by the condition $P(R_s)=0$. We focus on two different stellar models: (i) the constant density stars ($\rho={\rm const}$) studied by LMV, and (ii) the polytropic model $\rho=n m_b+Kn_0 m_b (\Gamma-1)^{-1}(n/n_0)^\Gamma, P=Kn_0 m_b(n/n_0)^\Gamma$, with $\Gamma=2.34, K=0.0195, m_b=1.66\times 10^{-24}\, {\rm g}$ and $n_0=0.1\,{\rm fm}^{-3}$ (this is the model that was considered in Ref.~\cite{Damour:1993hw} in the context of spontaneous scalarization). We have checked that nuclear-physics based EOS models would yield qualitatively similar results. Relativistic stellar configurations in GR correspond to $\Phi_c=0$, so that $\Phi=0$ everywhere. For each central density $\rho_c$, we used a shooting method to search for nonzero values of $\Phi_c$ such that $\Phi(r)\to 0$ as $r\to \infty$. For constant-density configurations, we find solutions with nonzero scalar field in the shaded regions of the $(\xi\,,M/R_s)$ diagram shown in the left panel of Fig.~\ref{fig:diagram}. The right panel of Fig.~\ref{fig:diagram} shows that the qualitative features of the existence diagram are the same for a polytropic EOS. These diagrams effectively reproduce previous results obtained many years ago in the context of spontaneous scalarization (see e.g.~\cite{Harada:1997mr}). It is remarkable that static solutions exist in the {\it same} region where the LMV instability operates (cf.~Fig.~1 in LMV). This provides strong evidence that these solutions (when they are stable) represent a plausible final state of the instability. In fact, the exact overlap between our own Fig.~\ref{fig:diagram} and Fig.~1 of LMV can be proved analytically. Focus for simplicity on constant-density stars and massless scalar fields (but our reasoning applies in general). The critical lines in Fig.~1 of LMV represent the curves where marginally stable modes exist. These modes are zero-frequency solutions of Eq.~(4) in LMV, where the potential is given by their Eq.~(6). On the other hand, the critical lines in our Fig.~\ref{fig:diagram} represent the boundary of regions where spherically symmetric, static solutions with nontrivial scalar field profiles cease to exist. These are solutions of the Einstein-Klein-Gordon equations with $\Phi=0$. As $\Phi\to 0$ the spacetime becomes arbitrarily close to that of a constant-density star and the Klein-Gordon equation reduces to Eq.~(4) in LMV, with potential given by their Eq.~(6). Furthermore, the same boundary conditions apply in both cases. Thus, the critical lines are obtained from the {\em very same} equations and they are indeed coincident, not just similar. \begin{figure*}[ht] \begin{center} \begin{tabular}{cc} \epsfig{file=fig2a.eps,width=7.55cm,angle=0}& \epsfig{file=fig2b.eps,width=7.3cm,angle=0} \end{tabular} \caption{Left: Gravitational mass as a function of the central baryonic density $\rho_c/\rho_0$, where $\rho_0=8\times 10^{14}$~g/cm$^3$ is a typical central density for neutron stars. The inset shows the (normalized) binding energy as a function of $\rho_c/\rho_0$. Right: Gravitational mass as a function of the radius for different values of the coupling. \label{fig:binding_energy}} \end{center} \end{figure*} Quite interestingly, the LMV instability threshold can be found {\it analytically} in the Newtonian limit $M/R_s\ll 1$ (i.e., in the bottom left corner of Fig.~\ref{fig:diagram}). The instability line defines the existence of static solutions with a small but nonvanishing massless scalar field. The relevant equation in this limit is $\Psi''-8\pi\,\xi( \rho-3P)\Psi=0$, where a prime denotes a derivative with respect to $r$, and we use the ansatz for the scalar field $\Phi=(\Psi/r)e^{-i\omega t}$ (i.e., we consider an s-wave). Assuming that $\rho$ is constant in the stellar interior (this assumption holds exactly for uniform-density stars and is a good approximation for most EOSs), a regular solution at the origin and at infinity that is also continuous at $R_s$ corresponds to \be 24M \xi=-\pi^2\,R_s\,.\label{analytics} \ee As shown in Fig.~\ref{fig:diagram}, this prediction is in very good agreement with the LMV results\footnote{The basic features of the instability in compact stars were understood by Ford (who studied unstable scalar fields as a possible mechanism to damp the effective value of the cosmological constant) as early as 1987 \cite{Ford:1987de}; see also \cite{Harada:1997mr}.}.

We have reconsidered a generic class of theories where a scalar field is nonminimally coupled to the Ricci scalar, that were recently shown to give rise to a semiclassical instability. We have pointed out an interesting relation between the semiclassical instability and the spontaneous scalarization effect in classical scalar-tensor theories. For certain values of the coupling parameter the scalar field can leave observable imprints on the equilibrium properties of relativistic stars. Our main finding is that the LMV instability may provide a ``natural'' seed mechanism to produce spontaneous scalarization, reinforcing the relevance of previous studies of compact stars in scalar-tensor or $f(R)$ theories of gravitation (see e.g.~\cite{Damour:1992we,Damour:1993hw,Babichev:2009fi}). We stress that corrections to GR due to scalar fields are a {\it generic} feature of a large class of unification theories. Our work suggests that strong-field modifications to GR compatible with weak-field tests may be astrophysically viable, with potentially observable consequences in the structure of compact stars. It will be important to explore the implications of vacuum amplification mechanisms for tests of strong-field gravity in compact objects (see e.g. \cite{Psaltis:2008bb} and references therein). {\bf \em Acknowledgements.} We thank Yanbei Chen for useful discussions. V.C. would like to thank Hideo Kodama, Akihiro Ishibashi and all the participants of the ExDip2010 workshop in KEK (Tsukuba, Japan) for useful discussions, and KEK for hospitality while this work was near completion. This work was supported by the {\it DyBHo--256667} ERC Starting and by FCT - Portugal through PTDC projects FIS/098025/2008, FIS/098032/2008 and CTE-AST/098034/2008, CERN projects FP/109306/2009, FP/109290/2009 and by an allocation through the TeraGrid Advanced Support Program under grant PHY-090003. E.B.'s and J.R.'s research was supported by NSF grant PHY-0900735. Computations were performed on the TeraGrid clusters TACC Ranger and NICS Kraken, the Milipeia cluster in Coimbra, Magerit in Madrid and LRZ in Munich.

1012.1343

1012

1012.2893_arXiv.txt

Superradiant instability turns rotating astrophysical black holes into unique probes of light axions. We consider what happens when a light axion is coupled to a strongly coupled hidden gauge sector. In this case superradiance results in an adiabatic increase of a hidden sector CP-violating $\theta$-parameter in a near horizon region. This may trigger a first order phase transition in the gauge sector. As a result a significant fraction of a black hole mass is released as a cloud of hidden mesons and can be later converted into electromagnetic radiation. This results in a violent electromagnetic burst. The characteristic frequency of such bursts may range from $\sim 100$ eV to $\sim 100$ MeV.

String theory landscape points towards ultimate unification of particle physics and geography. If the landscape is real it is very likely that in the short and midterm future particle physicists will be busy by mapping the compactification manifold where we happened to live. In the most optimistic scenario---if extra dimensions are large---this can be achieved directly at the LHC and future colliders. However, even if extra dimensions are small they are still likely to be populated by rich structures accessible for collider, laboratory, astrophysical and cosmological probes at relatively low energies. These structures escaped detection so far as a consequence of their distant geographic location in extra dimensions. The structures which are natural to expect in the landscape on theoretical grounds include hidden gauge sectors (``hidden valleys", \cite{Strassler:2006im}) and a plenitude of axion-like particles, dark photons and photini (the ``axiverse", \cite{Arvanitaki:2009fg}). Confinement scales of strongly coupled gauge sectors and axion masses are exponentially sensitive to the parameters of the compactification manifold and may take arbitrarily low values. For axions there is also an observational indication that they may be light. Indeed, not only the QCD contribution to the mass of the QCD axion is very small ($\sim 10^{-10}$~eV for the GUT scale axion decay constant), but also additional string corrections to the axion potential, expected to be there on general grounds, have to be suppressed at least by extra ten orders of magnitude. Consequently, one may expect the presence of even lighter axions with masses dominated by these corrections. A lot of efforts were concentrated recently on exploring the possibilities to discover these structures in the near future experiments and observations (see, e.g., \cite{Strassler:2006im}--\cite{Cheung:2010mc}). The focus of the current work will be the intriguing possibility to probe light axion particles with masses in the range $10^{-10}\div10^{-20}$~eV with ongoing observations of astrophysical black holes \cite{Arvanitaki:2009fg,Arvanitaki:2010sy}. This possibility is related to the famous Penrose process allowing extraction of kinetic energy from rotating black holes \cite{Penrose:1969pc}--\cite{Starobinskii}. If a boson with a Compton wavelength of order the size of a black hole is present, the Penrose process results in a superradiant instability \cite{Damour:1976kh}--\cite{Detweiler:1980uk}. A black hole releases its rotational energy by creating a cloud of a rotating Bose--Einstein condensate in the near horizon region. Compton wavelengths of axions in the above mass range match sizes of astrophysical black holes and Ref.~\cite{Arvanitaki:2010sy} studied opportunities provided by the superradiant instability for the discovery of these particless. The focus of Ref.~\cite{Arvanitaki:2010sy} was mainly on model-independent purely gravitational signatures of superradiance. Perhaps one of the most interesting conclusions is that Advanced LIGO may be able to discover the QCD axion if the decay constant is high. Instead, here we will consider model-dependent signatures arising when an axion causing the instability is coupled to a QCD-like hidden gauge sector (a hidden valley). The presence of these signatures depends on some of the details of the matter spectrum in the hidden valley. The payoff that we get for introducing a certain amount of model-dependence is that the resulting signals are electromagnetic rather than purely gravitational. Concretely, we are exploiting the following effect. From the point of view of a hidden gauge sector a growing cloud of axions surrounding a rotating black hole acts as a QCD laboratory with a slowly oscillating CP-violating $\theta$-parameter. The amplitude of oscillations grows and reaches order one at the late stages of the instability. As we discuss in section~\ref{phases}, it is a very common feature of gauge theory dynamics that under an adiabatic increase of $\theta$ the vacuum which is stable at $\theta=0$ first turns into an overheated metastable minimum and later on either becomes perturbatively unstable or even totally disappears as an extremum of the potential. Unfortunately, this possibility is not realized in the real-world QCD. However, it may be exhibited in one of hidden valleys. If this happens, the superradiant instability will trigger an avalanche phase transition in the hidden valley and a cloud of strongly interacting mesons will get produced. The total energy of the cloud depends on the axion decay constant, and for high scale axions may constitute $10^{-4}\div 10^{-3}$ of a black hole mass. Unlike axions, whose couplings to the rest of the matter are strongly suppressed by the decay constant, hidden sector mesons may have substantial interactions with the Standard Model fields. The net result is that a significant fraction of a black hole mass can be converted into electromagnetic energy on a characteristic time-scale set by a black hole size. This would give rise to violent transient astronomical sources---electromagnetic echoes of hidden valley avalanches---that can be looked for in a wide range of frequencies. The rest of the paper is organized as follows. In section~\ref{phases} we review how a vacuum structure with a non-trivial dependence on a $\theta$-parameter may arise in a QCD-like theory. As a fully calculable example we use a QCD-like theory with $N$ light quark flavors, where the relevant dynamics can be analyzed with the help of a low energy chiral Lagrangian. In section~\ref{avalanches} we discuss in details how superradiant instability may trigger an avalanche phase transition in a hidden valley. In section~\ref{EM} we estimate characteristic energies, frequencies and time-scales of the corresponding electromagnetic signal as a function of an axion mass. In the concluding section~\ref{conclusions} we present our conclusions and outline directions for the future work.

\label{conclusions} In this paper we outlined a proposal how rotating astrophysical black hole may open a new window into hidden QCD-like gauge sectors coupled to light axions. Thanks to axion superradiance rotating black holes may give rise to violent bursts of electromagnetic radiation with initial frequencies set by the confinement scale in the hidden sector. Clearly, a number of details need to be filled in to extract concrete predictions for the characteristics of these sources, allowing to distinguish them from astrophysical sources of a more conventional origin and to set limits on the existence of these phenomena at different frequencies. First, our proposal relies on a qualitative scenario for the evolution of superradiance instability developed in Ref.~\cite{Arvanitaki:2010sy}. Given that a black hole surrounded by an axionic cloud is a very rich non-linear dynamical system it is highly desirable to confirm this scenario by detailed numerical simulations. However, many features of our proposal can be studied without going into details of superradiance. Namely, from a purely phenomenological point of view what happens is that a black hole creates a fireball of photons and, if the frequencies are high enough, of electron/positron pairs. The initial temperature of the fireball is bounded from above by $\sim (f_a/R_g)^{1/2}$, where $R_g$ is the black hole size and $f_a$ is the axion decay constant. This is a rather common situation in astrophysics, for instance the initial stage of the fireball model for $\gamma$-ray bursts is very similar. There is a number of differences in our case which affect the further evolution of the fireball and should help to distinguish this kind of sources. For instance, according to the results of Ref.~\cite{Arvanitaki:2010sy}, an episode of superradiant instability may happen for a single black hole in a medium with a very small baryon contamination. In this case we expect the resulting burst to have spectrum close to thermal and have a shorter time scale than a typical $\gamma$-ray burst (here we refer to the case, when the frequency scale of the event corresponds to $\gamma$-rays). On the other hand, the electromagnetic fireball suggested here is immersed in a cloud of a dense liquid composed of strongly interacting hidden mesons, which may affect its evolution. When the episode of superradiance happens to a black hole which was just created as a result of supernova explosion, the whole setup becomes really similar to conventional $\gamma$-burst models. Even in this case the details may be different, though. For instance, superradiance is a relatively slow process. Furthermore, it can be shut off in a dense medium, so that it may be necessary for an environment to clean up a bit for the instability to start. As a consequence we expect a significant time delay between a black hole collapse and the creation of a fireball. Finally, supermassive black holes may give rise through superradiance to very powerful electromagnetic bursts at much lower (down to $\sim 100$~eV) frequencies than conventional $\gamma$-ray bursts. To conclude, superradiance opens an intriguing opportunity to test new physics with astrophysical black holes. Given that violent electromagnetic bursts are being observed in nature and the origin of some of them is unknown we believe a dedicated astrophysical study of the process proposed here is worthwhile.

1012.2893

1012

1012.5966_arXiv.txt

{We argue that the $\eta$-problem in supergravity inflation cannot be solved without knowledge of the ground state of hidden sectors that are gravitationally coupled to the inflaton. If the hidden sector breaks supersymmetry independently, its fields cannot be stabilized during cosmological evolution of the inflaton. We show that both the subsequent dynamical mixing between sectors as well as the lightest mass of the hidden sector are set by the scale of supersymmetry breaking in the hidden sector. The true cosmological $\eta$-parameter arises from a linear combination of the lightest mode of the hidden sector with the inflaton. Generically, either the true $\eta$ deviates considerably from the na\"{\i}ve $\eta$ implied by the inflaton sector alone, or one has to consider a multifield model. Only if the lightest mass in the hidden sector is much larger than the inflaton mass and if the inflaton mass is much larger than the scale of hidden sector supersymmetry breaking, is the effect of the hidden sector on the slow-roll dynamics of the inflaton negligible.}

\label{sec:intro} The construction of realistic models of slow-roll inflation in supergravity is a longstanding puzzle. Supersymmetry can alleviate the finetuning necessary to obtain slow-roll inflation --- if one assumes that the inflaton is a modulus of the supersymmetric ground state --- but cannot solve it completely. This is most clearly seen in the supergravity $\eta$-problem: if the inflaton is a lifted modulus, then its mass in the inflationary background is proportional to the supersymmetry breaking scale. Therefore, the slow-roll parameter $\eta \simeq V''/V$ generically equals unity rather than a small number \cite{Copeland:1994vg}. We will show here, however, that the $\eta$-problem is more serious than a simple hierarchy problem. In the conventional mode of study, the inflaton sector is always a subsector of the full supergravity theory presumed to describe our Universe. When the inflationary subsector of the supergravity is studied {\em an sich}, tuning a few parameters of the Lagrangian to order $10^{-2}$ will generically solve the problem. We will clarify that this split of the supergravity sector into an inflationary sector and other hidden sectors implicitly makes the assumption that all the other sectors are in a `supersymmetric' ground state: i.e. if the inflaton sector which must break supersymmetry is decoupled, the ground state of the remaining sectors is supersymmetric. If this is not the case, the effect on the $\eta$-parameter or on the inflationary dynamics in general can be large, even if the sypersymmetry breaking scale in the hidden sector is small. Blind truncation in supergravities to the inflaton sector alone, if one does not know whether other sectors preserve supersymmetry, is therefore an inconsistent approach towards slow-roll supergravity inflation. Coupling the truncated sector back in completely spoils the na\"{\i}ve solution found. This result, together with recent qualitatively similar findings for sequestered supergravities (where only the potential has a two-sector structure) \cite{Berg:2010ha}, provides strong evidence that to find true slow-roll inflation in supergravity one needs to know the global ground state of the system. The one obvious class of models where sector-mixing is not yet considered is the newly discovered manifest embedding of single field inflationary models in supergravity \cite{Kallosh:2010xz}. If these models are also sensitive to hidden sectors, it would arguably certify the necessity of a global analysis for cosmological solutions in supergravity and string theory. We will obtain our results on two-sector supergravities by an explicit calculation. The gravitational coupling between the hidden and the inflaton sectors is universal, which can be described by a simple $F$-term scalar supergravity theory. As in most discussions on inflationary supergravity theories, we will ignore $D$-terms as one expects its VEV to be zero throughout the early Universe \cite{Lyth:2009zz}. Including $D$-terms (which themselves always need to be accompanied by $F$-terms) only complicates the $F$-term analysis, which is where the $\eta$-problem resides. Furthermore, although true inflationary dynamics ought to be described in a fully kinetic description \cite{Achucarro:2010da}, we can already make our point by simply considering the mass eigenmodes of the system. In a strict slow-roll and slow-turn approximation the mass eigenmodes of the system determine the dynamics of the full system. Specifically we shall show the following for two-sector supergravities where the sectors are distinguished by independent R-symmetry invariant K\"{a}hler functions: \begin{itemize} \item Given a na\"{\i}ve supergravity solution to the $\eta$-problem, this solution is only consistent if the other sector is in its supersymmetric ground state. \item If it is not in its ground state, then the scalar fields of that sector cannot be static but \emph{must} evolve cosmologically as well. \item In order for the na\"{\i}ve solution to still control the cosmological evolution these fields must move very slowly. This translates in the requirement that the contribution to the first slow-roll parameter of the hidden sector must be much smaller than the contribution from the na\"{\i}ve inflaton sector, $\eps_{\textrm{hidden}} \ll \eps_{\textrm{na\"{\i}ve}}$. \item There are two ways to ensure that $\eps_{\textrm{hidden}}$ is small: Either the supersymmetry breaking scale in the hidden sector is very small or a particular linear combination of first and second derivatives of the generalized K\"ahler function is small. \begin{itemize} \item In the latter case, one finds that the second slow-roll parameter $\eta_{\textrm{na\"{\i}ve}}$ receives a very large correction $\eta_{\textrm{true}}-\eta_{\textrm{na\"{\i}ve}}\gg \eta_{\textrm{na\"{\i}ve}}$, unless the supersymmetry breaking scale in the hidden sector is small. This returns us to the first case. \item In the first case, one finds that the hidden sector always contains a light mode, because in a supersymmetry breaking (almost) stabilized supergravity sector there is always a mode that scales with the scale of supersymmetry breaking. This light mode will overrule the na\"{\i}ve single field inflationary dynamics. \end{itemize} \end{itemize} Thus for \emph{any} nonzero supersymmetry breaking scale in the hidden sector --- even when this scale is very small --- the true mass eigenmodes of the system are linear combinations of the hidden sector fields and the inflaton sector fields. We compute these eigenmodes. By assumption, the true value of the slow-roll parameter $\eta$ is the smallest of these eigenmodes. Depending on the values of the supersymmetry breaking scale and the na\"{\i}ve lowest mass eigenstate in the hidden sector, we find that \begin{enumerate} \item The new set of mass eigenmodes can have closely spaced eigenvalues, and thus the initial assumption of single field inflation is incorrect. Then a full multifield re-analysis is required. \item The relative change of the value of $\eta$ from the na\"{\i}ve to the true solution can be quantified and shows that for a supersymmetry breaking hidden sector, the na\"{\i}ve model is only reliable if the na\"{\i}ve lowest mass eigenstate in the hidden sector is much larger than the square of the scale of hidden sector supersymmetry breaking divided by the inflaton mass. This effectively excludes all models where the hidden sector has (nearly) massless modes. \item The smallest eigenmode can be dominantly determined by the hidden sector, and thus the initial assumption that the cosmological dynamics is constrained to the inflaton sector is incorrect. Again a full multifield re-analysis is required. \end{enumerate} One concludes that in general one needs to know/assume the ground states and the lowest mass eigenstates of {\em all} the hidden sectors to reliably find a slow-roll inflationary supergravity. The structure of our paper is the following. Section \ref{sec:sugra} reviews some definitions in supergravity and explains how sectors are coupled in supergravity. This leads directly to the first result that in a stabilized supergravity sector there always is a mode that scales with the scale of supersymmetry breaking. In section \ref{sec:inflation} we discuss the $\eta$-problem in a single sector theory and then consider the effect of a hidden sector qualitatively and quantitatively. The quantitative result is analysed in section \ref{sec:diagram} both in terms of effective parameters and direct supergravity parameters. As a notable example of our result, we show that if the hidden sector is the Standard Model, where its supersymmetry breaking is not caused by the inflaton sector but otherwise, spoils the na\"{\i}ve slow-roll solution in the putative inflaton sector. The paper is supplemented with two appendices in which some of the longer formulae are given.

In this paper we have studied the effect of hidden sectors on the finetuning of $F$-term inflation in supergravity, identifying a number of issues in the current methodology of finetuning inflation in supergravity. Finetuning inflationary models is only valid when the neglected physics does not affect this finetuning, in which case the inflationary physics can be studied independently. As shown in figures \ref{fig:etaplus} and \ref{fig:etamin} this assumption holds only under very special circumstances. The reason is that the everpresent gravitational couplings will always lead to a mixing of the hidden sectors with the inflationary sector, even in the case of the most minimally coupled action \eqref{eq:splitaction}. For a hidden sector vacuum that preserves supersymmetry, the sectors decouple consistently \cite{Choi:2004sx,deAlwis:2005tf,deAlwis:2005tg,Achucarro:2007qa,Achucarro:2008sy}. However, for a supersymmetry breaking vacuum the inflationary dynamics is generically altered, where the nature and the size of the change depends on the scale of supersymmetry breaking. For a hidden sector with a low scale of supersymmetry breaking, like the Standard Model, the cross coupling scales with the scale of supersymmetry breaking, and is therefore typically small. Yet, as shown in section \ref{sec:zeromode}, also the lightest mass of the hidden sector scales with the scale of supersymmetry breaking within that sector. This light mode is strongly affected by the inflationary physics and thus evolves during inflation. Therefore, any single field analysis is completely spoiled as discussed in section \ref{sec:inflationSM}. For massive hidden sectors, the problem is more traditional. For a small hidden sector supersymmetry breaking scale, one has a conventional decoupling as long as the lightest mass of the hidden sector is much larger than the inflaton mass. However, for large hidden sector supersymmetry breaking, this intuition fails. Then, the off-diagonal terms in the mass matrix \eqref{eq:massmatrix} will lead to a large correction of the $\eta$-parameter. To conclude, any theory that is working by only tuning the inflaton sector has made severe hidden assumptions about the hidden sector, which typically will not be easily met. Methodologically the only sensible approach is to search for inflation in a full theory, including knowledge of all hidden sectors.

1012.5966

1012

1012.5227_arXiv.txt

Galaxy interactions and mergers play a significant, but still debated and poorly understood role in the star formation history of galaxies. Numerical and theoretical models cannot yet explain the main properties of merger-induced starbursts, including their intensity and their spatial extent. Usually, the mechanism invoked in merger-induced starbursts is a global inflow of gas towards the central kpc, resulting in a nuclear starburst. We show here, using high-resolution AMR simulations and comparing to observations of the gas component in mergers, that the triggering of starbursts also results from increased ISM turbulence and velocity dispersions in interacting systems. This forms cold gas that are denser and more massive than in quiescent disk galaxies. The fraction of dense cold gas largely increases, modifying the global density distribution of these systems, and efficient star formation results. Because the starbursting activity is not just from a global compacting of the gas to higher average surface densities, but also from higher turbulence and fragmentation into massive and dense clouds, merging systems can enter a different regime of star formation compared to quiescent disk galaxies. This is in quantitative agreement with recent observations suggesting that disk galaxies and starbursting systems are not the low-activity end and high-activity end of a single regime, but actually follow different scaling relations for their star formation.

Evidence that interactions and mergers can strongly trigger the star formation activity of galaxies has been observed for more than two decades \citep[e.g.,][]{sanders88}. Nevertheless, both the underlying mechanisms the role of mergers in the cosmic budget of star formation remain largely unknown. Mergers have long appeared to potentially dominate the star formation history of the Universe. Starbursting galaxies with specific star formation rates much above the average, such as Ultraluminous Infrared Galaxies (ULIRGs), are almost exclusively major interactions and mergers -- in the nearby Universe (see \citealt{duc97}), not necessarily at high redshift. The decrease in the cosmic density of star formation between $z=1$ and $z=0$ \citep{lefloch05} may follow, and might even result from, the decrease in the galaxy merger rate. Other studies based on disturbed UV morphologies and/or kinematics also suggested high merger fractions among star forming galaxies at $z\sim1$ or even below \citep{hammer05}. However, disturbed morphologies and kinematics can also arise internally, without interactions, especially at high redshift when disk galaxies are wildly unstable, clumpy and irregular \citep{E07, FS09}. Merger-induced ULIRG-like starbursts may also be less important in the cosmic budget than the more moderate but more numerous objects with internally sustained star-formation, as is possibly the case for most high-redshift LIRGs \citep[e.g., ][]{daddi10a}. Recent studies aimed at accurately distinguishing the signatures of mergers from internal evolution actually suggest that major interactions and mergers account only for a small fraction of the cosmic star formation history \citep{jogee09, robaina09}. Moreover, observational estimates appear to be quite dependent on the chosen tracers of star formation. For instance, mergers are probably a more important trigger of dust-obscured star formation than of general star formation, and thus the fraction of mergers among objects with high infrared-traced star formation rate can be higher \citep[as discussed by][]{robaina09}. Numerical simulations are then required to understand the mechanisms leading to starbursting activity in mergers, but also to further probe the contribution of mergers to cosmic star formation since observational estimates remain uncertain. Section~2 reviews the standard knowledge on merger-induced star formation, which is mostly based on ``sub-grid'' modeling: all steps from the formation of cold/dense gas clouds to the formation of actual stars remain unresolved, and described with arbitrary recipes. This standard understanding cannot account for some general observationed features that we briefly review. Section~3 presents a new generation of high-resolution models in which the first steps of galactic-scale star formation, namely ISM turbulence and cold/dense gas cloud formation, are explicitly resolved using high-resolution codes. Based on this, we provide a substantially different explanation for interaction-triggered star formation and show that it could better account for recent observations of disk galaxies and starbursting mergers.

The results presented here were based mostly on low-redshift merger simulations. A recent set of high-redshift merger simulations with AMR is presented in Bournaud et al. (2011), and shows similar increase in the ISM velocity dispersions and clumpiness in mergers, with extended starbursts and high $\Sigma_{\mathrm SFR}/\Sigma_{\mathrm gas}$ ratios. Galaxy interactions and mergers play a significant but still debated and poorly understood role in the star formation history of galaxies. Numerical and theoretical models have significant difficulties in accounting for the properties of merger-induced starbursts, including their intensity and their spatial extent. Usually, the mechanism invoked to explain the triggering of star formation by mergers is a global inflow of gas towards the central kpc, resulting in a nuclear starburst. We show here, using high-resolution AMR simulations and comparing to observations of the gas component in mergers, that the triggering of starbursts also results from increased ISM turbulence and velocity dispersions in interacting systems. This results in the formation and collapse of dense and massive gas clouds in the regions of convergent flows and local shocks, these clouds being denser, more massive and/or more numerous than in quiescent disk galaxies. The fraction of dense cold gas largely increases, modifying the global density distribution of these systems, and efficient star formation results. Because the starbursting activity is not just from a global compacting of the gas to high average surface densities, but also from higher turbulence and fragmentation into massive and dense clouds, merging systems can enter a different regime of star formation compared to quiescent disk galaxies. This is in quantitative agreement with recent observations suggesting that disk galaxies and starbursting systems are not the low activity end and high activity end of a single regime, but actually follow different scaling relations for their star formation.

1012.5227

1012

1012.3771_arXiv.txt

We present an analysis of the properties of the lowest H$\alpha$-luminosity galaxies ($L_{H\alpha}\leq4\times10^{32}$ W; SFR$<0.02$ M$_\odot$yr$^{-1}$) in the Galaxy And Mass Assembly (GAMA) survey. These galaxies make up the the rise above a Schechter function in the number density of systems seen at the faint end of the H$\alpha$ luminosity function. Above our flux limit we find that these galaxies are principally composed of intrinsically low stellar mass systems (median stellar mass $=2.5\times10^8 M_\odot$) with only 5/90 having stellar masses $M>10^{10} M_\odot$. The low SFR systems are found to exist predominantly in the lowest density environments (median density $\sim0.02$ galaxy Mpc$^{-2}$ with none in environments more dense than $\sim1.5$ galaxy Mpc$^{-2}$). Their current specific star formation rates (SSFR; $-8.5 < $log(SSFR[yr$^{-1}] )<-12.$) are consistent with their having had a variety of star formation histories. The low density environments of these galaxies demonstrates that such low-mass, star-forming systems can only remain as low-mass and forming stars if they reside sufficiently far from other galaxies to avoid being accreted, dispersed through tidal effects or having their gas reservoirs rendered ineffective through external processes.

\label{sect:intro} Recent investigations into the evolutionary properties of galaxies have established that the distribution of galaxy colours exhibits a strong bimodality, separating red, quiescent early-type galaxies, with little ongoing star formation, from a blue sequence of star-forming spiral galaxies (e.g. \citealp{baldry06, blanton06, kauffmann03, driver06}). This bimodality can be interpreted in terms of star formation history and stellar mass content (e.g. \citealt{tinsley68}) with the stars in the most massive galaxies having formed quickly a long time ago and less massive galaxies still forming stars today. The current star formation of a galaxy can be traced through its H$\alpha$ luminosity as this emission gives a direct probe of the young massive stellar population \citep{kennicutt98}. The Galaxy And Mass Assembly survey (GAMA\footnote{http://www.gama-survey.org/}; \citealt{driver09,driver10}), has to date obtained optical spectra for $\sim120,000$ galaxies in the nearby Universe ($z<0.5$). GAMA's deep spectroscopic observations ($r<19.8$; two magnitudes fainter than the Sloan Digital Sky Survey; SDSS; \citealt{york00}) and wide-area sky coverage (three 48 square-degree regions) provide an excellent sample with which to analyse the current star formation rates % of a large sample of galaxies, down to very low stellar masses. Gunawardhana et al. (in prep.) have used early data from GAMA to establish that there is a significant increase in the number densities of H$\alpha$-emitting galaxies at very low luminosities (H$\alpha$ luminosities $<2.5\times10^{32}$ W, $H_0=100$ km s$^{-1}$ Mpc$^{-1}$) such that a Schechter Function is a poor fit at these luminosities. This has also been observed by \cite{westra10} with the Smithsonian Hectospec Lensing Survey. % Characterising the faint-end of the H$\alpha$ luminosity function is important as both \cite{lee07} and \cite{bothwell09} have found that there is a broadening of the star formation distribution to extremely low star formation rates (SFR) in low-mass as well as high-mass galaxies from their analysis of H$\alpha$-derived SFRs for galaxies within 11 Mpc. The broadening of the star formation distribution in high mass galaxies is due to the cessation of star formation. However, the broadening for low-mass galaxies is an open question with two potential resolutions: First, that the cause of this range is due to a change in the dominant physical process that regulates star formation. These galaxies are generally low mass and such galaxies have shallow potential wells. They are therefore subject to many processes that are negligible in higher-mass systems. Second, the possibility that, at low luminosities, the star formation rate is not being traced as closely by the H$\alpha$ luminosity due to the existence of fewer young stars, meaning that the initial mass function (IMF) is not fully sampled in these systems. \cite{meurer09} and \cite{lee09b} find that the H$\alpha$ luminosity underestimates the SFR relative to the FUV luminosity in dwarf galaxies and that this could be the result of a steeper initial mass function (IMF) for galaxies with lower SFRs, for which there is growing theoretical (e.g. \citealt{weidner05}) and observational evidence \citep{hoversten08,gunawardhana10}. In this paper we investigate the characteristics of the low H$\alpha$-luminosity galaxies that comprise the upturn in the H$\alpha$ luminosity function. We investigate their stellar masses to determine whether the low H$\alpha$-luminosities are from low-mass galaxies with high specific star formation rates (SSFRs) or massive galaxies with low SSFRs and find that these are generally low-mass systems with high SSFRs. We examine how much of the increase in number densities seen at low stellar masses ($<3\times10^{8}M_{\odot}$; \citealt{baldry08}) in the galaxy stellar mass function is due to these systems and find that they do not contribute significantly. We also use the wide area of the GAMA survey to investigate what environment these low star forming galaxies are found in and conclude that they are only found in low density environments. We describe the construction of the sample in Section \ref{sect:data}. In Section \ref{sect:properties} we describe the observed properties, star formation rates and environmental density of the sample, and discuss our findings in Section \ref{sect:discussion}. Throughout this paper we assume a Hubble constant of $H_0=70$ km s$^{-1}$ Mpc$^{-1}$ and an $\Omega_M=0.3$, $\Omega_\Lambda=0.7$ cosmology. All magnitudes are given in the AB system. We use a Salpeter IMF in our derivations of stellar masses and SFR for convenience but recognize that quantitative values may need to be scaled to a more realistic IMF.

\label{sect:discussion} We have investigated the properties of a sample of 90 of the faintest-H$\alpha$ luminosity galaxies from the $r-$band--selected GAMA survey. We find these galaxies to generally be low-stellar mass, blue, in low-density environments with a range of star formation histories. % This is the most distant sample of such low-mass, low star formation systems currently known and, with the area of the GAMA survey, the best yet available for probing their environmental dependence. We have shown that the low H$\alpha$ luminosity galaxies have similar properties to dwarf irregular galaxies in the Local Volume. Local Group galaxies are sufficiently nearby that their star formation histories have been studied in detail through resolved stellar population analysis. In this environment dwarf galaxies are the most numerous. The Local Group dwarf galaxies appear to have complex, past star formation histories, such that whatever star formation is currently occurring, there is always evidence for large intermediate age populations and for long episodes of low-to-moderate star formation intensity separated by short passive phases (e.g. \citealt{smecker-hane96,skillman03,tolstoy09}). % We find low-SFR, low-mass systems at higher redshifts than the Local Group/Volume samples yet with similar H$\alpha$ luminosities and wide ranges of star formation histories to the local samples. The continuous distribution of SSFRs we observe suggests that the galaxies cover a range of phases of star formation, some passive, some steady and some bursting. This is consistent with their being higher redshift analogues of local dwarf star-forming systems. As introduced in \S \ref{sect:intro}, understanding the causes of this range in star formation histories is ongoing. \cite{stinson07} simulate the collapse of isolated dwarf galaxies and find that star formation in low-mass galaxies can undergo a `breathing' mode where episodes of star formation trigger gas heating leading to galaxies with episodic star formation and significant intermediate age populations. \cite{quillen08} show that this `breathing' requires strong, delayed feedback in order to reproduce the observationally estimated episode times. Alternatively, we could have a steeper IMF as recently suggested for galaxies with lower SFRs (e.g. \citealt{weidner05,hoversten08,meurer09,lee09b,gunawardhana10}). We find these galaxies to be in low-density environments. That isolation has led to their remaining low mass and star forming, with the lowest mass galaxies generally found in the lowest density environments and those in denser environments tending to have higher stellar masses. These observations lead to the prediction that at high-redshift, when there has been less time for interactions, we would expect to see a higher proportion of low-mass, star-forming systems in dense environments since they will not have had time to be accreted, tidally-stripped or their gas reservoirs rendered ineffective through external processes.

1012.3771

1012

1012.1774_arXiv.txt

The emission from blazars is known to be variable at all wavelengths. The flux variability is often accompanied by spectral changes. Spectral energy distribution (SED) changes must be associated with changes in the spectra of emitting electrons and/or the physical parameters of the jet. Meaningful modeling of blazar broadband spectra is required to understand the extreme conditions within the emission region. Not only is the broadband SED crucial, but also information about its variability is needed to understand how the highest states of emission occur and how they differ from the low states. This may help in discriminating between models. Here we present the results of our SED modeling of the blazar S5 0716+714 during various phases of its activity. The SEDs are classified into different bins depending on the optical brightness state of the source.

\label{sec:intro} S5 0716+714 is a bright, high declination BL Lac object at a redshift, $z = 0.31 \pm$0.08 (Nilsson et al.\ 2008). This source has been extensively studied across the whole electromagnetic spectrum and exhibits strong variability on a wide range of timescales, ranging from minutes to years (e.g., Wagner et al.\ 1990; Heidt \& Wagner 1996; Villata et al.\ 2000; Raiteri et al.\ 2003; Montagni et al.\ 2006; Ostorero et al.\ 2006; Gupta et al.\ 2008a, b, 2009 and references therein). Nearly periodic oscillations of $\sim$15 minutes in the optical R band were detected in this source (Rani et al.\ 2010). The optical duty cycle of S5 0716+714 is nearly unity, indicating that the source is always in an active state in the visible (Wagner \& Witzel 1995). This blazar was recently shown to be a strong source in the high energy gamma-ray band by Fermi-LAT (Abdo et al. 2009).

\subsection{Limitations to the model} Although, as seen in Fig.\ 3, we were able to achieve reasonably good fits of the synchrotron emission for all six SEDs of the source, one should bear in mind that the one-zone BPL model is over-simplified in accounting for the radio-optical blazar emission. The applicability of a single-zone emitting region has been questioned by a number of authors (e.g., Vittorini et al.\ 2009, Raiteri et al.\ 2010). They showed that the BL Lac SED can be more successfully modeled with two synchrotron components (two different emitting populations). Furthermore, unfortunately, we do not have synchrotron peak measurements for the the source, which would significantly help to constrain the model. Another limitation to the model is that we do not correct our optical measurements for any possible host galaxy contribution. Some objects, such as BL Lac itself, have a significant contribution from starlight to the optical bands, which will modify the calculated synchrotron emission in this region, especially during the low states. We also do not make any corrections for Galaxy or internal absorption, which may again slightly affect the optical-UV part of the spectra. Last, but not least, we stress that our data is not strictly simultaneous but has been averaged over a period of months. As blazars are highly variable over timescales of a day or less, the time differences and averaging might have compromised somewhat the modeling. \subsection{Spectral Energy Distribution variation} To attempt a study of how the physical parameters related to emission region and synchrotron emission are changed when the BL Lac S5 0716+714 goes through various phases of its activity we used the simplest approach by fitting a single-zone SSC model. The observed results can be summarized as : \\ 1. No change between the R and B bands in the modeled SED are seen during different phases of activity. \\ 2. The Doppler boosting factor $\delta$ is higher during the optically bright states of the source compared to the dimmer phases of activity. \\ 3. The number density (N) of electrons emitting synchrotron photons is larger when the source is in lower states. \\ 4. The synchrotron peak frequency ($\nu_{peak}$) and peak intensity ($\nu F_{\nu}$$_{peak}$) are comparatively higher during the optical outburst phases of the BL Lac object. \\ This research was supported by CSIR Foreign travel grant {\bf Ref No. TG/5295/1-HRD} and 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. BR is very thankful to Margo Aller for providing the radio frequency data.

1012.1774

1012

1012.1859_arXiv.txt

Primordial non-Gaussianity introduces a scale-dependent variation in the clustering of density peaks corresponding to rare objects. This variation, parametrized by the bias, is investigated on scales where a linear perturbation theory is sufficiently accurate. The bias is obtained directly in real space by comparing the one- and two-point probability distributions of density fluctuations. We show that these distributions can be reconstructed using a bivariate Edgeworth series, presented here up to an arbitrarily high order. The Edgeworth formalism is shown to be well-suited for `local' cubic-order non-Gaussianity parametrized by $\gnl$. We show that a strong scale-dependence in the bias can be produced by $\gnl$ of order $10^5$, consistent with CMB constraints. On correlation length of $\sim100$ Mpc, current constraints on $\gnl$ still allow the bias for the most massive clusters to be enhanced by $20-30\%$ of the Gaussian value. We further examine the bias as a function of mass scale, and also explore the relationship between the clustering and the abundance of massive clusters in the presence of $\gnl$. We explain why the Edgeworth formalism, though technically challenging, is a very powerful technique for constraining high-order non-Gaussianity with large-scale structures.

One of the most intriguing unanswered questions in cosmology is whether or not the primordial seeds that grew into large-scale structures observed today were laid down as a Gaussian random field. In the simplest single-field inflation model of the early Universe, the initial distribution of the primordial seeds, or density fluctuations, is expected to be very close to Gaussian \cite{bartolo,chen}, but deviations from Gaussianity may be large in more complex models involving multiple fields \cite{rigopoulos,byrnes,chen2,langlois,bartolo2,sasaki} or a non-canonical Lagrangian \cite{alishahiha,arkani-hamed,chen3}. Therefore, a detection of a significant level of primordial non-Gaussianity is of great importance as it would effectively rule out a large class of single-field inflation and open an observational window to the early Universe. The observational signatures of primordial non-Gaussianity manifest across a large range of physical scales. On very large scales of order several gigaparsecs, non-Gaussianity can be detected, for instance, in the 3-point correlation function (bispectrum) of the cosmic microwave background (CMB) anisotropies (see \cite{bartoloreview,komatsu2} for recent reviews). In the simplest setting in which the bispectrum is parametrized by the constant $\fnl$, the prospect of constraining non-Gaussianity with the CMB seems very promising indeed. The Planck satellite\footnote{\texttt{http://planck.cf.ac.uk}} will most likely tighten the constraint on $\fnl$ to $\mc{O}$(a few). On smaller scales, the distribution of galaxy clusters can provide competitive constraints on non-Gaussianity, which changes the abundances and clustering properties of large-scale structures (see \cite{desjacquesreview,verde} and references therein). A particularly interesting large-scale-structure probe of non-Gaussianity was presented in the seminal work of Dalal \etal \cite{dalal}, who showed quantitatively that non-Gaussianity induces characteristic changes the clustering of density peaks corresponding to rare objects. Specifically, for a correlation length $r$, we can write \be \xi\sub{pk}(r) = b_L^2(r)\ff \xi(r),\ee where $\xi\sub{pk}$ denotes the correlation function of density peaks, $\xi$ is that of the underlying dark-matter distribution and $b_L$ is the \ii{bias} parameter (these parameters will be explained in detail later). Physically, the bias quantifies how the density peaks traces of the underlying matter distribution. If the density fluctuations are Gaussian distributed, it can be shown that the bias is almost constant (\ii{i.e.} scale-independent) to a good approximation \cite{kaiser}. The scale-\ii{dependence} of the bias induced by non-Gaussianity is the focus of this work. Scale-dependent bias from non-Gaussianity is a relatively young but rapidly developing topic. Whilst the dependence of the bias on $\fnl$ was investigated in \cite{dalal}, a number of authors have since examined the bias for higher-order non-Gaussianity \cite{desjacques}, non-local models \cite{schmidt} and, more recently, scale-dependent $\fnl$ \cite{shandera2} amongst others. The focus of previous works in this area has been the calculation of the bias in Fourier space whilst relying on either numerical simulations or some well-known mass functions. In this work, we show that it is possible to calculate the bias directly in real space by comparing the one- and two-point probability distribution functions (pdfs). We propose to reconstruct the pdfs by using the Edgeworth series in one and two variables (see \cite{blinnikov,kotz} for reviews). The Edgeworth formalism is a mathematically powerful way to capture the statistical essence of non-Gaussian distributions. In previous astrophysical applications, the Edgeworth series were invariably heavily truncated \cite{scherrer,bernardeau,juszkiewicz,amendola,loverde} yielding pdfs that may not be well-defined, non-negative distributions. In this work, we give a general algorithm which allows the Edgeworth series to be kept to arbitrarily high order. We shall see later that given a limited amount of statistical information on the density fluctuations, the Edgeworth formalism is particularly well suited for the reconstruction of non-Gaussian distributions in which the cubic-order non-Gaussianity parameter, $\gnl$, is non-zero. This parameter will be the main focus of our calculations. Once well-defined pdfs are reconstructed, the information on the non-Gaussian bias can then be easily extracted from the one- and two-dimensional pdfs.

\lab{conc} In this work, we have demonstrated an alternative method of calculating the bias in the clustering of rare objects in the presence of primordial non-Gaussianity. Our method is based on the reconstruction of the pdf of density fluctuations using the Edgeworth series in one and two variables. The bias obtained in this way is in real space, in contrast with previous works that examined the scale-dependence bias in Fourier space. A step-by-step guide to our method is presented in Section \ref{sectbias}. Some of the expressions involved (\eg \re{biedgeworth}) may seem complicated, but this is because they incorporate information on arbitrarily high-order correlations. As long as estimates on these high-order correlations are available, our formalism can, in principle, be used to study the observable signatures of high-order non-Gaussianity. In addition, the reconstruction algorithm is independent of the form of non-Gaussianity, hence making our method easily applicable to non-local forms of non-Gaussianity as well. The Edgeworth formalism is a powerful technique that captures all the statistical information of a probability distribution. However, previous astrophysical applications generally dealt with the lowest-order expansions, and therefore the reconstructed pdfs were often found not to be positive definite (in fact, at the lowest order the univariate pdf can never be positive definite). Results obtained from working with pdfs that are not positive definite are unreliable, especially in the context of large-scale structures which are particularly sensitive to the tail end of the pdf. In this work, we concentrate on the case of non-Gaussianity parametrized by positive $\gnl$, which yields pdfs (both uni- and bivariate) that \ii{are} positive definite. It may be surprising to some that the Edgeworth formalism is more easily applied to the case with $\gnl\neq0$ rather than the case with purely $\fnl$-type non-Gaussianity. The reason is that at leading order, $\fnl$ corresponds to the skewness of the distribution. As shown in our previous work \cite{me8}, this information alone cannot define a non-negative pdf. Our previous work also showed that the Edgeworth formalism for the case of pure $\fnl$ requires the knowledge of moments of order at least 5, for which there exist some observational constraints \cite{croton,ross}. The results for $\fnl$ are expected to be similar to that of $\gnl$. This degeneracy can, in theory, be broken by comparing the statistics of voids with that of massive clusters, as any asymmetry in the pdf must be due to the presence of odd-order cumulants. In practice, however, there is the obvious difficulty of determining the abundance and clustering properties of voids. See \cite{damico,me8,lamsheth} for recent progress. Our main results show that $\gnl$-type non-Gaussianity can significantly affect the clustering of massive clusters on large correlation scales ($\sim100$ Mpc, typical of inter-cluster distances). A strong scale dependence of the bias can be seen in Figure \re{figbias}, which summarises our main results for $\gnl$ up to $10^6$. It appears that current constraints on $\gnl$ still allow the bias for the most massive clusters to be enhanced by $20-30\%$ of the Gaussian value. Our findings are relevant to observations and $N$-body simulations in which the clustering of extremely massive objects are seen \cite{tian}. An interesting extension of this work is, therefore, a pdf reconstruction using moments observed in large surveys and simulations. It would then be important to include finite-volume effects \cite{kim,bernardeauuzan} which have been shown to systematically alter the cumulants and hence introduce spurious non-Gaussian effects. By using high-order moments and including finite-volume corrections, we expect to be able to extend the Edgeworth formalism to probe a much wider range of high-order non-Gaussianity. This is the subject of our future work. \mmm \sss \bbb \centerline{\bb{Acknowledgment}} \mmm SC is grateful to the referee for many insightful comments. SC supported by Lincoln College, Oxford. The code (in C++) for generating the bivariate Edgeworth expansion is available upon request. \bbb \liner \appendix

1012.1859

1012

1012.2966_arXiv.txt

We present the non-linear theory of shock acceleration applied to SNRs expanding into partially neutral plasma. Using this theory we show how the Balmer lines detected from young SNRs can be used to test the efficiency of shocks in the production of cosmic rays. In particular we investigate the effect of charge-exchange between protons and neutral hydrogen occurring in the precursor formed ahead of the shock. In this precursor the CR pressure accelerate the ionized component of the plasma and a relative velocity between protons and neutral hydrogen is established. On the other hand the charge-exchange process tends to equilibrate ions and neutrals resulting in the heating of both components. We show that even when the shock converts only a few per cent of the total bulk kinetic energy into CRs, the heating is efficient enough to produce a detectable broadening of the narrow Balmer lines emitted by the neutral hydrogen.

\label{sec:intro} Supernova remnants (SNRs) are thought to be the primary sources of Galactic cosmic rays (CRs). Diffusive shock acceleration (DSA) occurring around the forward shock is considered the most promising mechanism for the acceleration of particles up to very high energies. In order to explain the flux of CRs observed at the Earth, the shock should convert a fraction $\sim 10-30\%$ of the plasm kinetic energy into non thermal particles. Indeed, from the theoretical point of view, DSA can be very efficient in producing accelerated particles, but a clear confirmation of this prediction still lacks, even if many circumstantial evidences have been collected (see e.g. Ref. \refcite{morlino09}). A possible way to investigate the efficiency of CRs production is through the analysis of the Balmer lines associated with shocks of SNRs which propagate in a partially ionized plasma. The hydrogen lines of a Balmer shock consist of two superimposed components: the narrow component is emitted by neutral hydrogen after entering the shock front and the broad component by hot protons after undergoing charge exchange with incoming neutral hydrogen atoms. The width of the broad component reflects the proton temperature behind the shock front, while the width of the narrow component provides a direct measurement of the temperature upstream of the shock \cite{Chevalier80, Heng09}. We stress that, when the shock is not modified by the presence of CRs, the charge exchange process and the Balmer emission both occur only downstream of the shock. On the other hand, when CRs are accelerated efficiently, the shock structure is modified and a precursor is generated upstream: the CRs pressure slows down and compresses the ionized plasma. Even if the neutral component is not directly affected by the CR pressure, a relative velocity between ions and neutrals is now established. Hence charge exchange can also occur upstream of the shock, resulting in the pre-heating of neutral hydrogen. We do expect two remarkable consequences: 1) the narrow Balmer lines can be emitted also from the upstream region and 2) the typical width of these lines becomes larger with respect to the case without CRs, because the hydrogen temperature increases. A further consequence of efficient acceleration is that the temperature of the shocked ions downstream is lower with respect to the case with no acceleration. This occurs because a non negligible fraction of the shock kinetic energy is converted into CRs, rather than into thermal energy. As a consequence also the width of the broad Balmer lines is affected, being reduced with respect to the case with no acceleration. Remarkably all these aspects have been observed in some Balmer-dominated shocks:1) from a region of the Tycho remnant, the emission of narrow Balmer lines has been detected upstream of the shock \cite{Lee10}; 2) in several SNRs the narrow Balmer lines present a width incompatible with the typical temperature of the interstellar medium \cite{Sollerman03}; 3) in two different cases where the Balmer emission has been detected, i.e. RCW 86 \cite{Helder09} and SNR 0509-67.5 \cite{Helder10}, the width of the broad lines led to a downstream temperature which is lower than that estimated form the measurement of the shock proper motion, suggesting that a fraction of the total energy is converted into CRs. In order to quantify all these effects simultaneously, in this paper we present the solution for the stationary case of DSA when the shock propagates into a partially ionized plasma. We use a semi-analytical method, similar to those used in Ref.~\refcite{Amato06}, specialized for plane shock geometry.

\label{sec:disc} We consider a typical case of shock with the following parameters: shock speed $u_0=2000$ km/s, total upstream density $\rho_0= 1 \rm cm^{-3}$, neutral fraction 50\%, upstream temperature $T_0=10^4$ K, upstream magnetic field strength $B_1= 10 \mu$G and maximum momentum of accelerated protons fixed to $p_{\max}=10^4 m_pc$. The parameter which regulates the injection efficiency is taken $\xi_{\rm inj}= 3.95$ \cite{Amato06}. For these values we get $L_{prec}/L_{ce}\simeq 150$, hence we do expect an efficient role of charge exchange, while $L_{prec}/L_{ion}\simeq 10^{-4}$ hence the ionization can be neglected as stated above. In Fig.~\ref{fig:1} we show the velocity and the temperature profiles in the precursor for both protons and hydrogen. The dotted line shows the normalized CR pressure, which at the sub-shock position reaches the value $P_c/\rho_0 u_0^2=0.14$, hence we are dealing with a mildly efficient shock. The upstream temperatures of protons and hydrogen reach the values of $5 \cdot 10^5$ K and $4.7 \cdot 10^5$ K, respectively. Protons are slightly wormer than hydrogen because they also suffer the adiabatic compression due to the CR pressure. The narrow Balmer line width corresponding to the hydrogen temperature at ths subshock is 46 km/s, to be compared with 21 km/s resulting in the absence af a CR precursor. The downstream temperature of shocked ions is $6.7 \cdot 10^7$ K, which produces a broad Balmer line with a width of 1774 km/s; in the absence of acceleration the same width would be 2039 km/s. As already pointed out, this difference arises because when the acceleration is efficient a non negligible fraction of the total shock kinetic energy is channelled into CRs rather than into thermal energy. \begin{figure} \begin{center} \psfig{file=profile.eps,width=3.2in} \caption{Profiles of velocities and temperatures of hydrogen and protons in the precursor upstream of the shock. The thick-dotted line shows the CR pressure. All quantities are normalized as shown in the legend. The typical precursor length is $x_{\max}= D(p_{\max})/u_0$.} \label{fig:1} \end{center} \end{figure}

1012.2966

1012

1012.2157_arXiv.txt

} The introduction of any general talk about exoplanets always includes an account of the number of planets found. The number of planets discovered as of mid-September 2010 was 490\footnote{http://exoplanet.eu}. All exoplanets detected thus far, and recognized as such by the IAU, have been discovered within the past 18 years via a series of observational techniques, which have been quickly improving over time. That improvement is reflected in the histogram in Figure 1, where one can not only see how the number of exoplanets discovered per year has been steadily increasing, but also the number of techniques with which exoplanets are being discovered. \begin{figure} \center \includegraphics[width=10cm,angle=0,clip=true]{lopez_moralesmF1.jpg} \caption{\label{fig1} Histogram representing the number of exoplanets discovered per year since 1992 (last updated on September 2010; http://exoplanet.edu). On the top of the diagram are highlighted in chronological order the number of successful planet detection techniques. Each vertical dashed line indicates the year of the first discovery by each new technique, e.g. the first planet via the radial velocity technique was detected in 1995, and so on.} \end{figure} The first four exoplanets acknowledged as such were discovered between 1992 and 1994 via periodic timing variations in the very precise signal of pulsars \cite{Wolszczan1992,Wolszczan1994}. Then, in 1995, the first planet about a star like the Sun was discovered \cite{Mayor1995}. This planet, similar in mass to our Jupiter and orbiting its host star with a period of about four days, was discovered via the radial velocity technique. This technique continues to be the most successful, scoring close to 80$\%$ of the exoplanets discovered to date. Three other techniques have accomplished the detection of planets since then. In 1999 the discovery of the first planet via the transits technique was announced \cite{Charbonneau2000,Henry2000}. Then, in 2004, the micro-lensing technique produced its first detection \cite{Bond2004}. Finally, just two years ago, the first believed to be {\it bona-fide} imaged planets were detected via direct imaging \cite{Kalas2008,Marois2008}. With close to 500 planets discovered so far, just finding another planet is no longer news, and we are witnessing a split of the field of exoplanets into two main directions: on one hand work continues on trying to detect new planets, but now focusing on finding the first Earth analogs. On the other hand, many efforts are being put into trying to unveil the physical properties of those planets, specially the physical characteristics of their atmospheres. This is how a new field, which some call {\it Exoplanetology}, is now being born.

The detection and characterization of exoplanet amospheres is a brand new field which has started to bloom just in the past 3--4 years, thanks first to the advent of the SST, and most recently to the quick start off of successful detections with telescopes on the ground. Basically all atmospheres detected thus far are of hot Jupiters, but these detections are just the tip of the iceberg for what is yet to come. In the next few years, with more advanced instruments coming online, and the rapid improvement of analysis techniques, the characterization of exoplanet atmospheres will soon start to advance towards the lower mass planets regime, with the final goal in mind being the detection and characterization of Earth-like atmospheres. \small %

1012.2157

1012

1012.2820_arXiv.txt

{ 3C~345 is one of the archetypical active galactic nuclei, showing structural and flux variability on parsec scales near a compact unresolved radio core. During the last 2 years, the source has been undergoing a period of high activity visible in the broad spectral range, from radio through high-energy bands. We have been monitoring parsec-scale radio emission in 3C 345 during this period at monthly intervals, using the VLBA at 15, 24, and 43~GHz. Our radio observations are compared with gamma-ray emission detected by Fermi-LAT in the region including 3C~345 (1FGL J1642.5+3947). Three distinct gamma-ray events observed in this region are associated with the propagation of relativistic plasma condensations inside the radio jet of 3C~345. We report on evidence for the gamma-rays to be produced in a region of the jet of up to 40 pc (de-projected) in extent. This suggests the synchrotron self-Compton process as the most likely mechanism for production of gamma-rays in the source. }

The quasar \object{3C~345} is one of the best studied ``superluminal'' radio sources, with its parsec-scale radio emission monitored over the past 30 years. Substantial variability of the optical (\cite{1984Ap.....20..461B,1990A&A...239L...9K}) and radio (\cite{1996ASPC..110..208A, 1998A&AS..132..305T, 1999ApJ...521..509L}) emission has been observed, with a possible periodicity of 3.5--4.5 years and major flares occurring every 8--10 years. A new cycle of such enhanced nuclear activity began in early 2008 (\cite{2009ATel.2222....1L}). 3C~345 has also been known as a prominent variable source at high energies up to the X-ray band and only in the $\gamma$-ray regime had it not been clearly detected (\cite{2008A&A...489..849C}), possibly due to the low spatial resolution of previous $\gamma$-ray instruments and a lack of space instruments during periods of high source activity. The launch of the GLAST satellite (now \textit{Fermi}) equipped with the Large Area Telescope (LAT) survey instrument (\cite{2009ApJ...697.1071A}) enabled continuous monitoring of $\gamma$-ray emission originating from the vicinity of 3C~345. This paper presents first results from an analysis of \fermilat\ $\gamma$-ray monitoring data, combined with monthly radio observations made at 43.2~GHz (7~mm wavelength) with very long baseline interferometry (VLBI), using the VLBA\footnote{Very Long Baseline Array of the National Radio Astronomy Observatory, Socorro, USA} facility. Throughout this paper, a flat $\Lambda$CDM cosmology is assumed, with $H_0 = 71$\,km\,s$^{-1}$\,Mpc$^{-1}$ and $\Omega_\mathrm{M}$ = 0.27. At the redshift $z=0.593$ (\cite{2007AJ....134..102S}) of 3C~345, this corresponds to a luminosity distance $D_\mathrm{L} = 3.47$\,Gpc, a linear scale of 6.64 pc mas$^{-1}$, and a proper motion scale of 1$\,$mas$\,$year$^{-1}$ corresponding to 34.5$\, $c.

We have found a correspondence between the long-term \textit{Fermi}-LAT lightcurve of the region 1FGL~J1642.5$+$3947 and the radio emission of 3C~345, establishing the detection of $\gamma$-ray emission from 3C~345 by comparison of radio and $\gamma$-ray variability. More importantly, we find $\gamma$-ray emission to be related to the pc-scale jet of the source of up to 2~pc ($\sim$~40~pc de-projected adopting a viewing angle of $\theta \sim $~2.7$^\circ$; \cite{2005AJ....130.1418J}). We have been able to trace back the ejection of new superluminally moving and apparently accelerating features in the jet to be linked to $\gamma$-ray production. Even though a correspondence between the jet and $\gamma$-ray emission was found, there is no evidence for correlation between $\gamma$-ray emission and radio core flux density at 43~GHz. The observed radio properties of these features together with the observed $\gamma$-ray variability question existing jet models and suggest the synchrotron self-Compton (SSC) process as the most likely mechanism driving the production of $\gamma$-ray photons in the source. This conclusion is supported by the spectral energy distribution of 3C~345 being compatible with SSC exhibiting an IC to synchrotron peak ratio of only $\sim$5 (for the Oct. 2009 flare). Continued monitoring and more densely sampled VLBI observations could provide better confirmation of our results and provide an opportunity to localize more accurately the sites of individual flares in the jet. At this writing, the nuclear region of 3C~345 remains at high flux density and $\gamma$-ray flux levels are still elevated. The characteristics of the long-term activity of the source as observed over the last 30 years suggest that this continued activity may last for at least another year, giving a unique opportunity to trace this active state.

1012.2820

1012

1012.1362_arXiv.txt

{Using data from our Parkes \& ATCA \HI\ survey of six groups analogous to the Local Group, we find that the \HI\ mass function and velocity distribution function for loose groups are the same as those for the Local Group. Both mass functions confirm that the ``missing satellite" problem exists in other galaxy groups.}

1012.1362

1012

1012.1154_arXiv.txt

In the third part of the series presenting the Optical Gravitational Lensing Experiment (OGLE) microlensing studies of the dark matter halo compact objects (MACHOs) we describe results of the OGLE-III monitoring of the Large Magellanic Cloud (LMC). This unprecedented data set contains almost continuous photometric coverage over 8 years of about 35 million objects spread over 40 square degrees. We report a detection of two candidate microlensing events found with the automated pipeline and an additional two, less probable, candidate events found manually. The optical depth derived for the two main candidates was calculated following a detailed blending examination and detection efficiency determination and was found to be $\tau=(0.16 \pm 0.12) \times 10^{-7}$. If the microlensing signal we observe originates from MACHOs it means their masses are around $0.2~\msun$ and they compose only $f=3\pm2$ per cent of the mass of the Galactic Halo. However, the more likely explanation of our detections does not involve dark matter compact objects at all and rely on natural effect of self-lensing of LMC stars by LMC lenses. In such a scenario we can almost completely rule out MACHOs in the sub-solar mass range with an upper limit at $f<7$ per cent reaching its minimum of $f<4$ per cent at $M=0.1~\msun$. For masses around $M=10\msun$ the constraints on the MACHOs are more lenient with $f\sim20$ per cent. Owing to limitations of the survey there is no reasonable limit found for heavier masses, leaving only a tiny window of mass spectrum still available for dark matter compact objects.

The Milky Way's halo is probably one of the least known parts of our Galaxy. Numerous recent detections of tidal debris leftover after close encounters between smaller dwarf galaxies and our giant spiral confirm the presence of cold dark matter (CDM) substructures in the halo (\eg \citealt{Belokurov2006Streams}). However, the following question still remains unanswered: what actually is the dark matter (DM) or is not. Compact dark matter objects (MACHOs) would have been the most convenient explanation and in the last two decades this theory has been tested using various methods sensitive to different ranges of masses. In the high mass regime ($M>30~\msun$) wide halo binary objects were studied, but no signature of disturbance due to MACHOs was detected (see \citealt{Yoo2004MACHO} and \citealt{Quinn2009}). For detecting stellar-mass compact DM objects the technique of gravitational microlensing was suggested by \citet{Paczynski1986}. It employs a unique feature of gravitational lensing, namely its sensitivity to unseen objects when they bend and amplify the light of a distant source. The idea was simple: let us observe some distant background rich in stars (\eg Magellanic Clouds) and wait for their temporal brightening due to passage of a massive object located along the line-of-sight between us and the source. Several observing campaigns started after Paczynski's proposal: MACHO \citep{MACHO}, OGLE \citep{Udalski1993}, EROS \citep{EROS} and MOA \citep{MOA}, which for many years observed the Large and Small Magellanic Clouds (LMC and SMC, respectively). The MACHO collaboration was first to publish their results and claimed 20 per cent of mass of the galactic halo is composed of MACHOs with an average mass of 0.4 $\msun$ (\citealt{AlcockMACHOLMC}, \citealt{BennettMACHOLMC}). Based on detection of 10 candidates for microlensing events found in the central 15 sq. deg of the LMC observed over 5.7 years they derived the optical depth towards the LMC of $\tau_{\rm LMC} = (1.0 \pm 0.3) \times 10^{-7}$ \citep{BennettMACHOLMC}. In 2007 the EROS group analysed their data and published their conclusions, which contradicted the MACHO result quite severely. In their data comprising of 6.7 years of continuous observations of 84 deg$^2$ they found no candidates for microlensing events among their bright sample of stars \citep{TisserandEROSLMC}. This led them to the upper limit for the optical depth of $\tau_{\rm LMC} < 0.36 \times 10^{-7}$, which translated to $f = M_{\rm Macho}$/$M_{\rm halo} < 8$ per cent for MACHOs with masses $0.4~\msun$. The OGLE project monitored the LMC during its second (1996-2000) and third (2001-2009) phases (hereafter OGLE--II and OGLE--III, respectively), and the observations are still carried on in the current OGLE--IV phase. Our study of the OGLE--II data (\citealt{Wyrzykowski2009}, herafter Paper I) led to the detection of two candidates for microlensing events, OGLE-LMC-01 and OGLE-LMC-02, and $\tau_{\rm LMC} = 0.43\pm0.33\times 10^{-7}$, which is closer to the limit of EROS than the value of MACHO. Moreover, all detected microlensing signal can be attributed to self-lensing (i.e., when foreground LMC stars microlens backgound LMC stars), leaving no room for lensing due to DM halo objects in the sub-solar and solar mass range. However, the OGLE--II phase lasted only 4 years and covered only parts of the central bar of the LMC (about 4.7 deg$^2$), therefore its result is naturally limited and is subject to some uncertainty. Due to small number statistics, the upper limit on the MACHO presence in the Milky Way's halo was estimated only at $f<20$ per cent. In this paper we present yet another voice in this turbulent story of microlensing towards the Magellanic Clouds and report the results of the search for microlensing events in the OGLE--III data gathered towards the LMC. OGLE--III overwhelms most previous studies in terms of its duration (8 years) and coverage (40 deg$^2$) but more importantly, it uniformly covers the entire LMC bar region and much of LMC outskirts. The paper is organised as follows. In Section \ref{sec:data}, the OGLE--III LMC dataset is described. In the following section, the search algorithm is presented. Section \ref{sec:results} contains the results of the search and detailed description of all detected candidates for microlensing events. Next, the blending study is described and the optical depth is derived. The paper finishes with discussion of the results and conclusions. \begin{table*} \centering \caption{OGLE--III LMC fields.} \label{tab:fields} \begin{tiny} \begin{tabular}{cccrrc|cccrrc} \hline \noalign{\vskip5pt} Field & $\alpha_{J2000}$ & $\delta_{J2000}$ & \multicolumn{2}{c}{$N_{*\mathrm{good}}$[$10^3$]} & $\langle \log{N_{\mathrm{all_{CCD}}}}\rangle$ & Field & $\alpha_{J2000}$ & $\delta_{J2000}$ & \multicolumn{2}{c}{$N_{*\mathrm{good}}$[$10^3$]} & $\langle \log{N_{\mathrm{all_{CCD}}}}\rangle$ \\ & & & tpl & real & & & & & tpl & real & \\ \noalign{\vskip5pt} \hline \noalign{\vskip5pt} LMC100 & 5:19:02.2 & -69:15:07 & 712.2 & 973.9 & 5.13 & LMC158 & 4:30:59.9 & -70:26:01 & 20.0 & 20.8 & 3.63 \\ LMC101 & 5:19:03.1 & -68:39:19 & 279.4 & 323.2 & 4.84 & LMC159 & 5:25:11.4 & -68:03:58 & 129.1 & 137.6 & 4.47 \\ LMC102 & 5:19:03.4 & -68:03:48 & 125.9 & 134.7 & 4.52 & LMC160 & 5:25:20.9 & -68:39:24 & 264.4 & 296.0 & 4.74 \\ LMC103 & 5:19:02.9 & -69:50:26 & 540.6 & 687.4 & 5.02 & LMC161 & 5:25:32.5 & -69:14:59 & 453.6 & 561.8 & 4.99 \\ LMC104 & 5:19:02.4 & -70:26:03 & 272.1 & 313.2 & 4.84 & LMC162 & 5:25:43.3 & -69:50:24 & 817.7 & 1159.9 & 5.17 \\ LMC105 & 5:19:01.6 & -71:01:31 & 205.0 & 224.2 & 4.65 & LMC163 & 5:25:52.2 & -70:25:55 & 524.6 & 669.5 & 5.00 \\ LMC106 & 5:19:01.0 & -71:36:57 & 135.5 & 144.6 & 4.47 & LMC164 & 5:26:08.4 & -71:01:23 & 197.5 & 217.1 & 4.68 \\ LMC107 & 5:13:01.5 & -66:52:57 & 93.2 & 99.3 & 4.45 & LMC165 & 5:26:20.9 & -71:37:01 & 140.6 & 152.3 & 4.61 \\ LMC108 & 5:13:01.9 & -67:28:40 & 105.8 & 113.2 & 4.51 & LMC166 & 5:31:20.1 & -68:03:51 & 127.8 & 138.5 & 4.59 \\ LMC109 & 5:12:53.3 & -68:04:06 & 138.8 & 149.4 & 4.57 & LMC167 & 5:31:39.6 & -68:39:32 & 180.9 & 200.1 & 4.69 \\ LMC110 & 5:12:43.6 & -68:39:42 & 293.4 & 335.6 & 4.82 & LMC168 & 5:32:01.4 & -69:15:00 & 316.2 & 374.0 & 4.90 \\ LMC111 & 5:12:32.7 & -69:15:02 & 441.2 & 526.7 & 4.94 & LMC169 & 5:32:22.8 & -69:50:26 & 757.6 & 1054.8 & 5.14 \\ LMC112 & 5:12:21.5 & -69:50:21 & 481.7 & 594.0 & 4.97 & LMC170 & 5:32:48.1 & -70:25:53 & 588.6 & 769.3 & 5.04 \\ LMC113 & 5:12:10.9 & -70:25:48 & 289.7 & 334.2 & 4.85 & LMC171 & 5:33:10.6 & -71:01:30 & 235.8 & 268.3 & 4.81 \\ LMC114 & 5:11:58.9 & -71:01:22 & 109.4 & 116.2 & 4.44 & LMC172 & 5:33:34.4 & -71:36:54 & 185.7 & 205.9 & 4.72 \\ LMC115 & 5:07:09.7 & -66:52:59 & 133.8 & 143.0 & 4.47 & LMC173 & 5:37:29.3 & -68:03:50 & 126.9 & 135.2 & 4.44 \\ LMC116 & 5:07:00.9 & -67:28:29 & 117.3 & 124.4 & 4.41 & LMC174 & 5:37:59.8 & -68:39:26 & 155.5 & 169.2 & 4.63 \\ LMC117 & 5:06:55.3 & -68:03:58 & 260.4 & 298.2 & 4.82 & LMC175 & 5:38:32.3 & -69:15:01 & 268.7 & 305.3 & 4.79 \\ LMC118 & 5:06:25.4 & -68:39:25 & 383.6 & 463.1 & 4.94 & LMC176 & 5:39:01.6 & -69:50:30 & 357.0 & 414.9 & 4.86 \\ LMC119 & 5:06:02.5 & -69:15:02 & 578.7 & 723.4 & 5.01 & LMC177 & 5:39:38.0 & -70:25:49 & 459.3 & 577.8 & 5.01 \\ LMC120 & 5:05:39.8 & -69:50:28 & 339.4 & 399.4 & 4.89 & LMC178 & 5:40:14.1 & -71:01:27 & 260.6 & 292.4 & 4.77 \\ LMC121 & 5:05:14.4 & -70:25:59 & 200.6 & 223.5 & 4.74 & LMC179 & 5:40:52.3 & -71:36:58 & 167.3 & 180.7 & 4.57 \\ LMC122 & 5:04:52.9 & -71:01:25 & 139.3 & 150.6 & 4.56 & LMC180 & 5:40:51.5 & -72:12:28 & 111.6 & 120.0 & 4.51 \\ LMC123 & 5:01:18.0 & -66:53:00 & 113.5 & 121.3 & 4.48 & LMC181 & 5:43:35.7 & -68:03:58 & 97.2 & 103.1 & 4.40 \\ LMC124 & 5:01:00.3 & -67:28:27 & 136.9 & 147.6 & 4.56 & LMC182 & 5:44:16.0 & -68:39:32 & 141.0 & 152.9 & 4.59 \\ LMC125 & 5:00:36.1 & -68:03:54 & 162.5 & 177.3 & 4.65 & LMC183 & 5:45:02.8 & -69:14:59 & 173.3 & 189.9 & 4.66 \\ LMC126 & 5:00:02.4 & -68:39:31 & 270.9 & 309.7 & 4.82 & LMC184 & 5:45:43.2 & -69:50:33 & 243.2 & 277.0 & 4.78 \\ LMC127 & 4:59:33.6 & -69:14:54 & 296.1 & 340.8 & 4.84 & LMC185 & 5:46:30.8 & -70:25:51 & 350.9 & 413.7 & 4.90 \\ LMC128 & 4:59:03.6 & -69:50:24 & 184.2 & 203.4 & 4.71 & LMC186 & 5:47:21.2 & -71:01:24 & 205.6 & 225.9 & 4.67 \\ LMC129 & 4:58:24.6 & -70:26:07 & 151.9 & 164.6 & 4.58 & LMC187 & 5:48:12.6 & -71:36:52 & 141.6 & 154.2 & 4.57 \\ LMC130 & 4:57:50.8 & -71:01:20 & 118.1 & 126.2 & 4.46 & LMC188 & 5:48:26.6 & -72:12:27 & 68.0 & 71.9 & 4.28 \\ LMC131 & 4:55:28.6 & -66:52:46 & 133.3 & 142.8 & 4.49 & LMC189 & 5:50:37.9 & -68:39:26 & 81.7 & 86.1 & 4.28 \\ LMC132 & 4:55:00.6 & -67:28:36 & 107.1 & 114.0 & 4.44 & LMC190 & 5:51:33.2 & -69:14:55 & 107.7 & 114.4 & 4.40 \\ LMC133 & 4:54:29.2 & -68:03:47 & 186.6 & 204.2 & 4.64 & LMC191 & 5:52:20.1 & -69:50:24 & 137.7 & 147.9 & 4.48 \\ LMC134 & 4:53:49.2 & -68:39:18 & 158.5 & 171.4 & 4.58 & LMC192 & 5:53:24.1 & -70:25:51 & 143.3 & 154.0 & 4.46 \\ LMC135 & 4:53:05.2 & -69:14:51 & 140.0 & 150.6 & 4.55 & LMC193 & 5:54:21.7 & -71:01:34 & 83.4 & 88.1 & 4.24 \\ LMC136 & 4:52:23.7 & -69:50:25 & 117.3 & 125.1 & 4.49 & LMC194 & 5:55:29.7 & -71:36:59 & 44.8 & 46.8 & 3.99 \\ LMC137 & 4:51:30.2 & -70:26:01 & 87.0 & 92.3 & 4.40 & LMC195 & 5:56:00.0 & -72:12:25 & 24.9 & 25.8 & 3.73 \\ LMC138 & 4:49:34.7 & -66:53:07 & 52.2 & 54.8 & 4.17 & LMC196 & 5:56:54.7 & -68:39:29 & 51.7 & 54.2 & 4.16 \\ LMC139 & 4:49:05.2 & -67:28:30 & 52.8 & 55.5 & 4.21 & LMC197 & 5:58:02.7 & -69:15:06 & 50.2 & 52.5 & 4.06 \\ LMC140 & 4:48:18.2 & -68:04:05 & 81.9 & 86.8 & 4.37 & LMC198 & 5:59:02.5 & -69:50:35 & 43.0 & 44.8 & 3.95 \\ LMC141 & 4:47:26.7 & -68:39:36 & 85.6 & 90.7 & 4.38 & LMC199 & 6:00:14.7 & -70:26:00 & 39.2 & 40.8 & 3.90 \\ LMC142 & 4:46:31.9 & -69:15:08 & 112.7 & 120.4 & 4.45 & LMC200 & 6:01:27.5 & -71:01:36 & 35.1 & 36.6 & 3.87 \\ LMC143 & 4:45:43.1 & -69:50:19 & 78.7 & 82.9 & 4.29 & LMC201 & 6:02:45.9 & -71:37:04 & 46.9 & 49.2 & 4.06 \\ LMC144 & 4:44:40.2 & -70:26:01 & 55.8 & 58.5 & 4.16 & LMC202 & 6:03:28.3 & -72:12:34 & 42.9 & 44.9 & 4.00 \\ LMC145 & 4:43:47.5 & -66:52:43 & 30.1 & 31.3 & 3.91 & LMC203 & 6:03:29.9 & -72:48:04 & 40.9 & 42.8 & 3.95 \\ LMC146 & 4:43:03.0 & -67:28:17 & 39.4 & 41.2 & 4.03 & LMC204 & 6:03:14.6 & -68:39:25 & 55.3 & 58.1 & 4.13 \\ LMC147 & 4:42:07.8 & -68:03:55 & 46.9 & 49.1 & 4.11 & LMC205 & 6:04:32.9 & -69:15:04 & 36.7 & 38.4 & 4.01 \\ LMC148 & 4:41:06.8 & -68:39:27 & 49.1 & 51.4 & 4.14 & LMC206 & 6:05:40.3 & -69:50:27 & 38.3 & 40.0 & 3.99 \\ LMC149 & 4:40:05.1 & -69:14:57 & 52.3 & 54.8 & 4.16 & LMC207 & 6:07:04.2 & -70:25:55 & 35.9 & 37.4 & 3.95 \\ LMC150 & 4:39:05.3 & -69:50:16 & 44.8 & 46.8 & 4.08 & LMC208 & 6:08:30.4 & -71:01:27 & 44.1 & 46.2 & 4.04 \\ LMC151 & 4:37:51.6 & -70:25:45 & 38.1 & 39.8 & 4.05 & LMC209 & 6:10:07.0 & -71:37:00 & 37.5 & 39.1 & 3.91 \\ LMC152 & 4:37:54.1 & -66:52:52 & 25.3 & 26.5 & 3.76 & LMC210 & 6:10:55.7 & -72:12:37 & 34.5 & 36.1 & 3.94 \\ LMC153 & 4:37:01.7 & -67:28:30 & 27.3 & 28.5 & 3.85 & LMC211 & 6:11:22.0 & -72:48:04 & 31.3 & 32.7 & 3.88 \\ LMC154 & 4:35:59.1 & -68:04:02 & 26.6 & 27.8 & 3.91 & LMC212 & 6:11:04.0 & -69:14:50 & 39.9 & 41.7 & 4.01 \\ LMC155 & 4:34:49.4 & -68:39:32 & 33.0 & 34.5 & 3.95 & LMC213 & 6:12:17.9 & -69:50:37 & 32.7 & 34.1 & 3.81 \\ LMC156 & 4:33:32.7 & -69:15:00 & 34.3 & 35.9 & 3.96 & LMC214 & 6:13:58.2 & -70:26:08 & 31.6 & 32.9 & 3.88 \\ LMC157 & 4:32:23.8 & -69:50:26 & 25.6 & 26.6 & 3.74 & LMC215 & 6:15:36.4 & -71:01:28 & 32.0 & 33.3 & 3.88 \\ \hline \noalign{\vskip5pt} \multicolumn{9}{r}{total} & 19,424.4 & 22,740.0 & \\ \noalign{\vskip5pt} \hline \end{tabular} \end{tiny} \medskip \begin{flushleft} {\it Note:} Coordinates point to the centre of the field (centre of the mosaic), each being $35' \times 35'$. Number of ``good'' objects in the template is provided ($N>80$ and $\langle I \rangle < 20.4$ mag) together with the estimated number of real monitored stars (see Section \ref{sec:blending}). Mean number of all objects detected on a single CCD used for calculating the density of a field is given in the last column. \end{flushleft} \medskip \bigskip \bigskip \end{table*} \begin{figure*} \includegraphics[width=12cm]{fig1.eps} \caption{Positions of the OGLE--III LMC fields (black). Also shown are all OGLE--II fields (red rectangles). The three small filled squares show the positions of the HST fields used for our blending determination. Background image credit: ASAS all sky survey.} \label{fig:fields} \end{figure*}

In this study we analysed almost 8 years of observations of the Large Magellanic Cloud by OGLE--III. The data set with its volume, coverage and quality supersedes all previous determinations of the microlensing optical depth, including the one based on the OGLE--II data (Paper I). We detected two sound candidates for microlensing events and further two possible candidates. Neither of them, however, is likely to be caused by dark matter compact lenses from the halo of our Galaxy. The two best candidates can be explained as an expected signal from the self-lensing within the LMC. Of the remaining two, one is either a binary event or a some kind of chromatic outburst, whereas the other is a candidate for galactic disk lens. Such null detection for MACHO lensing led to estimating the upper limit on their contribution to the mass of the Halo of the Galaxy. The upper limit set at a level of 6-7 per cent at $M=0.4\msun$ leaves very little room for dark matter compact objects. Still, at the moment we can not exclude more heavy dark matter lenses, like black holes. Our survey puts a 20 per cent halo mass fraction limit on compact objects with masses of $M=10\msun$ and actually no limit on higher masses. This heavy mass end window should be now explored with more attention. As a side product of our analysis we also discovered that event MACHO-LMC-7, reported by the MACHO group and used in their final optical depth determination, exhibited couple of additional brightening episodes in the OGLE-III data, a feature which excludes it as a genuine microlensing event. With the OGLE project continuing now in its fourth phase we hope the sensitivity to extremely long events will improve significantly within next years. It should result in the increase in the statistics of potential black-hole lenses or allow us to rule out heavy dark matter compact objects as well and close that topic definitively.

1012.1154

1012

1012.1728.txt

We investigate the emission properties of polycyclic aromatic hydrocarbons (PAHs) in various metallicity environments with the \textit{Infrared Spectrograph} on board \textit{Spitzer}. Local giant H\2\ regions are used as references as they enable access to the distinct interstellar medium components that contribute to the mid-infrared spectrum of star-forming galaxies: photodissociation regions (PDRs), photoionized gas, stellar clusters, and embedded regions. Three objects are considered, NGC\,3603 in the Milky Way, 30\,Doradus in the Large Magellanic Cloud, and N\,66 in the Small Magellanic Cloud. From the variations of the PAH/14\mic\ ratio, we find that PAHs are destroyed in the ionized gas for a radiation field such that [Ne\3]/[Ne\2]$\gtrsim3$. From the variations of the PAH/Hu$\alpha$ ratio, we find that the PAH emission sources in the giant H\2\ regions follow the same photodestruction law regardless of metallicity. We then compare these results with observations of starburst galaxies, H\2\ galaxies, and blue compact dwarf galaxies (BCDs). While the integrated mid-infrared spectra of BCDs are reminiscent of a warm dusty ionized gas, we observe a significant contribution to the PAH emission in starburst galaxies that is not arising from PDRs.

\label{sec:intro} From a simplistic chemical viewpoint, the abundance of molecules in the interstellar medium (ISM) of galaxies should scale with the metallicity of the environment, either because of the presence of nucleosynthetic constituents or because of the need for catalysts (dust grains). Obviously, indirect consequences of the low metal abundance also impact molecule survival, such as the lack of certain coolants (such as [O\1] or CO) that results in a globally warmer ISM, and hard interstellar radiation field (ISRF) due to energetic photons arising from low-metallicity ionizing stars. Observations show that nearby star-forming galaxies with metallicities lower than $\approx1/6$\,Z$_\odot$\footnote{We choose from now on the solar abundances of Asplund et al.\ (2006) for reference.} ($12+\log {\rm(O/H)}\lesssim7.9$) seem to lack CO molecules in giant molecular clouds (GMCs; Arnault et al.\ 1988; Taylor et al.\ 1998) and polycyclic aromatic hydrocarbons (PAHs) in photodissociation regions (PDRs) (Engelbracht et al.\ 2005; Madden et al.\ 2006; Wu et al.\ 2006; O'Halloran et al.\ 2006; Rosenberg et al.\ 2006). Even H$_2$, although it is detected in PDRs of metal-poor star-forming galaxies (Vanzi et al.\ 2000; Hunt et al.\ 2010), is not seen in absorption in their diffuse ISM (Vidal-Madjar et al.\ 2000; Lebouteiller et al.\ 2004; Thuan et al.\ 2002, 2005). This result illustrates the indirect effect of low-metallicity since the paucity of diffuse H$_2$ is related to the combination of low H\1\ volumic density, dust grain abundance, and hard UV radiation (Vidal-Madjar et al.\ 2000). In the present paper, we study the emission properties of PAHs, which are carbon-based molecules suggested by Leger \& Puget (1984) and Allamandolla et al.\ (1985) to explain the mid-infrared (MIR) emission bands seen ubiquitously toward star-forming regions. PAH features are often used as a diagnostic tool, as it is thought that PAH emission bands could trace star-formation on timescales on the order of O-stars lifetime (F{\"o}rster Schreiber et al.\ 2004; Weedman \& Houck 2008) or B-stars lifetime (Peeters et al.\ 2004). It is therefore essential to understand the mechanisms of their formation and destruction as well as the physical conditions in which PAH emission occurs. The weakness of PAH emission in low-metallicity star-forming galaxies is thought to be due to a combination of several effects: insufficient chemical evolution of the environment leading to the lack of PAH progenitors (hydrocarbons), enhanced molecule destruction by energetic photons and by shocks, and smaller volume of PAH-dominated regions. The main difficulty in investigating these parameters is their cross-correlation. Although a trend of decreasing PAH emission with both metallicity and ionizing radiation hardness has been found in several studies (Madden 2000; Madden et al.\ 2006; Wu et al.\ 2006; Engelbracht et al.\ 2006), the dominant parameter remains unknown. The study of H\2\ regions within M\,101 by Gordon et al.\ (2008) suggests that the aromatic band emission tends to depend more on the ionization of the environment rather than on metallicity. The photodestruction mechanism could be isolated by investigating PAH emission at ionization fronts in reflection nebul\ae\ (e.g.; Sellgren et al.\ 1985) and across spatially-resolved H\2\ regions where the metallicity is known to be uniform (Madden et al.\ 2006; Lebouteiller et al.\ 2007; Bernard-Salas et al.\ 2011). The present study aims at better understanding how the PAH emission in galaxies depends on the ISM morphology, by comparing the MIR properties of star-forming galaxies with those of giant H\2\ regions. The latter enable access to the distinct physical regions contributing to the composite mid-infrared (MIR) spectrum of star-forming galaxies, such as warm photoionized gas, PDRs, molecular clouds, shock fronts, stellar clusters, or supernov\ae\ remnants. Three well studied giant H\2\ regions with different metallicities are used as templates. The objects are NGC\,3603 in the Milky Way with essentially a solar metallicity, 30\,Doradus (hereafter 30\,Dor) in the Large Magellanic Cloud (LMC) with 0.6\,Z$_\odot$, and N\,66/NGC\,346 in the Small Magellanic Cloud (SMC) with 0.2\,Z$_\odot$. Although these regions might not be representative of typical H\2\ regions in star-forming galaxies, it is possible to distinguish their ISM components in the MIR to compare to integrated spectra of galaxies. This paper is part of a series meant to investigate the MIR properties of nearby giant H\2\ regions. Lebouteiller et al.\ (2007) analyzed small spatial-scale variations of the MIR features in NGC\,3603 to characterize PAH destruction by energetic photons. A valuable advantage is that chemical abundances are remarkably uniform across the region, hypothesis confirmed a posteriori by Lebouteiller et al.\ (2008). Hence, it provided a unique possibility to isolate the metallicity parameter and study the specific effect of ionizing radiation hardness. The two other objects are being analyzed a similar way (Bernard-Salas et al.\ 2011; Whelan et al.\ 2011). We focus here on the brightest regions seen in the MIR to compare the PAH properties in the 3 giant H\2\ regions and to understand how PAH emission depends on both the physical and chemical conditions (Sect.\ \ref{sec:destr_ghiir}). We then view our results alongside integrated spectra of star-forming galaxies to determine general diagnostics of molecular and dust properties in various environments (Sect.\,\ref{sec:sfg}). The two specific issues we wish to address concern the influence of metallicity on the PAH-to-dust ratio, and whether the relative contribution of ionized gas versus PDRs in the integrated spectra of star-forming galaxies can explain the observed PAH deficiency in low-metallicity objects. Finally, we examine in Section \ref{sec:molh} the relation between PAH emission and molecular hydrogen.

In this paper, we first investigated the PAH emission properties of several sources within three giant H\2\ regions, NGC\,3603 in the Milky Way, 30\,Dor in the LMC, and N\,66 in the SMC. The sample spans an interesting range of metallicities (down to $0.2$\,Z$_\odot$) and physical conditions (ISRF hardness). The sample of giant H\2\ regions contain several ISM components contributing to the MIR spectra of star-forming galaxies, stellar clusters, individual PDRs, warm photoionized gas, as well as embedded MIR bright regions. The main results are: \begin{itemize} \item The PAH/14\mic\ ratio shows a strong dependence on the ISRF hardness as traced by [Ne\3]/[Ne\2]. The locations of the data points in the PAH/14\mic\ vs.\ [Ne\3]/[Ne\2] diagram agree well with the PDR vs.\ ionized gas nature of the sources. For soft ISRF, PAH/14\mic\ is uniform and dominated by PDRs. A threshold occurs around [Ne\3]/[Ne\2]$\approx3$ above which the dispersion of PAH/14\mic\ values greatly increases, with a decreasing ratio as the ISRF hardens. \item N\,66 is characterized by large PAH/14\mic\ values compared to positions in 30\,Dor with similar physical conditions. The difference is due to suppressed warm dust emission in N\,66, which we propose is an effect of the low dust-to-gas ratio in the ionized gas. \item The variations of the PAH/14\mic\ ratio are not dominated by varying ISRF conditions leading to different heating mechanism of PAHs and of the dust. A change in the PAH abundance is required, being consistent with the paradigm of photodestruction of PAH molecules in the ionized gas. \item PAH emission is compared to the H\1\ recombination line Hu$\alpha$ which traces the ionized gas. The variations of PAH/Hu$\alpha$ are directly related to the relative contribution of PDR vs.\ ionized gas. We find that all sources in the giant H\2\ regions follow the same trend as a function of the ISRF hardness. We use these results as a reference for discussing observations of star-forming galaxies. \item By investigating the dependence of PAH/Hu$\alpha$ on the iron abundance, we do not find any evidence for enhanced PAH destruction by shocks. \end{itemize} The results on the giant H\2\ regions are compared to those of star-forming galaxies (BCDs, starburst galaxies, and H\2\ galaxies), with the following findings: \begin{itemize} \item The PAH/14\mic\ ratio in star-forming galaxies follows on first order the trend of the giant H\2\ regions. Starburst galaxies globally lie with the PDRs in giant H\2\ regions while BCDs lie with sources in giant H\2\ regions dominated by ionized gas. \item PAH/Hu$\alpha$ in starburst galaxies is larger than PDR values in giant H\2\ regions which is likely due to the PAH excitation by old stellar populations in the diffuse ISM. The PAH emission component in the galaxies of the sample is thus not dominated by a collection of PDRs similar to the ones observed in giant H\2\ regions. \item We find a metallicity dependence of the 14\mic/Hu$\alpha$ ratio, which reflects the warm dust abundance in the photoionized gas. \end{itemize} Low-metallicity star-forming galaxies are particularly challenging environments for detecting PAH emission because of the low carbon abundance which limits the formation of hydrocarbon carriers, and of the hard and intense radiation field which is able to destroy PAH molecules. Furthermore, dwarf galaxies do not contain significant amounts of GMCs, implying little integrated volume of PDRs. Our findings suggest that metal-poor star-forming galaxies such as BCDs have PAH/Hu$\alpha$ ratios that agree on first order with what is expected from typical physical conditions similar to the ionized gas in H\2\ regions. Even though PAHs could be formed in large amounts over the galaxy history, molecules are mostly photodissociated in the ionized gas. In more metal-rich galaxies such as starburst galaxies, there is a sign of significant diffuse PAH emission which could severely bias star-formation rate indicators based on PAH emission.

1012.1728

1012

1012.4547_arXiv.txt

Firstly, we give a historical overview of attempts to incorporate magnetic fields into the Smoothed Particle Hydrodynamics method by solving the equations of Magnetohydrodynamics (MHD), leading an honest assessment of the current state-of-the-art in terms of the limitations to performing realistic calculations of the star formation process. Secondly, we discuss the results of a recent comparison we have performed on simulations of driven, supersonic turbulence with SPH and Eulerian techniques. Finally we present some new results on the relationship between the density variance and the Mach number in supersonic turbulent flows, finding $\sigma^{2}_{\ln \rho} = \ln (1 + b^{2} \mathcal{M}^{2})$ with $b=0.33$ up to Mach~20, consistent with other numerical results at lower Mach number \citep{ls08} but inconsistent with observational constraints on $\sigma_{\rho}$ and $\mathcal{M}$ in Taurus and IC5146.

Magnetic fields and turbulence are thought to be two of the most important ingredients in the star formation process, so it is critical that we are able to model the effect of both in star formation simulations. Smoothed Particle Hydrodynamics (SPH, for recent reviews see \citealt{monaghan05,price04}), since it is a Lagrangian method where resolution follows mass, is a very natural method to use in order to model star formation. However the introduction of magnetic fields in SPH has a somewhat troubled history, and there are perceptions that the explicit use of artificial viscosity terms in order to treat shocks is too crude to enable accurate modelling of supersonic turbulence. We will discuss both of these aspects in this talk.

In summary: \begin{itemize} \item Being a television presenter is easier than getting MHD in SPH to work \item Development of MHD in SPH would proceed faster in the absence of arguments over basic misunderstandings of SPH which require lengthy comparison projects to resolve. \item Simulations with supercritical field strengths already tell us that magnetic fields can significantly change star formation even with weak fields. \item SPH and grid codes agree very well on the statistics of supersonic turbulence provided comparable resolutions are used (for power spectra this means $n_{cells} \approx n_{particles}$, but SPH was found to be much better at resolving dense structures). \item We have constrained the density variance -- Mach number relation in supersonic turbulence to $\sigma_{s}^{2} = \ln( 1 + b^{2} \mathcal{M}^{2})$, with $b\approx0.33$ up to Mach~20, but observed density variances suggest $b\approx0.5$, higher than can be produced with purely solenoidally-driven turbulence -- rather, some form of compressive driving or additional physics such as gravity is required to explain the observations. \end{itemize}

1012.4547

1012

1012.3659_arXiv.txt

\noindent In a previous paper, the observational, mapping and source-extraction techniques used for the Tenth Cambridge (10C) Survey of Radio Sources were described. Here, the first results from the survey, carried out using the Arcminute Microkelvin Imager Large Array (LA) at an observing frequency of 15.7~GHz, are presented. The survey fields cover an area of $\approx 27$~deg$^{2}$ to a flux-density completeness of 1~mJy. Results for some deeper areas, covering $\approx 12$~deg$^{2}$, wholly contained within the total areas and complete to 0.5~mJy, are also presented. The completeness for both areas is estimated to be at least 93~per~cent. The 10C survey is the deepest radio survey of any significant extent ($\gtrsim 0.2$~deg$^{2}$) above 1.4~GHz. The 10C source catalogue contains 1897 entries and is available online from the survey website (www.mrao.cam.ac.uk/surveys/10C). The source catalogue has been combined with that of the Ninth Cambridge Survey to calculate the 15.7-GHz source counts. A broken power law is found to provide a good parameterisation of the differential count between 0.5~mJy and 1~Jy. The measured source count has been compared to that predicted by \citet{dezotti2005} -- the model is found to display good agreement with the data at the highest flux densities. However, over the entire flux-density range of the measured count (0.5~mJy to 1~Jy), the model is found to under-predict the integrated count by $\approx 30$ per cent. Entries from the source catalogue have been matched to those contained in the catalogues of the NRAO-VLA Sky Survey and the Faint Images of the Sky at Twenty Centimetres survey (both of which have observing frequencies of 1.4~GHz). This matching provides evidence for a shift in the typical 1.4-to-15.7-GHz spectral index of the 15.7-GHz-selected source population with decreasing flux density towards sub-mJy levels -- the spectra tend to become less steep. Automated methods for detecting extended sources, developed in the earlier 10C paper, have been applied to the data; $\approx 5$~per~cent of the sources are found to be extended relative to the LA synthesised beam of $\approx 30$~arcsec. Investigations using higher-resolution data showed that most of the genuinely extended sources at 15.7~GHz are classical doubles, although some nearby galaxies and twin-jet sources were also identified.

\subsection{Background} The Ninth Cambridge (9C) Survey of Radio Sources \citep{waldram2003,waldram2010}, carried out using the Ryle Telescope (RT) at an observing frequency of 15.2 GHz, was a milestone in the exploration of the high-radio-frequency sky, as the first survey of significant extent and depth at such a high radio frequency. Since the publication of the first 9C paper, extensive survey work has been carried out using the Australia Telescope Compact Array at 20~GHz \citep[ATCA;][]{ricci2004,sadler2006,massardi2008,massardi2010, murphy2010}. The two surveys are complementary, with 9C probing deeper flux-density levels (down to 5.5~mJy) and the ATCA surveys covering shallower and wider areas (most recently, the whole southern sky). It is well known that high-frequency radio surveys are highly time-consuming. The scaling of interferometer primary-beam areas with frequency ($\propto \nu^{-2}$), and the typical synchrotron spectra of radio sources ($\propto \nu^{-0.7}$) conspire so that the time required to carry out a survey of equivalent depth and sky-coverage, using a telescope of fixed aperture diameter, scales as $\nu^{3.4}$. Things are somewhat better if it is assumed that the available bandwidth scales linearly with frequency. However, the fact that the noise temperatures of the available front-end, low-noise amplifiers used in interferometers tends to increase with frequency must also be taken into account. For these reasons, relatively little survey work has been undertaken at high radio frequencies and the knowledge of the source population remains poor. Nevertheless, familiarity with the properties of this population is important for the interpretation of the results from observations of the Cosmic Microwave Background (CMB), such as those made by \textit{Planck} \citep{tauber2010}. At mm wavelengths, foreground radio sources are the dominant source of contamination of small-scale CMB anisotropies \citep{dezotti1999}. \citet{waldram2003,waldram2010} have demonstrated that extrapolation of the flux densities of sources at low frequencies cannot be relied upon to predict their high-frequency properties, which emphasises the value of survey work at the frequencies of interest ($\gtrsim 10$~GHz) for CMB work. Samples of bright sources selected at high radio frequencies have significant proportions with flat or rising spectra \citep[see, for example,][]{taylor2001,davies2009}. In the main, these sources are believed to be blazars with synchrotron self-absorbed spectra; the self-absorbed components of such sources are often highly variable \citep[see, for example,][]{franzen2009}. High-frequency-selected samples also include appreciable numbers of sources with convex spectra, peaking at GHz frequencies \citep[see, for example,][]{bolton2004}. Some of these GHz peaked spectrum (GPS) sources \citep[see][for a review]{odea1998} are believed to be associated with young objects, which later expand into powerful radio sources, though many are dominated by emission from a strongly-beamed self-absorbed component \citep{bolton2006}. Surveys such as the 9C provide flux-density-limited samples, which are useful for gaining further understanding of the evolution of such objects. \subsection{This work} Since the 9C survey was carried out, the RT has been transformed, by the installation of new front-end receivers and back-end electronics (including a new correlator), into the Arcminute Microkelvin Imager (AMI) Large Array (LA) \citep[see][for a detailed description of the telescope]{zwart2008}. The LA is a radio synthesis-telescope, located $\approx 19$~m above sea level near Cambridge. It is used to observe at a centre frequency of 15.7~GHz and has a usable bandwidth of 4.5~GHz. At this frequency, the telescope has a full-width-at-half-maximum (FWHM) primary beam of $\approx 5.5$~arcmin and a resolution of $\approx 30$~arcsec. The LA has been used to carry out the Tenth Cambridge (10C) Survey of Radio Sources. As part of this survey, the improved flux sensitivity of the LA, compared with the RT, has been used to explore the 15-GHz-band sky to sub-mJy levels. In a previous paper \citep[hereafter Paper~I]{franzen2010} detailed technical information regarding the survey strategy, mapping and source-extraction techniques for the 10C survey was provided. In this paper, the first results from 10 fields, including the 15.7-GHz source count, are presented. Throughout this paper any equatorial coordinates use equinox J2000 and spectral indices are defined using the convention that $S \propto \nu^{-\alpha}$.

The AMI LA has been used to carry out the 10C survey, the deepest radio source survey of any significant extent ($\gtrsim 0.2$~deg$^{2}$) above 1.4~GHz. The resulting deep 15.7-GHz source counts are useful for the interpretation of CMB data, for which foreground radio sources are an important contaminant. The source catalogue also provides an invaluable resource for the study of faint high-frequency-selected radio sources. The survey covers $\approx 27$~deg$^{2}$ complete to 1~mJy and $\approx 12$~deg$^2$ (wholly contained within the larger area) complete to 0.5~mJy. The number of sources with $S > 25$~mJy, appearing in the survey catalogue, is biased low; several of the survey fields were chosen to minimise the number of such sources appearing within the survey areas. In total, 1897 sources appear in the 4.62-$\sigma$ catalogue; the faintest being $\approx 100$~$\mu$Jy. A number of the key conclusions from the work are listed below. \begin{noindlist} \item[(1)] The 10C differential source count was parameterised using a broken power law. This was found to provide a significantly better fit to the data than a simpler, single-power-law parameterisation. The best-fit broken power law is \[n(S) \equiv \frac{\mathrm{d}N}{\mathrm{d}S} \approx \left\{ \begin{array}{l l} ~~24 \left(\frac{S}{\mathrm{Jy}}\right)^{-2.27}~\mathrm{Jy^{-1}~sr^{-1}}~\mathrm{for}~2.8 \leq S \leq 25~\mathrm{mJy} \\ 376 \left(\frac{S}{\mathrm{Jy}}\right)^{-1.80}~\mathrm{Jy^{-1}~sr^{-1}}~\mathrm{for}~0.5 \leq S < 2.8~\mathrm{mJy.} \end{array} \right. \] \newline \item[(2)] After having applied corrections to the individual source flux densities measured as part of the 9C survey to account for the small difference in frequencies between the 9C and 10C surveys, the 9C and 10C data were combined. The addition of the 9C data allowed the calculation of the best estimate of the 15.7-GHz differential source count by improving the source-count statistics for $5.5 \leq S < 25$~mJy and by providing data for a complete sample of sources with $25~\mathrm{mJy} \leq S \leq 1~\mathrm{Jy}$. Again, a broken-power-law parameterisation was found to offer a significantly improved fit to the data compared with that provided by a single power law. The best-fit broken power law is \[n(S) \equiv \frac{\mathrm{d}N}{\mathrm{d}S} \approx \left\{ \begin{array}{l l} ~~48 \left(\frac{S}{\mathrm{Jy}}\right)^{-2.13}~\mathrm{Jy^{-1}~sr^{-1}}~\mathrm{for}~2.2~\mathrm{mJy} \leq S \leq 1~\mathrm{Jy} \\ 340 \left(\frac{S}{\mathrm{Jy}}\right)^{-1.81}~\mathrm{Jy^{-1}~sr^{-1}}~\mathrm{for}~0.5 \leq S < 2.2~\mathrm{mJy.} \end{array} \right. \] \newline \item[(3)] The model counts by \citet{dezotti2005} are found to display good agreement with the 9C and 10C data at the high flux-density-end of the measured count. However, with decreasing flux density the model first over-predicts and then, below about 5~mJy, under-predicts the measured count. By integrating the model differential source count, the model was found to under-predict the total number of sources, with flux densities between 0.5~mJy and 1~Jy, per unit area by approximately 30~per~cent. This deficit, over the entire flux-density range, is attributable to the model underestimating the count at the lowest flux densities. \newline \item[(4)] Entries from the 10C source catalogue were matched to those contained in the catalogues of the NVSS and FIRST survey (both of which have observing frequencies of 1.4~GHz). The matching revealed a shift in the typical 1.4-to-15.7-GHz spectral index of the 15.7-GHz-selected source population with decreasing flux density towards sub-mJy levels. When matching to NVSS, 49 per cent of sources with $5.0 \leq S_{10C} < 25.0$~mJy were found to have $\alpha^{15.7}_{1.4} > 0.81$. However, for sources with $0.5 \leq S_{10C} < 0.6$~mJy the corresponding figure is 18 per cent. The observed trend is in contrast to that measured for sources with higher flux densities by \citet{waldram2010}. They found that the typical spectral index became steeper for sources with decreasing flux densities between $\approx 100$~mJy and $\approx 10$~mJy. A similar effect, to that measured using the 10C data, has been observed as part of lower-frequency surveys -- in lower-frequency-selected samples significant numbers of flatter spectrum sources start to enter at sub-mJy levels \citep[see, for example,][and references therein]{prandoni2006}. \newline \item[(5)] Automated techniques for identifying extended sources, described in Paper~I, have been applied to the data. The proportion of sources that are extended relative to the LA synthesised beam of $\approx 30$~arcsec is $\approx 5$~per~cent; this is similar to the proportion of $\approx 6$ per cent measured by \citet{waldram2003}. A subset of 39 extended or overlapping sources, thought likely to be genuinely extended on the basis of their 15.7-GHz morphologies, were investigated further using higher-resolution data. These data confirmed that at least 29 of these sources are genuinely extended; most were identified as classical doubles but three are nearby galaxies and four are twin-jet sources. \end{noindlist}

1012.3659

1012

1012.1376_arXiv.txt

We examine the recoil velocity induced by the superposition of the magnetic dipole and quadrupole radiation from a pulsar/magnetar born with rapid rotation. The resultant velocity depends on not the magnitude, but rather the ratio of the two moments and their geometrical configuration. The model does not necessarily lead to high spatial velocity for a magnetar with a strong magnetic field, which is consistent with the recent observational upper bound. The maximum velocity predicted with this model is slightly smaller than that of observed fast-moving pulsars.

The surface magnetic field strength $B_{s}$ of a pulsar is conventionally estimated by matching the rotational energy loss rate with the magnetic dipole radiation rate, that is, $B_{s} \approx $ $(3 c^{3} I P \dot{P} )^{1/2} /(2\sqrt{2} \pi R_{s} ^{3})$, where $I$ is the inertial moment, $R_{s}$ is the stellar radius, $P$ is the spin period, and $\dot{P}$ is the time derivative of the spin period. The precision of this approximation is only at the order of magnitude level because actual energy loss is not well described by magnetic dipole radiation in a vacuum. A more realistic model with current flows and radiation losses is required, but has not yet been established. A simple estimate provides $B_{s} \approx 10^{12}$G for typical radio and X-ray pulsars, and $B_{s} \approx 10^{13}$-$ 10^{15}$G for magnetars, although the level of the approximation must be noted. Dynamo action in a rapidly rotating proto-neutron star with $P \approx 1 $ ms is proposed as a mechanism for this amplification by 2-3 orders of magnitude (see e.g., \cite{ThDu93, BoReUr03, BoUrBe06}). Actual upper limit of $B_{s} $ generated in the convective proto-neutron star is estimated as $10^{15}-10^{16}$G, beyond which all sorts of instabilities are suppressed by strong magnetic fields \citep{MiPoUr02}. Recent numerical simulations of dynamo action can be used to study the large-scale fields in fully convective rotating stars. For example, non-axisymmetric fields are generated in the case of uniform rotation \citep{ChKur06}, while mostly axisymmetric fields with a mixture of the first few multipoles are formed in the case of a differentially rotating star\citep{DoStBr06}. The results may not directly apply to pulsars or magnetars, but suggest that the magnetic field configuration of neutron stars may not be an ordered dipole. If there are higher-order multipoles, these will also contribute to the radiation loss. The upper bounds on their surface magnetic fields are rather loose. See \cite{Krol91} for a discussion of the magnetic fields of millisecond pulsars. The magnetic field strength $B_{lm}$ relevant to the multipole moment of order $(l,m)$ is limited to $ B_{lm} \le B_{s} /(m R_{s}\Omega /c)^{l-1} $, where $\Omega=2\pi/P$ is angular velocity, and the radiation of each multipole $ L_{lm} \sim c ( m R_{s} \Omega /c) ^{2l+2} $ $ (B_{lm} R_{s} ) ^{2} $ is assumed to be smaller than that of a dipole. Thus, a model with complex magnetic configuration at surface $ B_{lm} \ge B_{s} ~( l > 1 ) $ is allowed because of the small factor $ R_{s} \Omega /c \ll 1 $ for observed stars. Some proto-neutron stars are conjectured to be born in hypothetical extreme state of rapid rotation $P \approx 1 $ ms with an ultra strong magnetic field $B_{s} \approx 10^{15}$G. Is there any remaining evidence of this stage? The proper motion can possibly be used as a probe. Several kick mechanisms operative at the core bounce of a supernova explosion have been proposed to date: anisotropic emissions of neutrinos (e.g., \cite{ArLa99,FrKu06}), hydrodynamical waves (e.g., \cite{Scetal06}), and MHD effects (e.g., \cite{SaKoYa08}). These mechanisms operate on a dynamical timescale of the order of milliseconds or the cooling timescale of $\sim10$ s. If the strong magnetic fields are generated on a longer timescale, some natal kick mechanisms involved the magnetic-field-driven anisotropy do not work effectively. Recoil driven by electromagnetic radiation, which is operative on a longer spindown timescale of $\sim 10^3 (B/10^{15} {\rm G})^{-2} (P_i/1{\rm ms})^{2}$ s, has been proposed as a post-natal kick mechanism\citep{HaTa75} (see also \cite{LaChCo01} for the corrected expression). In their model, an oblique dipole moment displaced by a distance $s$ from the stellar center rotates. This causes the radiation of higher order multipoles, whose superposition is generally asymmetric in the spin direction, leading to the kick velocity. In the off-center model, the quadrupole field $B_{2}$ of order $B_{2} \sim (s/R_{s}) \times B_{1} \sim B_{1} $ is involved. It is interesting to study the case where $ B_{2} \gg B_{1} $, because the constraint of the higher order component by the radiation is very weak, for example, $B_{2} \sim (c/( R_{s}\Omega)) \times B_{1} \gg B_{1} $. In this paper, we revisit the kick velocity induced by electromagnetic radiation from a magnetized rotating star with both dipole and quadrupole fields, in which a larger quadrupole field $ B_{2} \gg B_{1} $ at the surface is allowed. This paper is organized as follows. In Section 2, we present the radiation from rotating dipole and quadrupole moments in vacuum. Their field strength and inclination angle with respect to the spin axis are arbitrary. We also compare our model with the off-centered dipole model. In Section 3, we evaluate the maximum kick velocity as a recoil of momentum radiation. Section 4 presents our conclusions.

Magnetic field strength itself is critical in most kick mechanisms. For example, $ B_{s} > 10^{15}$ G at the surface is required in asymmetric neutrino emission (e.g., \cite{ArLa99}), as well as in asymmetric magnetized core collapse (e.g., \cite{SaKoYa08}). The resultant velocity increases with the field strength because the asymmetry arises from the magnetic field. Magnetars are therefore expected to have high velocity if one of these mechanisms is operative. Recent observations do not support the high velocity. Rather, the upper limit of the transverse velocity $v_{\bot}$ has been reported, although there is uncertainty in the value. For example, $v_{\bot} \sim 210$ km s$^{-1}$ for AXP XTEJ1810-197\citep{Helfandetal07}, $v_{\bot} < 1300$ km s$^{-1}$ for SGR 1900+14 \citep{Kaplanetal09, Lucaetal09} and $v_{\bot} < 930$ km s$^{-1}$ for AXP 1E2259+586\citep{Kaplanetal09}. For fast moving pulsars, PSR2224+45 ($v_{\bot} > 800$ km s$^{-1}$ \cite{CoRoLu93}) and B1508+55 ($v_{\bot} \sim 1000$ km s$^{-1}$ \cite{Chatterjee05}) have been reported. These magnetic fields are quite ordinary, $B_{s} =2.7\times 10^{12}$G and $2.0\times 10^{12}$G, respectively. Thus, there is no clear correlation between the field strength and the velocity in the present sample. The electromagnetic rocket mechanism considered in \cite{HaTa75} and in this paper does not depend on field strength if the spin evolution is determined from the radiation loss. In our model, the ratio of dipole and quadrupole moments is important. The condition for high velocity is that the quadrupole field is large enough in magnitude for the radiation loss to be of the same order as the dipole field. The velocity also depends on the geometrical configuration of the multipole moments, that is, each inclination angle from the spin axis and the angle between the axes of symmetry of the moment. Assuming that the directions of moments are random, and that they are equally likely to be oriented in any direction, it is found that the mean velocity with respect to the configuration is not so large, $ \sim 120 (P_{i}/1{\rm ms})^{-2}$ km s$^{-1}$, for the optimized dipole-quadrupole ratio. The maximum velocity is realized for a specific configuration in which the inclination angle of the quadrupole moment is small, and the meridian plane in which the quadrupole moment lies is perpendicular to the plane of the dipole. The velocity increases up to $ \sim 930 (P_{i}/1{\rm ms})^{-2}$ km s$^{-1}$. This value is slightly smaller than the maximum observed velocity of a pulsar. The configuration is unknown, and is closely related to the origin of the magnetic field, dynamo or fossil. Nevertheless, \cite{BoUrBe06} reported interesting results within the mean-field dynamo theory. They argued that strong large-scale and weak small-scale fields are generated only in a star with a very short initial period, that is, the Rossy number is small, and that the maximum strength decreases and small-scale fields become dominant with the decrease of the initial period. Thus, magnetars may have an ordered dipole with a strong field, while some pulsars may have rather irregular fields with higher multipoles. Through the superposition of higher multipoles, pulsars in general come to have a larger radiation recoil velocity than magnetars. Finally, if the kick velocity of pulsars and magnetars is governed by the same mechanism, it either should not simply depend on magnetic field, or should depend on only the configuration. The latter possibility was explored here. Present argument is recognized as the order of magnitude level due to the rotating model in vacuum. Further improvement of the magnetosphere will be of importance to explore the idea.

1012.1376

1012

1012.2834_arXiv.txt

A new class of core-collapse supernovae (SNe) has been discovered in recent years by optical/infrared surveys; these SNe suggest the presence of one or more extremely dense ($\sim {10}^{5-11}~{\rm cm}^{-3}$) shells of circumstellar material (CSM) on ${10}^{2-4}$~AU scales. We consider the collisions of the SN ejecta with these massive CSM shells as potential cosmic-ray (CR) accelerators. If $\sim 10$\% of the SN energy goes into CRs, multi-TeV neutrinos and/or GeV-TeV gamma rays almost simultaneous with the optical/infrared light curves are detectable for SNe at $\lesssim 20-30$~Mpc. A new type of coordinated multi-messenger search for such transients of duration $\sim 1-10$~months is required; these may give important clues to the physical origin of such SNe and to CR acceleration mechanisms.

Introduction} The much-anticipated era of multi-messenger astronomy is coming. GeV-TeV gamma rays, which are powerful tracers of cosmic rays (CRs), are detected by \textit{Fermi} and ground-based Cherenkov detectors. Neutrinos, which can uniquely identify the production of hadronic CRs and probe dense sources, are detectable by the nearly-completed IceCube and the planned KM3Net~\cite{Ahr+04}. They are useful for studying sources especially when photons cannot escape directly. Violent explosions of massive stars, such as supernovae (SNe) and gamma-ray bursts (GRBs) may be prodigious neutrino/gamma-ray sources~\cite{HH02}. High-energy observations should reveal their physical origin, nonthermal processes, and extreme environments. TeV-PeV neutrino detections from extragalactic sources may be possible for GRBs~\cite{WB97,Mur08} and certain other kinds of SNe, e.g., those with hidden relativistic jets~\cite{RMW04}, relativistic pulsar winds~\cite{GS87}, or semi-relativistic external shocks~\cite{Wan+07}. But are SNe with more ordinary explosions also detectable? In recent years, blind surveys for optical transients have discovered ultra-bright SNe such as SN 2008am~\cite{Cha+11}, 2008iy~\cite{Mil+10}, 2006gy~\cite{Ofe+07}, 2005ap~\cite{Qui+07} and 2003ma~\cite{Res+09}, which are much more luminous than ordinary SNe. Their physical origin is not settled, but a plausible interpretation is strong shock dissipation by collision with a massive ($M_{\rm sh}\sim 1-30 M_{\odot}$) circumstellar material (CSM) shell at $R \sim {10}^{2-3}$~AU~\cite{SM07,WBH07}, which is more or less analogous to the mechanism of type IIn SNe, though some may be pair-instability SNe~\cite{Gal+09}. The local rate of these ultra-bright SNe may be $\sim {10}^{3-4}$ times larger than the local apparent rate of classical long GRBs~\cite{Mil+09}. But SNe with CSM do not have to be optically ultra-bright, and the existence of massive CSM shells may be even more common. The existence of dense CSM shells is indicated from other SNe, e.g., SN 2006jc, 2005ip and PTF 09UJ~\cite{Imm+08,Smi+09}. Also, massive CSM eruptions were observed in luminous blue variables such as $\eta$ Carinae~\cite{Smi+03}, and are supported by the discovery of a self-obscured SN~\cite{Koz+10}. These observations motivate us to investigate the system of SN ejecta crashing into massive CSM shells at various radii. We consider the possibility of CR acceleration at strong shocks formed by the crashes, and suggest a new class of high-energy transients that are neither bursts like GRBs nor persistent sources like SN remnants. Given their rate and timescales, new types of searches at intermediate time scales, with coordinated observations of these extragalactic SNe at $\lesssim 100$~Mpc, are required. We adopt the conventional notation $Q_x = Q/{10}^x$.

We have shown that SNe crashing into dense CSM are interesting targets for current and near-future high-energy detectors. Importantly, new types of coordinated multi-messenger searches are required to detect such $0.1-1$~yr transients. CR acceleration in dense surroundings is also motivated by recent gamma-ray observations~\cite{Abd+10,Wal+10}. Both detections and non-detections of SN-CSM emission will be useful, since physical mechanisms in such extreme environments are uncertain, especially when $\tau_T \gtrsim 1$. In particular, neutrinos have the benefit of probing hadronic CR accelerators, and they can escape earlier than thermal photons, which may be delayed by diffusion. In addition, the detections of signals would support the SN-CSM scenario~\cite{SM07,WBH07} rather than the pair-instability scenario~\cite{Gal+09}, useful for revealing the origin of bright SNe. One may expect synchrotron emission in the infrared-to-gamma-ray bands, as electrons of energy $E_e$ emit photons with $\sim 44~{\rm keV}~E_{e,\rm TeV}^2 B$. But the emission should be reprocessed to energies $\lesssim 10$~keV and/or strongly thermalized when $\tau_T \gg 1$, so we basically expect thermal emission observed from such SNe. But the synchrotron emission can be seen in the sufficiently hard x-ray range, if the collision happens at $\tau_T < 1$, as in Model B. The unabsorbed energy flux is estimated to be $\sim 7 \times {10}^{-13}~{\rm erg}~{\rm cm}^{-2}~{\rm s}^{-1}~f_{\rm syn} {\rm min}[1,f_{pp}] \epsilon_{\rm cr,-1} \mathcal{E}_{\rm ej,51} d_1^{-2} t_{s,7.8}^{-1}$, where $f_{\rm syn} \leq 1$ is the efficiency of synchrotron cooling. In Model B, the synchrotron cooling is more important than the synchrotron self-Compton and external Compton cooling at sufficiently high energies, so that pionic gamma rays are dominant in the GeV-TeV range while the synchrotron component, whose high-energy photon index is $(q+2)/2 \sim 2$, is relevant below $\sim 100$~MeV. (But other cooling processes such as the Coulomb interaction becomes more important especially at $E_e \lesssim 0.1-1$~GeV.) Hence, this signal will be an interesting target for the near-future x-ray monitor \textit{NuStar} (softer x rays can be masked by strong thermal bremsstrahlung emission that seems to be observed~\cite{Mil+10}). The radio emission is suppressed by the Razin effect, free-free absorption, and synchrotron self-absorption when the collision radius is small enough, though it may be expected at very large radii. For more quantitative theoretical studies, hydrodynamical simulations with radiation transfer and CR back-reaction are desirable. But our results are enough for the purpose of this work, which, in part, is to motivate new searches starting now. The relevant quantities, $\mathcal{E}_{\rm cr}$, $E_{p,s}^{\rm max}$ and $q$, have uncertainties, but we could see a source up to $d \sim 30-60$~Mpc if $\epsilon_{\rm cr}$ is larger than 0.1. The spectral index also affects the results. Although we adopt $q=2$, steeper indices lead to lower muon yields from SN-CSM neutrinos. But harder indices may also be expected since the gas may be radiation-dominated and/or the shock may be CR-mediated~\cite{BE87}. Since $E_{p,s}^{\rm max}$ depends on shock velocities, the non-uniformity of the ejecta affects the RS velocity and ratio of the RS dissipation to the FS dissipation~\cite{Che82}, but its pre-collision velocity distribution is uncertain since the ejecta may sweep the CSM before the collision. Also, the shock evolution may be radiative rather than adiabatic~\cite{Mar+10}. Plasma effects can modify the results, via e.g., wave damping by neutral particles or radiation. Especially, $E_{p,s}^{\rm max}$ may be limited by the size of the ionization region, since damping occurs in a time $\sim \frac{1}{n_{n} <\sigma_{i-n} v_{\rm th}>} \sim {10}^{0.5}~{\rm s}~n_{n,7.5}^{-1} T_{0}^{-0.4}$ in the neutral region~\cite{KC71}. Although the CSM gas would be initially neutral, ionization in the downstream and upstream region (not far from the shock) seems expected observationally~\cite{Smi+09} and theoretically~\cite{HS84} since the post-shock temperature is high. We have considered extragalactic SNe with dense and massive CSM shells. Such collisions may happen even for GRBs, where the jet breakout emission can be expected. The high-energy neutrino and gamma-ray emission is also possible, which would be more or less analogous to the (sub-)photospheric emission~\cite{Mur08,RMW04}. In the Galaxy, $\eta$ Carinae is a promising candidate that showed violent mass eruptions~\cite{Smi+03}. The radius of its massive nebula is larger than our typical values that are required for bright SNe. But the neutrino detection seems possible if the star explodes, since $N_{\mu, > \rm TeV} \sim 2 \times {10}^5 \mathcal{E}_{\rm cr,50} (M_{\rm sh}/10 M_{\odot}) R_{\rm sh,17}^{-2} V_{s,3.5}^{-1}$ (where $f_{pp}<1$). For smaller CSM mass, though the radiation might push the CSM shell, SN dynamics are not largely affected as in ordinary SNe~\cite{BK00}, where $\mathcal{E}_{\rm cr}$, $E_{p,s}^{\rm max}$ and $f_{pp}^{\rm sh}$ are much smaller. Detections would be challenging, but it may be interesting for a Galactic event. After this work was submitted and put onto the arXiv (arXiv:1012.2834), we became aware of Ref.~\cite{KSW11}, which is closely related to ours, and which supports our claims that these unusual SNe are interesting and that their high-energy emission is an important probe. We thank K. Ioka, C. Kochanek, C. Rott and R. Yamazaki for discussions. This work is supported by JSPS and CCAPP (KM), E. C. Howald Presidental Fellowship (BCL), Sloan Fellowship and NSF Grant AST-0908816 (TAT), and NSF CAREER Grant PHY-0547102 (JFB).

1012.2834

1012

1012.5413_arXiv.txt

Using \emph{ab initio} simulations we investigate whether water ice is stable in the cores of giant planets, or whether it dissolves into the layer of metallic hydrogen above. By Gibbs free energy calculations we find that for pressures between 10 and 40 Mbar the ice-hydrogen interface is thermodynamically unstable at temperatures above approximately 3000~K, far below the temperature of the core-mantle boundaries in Jupiter and Saturn. This implies that the dissolution of core material into the fluid layers of giant planets of giant planets is thermodynamically favoured, and that further modelling of the extent of core erosion is warranted.

According to core accretion models \citep{mizuno-78}, giant gas planets such as Jupiter and Saturn formed via the accumulation of an protocore of rock and ice which gained solid material until it reached sufficient size to begin accreting the gaseous component of the protosolar nebula. The existence of solid ice in the outer solar system promotes the rapid growth of the more massive protocores allowing the accretion of large quantities of gas necessary for Jupiter-sized planets. Giant planets thus have a dense core of rock and ice surrounded by a H-He envelope. It is not known, however, whether the initial dense core remains stable following the accretion of the H-He outer layer or whether the core erodes into the fluid hydrogen-rich layers above \citep{stevenson-pss-82,guillot-book}. The gravitational moments of Jupiter and Saturn, which have been measured by prior planetary missions and will be determined for Jupiter with unprecedented accuracy by the upcoming Juno mission, may be used in combination with interior models \citep{militzer-apj-08,guillot-book,hubbard-ass-05,saumon-apj-04,nettelmann} to estimate the mass of the present-day core, but it is unclear whether these masses correspond to the primordial core mass. It has been suggested \citep{guillot-book,saumon-apj-04} that the present-day core mass of Jupiter may be too small to explain its formation by core accretion within the relatively short lifetime of the protosolar nebula \citep{pollack-icarus-96}, although a more recent Jupiter model \citep{militzer-apj-08} predicted a larger core of 14--18 Earth masses which is consistent with core accretion. Furthermore, direct measurements of Jupiter's atmosphere suggest a significant enhancement in the concentration of heavy $(Z > 3)$ elements \citep{niemann-science-96}, but it is unknown to what extent this should be attributed to a large flux of late-arriving planetesimals versus the upwelling of core material. Determining the extent of core erosion is thus a major priority for understanding the interiors of giant planets and the process by which they were formed. In this work we focus on water ice, presumed to be a major constituent of the core, and consider the question of whether it has significant solubility in fluid metallic hydrogen at conditions corresponding to the core-mantle boundaries of giant gas planets. Water ice is the most prevalent of the planetary ices (water, methane and ammonia) which may be assumed to make up the outermost layers of a differentiated rock-ice core \citep{hubbard-science-81}. At the conditions of temperature and pressure prevalent at giant planet cores, water ice is predicted \citep{cavazzoni,french-prb-09} to be in either in a fully atomic fluid phase in which oxygen and hydrogen migrate freely and independently, or in a superionic phase in which oxygen atoms vibrate around defined lattice sites while hydrogen atoms migrate freely. Assuming the existence of a core-mantle boundary at which water ice and the fluid H-He phase are in direct contact, the relevant question is the extent to which the system may lower its Gibbs free energy by the redistribution of the atoms of the ice phase into the fluid hydrogen. The extreme pressure and temperature conditions prevalent at giant planet core-mantle boundaries (8000--12000K and 8--18 Mbar for Saturn, 18000--21000K and 35--45 Mbar for Jupiter) are not yet obtainable in the laboratory, thus \emph{ab initio} simulations provide the best available guide to determining the extent of core solubility.

The consequences of core erosion for planetary evolution models have been previously considered by \citet{stevenson-pss-82} and later \citet{guillot-book}. The effects of core erosion can potentially be detected either by orbital probes such as Juno or by atmospheric entry probes, since the redistribution of core material throughout the planet will manifest itself both by a smaller core (detectable from gravitational moments) and a higher concentration of heavy elements in the atmosphere than would be expected in a planet without core erosion. Once core material has dissolved into the metallic H layers, the rate at which core material can be redistributed throughout the planet is expected to be limited by double diffusive convection \citep{turner,huppert}. Since the higher density due to compositional gradients of the lower material interferes with the convection process, convection may be slowed significantly. \citet{guillot-book} modelled the effect of core erosion under the assumptions of fully soluble 30 Mbar core in each planet. Under their assumptions up to 19 Earth masses could have been redistributed from Jupiter's core but only 2 Earth masses from Saturn's, the difference being Jupiter's higher temperatures. While this prediction is subject to significant uncertainty in many aspects of the model, it does suggest that a redistribution of a significant fraction of the initial protocore is possible, at least in Jupiter. Further refinement of models for the upconvection of core material and its observable consequences may be fruitful. The effect of core erosion on the heat transport and mass distribution properties of Jupiter and Saturn should also be taken into account in future static models of these planets' interiors. These calculations can be expanded in several ways. We have neglected the presence of helium in the hydrogen-rich mantle, however due to the large magnitude of $\Delta G_{sol}$ and the chemical inertness of helium we do not expect the presence of helium in the mantle to significantly affect the solubility behavior. We have explicitly made the assumption that ice and hydrogen are in direct contact, an assumption which might fail in one of two ways. If hydrogen and helium are immiscible at the base of the atmosphere then the core may make direct contact with a helium-rich layer. This, however, is unlikely in the context of the calculations of \citet{morales-pnas-09} who predict hydrogen-helium immiscibility only far away from the cores of Jupiter and Saturn. The other possibility is that the ice layers of the core may be gravitationally differentiated, leaving ice beneath layers of the less dense planetary ices methane and ammonia. Since the bonding in these is similar to that in water ice one could assume that they show a similar solubility behavior, but this analysis is the subject of future work. We have also considered only the structure of the present-day planet. As suggested by \citet{slav}, a proper consideration of core solubility must also include solubility during the formation process, as dissolution of the icy parts of the core into the accreting hydrogen during the formation may result in the amount of ice on the core itself being small by the time the planet reaches its final size. A treatment of the formation processes for Jupiter and Saturn using ice solubilities derived from \emph{ab initio} calculations may be valuable. Our calculations strongly suggest that icy core components are highly soluble in the fluid mantle under the conditions prevalent at the core-mantle boundaries of Jupiter and Saturn. Since many recently-discovered exoplanets are more massive and hence internally hotter than Jupiter, it can be expected that any initial icy cores in these exoplanets will also dissolve. The presence of core erosion may allow models predicting a small present-day Jovian core to be made consistent with the large initial core required by core erosion, however models of the interior mass distribution of the planet will need to be revised to take the inhomogeneous composition of the lower layers implied by convection-limited core redistribution into account. Improved models which include core redistribution processes, combined with the data from the Juno probe, may assist in understanding the history and present structure of Jupiter and other planets in our own and in other solar systems.

1012.5413

1012

1012.2375_arXiv.txt

{Most extrasolar planets currently known were discovered by means of an indirect method that measures the stellar wobble caused by the planet. We previously studied a triple system composed of a star and a nearby binary on circular coplanar orbits. We showed that although the effect of the binary on the star can be differentiated from the stellar wobble caused by a planet, because of observational limitations the two effects may often remain indistinguishable. Here, we develop a model that applies to eccentric and inclined orbits. We show that the binary's effect is more likely to be mistaken by planet(s) in the case of coplanar motion observed equator-on. Moreover, when the orbits are eccentric, the magnitude of the binary's effect may be larger than in the circular case. Additionally, an eccentric binary can mimic two planets with orbital periods in the ratio 2/1. However, when the star's orbit around the binary's center of mass has a high eccentricity and a reasonably well-constrained period, it should be easier to distinguish the binary's effect from a planet. }

In \citet{Morais&Correia2008}, we studied a triple system consisting of a star perturbed by a binary system. Our aim was to derive an expression for the star's radial velocity in order to determine whether the binary's effect could be mistaken for that of a planet companion to the star. This question is relevant if one or even both binary components are unresolved which can happen for instance when these are faint but very common M stars. About $76\%$ of nearby main sequence stars are M-type \citep{Starrysky2001}. Moreover, in the solar neighborhood, over $50\%$ of G and K-stars \citep{Duquennoy&Mayor1991,Halbwachs_etal2003,Eggenberger_etal2004} and about $30\%$ of M-stars \citep{Fischer&Marcy1992,Delfosse_etal2004} belong to binary or even multiple systems. We previously showed that the binary system's effect on the star is a sum of two periodic signals with very close frequencies. These signals, if detected, could be associated with two planets on circular orbits around the star. However, these planets would have very close orbits and we should be able to discard them as unstable systems. Moreover, we derived expressions for the frequencies and amplitudes of these signals, thus were able to identify the hidden binary's parameters. As explained in \citet{Morais&Correia2008}, our results do not agree with previous work by \citet{Schneider&Cabrera2006} who studied the case of an equal-mass binary and concluded that its effect is a single periodic signal that mimics a planet on an eccentric orbit. Nevertheless, in \citet{Morais&Correia2008} we saw that there are realistic situations where we can identify one of the periodic signals but not the other. In this case, we may mistake the binary's effect for that of a planet on a circular orbit around the star. However, we can still apply our theory to compute the hidden binary's parameters and then try to detect these objects. Our work was based on a full three-body model but assumed coplanar motion and initial circular orbits for the star and binary. However, general triple-star systems are likely to have non-coplanar motion and eccentric orbits. Therefore, in this article, we extend our study to the case of three-dimensional non-circular motion of the star and binary system. In Sect.~2, we present the model and in Sect.~3 we discuss the circumstances that lead to the binary being mistaken by one or more planets. In Sect.~4, we compare the theoretical predictions with results obtained from simulations of hypothetical triple systems, we discuss the triple system HD 188753, and we (re)analyze the exoplanets discovered within binary systems. Finally, in Sect.~5 we present our conclusions.

We have studied a triple system composed of a star and a binary system. The star and binary system have eccentric and inclined orbits. This is an extension of earlier work where we assumed that the star and binary system are on circular coplanar orbits \citep{Morais&Correia2008}. We demonstrated that if we are unaware of the binary system’s presence (one or even both components may be unresolved for instance because they are faint M stars) we may then be led to believe that the star has one or even two planet companions. Although the radial velocity variations due to a binary are distinct from those due to planet(s), in practice, the measured effect depends on the instrument's precision and the observation time span. We have shown that because of the limited instrumental precision, we may only detect periodic terms with well separated frequencies hence we mistake these for planet(s). We also saw that, if the observation time span is shorter than the star's long-period motion around the binary's center of mass, we may not be able to resolve terms with nearby frequencies, which means that we cannot distinguish fake planet(s) from a binary. We have also demonstrated that the binary's effect is more likely to be mistaken for planet(s) when the radial velocity oscillations are composed of large dominant periodic terms with well separated frequencies. This is more likely to happen in the case of coplanar orbits observed equator-on. Moreover, we have seen that when the orbits are eccentric, the magnitude of the binary's effect can increase with respect to the circular case. Nevertheless, our model is valid for any triple system's configuration. We have presented an example of a binary system that affects a nearby star's motion. When the star's long-period motion has low to medium eccentricity, the binary can mimic planet(s) of 10 to 50 Earth masses, which are near the current detection limit. However, if the long-period motion has a high eccentricity we are more likely to detect multiple signals with very close frequencies and therefore reject the planet hypothesis. An exception occurs when the long period motion has a high eccentricity but its period is poorly constrained (due to a short observational time span) in which case we may mistake the binary for a planet of about 100 Earth masses. We showed that when the binary has an eccentric orbit it can mimic two planets with periods approximately in the ratio 2/1. Therefore we may be misled to think that we found planets in the 2/1 orbital resonance when we have an eccentric hidden binary. This is somehow analogous to the case studied by \citet{Anglada-Escude2010} where two planets on circular orbits in the 2/1 orbital resonance can be mistaken for a single planet with an eccentric orbit. However, our scenario is probably more realistic since planets in orbital resonances are not likely to have circular orbits. We propose that new planet detections in close binary systems, especially Earth-sized objects that are the targets of the planned search program EXPRESSO/CODEX \citep{EXPRESSOCODEX2007}, be checked carefully because they could indeed be artifacts caused by a hidden binary. This could be done by comparing fits to the data using (1) a model composed of the binary star system and a planet with (2) those obtained for a model composed of a hierarchical triple star system. If the fits of (2) are at least as good as (1), then the hidden binary hypothesis should be considered.

1012.2375

1012

1012.3079_arXiv.txt

Mean motion resonances are a common feature of both our own Solar System and of extrasolar planetary systems. Bodies can be trapped in resonance when their orbital semi-major axes change, for instance when they migrate through a protoplanetary disc. We use a Hamiltonian model to thoroughly investigate the capture behaviour for first and second order resonances. Using this method, all resonances of the same order can be described by one equation, with applications to specific resonances by appropriate scaling. We focus on the limit where one body is a massless test particle and the other a massive planet. We quantify how the the probability of capture into a resonance depends on the relative migration rate of the planet and particle, and the particle's eccentricity. Resonant capture fails for high migration rates, and has decreasing probability for higher eccentricities, although for certain migration rates, capture probability peaks at a finite eccentricity. More massive planets can capture particles at higher eccentricities and migration rates. We also calculate libration amplitudes and the offset of the libration centres for captured particles, and the change in eccentricity if capture does not occur. Libration amplitudes are higher for larger initial eccentricity. The model allows for a complete description of a particle's behaviour as it successively encounters several resonances. Data files containing the integration grid output will be available on-line. We discuss implications for several scenarios: (i) Planet migration through gas discs trapping other planets or planetesimals in resonances: We find that, with classical prescriptions for Type~I migration, capture into second order resonances is not possible, and lower mass planets or those further from the star should trap objects in first-order resonances closer to the planet than higher mass planets or those closer to the star. For fast enough migration, a planet can trap no objects into its resonances. We suggest that the present libration amplitude of planets may be a signature of their eccentricities at the epoch of capture, with high libration amplitudes suggesting high eccentricity (e.g., HD~128311). (ii) Planet migration through a debris disc: We find the resulting dynamical structure depends strongly both on migration rate and on planetesimal eccentricity. Translating this to spatial structure, we expect clumpiness to decrease from a significant level at $e\lesssim 0.01$ to non-existent at $e\gtrsim 0.1$. (iii) Dust migration through PR drag: We predict that Mars should have its own resonant ring of particles captured from the zodiacal cloud, and that the capture probability is $\lesssim 25\%$ that of the Earth, consistent with published upper limits for its resonant ring. To summarise, the Hamiltonian model will allow quick interpretation of the resonant properties of extrasolar planets and Kuiper Belt Objects, and will allow synthetic images of debris disc structures to be quickly generated, which will be useful for predicting and interpreting disc images made with ALMA, Darwin/TPF or similar missions.

Mean motion resonances (MMRs) occur when two objects' orbital periods are close to a ratio of two integers, and a particular combination of orbital angles, the resonant argument, is librating. Examples in the Solar System include Neptune and Pluto (3:2 resonance) and the inner Galilean moons of Jupiter (4:2:1 Laplace resonance). There are also now numerous examples of suspected or confirmed MMRs in extrasolar planetary systems (e.g., GJ~876~b and c in a 2:1 resonance, \citealt{2001ApJ...551L.109L}). Mean motion resonances also occur between planets and small dust particles, as seen in the Earth's resonant dust ring~\citep{1994Natur.369..719D}. Some extrasolar debris discs, such as Vega, show evidence of non-axisymmetric clumps \citep{1998Natur.392..788H,2002ApJ...569L.115W}, and several authors have attempted to model these as arising from a planet's resonant perturbations \citep[e.g.,][]{2003ApJ...588.1110K,2003ApJ...598.1321W}. Although resonant orbits occupy only a small volume of phase space, they are common because of a locking mechanism which can preserve the resonance once attained. If a particle's or planet's orbital semi-major axis changes due to non-conservative forces, the bodies can approach a resonance and then remain trapped there even under further action of the non-conservative forces. The associated orbital angular momentum change then drives an eccentricity change, while the semi-major axis ratio remains approximately fixed at the resonance. There are many mechanisms by which such a semi-major axis change can be driven. Early work looked at the tidal evolution of satellite orbits \citep{1965MNRAS.130..159G}. In a protoplanetary disc, planets can migrate by tidal interaction with the gas disc \citep[see][for a recent review]{2009AREPS..37..321C}, and small planetesimals by aerodynamic drag \citep{1977MNRAS.180...57W}. In a gas-depleted debris disc, planets can migrate by gravitational scattering of planetesimals \citep{1984Icar...58..109F,2009Icar..199..197K}. Interplanetary dust drifts towards the Sun under the influence of Poynting-Robertson (PR) drag \citep{1979Icar...40....1B}, and large bodies can be moved more slowly by the Yarkovsky effect \citep{2006AREPS..34..157B}. At the end of a star's main-sequence lifetime, planetesimals can experience aerodynamic drag as the star loses mass \citep{2010ApJ...715.1036D}. Moreover, for a planet orbiting the secondary component in a binary system, formation of a disc following mass transfer from the primary to the secondary could trigger renewed planet migration \citep{2010arXiv1001.0581P}. Resonance capture has been studied by several authors, going back to \cite{1965MNRAS.130..159G}. The regime of adiabatic migration, where the migration timescale is much longer than the resonant argument's libration timescale, has been studied extensively analytically using a Hamiltonian model \citep[e.g.,][]{1982CeMec..27....3H,1984CeMec..32..127B}. With adiabatic migration, capture is certain if the particle has an eccentricity below a critical value, and probabilistic with decreasing probability as eccentricity increases beyond this. Rapid migration was studied using full N-body models by \citet[henceforth W03]{2003ApJ...598.1321W} for the case of a planet migrating into a planetesimal disc, and using the Hamiltonian model by \citet[henceforth Q06]{2006MNRAS.365.1367Q} for general migration scenarios. Q06 obtained capture probability as a function of migration rate and eccentricity for the Hamiltonian containing a single resonant term, and went on to consider the role played by additional resonant terms in affecting capture probability. Such terms can be important when the planet is eccentric. In this paper we extend this work in a different direction, and using the Hamiltonian model with a single resonant term we calculate capture probabilities, libration amplitudes and offsets for particles that are captured, and eccentricity jumps for those that pass through the resonance without capture. We validate the model against the numerical integrations of W03 and Q06. We also discuss the application of the model to various migration scenarios, and discuss previous studies in its light, in particular those of \cite{2008A&A...480..551R} who investigated capture of planetesimals by a migrating planet, and \cite{1994Natur.369..719D} who studied the formation of the Earth's resonant dust ring, both using N-body integrations. We should like to emphasize the role eccentricity can play in affecting capture probabilities and libration amplitudes, which, while long understood in the Solar System \citep[e.g.,][]{1999ssd..book.....M}, is sometimes neglected in studies of extrasolar planets and discs. The Hamiltonian model we use has several advantages over N-body simulations: (1) it allows some results to be derived analytically; (2) it is faster to integrate numerically than the 3-body problem; (3) all resonances of the same order reduce to a Hamiltonian of the same form, with fewer free parameters than the three-body problem. Once a suite of numerical integrations of the Hamiltonian model is performed, it can be applied to any system, without the need for running a different N-body integration every time the system parameters are changed. The plan of this paper is as follows: In \S\,2 we describe the dimensionless Hamiltonian model. Readers interested in the details may read the Appendix which contains the mathematical derivation. In \S\,3 we summarise how physical parameters relate to the dimensionless parameters for test particles. In \S\,4 we describe the results of our numerical integrations. In \S\,5 we compare the Hamiltonian model to N-body simulations. In \S\,6 we discuss applications. \S\,7 summarises our work.

We have systematically investigated the Hamiltonian model for capture of a test particle into first and second order resonances with a planet. The model reduces the full complexity of the restricted three body problem to only one degree of freedom, the resonant angle $\theta$, and two parameters, proportional to the particle's eccentricity and the migration rate. Only the relative migration rate of the two bodies, in the form $\dot{a}_1/a_1-\dot{a}_2/a_2$, is important. External and internal resonances behave the same way, but with different proportionality constants. We confirmed previous work showing that capture into resonance is certain at low (rescaled) eccentricities and migration rates, possible with decreasing probability at high eccentricities, and impossible at low eccentricities and high migration rates. As eccentricity increases, the transition from certain capture at low migration rate to impossible capture at high migration rate broadens. At higher eccentricities, capture is possible with faster migration than at lower eccentricities. This effect is more pronounced for second order resonances. We have also found the libration amplitudes and centres of the resulting resonant motion. The libration centres are offset from $\theta=\pi$, the centre in the absence of migration, by an amount increasing with the migration rate, agreeing with previous work. In addition, we have found that the offset is almost independent of eccentricity for first order resonances, except at high eccentricities. Libration amplitudes are small if capture occurs at low eccentricity and low migration rate. They are somewhat larger for migration rates just less than critical. Very large libration amplitudes ($\gtrsim 90^\circ$) can be attained if the initial eccentricity was very high. We have also found the jumps in eccentricity when a particle encounters a resonance but is not captured. In the case of slow migration, the eccentricity is always driven down, in agreement with adiabatic theory. When the migration is fast and the eccentricity low, eccentricity jumps up. However, when the migration is fast and eccentricity high, eccentricity can jump either up or down. We checked the results of the model against the N-body simulations of \cite{2003ApJ...598.1321W}, finding excellent quantitative agreement for capture probabilities, and qualitative agreement for libration amplitudes and offsets. We found that accounting for the particles' eccentricity is necessary to fully explain the dependence of capture probability on migration rate. While the effect of eccentricity has long been understood in the context of the Solar System, its implications for extrasolar planets are less well explored, and for low mass planets in particular its effects are important. We have applied our model to several situations in which planet or particle migration is likely to occur. We can easily determine whether capture occurs for planets migrating in Type~I or Type~II regimes. The pre-capture eccentricity can be constrained by the present libration amplitude, with higher eccentricities giving higher libration amplitudes. For planets migrating through a gas disc, a non-zero eccentricity prior to capture can lead to large libration amplitudes such as those seen in the HD~128311 system. We find that, if planetary eccentricity can be raised to $e\gtrsim 0.1$ for Jupiter-mass planets, for example by planet-disc interaction, the resonant capture process by itself will result in high libration amplitudes without the need to invoke extra mechanisms such as turbulent torques. For lower mass planets the necessary eccentricity is lower. Also, very low libration amplitudes (e.g., GJ~876) suggest a low eccentricity during the capture process. We also predict, based on classical formulae for migration rates, that lower mass planets will be found in higher $j$ first order resonances than higher mass planets, and planets migrating in the type~I regime will be moving too fast to capture smaller particles into any second-order resonance. The model may also be useful for population synthesis of multi-planet systems where it is necessary to account for planet--planet interactions without recourse to full N-body integrations. We then discussed debris disc structure. A planet that has migrated into a disc will impose clumpy structure on a dynamically cold disc, but the clumps will be at a lower level and smeared out for migration into excited discs. For dust migrating under PR drag, the model can explain the structure of the Earth's resonant ring reasonably well. We predict that Mars has a dust ring at a level $\lesssim 25\%$ that of Earth's, consistent with observed upper limits. The data on capture probabilities, libration amplitudes and offsets, and eccentricity jumps, are available on-line at \url{http://www.ast.cam.ac.uk/~ajm233/} and at the journal website and may be used provided that this work is cited.

1012.3079

1012

1012.4373_arXiv.txt

In the faint star KIC\,9700322 observed by the {\it Kepler} satellite, 76 frequencies with amplitudes from 14 to 29000\,ppm were detected. The two dominant frequencies at 9.79 and 12.57\,d$^{-1}$ (113.3 and 145.5 $\mu$Hz), interpreted to be radial modes, are accompanied by a large number of combination frequencies. A small additional modulation with a 0.16\,d$^{-1}$ frequency is also seen; this is interpreted to be the rotation frequency of the star. The corresponding prediction of slow rotation is confirmed by a spectrum from which $v \sin i = 19 \pm 1$\,km\,s$^{-1}$ is obtained. The analysis of the spectrum shows that the star is one of the coolest $\delta$\,Sct variables. We also determine T$_{\rm eff}$\,=\,6700~$\pm$~100 K and $\log g$\,=\,3.7~$\pm$~0.1, compatible with the observed frequencies of the radial modes. Normal solar abundances are found. An $\ell=2$ frequency quintuplet is also detected with a frequency separation consistent with predictions from the measured rotation rate. A remarkable result is the absence of additional independent frequencies down to an amplitude limit near 14\,ppm, suggesting that the star is stable against most forms of nonradial pulsation. A low frequency peak at 2.7763\,d$^{-1}$ in KIC\,9700322 is the frequency difference between the two dominant modes and is repeated over and over in various frequency combinations involving the two dominant modes. The relative phases of the combination frequencies show a strong correlation with frequency, but the physical significance of this result is not clear.

The {\it Kepler Mission} is designed to detect Earth-like planets around solar-type stars \citep{Koch2010}. To achieve that goal, {\it Kepler} is continuously monitoring the brightness of over 150\,000 stars for at least 3.5\,yr in a 105 square degree fixed field of view. Photometric results show that after one year of almost continuous observations, pulsation amplitudes of 5\,ppm are easily detected in the periodogram for stars brighter than $V = 10$\,mag, while at $V = 14$\,mag the amplitude limit is about 30\,ppm. Two modes of observation are available: long cadence (29.4-min exposures) and short-cadence (1-min exposures). With short-cadence exposures \citep{gilliland2010} it is possible to observe the whole frequency range seen in $\delta$\,Sct stars. Many hundreds of $\delta$\,Sct stars have now been detected in {\it Kepler} short-cadence observations. This is an extremely valuable homogeneous data set which allows for the exploration of effects never seen from the ground. Ground-based observations of $\delta$\,Sct stars have long indicated that the many observed frequencies, which typically span the range $5-50$\,d$^{-1}$, are mostly $p$~modes driven by the $\kappa$-mechanism operating in the He\,\textsc{ii} ionization zone. The closely-related $\gamma$\,Dor stars lie on the cool side of the $\delta$\,Sct instability strip and have frequencies below about 5\,d$^{-1}$. These are $g$~modes driven by the convection-blocking mechanism. Several stars exhibit frequencies in both the $\delta$\,Sct and $\gamma$\,Dor ranges and are known as hybrids. \cite{Dupret2005} have discussed how the $\kappa$ and convective blocking mechanisms can work together to drive the pulsations seen in the hybrids. The nice separation in frequencies between $\delta$\,Sct and $\gamma$\,Dor stars disappears as the amplitude limit is lowered. {\it Kepler} observations have shown that frequencies in both the $\delta$\,Sct and $\gamma$\,Dor regions are present in almost all of the stars in the $\delta$\,Sct instability strip \citep{Grigahcene2010}. In other words, practically all stars in the $\delta$\,Sct instability strip are hybrids when the photometric detection level is sufficiently low. \begin{figure*} \centering \includegraphics[bb=120 200 720 550,width=14cm]{fig1.eps} \caption{A sample of the {\it Kepler} light curve covering 0.8\,d. The multifrequency fit is shown as a solid curve. The pattern shown here roughly repeats every 0.72\,d.} \label{lcurve} \end{figure*} \begin{figure*} \centering \includegraphics[width=8cm]{fig2a.ps} \includegraphics[width=8cm]{fig2b.ps} \caption{Two portions of observed spectrum matched to a model with $T_{\rm eff}$\,=\,6700, $\log g$\,=\,3.7 (red line). In the left panel we show the region around H$\beta$ that is sensitive to temperature, and in the right panel the region around Mg{\sc i} triplet sensitive to gravity.} \label{spectrum} \end{figure*} Statistical analyses of several $\delta$\,Sct stars observed from the ground have already shown that the photometrically observed frequencies are not distributed at random, but that the excited nonradial modes cluster around the frequencies of the radial modes over many radial orders. The observed regularities can be partly explained by modes trapped in the stellar envelope \citep{BregerLenzPamyatnykh2009}. This leads to regularities in the observed frequency spectra, but not to exact equidistance. In examining the {\it Kepler} data for $\delta$\,Sct stars we noticed several stars in which many exactly equally-spaced frequency components are present. There are natural explanations for nearly equally spaced frequency multiplets such as harmonics and non-linear combination frequencies. In some of these stars, however, these mechanisms do not explain the spacings. In these stars there is often more than one exact frequency spacing and these are interleaved in a way which so far defies any explanation. Some examples of equally-spaced frequency components which remain unexplained are known from ground-based observations. The $\delta$\,Sct star 1\,Mon has a frequency triplet where the departure from equidistance is extremely small: only $0.000079 \pm 0.000001$\,d$^{-1}$ (or $0.91 \pm 0.01$\,nHz), yet the frequency splitting cannot be due to rotation because $\ell = 0$ for the central component and $\ell = 1$ for the other two modes \citep{BregerKolenberg2006,BalonaBartlett2001}. In the $\beta$\,Cep star 12\,Lac, there is a triplet with side peaks spaced by 0.1558 and 0.1553 d$^{-1}$. The probability that this is a chance occurrence is very small, yet photometric mode identification shows that two of these modes are $\ell = 2$ and the third is $\ell = 1$. This is therefore not a rotationally split triplet either \citep{Handler2006}. One solution to these puzzling equally-spaced frequencies could be non-linear mode interaction through frequency locking. \cite{Buchler1997} show that frequency locking within a rotationally split multiplet of a rapidly rotating star (150 to 200 km\,s$^{-1}$) could yield equally-spaced frequency splitting, which is to be contrasted to the prediction of linear theory where strong departures from equal splitting are expected. In this paper we present a study of the $\delta$\,Sct star KIC\,9700322 (RA = 19:07:51, Dec = 46:29:12 J2000, Kp = 12.685). There are two modes with amplitudes exceeding 20000\,ppm and several more larger than 1000\,ppm. The equal frequency spacing is already evident in these large amplitude modes. This star does not fall in the unexplained category discussed above. It is, however, a remarkable example of a star in which combination frequencies are dominant. The star has a large pulsational amplitude which can easily be observed from the ground. It was found to be variable in the "All Sky Automated Survey" \citep{Pigulski2008}, where it is given the designation ASAS 190751+4629.2. It is classified as a periodic variable (PER) with a frequency of 7.79 d$^{-1}$. This is the 2 d$^{-1}$ alias of the main frequency (9.79 d$^{-1}$), which is determined below from the {\it Kepler} data. The {\it Kepler} data is, of course, not affected by daily aliasing. It was also examined during the "Northern Sky Variability Survey" \citep{Wozniak2004} with up to two measurements per night. Due to the short periods of the star, the 109 points of NSVS 5575265 were not suitable for a comparison with our results.

Although this star was selected because of its very clear exactly equal frequency spacing, it turns out that the frequency spacing is explained as simple combination frequencies arising from non-linearities of the oscillation. This is different from another class of $\delta$\,Sct stars in the {\it Kepler} database which also show exact frequency spacings, but in a manner which is not at present understood. Examples of this strange class will be presented in a separate paper. What makes KIC\,9700322 interesting is the remarkable way in which the large number of frequencies are related to the two main frequencies, $f_1$ and $f_2$. This behaviour is very similar to the high amplitude $\delta$\,Sct star KIC\,9408694, also discovered in the {\it Kepler} database. The frequency patterns together with their amplitudes permit us to identify the different frequencies and to provide physical interpretations. \subsection{The dominant radial modes} \label{sec:radmodes} The period ratio of $f_1$ and $f_2$ is 0.779. This is close to the expected period ratio for fundamental and first overtone radial pulsation. The pulsation amplitudes of $\delta$\,Sct stars increase with decreasing rotation, e.g., see Fig.\,5 of \cite{Breger2000}. Furthermore, high amplitudes occur mainly in slowly rotating, radial pulsators: in fact, the high-amplitude $\delta$\,Sct (HADS) subgroup is defined on the basis of peak-to-peak amplitudes in excess of 0.3 mag. Nevertheless, a rigid separation between radial HADS and lower-amplitude nonradial $\delta$\,Sct stars does not exist. Dominant radial modes with amplitudes smaller than 0.3 mag have previously been found. Examples of EE Cam \citep{bregerrucinskireegen2007} and 44 Tau \citep{LenzPamyatnykhZdravkovBreger2010}. The situation might be summarized as follows: Dominant radial modes occur only in slowly rotating stars. Since KIC\,9700322 is sharp-lined ($v \sin i = 19$\,km\,s$^{-1}$) and presumably also a slow rotator, it follows this relationship. The presence of two dominant radial modes with amplitudes less than the HADS limits of peak-to-peak amplitudes of 0.3 mag is not unusual. The measured $v \sin i$ value supports the interpretation of the observed 0.1597\,d$^{-1}$ peak as the rotational frequency. In fact, both dominant modes have very weak side lobes with spacings of exactly the rotational frequency. The side lobes are very weak: for $f_1$ and $f_2$ the amplitudes are only 0.0018 and 0.0011 of the central peak amplitudes. We interpret this as a very small modulation of the amplitudes with rotation. An alternate explanation in terms of rotational splitting of nonradial modes is improbable because rotational splitting does not lead to exact frequency separation unless there is frequency locking due to resonance. Also, the extreme amplitude ratios tend to favour the interpretation in terms of amplitude modulation. Based on this mode identification assumption we investigated representative asteroseismic models of the star. We have used two independent numerical packages: the first package consisted of the current versions of the Warsaw-New Jersey stellar evolution code and the Dziembowski pulsation code \citep{wd1977,wdgd1992}. The second package is composed by the evolutionary code \cesam\ \citep{Morel97}, and the oscillation code \filou\ \citep{Sua02thesis, Sua06rotcel}. Both pulsation codes consider second-order effects of rotation including near degeneracy effects. The period ratio between the first radial overtone and fundamental mode mainly depends on metallicity, rotation and stellar mass. Moreover, the radial period ratio also allows for inferences on Rosseland mean opacities as shown in \cite{LenzPamyatnykhZdravkovBreger2010}. Indeed, an attempt to reproduce the radial fundamental and first overtone mode at the observed frequencies with the first modelling package revealed a strong dependence on the choice of the chemical composition and the OPAL vs. OP opacity data \citep{IglesiasRogers1996,Seaton2005}. The best model found in this investigation was obtained with OP opacities and increased helium and metal abundances. Unfortunately, this model ($T_{\rm eff} = 7400$\,K, $\log L/{\rm L}_\odot = 1.27$, $\log g = 3.87$, 2M$_\odot$) is much hotter than observations indicate. The disagreement in effective temperature indicates that this model is not correct despite the good fit of the radial modes. \begin{figure*} \centering \includegraphics[bb=20 17 800 560,width=170mm,clip]{fig7.eps} \caption{Top panel: Amplitude spectrum of those frequencies which are numerical combinations of the two dominant modes only. This demonstrates the richness of the combinations. Notice that the patterns are more complex than a simple Fourier sum. The middle panel demonstrates that some of those combination frequencies show further splittings with $f_3$ (rotation). The bottom panel shows an additional set of five frequencies together with the combinations of the set with $f_1$ and $f_2$. } \label{pattern} \end{figure*} As an additional test, by adopting the radial linear nonadiabatic models developed by \citet{mp98} and \citet{m04}, we are able to reproduce the values of the two dominant frequencies with pulsation in the fundamental and first overtone modes, but with a lower period ratio (0.770) than observed. The best fit solution obtained with these models, for an effective temperature consistent with the spectroscopic determination and assuming solar chemical composition, corresponds to: $M=1.65 M_{\odot}$, $\log L/{\rm L}_\odot = 1.1$, $T_{\rm eff}=6700$ K, $\log g = 3.83$. We notice that for this combination of stellar parameters, both the fundamental and the first overtone mode are unstable in these models. Moreover, looking at the Main Sequence and post-Main Sequence evolutionary tracks in the gravity versus effective temperature plane, as reported in Fig.~4 of \citet{c10}, the solution $T_{\rm eff}$ = 6700 K, $\log g = 3.83$ is consistent with a $1.65 M_{\odot}$ stellar mass. However, as already noted, the period ratio in our models is lower than the observed value. To resolve this discrepancy, the possibility of low metallicity and rotation effects was examined in more detail with the second modeling package. Models between $T_{\rm eff} = 6200$\,K and $8600\,$K with masses between 1.2 and 1.76$\msun$, were found to represent a good fit of $f_1$ and $f_2$ as radial fundamental and first overtone, respectively. The best fit with the observations was found for $M=1.2\,\msun$ models computed with $\amlt=0.5$, ${\rm d}_{ov}=0.1$, and a metallicity of [Fe/H]\,=\,-0.5~dex. Such a low value for the convection efficiency is in good agreement with the predictions by \citet{Casas06} for $\delta$~Sct stars, based on their non-adiabatic asteroseismic analysis. All these parameters roughly match the general characteristics of the $\delta$\,Sct stars with dominant radial modes and large amplitudes, despite being in the limit in metallicity. The $P_1/P_0$ period ratios predicted by these models (which simultaneously fit $P_0$) are near 0.775, which is lower than the observed ratio, 0.779. A period ratio of 0.775 is also obtained by adopting the radial linear nonadiabatic models by \citet{m04} at $Z=0.006$, according to which the best fit solution with effective temperature consistent with the spectroscopic determination, corresponds to $Z=0.006$, $Y=0.25$, $M=1.5$, $\log L/L_{\odot}=1.04$, $T_e=6700$ K, $\log g = 3.83$. Again the fundamental and first overtone modes are predicted to be simultaneoulsy unstable for this parameter combination. We explored the possibility that such a discrepancy might be due to rotation effects, particularly second-order distortion effects, as discussed by \citet{Sua06pdrotI} and \citet{Sua07pdrotII}. These investigations analyze theoretical Petersen Diagrams including rotation effects (Rotational Petersen Diagrams, hereafter RPDs), and show that $P_1/P_0$ ratios increase as stellar surface rotation increases. The rotation rate derived from observations is slightly below 25\,km\,s$^{-1}$ (see section\,\ref{sec:quintuplet}). At such rotation rates near degeneracy effects on the period ratio are small (less than 0.001 in $P_1/P_0$). However, when non-spherically symmetric components of the centrifugal force are considered, near-degeneracy effects may be larger, around 0.0025, causing the presence of wriggles in the RPDs (see Fig. 5 in \citet{Sua07pdrotII} and Fig. 6 in \citet{Pamyatnykh2003}). Such effects are even more significant for rotational velocities higher than 40 - 50\,km\,s$^{-1}$. Consequently, near-degeneracy effects may help to decrease the discrepancy between the observed period ratio and the slightly lower values predicted by the models. If the star had a low metal abundance (close to Pop.~II), a detailed analysis of RPDs might have provided an independent estimate of the true rotational velocity (and thereby of the angle of inclination). However, the spectroscopic analysis indicates that the star has a solar abundance. KIC 9700322 therefore represents a challenge to asteroseismic modeling, since it appears impossible to reproduce all observables simultaneously with standard models. \subsection{The combination frequencies} We have already shown that the 50+ detected frequency peaks can be explained by simple combinations of the two dominant modes and the rotational frequency. Several different nonlinear mechanisms may be responsible for generating combination frequencies between two independent frequencies, $\nu_1$ and $\nu_2$. For example, any non-linear transformation, such as the dependence of emergent flux variation on the temperature variation ($L = \sigma T^4$) will lead to cross terms involving frequencies $\nu_1 + \nu_2$ and $\nu_1 - \nu_2$ and other combinations. The inability of the stellar medium to respond linearly to the pulsational wave is another example of this effect. Combination frequencies may also arise through resonant mode coupling when $\nu_1$ and $\nu_2$ are related in a simple numerical way such as $2\nu_1 \approx 3\nu_2$. The interest in the combination frequencies derives from the fact that their amplitudes and phases may allow indirect mode identification. For nonradial modes, some combination frequencies are not allowed depending on the parity of the modes \citep{Buchler1997} which could lead to useful constraints on mode identification. Since $f_1$ and $f_2$ in KIC\,9700322 are both presumably radial, there are no such constraints. The identification of $f_1$ and $f_2$ with radial modes allows us to investigate the properties of the Fourier parameters of the combination modes with the aim to disentangle less obvious cases and/or solutions with a smaller number of combination terms. \cite{Buchler1997} show that a resonance of the type $f_c = n_1f_1 + n_2f_2$ leads to a phase $\phi_r = \phi_c - (n_1\phi_1 + n_2\phi_2)$. In the same way we may define the amplitude ratios $A_r = A_c/(A_1A_2)$. To investigate how $\phi_r$ and $A_r$ behave with frequency, we first need the best estimate of the parent frequencies. We obtained these by non-linear minimization of a truncated Fourier fit involving $f_1, f_2$ and all combination frequencies up to the 4$^{\rm th}$ order. The best values are $f_1 = 9.792514$ and $f_2 = 12.568811$\,d$^{-1}$. The resulting amplitude and phases are shown in Table\,\ref{cfreq} together with the values of $\phi_r$ and $A_r$. The phases were calculated relative to BJD\,245\,5108.3849 which corresponds to the midpoint of the observations. \begin{table} \caption{Best fitting amplitudes, $A_c$ (ppm), and phases $\phi_c$ (radians), for the parent frequencies and their combination frequencies up to fourth order. Relative phases, $\phi_r$ and amplitudes, $A_r$ are also shown.} \label{cfreq} \begin{flushleft} \begin{tabular}{rrrrrr} \hline $(n_1, n_2)$ & $f_c$ & $A_c$ & $\phi_c$ & $\phi_r$ & $A_r$ \\ \hline ( 1, 0) & 9.792514 & 27271 & 2.710 & & \\ ( 0, 1) & 12.568811 & 29443 & -1.313 & & \\ ( 2, 0) & 19.585028 & 2225 & -0.110 & 0.751 & 0.002553 \\ ( 1, 1) & 22.361325 & 4898 & 1.954 & 0.557 & 0.005619 \\ ( -1, 1) & 2.776297 & 2636 & 0.826 & -1.431 & 0.003024 \\ ( 0, 2) & 25.137622 & 2663 & -2.317 & 0.310 & 0.003055 \\ ( 3, 0) & 29.377542 & 192 & -2.749 & 1.684 & 0.000222 \\ ( 2, 1) & 32.153839 & 476 & -0.835 & 1.339 & 0.000547 \\ ( 2, -1) & 7.016217 & 633 & 0.127 & -0.324 & 0.000726 \\ ( 1, 2) & 34.930136 & 536 & 0.202 & 0.119 & 0.000615 \\ ( -1, 2) & 15.345108 & 502 & 0.525 & -0.418 & 0.000577 \\ ( 0, 3) & 37.706433 & 329 & -0.658 & -3.000 & 0.000377 \\ ( 4, 0) & 39.170056 & 12 & 0.844 & 2.567 & 0.000014 \\ ( 3, 1) & 41.946353 & 22 & -2.748 & 2.999 & 0.000026 \\ ( 3, -1) & 16.808731 & 69 & -2.521 & 0.598 & 0.000080 \\ ( 2, 2) & 44.722650 & 114 & 0.869 & -1.924 & 0.000132 \\ ( -2, 2) & 5.552594 & 87 & -2.717 & -0.951 & 0.000101 \\ ( 1, 3) & 47.498947 & 84 & -1.941 & -0.710 & 0.000097 \\ ( -1, 3) & 27.913919 & 205 & 0.199 & 0.568 & 0.000236 \\ ( 0, 4) & 50.275244 & 84 & -1.061 & -2.088 & 0.000097 \\ \hline \end{tabular} \end{flushleft} \end{table} \begin{figure} \centering \includegraphics[scale=0.6]{fig8.ps} \caption{The relative amplitudes, $A_r$, and phases, $\phi_r$ (radians) as a function of frequency (d$^{-1}$). The dotted lines show the location of the parent frequencies.} \label{comb} \end{figure} Fig.\,\ref{comb} shows how $A_r$ and $\phi_r$ vary with frequency. From the figure we note that $A_r$ is largest for $f_1 + f_2$, $2f_1$, $2f_2$ and $f_2-f_1$ and very small for the rest. It is also interesting that $\phi_r$ is a relatively smooth function of frequency, being practically zero in the vicinity of the parent frequencies, decreasing towards smaller frequencies and increasing towards higher frequencies. This result is almost independent of the choice of $f_1$ and $f_2$. The standard deviation of $f_1$ and $f_2$ is 0.0001~d$^{-1}$ using the \cite{Montgomery1999} formula. One may arbitrarily adjust $f_1$ and $f_2$ in opposite directions by this value, and using the corresponding calculated values of the combination frequencies, fit the data to obtain new phases. The resulting $\phi_r$ versus frequency remains monotonic, but the slope does change. The smooth monotonic nature of the $\phi_r$ versus frequency diagram remains even for a change of ten times the standard deviation in opposite directions for $f_1$ and $f_2$ and for arbitrary changes in epoch of phase zero. The result is clearly robust to observational errors, but it is not clear what physical conclusions may be derived from this result. The behaviour is certainly not random and must have a physical basis. Note that for simple trigonometric products, $\phi_r$ will always be zero. Finally, we note that the amplitudes of the combination modes relative to the amplitudes of their parents can be compared with values detected in the star 44\,Tau \citep{BregerLenz2008}. They agree to a factor of two or better, suggesting that KIC\,9700322 is not unusual in this regard, just more accurately studied because of the {\it Kepler} data. \subsection{The quintuplet} \label{sec:quintuplet} In addition to the quintuplet structure around the two dominant modes another quintuplet with different properties is present in KIC 9700322 (see the listing of $f_4$ to $f_8$ in Table\,\ref{quin}). The average spacing between the frequencies in this quintuplet is slightly smaller than the rotational frequency (0.1338\,d$^{-1}$ vs. 0.1597\,d$^{-1}$). This makes this quintuplet different from the quintuplet structures found around the two dominant modes, which exhibit a spacing that corresponds exactly to the rotation frequency. Moreover, the distribution of amplitudes within the third quintuplet is fundamentally different to the patterns around $f_1$ and $f_2$. The given characteristics support an interpretation of the quintuplet as an $l$ = 2 mode. The location of the quintuplet near the centre in between the radial fundamental and first overtone mode rules out pure acoustic character. Consequently, the observed quintuplet consists of mixed modes with considerable kinetic energy contribution from the gravity-mode cavity. For such modes theory predicts a smaller (and more symmetrical) rotational splittings compared to acoustic modes due to different values of the Ledoux constant $C_{nl}$. Using the framework of second order theory \citep{wdgd1992} we determined the equatorial rotation rate which provides the best fit of the observed quintuplet with an $\ell=2$ multiplet. The best results were obtained for an equatorial rotation rate of 23\,km\,s$^{-1}$. This is only slightly higher than the observed $v \sin i$ value of 19\,km\,s$^{-1}$, and therefore indicates a near-equator-on-view. The Ledoux constant, $C_{nl}$, of the $\ell=2$ quintuplet is 0.164. For quadrupole modes $C_{nl}$ ranges between $\approx$0.2 for pure gravity modes to smaller values for acoustic modes. With (1 - $C_{nl}$) = 0.836 this leads to a rotational frequency, $\nu_{rot}=\frac{\Omega}{2\Pi}$, of around 0.16 d$^{-1}$. Consequently, this theoretical result confirms the interpretation of $f_3$ as a rotational feature and of the quintuplet as $l$ = 2 modes. Further support is provided by the fact that we see various combinations of the quintuplet with $f_1$ and $f_2$. Moreover, the location of the quintuplet allows us to determine the extent of overshooting from the convective core. In the given model we obtained $\alpha_{ov}=0.13$ but the uncertainties elaborated in Section~\ref{sec:radmodes} currently prevent an accurate determination. \subsection{Further discussion} A remarkable aspect of the star is the fact that so few pulsation modes are excited with amplitudes of 10\,ppm or larger. In the interior of an evolved $\delta$\,Sct star, even high-frequency $p$~modes behave like high-order $g$~modes. The large number of spatial oscillations of these modes in the deep interior leads to severe radiative damping. As a result, nonradial modes are increasingly damped for more massive $\delta$\,Sct stars, which explains why high-amplitude$\delta$\,Sct stars pulsate in mostly radial modes and why in even more massive classical Cepheids nonradial modes are no longer visible. In general, we do not expect the frequencies in the $\delta$\,Sct stars observed by {\it Kepler} to be regularly spaced because, unlike ground-based photometry, the observed pulsation modes are not limited to small spherical harmonic degree, $l$. For the very low amplitudes detected by {\it Kepler} we may expect to see a large number of small-amplitude modes with high $l$. The observed amplitudes decreases very slowly with $l$ and, all things being equal, a large number of modes with high $l$ might be expected to be seen in $\delta$\,Sct and other stars \citep{Balona1999}. The $\delta$\,Sct stars HD\,50844 \citep{Poretti2009} and HD\,174936 \citep{Hernandez2009} observed by {\it CoRoT} show many hundreds of closely-spaced frequencies and may be examples of high-degree modes. The relatively small number of independent frequencies detected in KIC\,9700322 stands in strong contrast to the two stars observed by {\it CoRoT}. It should be noted that, unlike many $\delta$\,Sct stars observed by {\it Kepler}, KIC\,9700322 does not have any frequencies in the range normally seen in $\gamma$\,Dor stars. The only strong frequencies in this range are a few combination frequencies. Although we have identified significant frequencies below 0.5\,d$^{-1}$, it is not possible at this stage to verify whether these are due to the star or instrumental artefacts. At present, we do not understand why low frequencies are present in so many $\delta$\,Sct stars. Regularities in the frequency spacing due to combination modes have already been observed from the ground even in low amplitude $\delta$\,Sct stars. An example is the star 44\,Tau \citep{BregerLenz2008}. Fig.\,2 of \cite{BregerLenzPamyatnykh2009} demonstrates that all the observed regularities outside the $5-13$\,d$^{-1}$ range are caused by combination modes. For combination modes the frequency spacing must be absolutely regular within the limits of measurability. This is found for KIC\,9700322.

1012.4373

1012

1012.1626_arXiv.txt

This splinter session was devoted to reviewing our current knowledge of correlated X-ray and radio emission from cool stars in order to prepare for new large radio observatories such as the EVLA. A key interest was to discuss why the X-ray and radio luminosities of some cool stars are in clear breach of a correlation that holds for other active stars, the so-called G\"udel-Benz relation. This article summarizes the contributions whereas the actual presentations can be accessed on the splinter website\footnote{http://cxc.harvard.edu/cs16xrayradio/}.

\begin{figure}[t!] \includegraphics[angle=0,width=9cm]{forbrich_j_fig1.eps} \caption{Radio vs. X-ray correlation. Different symbols in the upper right part of the figure refer to different types of active stars. The letters in the lower left show the loci of solar flares (luminosities averaged over the duration of the X-ray flare; see \citealt{benz94} for details). The red and blue triangles show examples of stellar flares from M dwarfs \citep{guedel96} and RS CVn binaries \citep{osten04}.} \label{fig1} \end{figure} X-rays and radio emission are excellent diagnostic probes to study energy release in magnetized stellar coronae. Solar observations have been key to deciphering the plethora of phenomena seen in these wavelength ranges. In brief, X-rays trace the presence of dense, hot (million-degree) plasma trapped in closed coronal magnetic fields, heated by processes that are still not fully understood. Radio observations, in contrast, probe both thermal atmospheric components from the chromosphere to the corona, and populations of non-thermal, accelerated electrons, typically residing in low-density, open or closed coronal magnetic fields. A solar-stellar analogy is, however, complicated by phenomenology in {\it magnetically active stars} that is, at first sight, not present in the Sun. X-ray emission becomes much stronger toward more active stars, most likely as a result of increased surface coverage with active regions; however, the average, ``characteristic'' temperature of the corona increases along with the coronal luminosity (e.g., \citealt{schrijver84, guedel97}), a trend that requires additional physical explanation. At radio wavelengths, magnetically active stars also show a different face. Solar radio emission in the 1-10~GHz range is dominated by bremsstrahlung from various chromospheric/transition region levels and by optically thick gyroresonance emission from coronal layers above magnetic active regions. In contrast, observed radio brightness temperatures and radio spectra from active stars indicate gyrosynchrotron radiation from electrons with much higher energies than typically present in the solar atmosphere. The Sun occasionally features gyrosynchrotron emission, often accompanied by a variety of coherent radiation types (see Sect.~\ref{sect_benz}), but such radiation is confined to episodes of flaring. Magnetically active stars reveal further radio properties lacking any clear solar analogy such as extremely large coronal structures with size scales of order a stellar radius and more (e.g., \citealt{benz98, mutel98, peterson10}). In contrast, X-ray coronae tend to be rather compact even in extremely active stars (e.g., \citealt{walter83, ottmann93}); this is a consequence of the X-ray brightness scaling with the square of the electron density combined with relatively small pressure scale heights. Radio and X-ray sources are therefore not necessarily co-spatial and probe rather different plasmas or particle populations in stellar atmospheres, different atmospheric layers and structures, and perhaps even different energy sources. There should be little reason to expect that the two types of radiation are correlated in stars. It therefore came as a surprise when such a correlation was uncovered for the steady radio and X-ray luminosities of RS CVn close binary stars. \citet{drake89} found that the soft X-ray and radio (6~cm) luminosities are correlated over a few orders of magnitude albeit somewhat different from a linear trend ($L_{\rm R} \propto L_{\rm X}^{1.37\pm 0.13}$). They suggested a self-consistent scheme in which the radio emission originates from the tail of the Maxwellian electron distribution of a very hot ($>$50~MK) plasma through the gyrosynchrotron process; such a plasma component is suggested from X-ray observations. This model is elegant as it links X-rays and radio emission by suggesting a common source. However, serious difficulties remain. Gyrosynchrotron spectra from a thermal plasma reveal a steep decline toward higher frequencies not observed in any magnetically active star; an acceptable spectral fit requires an extraordinary setup of the coronal magnetic field, such as a field strength decreasing with radius as $r^{-1}$ \citep{chiuderi93, beasley00}. Non-thermal (power-law) electron distributions, in contrast, readily produce shallow spectra as observed, especially if the ``aging'' of an injected electron distribution, leading to spectral modifications due to synchrotron and collisional losses, is taken into account \citep{chiuderi93}. A closer inspection of stars more akin to the Sun is in order. To that end, \citet{guedel93a} studied X-ray and radio luminosities for M dwarfs, followed by other spectral types including G dwarfs (e.g., \citealt{benz94, guedel95}). Again, active stars of all late-type spectral classes followed a similar trend, best described by a proportionality, $L_{\rm X}/L_{\rm R} \approx 10^{15.5\pm 0.5}$~Hz (in the following referred to as the GB relation). Combining these samples with samples of RS CVn binaries, Algol binaries, FK Com-type stars and also pre-main sequence weak-lined T Tauri stars, a coherent trend is found over 5-6 orders of magnitude in $L_{\rm R}$ and $L_{\rm X}$ (Figure~\ref{fig1}). It is important to note that the $L_{\rm X}/L_{\rm R}$ ratio is by no means universal. It has been demonstrated exclusively for {\it magnetically active stars} but does clearly not apply to inactive stars like the Sun; such stars keep appreciable levels of quasi-steady soft X-ray emission but are {\it not} sources of continuous radio emission of the gyrosynchrotron type. In fact, present-day radio observatories still cannot systematically detect nearby cool stars except for extremely active examples. Active stars stand out by two properties mentioned above - their extremely hot plasma seen in X-rays, and their non-thermal electron populations evidenced by their radio emission. Let us assume that the energy initially contained in the accelerated electrons eventually heats the coronal plasma. If the corona releases energy at a rate $\dot{E}$ by injecting accelerated electrons at an energy-dependent rate $\dot{n}_{\rm in}(\epsilon)$, then \begin{equation}\label{equil} \dot{E} = {1\over a}\int_{\epsilon_0}^{\infty} \dot{n}_{\rm in}(\epsilon)\epsilon d\epsilon = {1\over b}L_{\rm X} \end{equation} where $a$ is the fraction of the total energy that is channeled into particle acceleration and $b$ is the fraction of the total energy that is eventually radiated as soft X-rays. Eq.~\ref{equil} assumes an equilibrium between energy injection and energy loss. After introducing radiation processes into Eq.~\ref{equil}, one finds \citep{guedel93b} \begin{equation}\label{correl} L_{\rm R} = 3\times 10^{-22}B^{2.48}{a\over b}\tau_0 (\alpha + 1) L_{\rm X} \label{luminosities} \end{equation} i.e., a proportionality if several parameters on the right-hand side take characteristic, constant values, in particular $B$ and the ratio $a/b$ ($\alpha$ is the power-law index for the energy dependence of the electron lifetime). Conversely, comparing Eq.~\ref{correl} with observations, one finds (e.g., for $a/b\approx 1$) the time scale $\tau_0$ for electron trapping (i.e., the lifetime of the population), concluding that the radiation must decay on time scales of minutes to hours in most cases. This necessitates frequent or quasi-continuous replenishment of the corona by high-energy electrons. This mechanism is what the ``standard model'' for a solar flare would predict. The standard solar flare model, the chromospheric evaporation scenario, posits that electrons initially accelerated in reconnecting magnetic fields propagate to chromospheric layers where they heat and ablate material which escapes into closed magnetic loops and cools by X-ray radiation. The best observational evidence for this model is the ``Neupert Effect'', stating that the time derivative of the flare X-ray light curve resembles the radio (or hard X-ray or U-band) light curve, $dL_{\rm X}/dt\propto L_{\rm R}$. This prediction follows from assuming that $L_{\rm X}$ roughly scales with the thermal energy content in the hot plasma accumulated from the high-energy electrons, while radio emission scales with the number of such electrons present at any given time. The Neupert Effect is frequently observed in solar flares (e.g., \citealt{dennis93}), but has also frequently been seen in stellar flares, both extremely large events and the smallest yet discerned in stellar X-rays (e.g., \citealt{guedel96, guedel02, osten04}). We need one further ingredient, relating flares to the observed quiescent emission. During the past decade, a number of studies have shown that the occurrence rate of stellar flares in X-rays is distributed as a power law in radiated energy (a concept familiar to solar physics), $dN/dE \propto E^{-\alpha}$, with $\alpha \ga 2$ (e.g., \citealt{audard00, kashyap02}). In that case, and assuming that the power law continues toward smaller energies, the energy integration, $ \int_{\rm E_0}^{\infty}E(dN/dE)dE$ diverges for $E_0\rightarrow 0$, i.e., a lower cut-off is required. More relevant here, the entire apparently steady emission level could be explained by the large number of small flares that superpose to a quasi-steady emission level while not recognizable individually in light curves. Assembling the above pieces, we then suggest to solve the $L_{\rm X}-L_{\rm R}$ puzzle as follows: Radio and X-ray emission correlate in magnetically active stars because the radiation we perceive as ``quiescent'' emission is made up of contributions from numerous small flares; each of these flares heats plasma by transforming kinetic energy from accelerated electrons; a portion of the latter is evident from their radio emission, while the heated plasma is observed by its X-ray emission. The $L_{\rm X}/L_{\rm R}$ ratio therefore reflects the energy loss ratio of individual flares. As a check, we consider whether {\it solar and stellar flares} reveal radiative output ratios similar to those of the ``quiescent'' radiation. Average X-ray and radio luminosities for a range of solar flares (with specified total flare durations) as well as a sample of stellar flares are overplotted in Fig.~\ref{fig1}. Indeed, the solar flares continue the trend seen in magnetically active stars \citep{benz94} and the stellar flares show luminosity ratios in perfect agreement with the trend for quiescent emission. These observations support a picture in which flares are at the origin of coronal heating, of the steady radiation in magnetically active stars, and consequently of the $L_{\rm X}/L_{\rm R}$ correlation. \vspace*{-5mm}

In short, we are not headed toward a divorce. Instead, the X-ray and radio luminosities of cool stars appear to be in an ``open relationship'' where a lot depends on the type of radio emission that is present. Until now, observational data have not always been sensitive enough to unambiguously identify dominant emission mechanisms, for example in the case of YSOs. Clearly, the $L_{\rm X}/L_{\rm R}$ relation does not apply to all stars. As mentioned above, inactive stars violate this relation as they do not show non-thermal radio (gyrosynchrotron) emission. The ``non-flaring'' Sun is an example. Either, there are additional heating mechanisms at work in these stars that do not involve high-energy electrons, or the energy transformation process is more efficient in heating the plasma to sufficiently high temperatures so as to become visible in X-rays. The other important class of stars violating the relation are brown dwarfs, but coherent radio radiation mechanisms may matter here. Similarly, a number of protostellar objects do not follow the standard trend although they are magnetospheric radio and X-ray sources. Here, however, the measurement of either of the luminosities is difficult. Radio gyrosynchrotron emission may be attenuated by overlying ionized winds (e.g., the jets easily detected as thermal radio sources); X-ray emission could be partially attenuated by neutral gas masses, such as neutral winds, accretion streams, or molecular outflows. If only part of the coronal emission is (fully) attenuated, e.g., by accretion streams, then an assessment of the intrinsic luminosities becomes impossible. Further observations, particularly deeper radio observations with new instruments such as EVLA, as well as large X-ray surveys with \emph{Chandra} and XMM-\emph{Newton} will help shed light on the limits of the radio/X-ray correlation in cool stars. \vspace*{-6mm}

1012.1626

1012

1012.2625_arXiv.txt

We reexamine the model--independent data analysis methods for extracting properties of Weakly Interacting Massive Particles (WIMPs) by using data (measured recoil energies) from direct Dark Matter detection experiments directly and, as a more realistic study, consider a small fraction of residue background events, which pass all discrimination criteria and then mix with other real WIMP--induced signals in the analyzed data sets. In this talk, the effects of residue backgrounds on the determination of the WIMP mass as well as the spin--independent WIMP coupling on nucleons will be discussed.

In our earlier work on the development of model--independent data analysis methods for extracting properties of Weakly Interacting Massive Particles (WIMPs) by using measured recoil energies from direct Dark Matter detection experiments directly% \cite{DMDDf1v,DMDDmchi,DMDDfp2,DMDDidentification-DARK2009}, it was assumed that the analyzed data sets are background--free, i.e., all events are WIMP signals. Active background discrimination techniques should make this condition possible. For example, by using the ratio of the ionization to recoil energy, the so--called ``ionization yield'', combined with the ``phonon pulse timing parameter'', the CDMS-II collaboration claimed that electron recoil events can be rejected event--by--event with a misidentification fraction of $< 10^{-6}$.% \cite{Ahmed09b} The CRESST collaboration demonstrated also that the pulse shape discrimination (PSD) technique can distinguish WIMP--induced nuclear recoils from those induced by backgrounds by means of inserting a scintillating foil, which causes some additional scintillation light for events induced by $\alpha$-decay of $\rmXA{Po}{210}$ and thus shifts the pulse shapes of these events faster than pulses induced by WIMP interactions in the crystal% \cite{CRESST-bg}.% \footnote{ More details about background discrimination techniques and status see also e.g., Refs.~\refcite{bg-papers}. } However, as the most important issue in all underground experiments, possible residue background events which pass all discrimination criteria and then mix with other real WIMP--induced events in our data sets should also be considered. Therefore, as a more realistic study, we take into account small fractions of residue background events mixed in experimental data sets and want to study how well the model--independent methods could extract the {\em input} WIMP properties by using these ``impure'' data sets and how ``dirty'' these data sets could be to be still useful. In this article, I focus on two properties of WIMP Dark Matter: the mass $\mchi$ and the spin--independent (SI) coupling on nucleons $f_{\rm p}$. More detailed discussions can be found in Refs.~\refcite{DMDDbg-mchi,DMDDbg-fp2}.

In this article we reexamine the data analysis methods introduced in Refs.~\refcite{DMDDmchi,DMDDfp2} for determining the mass of Dark Matter particle and its spin--independent coupling on nucleons from measured recoil energies of direct detection experiments directly, by taking into account small fractions of residue background events, which pass all discrimination criteria and then mix with other real WIMP--induced events in the analyzed data sets. Our simulations show that, with a background ratio of $\sim$ 10\% -- 20\% in data sets of only $\sim$ 50 total events, while the 1$\sigma$ statistical uncertainty band of the reconstructed WIMP mass can cover the true value pretty well, especially for an input mass of $\sim$ 100 GeV, the reconstructed SI WIMP coupling on nucleons would be $\sim$ 10\% -- 15\% overestimated and the deviation would be the largest once the WIMP mass is between 50 and 100 GeV.

1012.2625

1012

1012.5005_arXiv.txt

We review the main observational and theoretical facts about acceleration of Galactic cosmic rays in supernova remnants, discussing the arguments in favor and against a connection between cosmic rays and supernova remnants, the so-called supernova remnant paradigm for the origin of Galactic cosmic rays. Recent developments in the modeling of the mechanism of diffusive shock acceleration are discussed, with emphasis on the role of 1) magnetic field amplification, 2) acceleration of nuclei heavier than hydrogen, 3) presence of neutrals in the circumstellar environment. The status of the supernova-cosmic ray connection in the time of Fermi-LAT and Cherenkov telescopes is also discussed.

The possibility that the bulk of cosmic rays (CRs) observed at Earth may be generated in supernova remnants (SNRs) dates back to the '30s \cite{Baade:1934p886} and with time it has acquired the rank of a paradigm, mainly because of the fact that on energetic grounds \cite{Ginzburg:1964p929} supernovae are the only class of sources in the Galaxy that can provide enough energy as to explain the CR flux observed at Earth. The basic requirements for the SNR paradigm are: 1) that SNRs may accelerate with typical efficiency of $\sim 10-20\%$; 2) that the spectrum of individual elements and the consequent all-particle spectrum are reproduced, including the presence of a knee at $\sim 3\times 10^{15}$ eV; 3) that the chemical abundances of nuclei are well described; 4) that the multifrequency observations of individual remnants, from the radio to the gamma ray band, are well described; 4) that the anisotropy induced by the spatial distribution of SNRs in the Galaxy is compatible with observations. The acceleration mechanism that is usually assumed to work in SNRs is diffusive shock acceleration (DSA) \cite{Bell:1978p1344,Blandford:1987p881}, but the energetic requirement that at least $\sim 10-20\%$ of the kinetic energy of the supernova shell is converted to CRs leads to realize immediately that the standard test-particle version of the theory describing this process is not applicable to the description of CR acceleration: the reaction of accelerated particles onto the accelerator cannot be neglected and in fact it is responsible for spectral features (such as spectral concavity) that may represent potential signatures of CR acceleration. There is another, possibly more important reason why DSA must include the reaction of accelerated particles: the standard diffusion coefficient typical of the interstellar medium (ISM) only leads to maximum energies of CRs in the range of $\sim$ GeV, rather than $\sim 10^{6}$ GeV (around the knee) required by observations. The only way that the mechanism can play a role for CR acceleration is if the accelerated particles generate the magnetic field structure on which they may scatter \cite{Lagage:1983p1348}, thereby reducing the acceleration time and reach larger values of the maximum energy. A non linear version of DSA including the dynamical reaction of accelerated particles on the shock was developed by many authors (see Ref. \refcite{Malkov:2001p765} for a review) and more recently completed with the inclusion of self-generation of magnetic field \cite{Amato:2005p112,Amato:2006p139,Caprioli:2008p123} and acceleration of nuclei other than Hydrogen \cite{Caprioli:2010p789}. From the observational point of view the detection of narrow rims in the X-ray emission of several SNRs (Ref. \refcite{Parizot:2006p933} and references therein) has provided an important support to the idea that CRs may amplify the magnetic field close to the shock surface, thereby leading to reach higher energies: these rims are in fact most easily interpreted as the result of the synchrotron emission on a time scale comparable with the loss length of the highest energy electrons \cite{Volk:2005p968}. A simple estimate leads to magnetic fields of order $\sim 100-1000 \mu G$ downstream of the shock, which are hard to interpret as the result of solely the compression of the field at the shock surface since the ISM magnetic field is typically $\sim 1-10 \mu G$ and the compression at a strong shock is a factor $\sim 4$. It is important to realize that the CR induced magnetic field amplification occurs upstream of the shock; the perpendicular components of the field are further compressed at the shock surface. Moreover the overall structure of the rims, at least in the case of SN1006 appears to be inconsistent with the absence of magnetic field amplification upstream of the shock \cite{Morlino:2010p168}. However alternative explanations, involving acceleration at perpendicular shocks and magnetic field amplification due to turbulent eddies downstream of the shock \cite{Giacalone:2007p962} may still beb plausible explanations of the data. The rigidity dependent nature of DSA leads to predict higher energies for accelerated nuclei (if they get fully ionized), so that the knee may result from the superposition of the spectra of accelerated nuclei, if the magnetic field is amplified to sufficiently high levels. Unfortunately the problem of injection, hard enough for protons and electrons, becomes even harder for nuclei especially for those that may result from sputtering of dust grains \cite{Ellison:1997p609}. Some phenomenological attempts to calculating the contribution of nuclei in the context of non linear theory of DSA to the all-particle spectrum observed at Earth have recently been carried out \cite{Berezhko:2007p1010,Ptuskin:2010p1025,Caprioli:2010p789}. A crucial step towards confirming or rejecting the SNR paradigm might be made through gamma ray observations both in the TeV energy range, by using Cherenkov telescopes, and in the GeV energy range accessible to the Fermi gamma ray telescope. Gamma radiation can be produced mainly as a result of inverse Compton scattering (ICS) of relativistic electrons on the photon background and in inelastic proton-proton scatterings with production and decay of neutral pions. The present observational situation is rather puzzling and deserves some special discussion (see \S \ref{sec:gamma}): most gamma ray spectra observed by Fermi (see \cite{funk,tanaka} for reviews) (with some important exceptions, e.g. RX J1713-3946) hint to rather steep spectra of accelerated particles ($\propto E^{-\gamma}$, with $\gamma\sim 2.4-3$), which are not easy to accomodate in the context of NLDSA that predicts flat spectra, possibly even flatter than $E^{-2}$ at high enough energy \cite{Caprioli:2009p145}. Flat injection spectra would also lead to require a steep dependence of the Galactic diffusion coefficient on energy, which in turn is known to result in exceedingly large anisotropy \cite{Ptuskin:2006p620}. The most likely explanation for this discrepancy might lie in a rather subtle detail of the DSA theory, namely that the velocity relevant for particle acceleration is the velocity of waves with respect to the plasma. Usually the wave speed is negligible compared with the plasma velocity in the shock frame, but in the presence of magnetic field amplification this condition might be weakly violated. This is very bad news in that the spectral changes induced by this effect depend not only on the wave speed but on the wave polarization as well. In \cite{Caprioli:2010p789,Ptuskin:2010p1025} the authors show that there are situations in which the spectral steepening can indeed be sufficient to explain the observed spectrum of CRs and required by Fermi data on some SNRs.

We provided a short review of the main arguments in favor and against the so called SNR paradigm for the origin of Galactic CRs. There is a train of recent evidence that efficient CR acceleration takes place in several SNRs. Probably the most striking findings are the detection of narrow X-ray rims interpreted as evidence for magnetic field amplification, and numerous observational results on Balmer dominated shocks, which also lead to conclude that CRs are being accelerated effectively. On the other hand, there are a few pieces of the puzzle that do not appear to be in place: the non-detection of X-ray lines from SNR RX J1713-3946 is certainly problematic for a scenario in which the detected gamma rays are of hadronic origin. Fermi-LAT gamma ray observations of several SNRs with steep spectra also seems to be at odds with the standard predictions of NLDSA which is expected to describe efficient CR acceleration in SNR shocks. Steep spectra are also suggested by observations of CR anisotropy. We have discussed how the last two issues mentioned above could be understood by taking into account the finite velocity of the waves scattering particles close to the shock.

1012.5005

1012

1012.2280_arXiv.txt

A theory of exponential modified gravity which explains both early-time inflation and late-time acceleration, in a unified way, is proposed. The theory successfully passes the local tests and fulfills the cosmological bounds and, remarkably, the corresponding inflationary era is proven to be unstable. Numerical investigation of its late-time evolution leads to the conclusion that the corresponding dark energy epoch is not distinguishable from the one for the $\Lambda$CDM model. Several versions of this exponential gravity, sharing similar properties, are formulated. It is also shown that this theory is non-singular, being protected against the formation of finite-time future singularities. As a result, the corresponding future universe evolution asymptotically tends, in a smooth way, to de Sitter space, which turns out to be the final attractor of the system.

} Modified gravity is getting a lot of attention from the scientific community, owing in particular to the remarkable fact that it is able to describe early-time inflation as well as the late-time (dark energy) acceleration epoch in a unified way. This approach appears to be very economical, as it avoids the introduction of any extra dark component (inflaton % or dark energy of any kind) for the explanation of both inflationary epochs. Moreover, it may be expected that, with some additional effort, it will be able to provide a reasonable resolution of the dark matter problem as well as of reheating, two other important issues in the description of the evolution of our universe. Furthermore, as a by-product, modified gravity has the potential to lead to a number of interesting applications in high-energy physics. In particular, $R^2$ gravity, which provides a simple example of modified gravity, may serve for the unification of all fundamental interactions, including quantum gravity, in an asymptotically-free theory \cite{bos92}. Modified gravity may also be used to provide the scenario for the resolution of the hierarchy problem of high energy physics \cite{cenoz0601}. Finally, the corresponding string/M-theory approach modifies gravity already in the low-energy effective-action approximation, so that a theory of the kind considered appears to be quite natural from very fundamental considerations (see, for instance, \cite{Nojiri:2003rz}). Presently, a number of viable $F(R)$ gravities leading to a unified description as explained have been identified (for a recent review, see \cite{F(R)}, and for a description of the observable consequences of such models, see \cite{F(R)a}). It goes without saying that all those models are constrained to obey the known local tests, as well as cosmological bounds. However, this might not be such a severe problem, since already the first model proposed \cite{Nojiri2003} which unified inflation with dark energy already satisfied many of these local tests. The real internal problem of $F(R)$ gravity is related with its being a higher-derivative theory, which renders it highly non-trivial. This means that it is hard, in fact, to explicitly work with such theories and to get observable predictions from them. The main aim of this paper is to propose a reasonably simple but indeed viable version of $F(R)$ gravity which consistently describes the unification of the inflationary epoch with the dark energy stage, while satisfying the known local tests and cosmological bounds. Specifically, to address the issues above, we here propose exponential gravity, which on top of being simple is moreover free from any kind of finite-time future singularity and exhibits other very interesting properties, as we will see. The paper is organized as follows. In the next section we briefly review $F(R)$ gravity as well as the corresponding FRW cosmological equations. Special attention is paid to de Sitter and spherically-symmetric solutions. Sect.~\ref{SectIII} is devoted to the discussion of the viability conditions in $F(R)$ gravity. These conditions are investigated for the simple and realistic theory of exponential gravity, proposed as a dark energy model, in Sect.~\ref{SectIV}. In Sect.~\ref{SectV} we carry out a detailed analysis of our explicit proposal: exponential gravity which describes in a natural, unifying way both early-time inflation and late-time acceleration. It is there demonstrated, too, that the model leads to a satisfactory, graceful exit from inflation (the de Sitter inflationary solution being unstable). In Sect.~\ref{SectVI} we show that the theory does not lead to any sort of finite-time future singularities. A careful numerical investigation of late-time cosmological dynamics is carried out in Sect.~\ref{SectVII}. It will be demonstrated there that exponential gravity makes specific predictions which are not distinguishable from those of the $\Lambda$CDM model in the dark energy regime. The asymptotic behavior of the theory at late times is investigated in Sect.~\ref{SectVIII}. Section \ref{detailed} is devoted to a somehow different model, a variant of exponential gravity which unifies unstable inflation with the dark energy epoch and which is protected against future singularities by construction. This opens the window to other variations of the basic model sharing all its good properties. In the discussion section \ref{SectX}, a final summary and outlook are provided, and there is an Appendix on the Einstein frame. \setcounter{equation}{0}

} In summary, we have investigated in this paper some models corresponding to the quite simple exponential theory of modified $F(R)$ gravity which are able to explain the early- and late-time universe accelerations in a unified way. The viability conditions of the models have been carefully investigated and it has been demonstrated that the theory quite naturally complies with the local tests as well as with the observational bounds. Moreover, the inflationary era has been proven to be unstable and graceful exit from inflation has been established. A numerical investigation of the dark energy epoch shows that the theory is basically non-distinguishable from the latest observational predictions of the standard $\Lambda$CDM model in this range. Special attention has been paid in the paper to the occurrence of finite-time future singularities in the theory under consideration. It has been shown that it is indeed protected against the appearance of such singularities. Moreover, its evolution turns out to be asymptotically de Sitter (it has a late-time de Sitter universe as an attractor of the system). Hence, the future of our universe, according to such modified gravity, is eternal acceleration. We have also demonstrated that slight modifications of the theory may lead to other non-singular exponential gravities with similar predictions, what points towards a sort of stable class of well-behaved theories. Very nice properties of exponential gravity are its extreme analytic simplicity, as well as the noted singularity avoidance. In this respect, the theory considered seems to be a very natural candidate for the study of cosmological perturbations and structure formation, which are among the most basic issues of evolutional cosmology. However, the theory remains in the class of higher-derivative gravities, which is not yet well understood, even concerning its canonical formulation \cite{woodard}. In this respect, the covariant perturbation theory developed in \cite{dunsby} could presumably be applied for such investigation. This will be pursued elsewhere.

1012.2280