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JSALT2024 Evaluating LLMs for Research Astronomy 7

951172

2011PASP..123..138V

1012

1012.4570_arXiv.txt

\label{sect:intro} As interstellar clouds collapse to form new stars, part of the gas and dust are transported from the infalling envelope to the rotating disk out of which new planetary systems may form \citep{Shu87}. Water has a pivotal role in these protostellar and protoplanetary environments \citep[see reviews by][]{Cernicharo05,Melnick09}. As a dominant form of oxygen, the most abundant element in the universe after H and He, it controls the chemistry of many other species, whether in gaseous or solid phase. It is a unique diagnostic of warm gas and energetic processes taking place during star formation. In cold regions, water is primarily in solid form, and its presence as an ice may help the coagulation process that ultimately produces planets. Asteroids and comets containing ice have likely delivered most of the water to our oceans on Earth, where water is directly associated with the emergence of life. Water also contributes to the energy balance as a gas coolant, allowing clouds to collapse up to higher temperatures. The distribution of water vapor and ice during the entire star and planet formation process is therefore a fundamental question relevant to our own origins. The importance of water as a physical diagnostic stems from the orders of magnitude variations in its gas phase abundance between warm and cold regions \citep[e.g.,][]{Cernicharo90,vanDishoeck96,Ceccarelli96,Harwit98,Ceccarelli99,Snell00,Nisini02,Maret02,Boonman03,vanderTak06}. Thus, water acts like a `switch' that turns on whenever energy is deposited in molecular clouds and elucidates key episodes in the process of stellar birth, in particular when the system exchanges matter and energy with its environment. This includes basic stages like gravitational collapse, outflow injection, and stellar heating of disks and envelopes. Its unique ability to act as a natural filter for warm gas and probe cold gas in absorption make water a highly complementary diagnostic to the commonly used CO molecule. Studying water is also central to understanding the fundamental processes of freeze-out, grain surface chemistry, and evaporation \citep[e.g.,][]{Hollenbach09}. Water is the dominant ice constituent and can trap various molecules inside its matrix, including complex organic ones \citep[e.g.,][]{Gibb04,Boogert08}. When and where water evaporates back into the gas is therefore relevant for understanding the wealth of organic molecules observed near protostars \citep[see review by][]{Herbst09}. Young stars also emit copious UV radiation and X-rays which affect the physical structure as well as the chemistry, in particular that of hydrides like water \citep{Stauber06}. Moreover, the level of deuteration of water provides an important record of the temperature history of the cloud and the conditions during grain surface formation, and, in comparison with cometary data, of its evolution from interstellar clouds to solar system objects \citep[e.g.,][]{Jacq90,Gensheimer96,Helmich96,Parise03}. Finally, water plays an active role in the energy balance of dense gas \citep[e.g.,][] {Goldsmith78,Neufeld93,Doty97}. Because water has a large dipole moment, its emission lines can be efficient coolants of the gas, contributing significantly over the range of physical conditions appropriate in star-forming regions. The large dipole moment can also lead to heating through absorption of infrared radiation followed by collisional de-excitation \citep{Ceccarelli96}. It is therefore important to study the gaseous water emission and absorption, and compare its cooling or heating efficiency quantitatively with that of other species. Interstellar water was detected more than 40 years ago through its 22 GHz maser emission towards Orion \citep{Cheung69}. While this line remains a useful beacon of star-formation activity, its special excitation and line formation conditions limit its usefulness as a quantitative physical and chemical tool. Because observations of thermally excited water lines from Earth are limited, most information on water has come from satellites. The {\it Submillimeter Wave Astronomy Satellite} \citep[SWAS,][]{Melnick00} and the Odin satellite \citep[e.g.,][]{Hjalmarson03,Bjerkeli09} observed only the 557 GHz ground-state line of ortho-H$_2$O at spatial resolutions of $\sim 3.3'\times 4.5'$ and $126''$, respectively. The 557 GHz line was detected in just the brightest objects by these missions and other lines of water (including those of para-H$_2$O) and related species could not be observed. The Short (SWS) and Long Wavelength Spectrometers (LWS) on the {\it Infrared Space Observatory} (ISO) covered a large number of pure rotational and vibration-rotation lines providing important insight in the water excitation and distribution, but with limited spectral and spatial resolution and mapping capabilities \citep[see review by][]{vanDishoeck04}. The {\it Spitzer Space Telescope} has detected highly excited pure rotational H$_2$O lines at mid-infrared wavelengths from shocks \citep{Melnick08,Watson07} and from the inner few AU regions of protoplanetary disks \citep{Salyk08,Carr08,Pontoppidan10}. Even higher-lying vibration-rotation H$_2$O lines have been observed at near-infrared wavelengths from the ground \citep{Salyk08}. However, these data do not provide any information on the cooler water reservoir where the bulk of the disk mass is. \begin{figure} \includegraphics[angle=0,width=0.6\textwidth]{f1.ps} \caption{Energy levels of ortho- and para-H$_2$O, with HIFI transitions (in GHz, red) and PACS transitions (in $\mu$m, blue) observed in WISH indicated.} \label{fig:energy} \end{figure} The 3.5m {\it Herschel Space Observatory} with its suite of instruments \citep{Pilbratt10}\footnote{Herschel is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA} is particularly well suited to address the distribution of cold and warm water in star- and planet-forming regions, building on the pioneering results from previous missions. Its wavelength coverage of 55--671 $\mu$m (0.45--5.4 THz; 15--180 cm$^{-1}$)\footnote{This paper follows the common usage of giving frequencies in GHz for lines observed with HIFI and wavelengths in $\mu$m for lines observed with PACS} includes both low- and high-excitation lines of water, enabling a detailed analysis of its excitation and abundance structure (Fig.~\ref{fig:energy}). Compared with the ISO-LWS beam ($\sim$80$''$), Herschel has a gain of a factor of $\sim 8$ in spatial resolution at similar wavelengths, and up to $10^4$ in spectral resolution. Compared to observatories with high spectral resolution heterodyne instruments, its diffraction limited beam of 37$''$ at 557 GHz is a factor of 3--5 smaller than that of SWAS or Odin, with a gain of $>10$ in sensitivity because of the bigger dish and improved detector technology. The Heterodyne Instrument for the Far-Infrared (HIFI; \citealt{deGraauw10}) has spectral resolving powers $R=\lambda/\Delta \lambda>10^7$ up to 1900 GHz, allowing the kinematics of the water lines to be studied. The Photoconducting Array Camera and Spectrometer (PACS; \citealt{Poglitsch10}) $5\times 5$ pixel array provides instantaneous spectral mapping capabilities at $R=1500-4000$ of important (backbone) water lines in the 55--200 $\mu$m range. These enormous advances in sensitivity, spatial and spectral resolution, and wavelength coverage provide a unique opportunity to observe both cold and hot water in space, with no other space mission with similar capabilities being planned. The goal of the `Water In Star-forming regions with Herschel' (WISH) Key Program (KP) is to use gas-phase water as a physical and chemical diagnostic and follow its abundance throughout the different phases of star and planet formation. A comprehensive set of water observations is carried out with HIFI and PACS toward a large sample of young stellar objects (YSOs), covering a wide range of masses and luminosities -- from the lowest to the highest mass protostars-- and a large range of evolutionary stages -- from the first stages represented by the pre-stellar cores to the late stages represented by the pre-main sequence stars surrounded only by disks. Lines of H$_2$O, H$_2^{18}$O, H$_2^{17}$O and the chemically related species O, OH, and H$_3$O$^+$ are targeted\footnote{Both H$_2$O and `water' are used to denote the main H$_2^{16}$O isotopologue throughout this paper}. In addition, a number of hydrides which are diagnostic of the presence of X-rays and UV radiation are observed. Selected high-frequency lines of CO, $^{13}$CO, and C$^{18}$O as well as dust continuum maps are obtained to constrain the physical structure of the sources independent of the chemical effects. Together with the atomic O and C$^+$ lines, these data also constrain the contributions from the major coolants. The Herschel data are complemented by ground-based maps of longer wavelength continuum emission and lines of HDO, CO, C and other molecules to ensure a self-consistent data set for analysis. In terms of water, the WISH program is complementary to the `CHEmical Survey of Star-forming regions' (CHESS; \citealt{Ceccarelli10chess}) and `Herschel observations of EXtraOrdinary Sources' (HEXOS; \citealt{Bergin10}) HIFI KPs which survey the entire spectral range (including many lines of water) but only for a limited number of sources and with smaller on-source integration times per frequency setting. The `Dust, Ice and Gas in Time' (DIGIT) (PI N.J.\ Evans) and `Herschel Orion Protostar Survey' (HOPS) (PI T.\ Megeath) KPs complement WISH by carrying out full PACS spectral scans for a larger sample of low-mass embedded YSOs \citep[e.g.,][]{vanKempen10dkcha}. The `PRobing InterStellar Molecules with Absorption line Studies' (PRISMAS) targets the water chemistry in the diffuse interstellar gas \citep{Gerin10} whereas `Water and related chemistry in the Solar System' (HssO) observes water in planets and comets \citep{Hartogh10}. Together, these Herschel data on water and related hydride lines will provide a legacy for decades to come. In the following sections, background information on the water chemistry and excitation relevant for interpreting the data is provided (\S 2). The observational details and organization of the WISH program are subsequently described in \S 3, with the specific goals and first results of the various subprograms presented in \S 4. A discussion of the results across the various evolutionary stages and luminosities is contained in \S 5 together with implications for the water chemistry, with conclusions in \S 6. Detailed information can be found at the WISH KP website {\tt http://www.strw.leidenuniv.nl/WISH/}. This website includes links to model results, modeling tools, complementary data and outreach material. These data and analysis tools will be useful not only for {\it Herschel} but also for planning observations with the Atacama Large Millimeter/submillimeter Array (ALMA) and future far-infrared missions.

The goal of the WISH program is to use water and related species as physical and chemical probes of star-forming regions over a range of luminosities and evolutionary stages. The initial results presented in \S 4 and discussed in \S 5 demonstrate that many of the questions asked in the individual subprograms can indeed be addressed by the WISH data, thanks to the combination of excellent sensitivity, fully resolved line profiles and spatial information, thereby validating the observational strategy. Initial surprises include a near absence of gaseous water in pre-stellar cores and disks, the dominance of shocks rather than hot cores in controlling the bright water emission from protostars, and the detection of all of the ions and hydrides involved in the water chemistry schemes. For pre-stellar cores our observations are in agreement with models that suggest near total freeze-out of water in the form of ice, whereas in disks the inferred water abundance is at levels significantly lower than predicted. So far, water has not yet been found to be the major coolant in any of these regions. All remaining observations are expected to be taken within the next year. Quantitative analysis will require the further development of multidimensional models of protostellar envelopes, outflows and disks. Together with the results from related {\it Herschel} key programs, they will greatly enhance our understanding of water in the galactic interstellar medium and solar system, and provide a true legacy to follow the water trail from the most diffuse gas to dense cores and disks, and eventually comets and planets in our own Solar system.

1012.4570

Water In Star-forming regions with Herschel (WISH) is a key program on the Herschel Space Observatory designed to probe the physical and chemical structures of young stellar objects using water and related molecules and to follow the water abundance from collapsing clouds to planet-forming disks. About 80 sources are targeted, covering a wide range of luminosities—from low (&lt; 1 L<SUB>⊙</SUB>) to high (&gt;10<SUP>5</SUP> L<SUB>⊙</SUB>)—and a wide range of evolutionary stages—from cold prestellar cores to warm protostellar envelopes and outflows to disks around young stars. Both the HIFI and PACS instruments are used to observe a variety of lines of H<SUB>2</SUB>O, H<SUB>2</SUB><SUP>18</SUP>O and chemically related species at the source position and in small maps around the protostars and selected outflow positions. In addition, high-frequency lines of CO, <SUP>13</SUP>CO, and C<SUP>18</SUP>O are obtained with Herschel and are complemented by ground-based observations of dust continuum, HDO, CO and its isotopologs, and other molecules to ensure a self-consistent data set for analysis. An overview of the scientific motivation and observational strategy of the program is given, together with the modeling approach and analysis tools that have been developed. Initial science results are presented. These include a lack of water in cold gas at abundances that are lower than most predictions, strong water emission from shocks in protostellar environments, the importance of UV radiation in heating the gas along outflow walls across the full range of luminosities, and surprisingly widespread detection of the chemically related hydrides OH<SUP>+</SUP> and H<SUB>2</SUB>O<SUP>+</SUP> in outflows and foreground gas. Quantitative estimates of the energy budget indicate that H<SUB>2</SUB>O is generally not the dominant coolant in the warm dense gas associated with protostars. Very deep limits on the cold gaseous water reservoir in the outer regions of protoplanetary disks are obtained that have profound implications for our understanding of grain growth and mixing in disks.

12137001

2010EPJH...35..309C

1012

1012.5068.txt

By 1785 de Coulomb found \cite{Cou1785} that electroscopes can spontaneously discharge due to the action of the air and not because of defective insulation. After dedicated studies by Faraday around 1835 \cite{Far1835}, Crookes observed in 1879 \cite{Cro1879} that the speed of discharge of an electroscope decreased when pressure was reduced. The explanation of this phenomenon came in the beginning of the 20th century and paved the way to one of mankind's revolutionary scientific discoveries: cosmic rays. %Several reviews have been written on the history of the research on cosmic rays, see for example \cite{janossy,rossi,hillas,wilson,xu}. Early seminal reviews had been published in the 20's and 30's \cite{wigandr,millikanr,comptonr}. Several reviews have been written on the history of the research on cosmic rays \cite{montgo,janossy,leprince,rossi,hillas,wilson,wigandr,millikanr,comptonr,swann,xu,ginz,puppi}.

%The discovery of cosmic rays, a milestone in science, comprised scientists in Europe and the US and took place during a period characterised by nationalism and lack of communication. Many scientists that took part in this research a century ago, either alone or as a two person group, were intrigued by the penetrating radiation and tried to understand the origin of it. Many important contributions have been forgotten and in particular that of Pacini. From Hess' pioneering balloon flight in 1912 it took until 1926 before the extraterrestrial nature of the penetrating radiation was generally accepted, partly because of the World War I, but partly for lack of communication between different countries in Europe and between Europe and the United States. Several historical, political and personal facts might have contributed to the substantial disappearance of Pacini from the history of science. %The nominations for a Nobel Prize in physics for the discovery of cosmic rays started only in the 1930Õs and the Prize was given in 1936, two years after the death of Pacini. The work behind the discovery of cosmic rays, a milestone in science, comprised scientists in Europe and the New World and took place during a period characterised by lack of communication and by nationalism caused primarily by the World War I. The many scientists that took part in this research starting a century ago, either alone or as a two-person group, were fascinated by the penetrating radiation and wanted to understand the origin and properties of it. It took from the turn of the century until 1926 before the extraterrestrial nature of the penetrating radiation was generally accepted. In the work that culminated with high altitude balloon flights, many important contributions have been forgotten and in particular those of Pacini in 1909-1911. Several historical, political and personal facts might have contributed to the lack of references to the work of Pacini in the history of cosmic rays. The nominations for a Nobel Prize in physics for the discovery of cosmic rays started only in the 1930s. Pacini was never nominated and died two years before the Nobel Prize was given to Hess in 1936. %

1012.5068

The discovery of cosmic rays, a milestone in science, comprised of scientists in Europe and North America and took place during a period characterised by nationalism and lack of communication. Many scientists that took part in this research a century ago were intrigued by the penetrating radiation and tried to understand the origin of it. Several important contributions to the discovery of the origin of cosmic rays have been forgotten in particular that of Domenico Pacini, who in June 1911 demonstrated by studying the decrease of radioactivity with an electroscope immersed in water that cosmic rays could not come from the crust of the Earth. Several historical, political and personal facts might have contributed to the substantial disappearance of Pacini from the history of science.

12213563

2011MNRAS.413.1054M

1012

1012.3020_arXiv.txt

\label{sec:intro} The strong cosmic evolution of the most powerful radio sources was first deduced by \citet{Longair1966} from low frequency radio source counts. This evolution has since been confirmed by a number of investigations which imply that the co-moving space density of high luminosity radio sources has decreased by a factor of approximately 1000 between $z\sim2$ and $z\sim0$ (e.g. \citealp*{LRL1983}; \citealp{DP1990,Willott2001}). Beyond $z\sim 2$ the evolution remains uncertain \citep[e.g.][]{JarvisRawlings00,Jarvis01b} but appears to undergo a gradual decline \citep[e.g.][]{Wall05}. The evolution of the low-power radio sources is less well constrained, but early studies of the radio source counts indicated that the low luminosity sources could not be evolving as strongly as their high luminosity counterparts (\citealt{Longair1966}; \citealt*{DLZ1970}). Many models of the evolution of radio sources thus divide the radio source population into two independently evolving components, a strongly evolving high luminosity component and a lower luminosity component with much weaker evolution (e.g. \citealt{JW1999,Willott2001}). Radio galaxies can be classified into two morphologically distinct groups identified as Fanaroff-Riley class I (FRI) and Fanaroff-Riley class II (FRII) sources \citep{FR1974}. The FRI sources have higher surface brightness close to the centre of their radio lobes whereas the FRII sources have highly collimated large-scale jets and bright emission hotspots at the edges of their radio lobes. The FRII sources are typically more luminous than the FRI sources with the division in luminosity falling at roughly $L_{\rm{1.4GHz}}=10^{25}~\rm{W~Hz^{-1}sr^{-1}}$. This dividing luminosity may be dependent on the optical luminosity of the host galaxy, a higher optical host luminosity results in a higher FRI/FRII division luminosity \citep{LO1996}. It has been suggested that the FRI and FRII sources might respectively be associated with the slowly and rapidly evolving components of the radio source population \citep{JW1999}. However \citet{Rigby} find evidence that high luminosity FRI sources evolve as rapidly as FRII sources of comparable radio power. Furthermore, \citet*{Gendre2010} performed a detailed comparison of the radio luminosity function of FRI and FRII sources at a number of redshifts. This analysis revealed that both populations experience luminosity dependent number density enhancements at higher redshifts, z$\approx$0.8$\sim$2.5, and that there were no significant differences between the enhancements measured for these two populations. These results would seem to suggest that both types of FR sources experience a common evolutionary history and thus cannot fully account for the observed dichotomy in cosmic evolution. Recent evidence suggests that a more important division in the radio source population, distinct from the morphological classification, may be related to different modes of black hole accretion. Radio galaxies have also been classified according to the presence or absence of narrow high-excitation emission lines in the spectra of their optical host galaxies \citep{HL1979,Laing1994,JR1997,Willott2001}. Objects without high-excitation emission lines are referred to as low-excitation radio galaxies (LERG). Low luminosity FRI galaxies are predominantly LERG's and most objects with high excitation emission lines are associated with the more powerful FRII sources. However the relationship between FR class and the emission line classification scheme is not one-to-one as many FRII galaxies have been found to be low-excitation radio galaxies \citep[e.g.][]{Evans2006}. It has been argued that these two groups correspond to different AGN phenomena powered by fundamentally different modes of accretion. The LERG sources are powered by a radiatively inefficient accretion of the hot gas in the intergalactic medium, referred to as ``radio'' mode accretion. While the HERG are powered by radiatively efficient accretion of cold gas, referred to as quasar mode accretion (\citealt{Evans2006}; \citealt*{Hardcastle2006,Hardcastle2007}; \citealt{Herbert2010}). \citet{Smol2009} argue that the observed bimodal evolution of radio sources may be caused by differences in the evolution of these black hole accretion modes. Understanding the relationships between different classes of radio galaxies, and their cosmic evolution, is important for understanding models of galaxy formation and evolution. Semi-analytic models of hierarchical galaxy formation predict that galaxies in the local universe should be more massive and have higher rates of star formation than is observed \citep{Bower2006}. There is increasing observational evidence that radio activity from AGN may disrupt cooling flows in luminous early-type galaxies thus slowing or preventing significant accretion of gas and further star formation \citep{Fabian,Birzan2004,RJ2004}. This negative AGN feedback has been successfully incorporated into models of galaxy formation which are then able to better reproduce several features of the observed universe including the exponential cut-off in the bright end of the galaxy luminosity function \citep{croton2006,Bower2006}. \citet{Best2006} implied that the low luminosity radio sources may be the predominant contributors to this negative feedback effect. Current studies of the cosmic evolution of low power radio sources seem to imply that these sources are not evolving at all or experience only mild negative evolution. \citet{Clewley2004} found that the low luminosity radio sources with $L_{\rm{325MHz}}<10^{25}\rm{~W~Hz^{-1}sr^{-1}}$ exhibit no evolution out to redshift $z\sim0.8$, \citet{Waddington} also found no evidence of evolution in their study of 72 sources. \citet{Sadler} find evidence of mild evolution out to $z\sim 0.7$ in their comparison of the radio luminosity function of sources in the 2dF-SDSS Luminous Red Galaxy (2SLAQ) and QSO survey with that of sources in the 6dF Galaxy survey (6dFGS). This was consistent with pure luminosity evolution of the form $(1+z)^{2.0\pm0.3}$ which ruled out the no evolution scenario at the 6$\sigma$ level. \citet*{Donoso} also find evidence of mild positive evolution in the $z=0.1\sim0.55$ redshift range. Using radio sources in the VLA-COSMOS survey \citet{Smol2009} find evidence that this mild evolution continues out to $z\sim 1.3$. All of these studies imply that the level of evolution in the low power radio sources is significantly less than that taking place in the high luminosity sources. In this study we use low frequency radio sources detected in the XMM-Large Scale Structure survey field to investigate the co-moving space density of low-luminosity radio sources. The low-frequency selection is preferred over high-frequency (e.g. $\geq 1.4$~GHz), as it provides an orientation independent selection, as the low-frequency detects the optically-thin lobe emission, whereas high-frequency surveys contain a higher fraction of pole-on sources where the optically thick core dominates \citep[e.g.][]{JM02}. In section 2 and 3 we discuss the radio and optical observations used in our analysis. Section 4 outlines the method we used to cross match the radio sources with their optical counterparts. In section 5 we describe the method used to obtain redshifts for our sample. In the section 6 and 7 we discuss the $V/V_{\rm{max}}$ statistic used to quantify the evolution of our sources. In section 8 we determine the radio luminosity function of the low luminosity sources in our observed field and use it to investigate their evolution. We present our conclusions in section 9. We assume the $H_{0}$=70~km.s$^{-1}$.Mpc$^{-1}$ and a $\Omega_{M}$=0.3 and $\Omega_{\Lambda}$=0.7 cosmology throughout this paper.

We have used the non-parametric $V/V_{\rm{max}}$ test and the radio luminosity function to investigate the the cosmic evolution of low power radio sources in the XMM-LSS field. Previous investigations of the evolution of low-power radio sources have found evidence of mild evolution taking place in the redshift range $z=0-1.3$. \citet{Donoso} find that the co-moving density of low luminosity sources ($L_{\rm{1.4GHz}}<10^{25}~\rm{W~Hz^{-1}sr^{-1}}$) in the MegaZ-Luminous Red Galaxy catalogue increases by a factor of 1.5 between $z=0.1-0.55$. Similarly \citet{Sadler} find evidence that low luminosity sources experience mild evolution with an increase in their number density by a factor of $\sim$2 at z=0.55. This is significantly less than the strong evolution detected in the high power radio sources whose number densities are enhanced by a factor of 10 in the same cosmic timeframe. The results of \citet{Sadler} seem to rule out the no-evolultion scenario. \citet{Gendre2010} find evidence that higher luminosity, $L_{\rm{1.4GHz}} \geq 10^{24.5}~\rm{W~Hz^{-1}sr^{-1}}$ FRI \citet{Smol2009} find mild evolution of the low power AGN in the VLA-COSMOS survey out to $z\sim1.3$. Our results are broadly consistent with these previous works and imply density enhancements by a factor $\sim1.5$ at z=0.8. This is slightly less than the evolution implied in \citet{Sadler}, but in fairly good agreement with the estimates in \citet{Smol2009} and \citet{Donoso}. We also find evidence that the low power radio AGN are evolving more slowly than their high power counterparts in the redshift range $z=0-0.8$ and tentative evidence that this separation in evolutionary behaviour persists to $z=1.2$. The division in the FRI/FRII classification system and the LERG/HERG system occurs at roughly the same luminosity threshold of $L_{\rm{1.4GHz}}\sim10^{25}$. As our analysis categorises our radio sources solely on their luminosity we are unable to determine whether the separation in evolutionary behaviour is due to the FRI/FRII dichotomy or to differences in black hole accretion modes in HERGS and LERGS. Our results demonstrate that using the $K-$band (or similar wavelength), combined with radio surveys, is a viable route to investigating the evolution of the radio source population, at least up to $z \sim 1.2$. In the future, all-sky radio surveys such as those carried out with the Low-Frequency Array \citep[LOFAR; ][]{Morganti2010} and the Australian Square Kilometre Array Pathfinder (ASKAP) telescope, combined with the UKIDSS large area survey and the VISTA Hemisphere Survey, as well as Wide-field Infrared Survey Explorer (WISE), will enable us to pin down the evolution of the radio source population to a much higher degree of accuracy.

1012.3020

We investigate the cosmic evolution of low-luminosity (L<SUB>1.4 GHz</SUB> &lt; 10<SUP>25</SUP> W Hz<SUP>-1</SUP> sr<SUP>-1</SUP>) radio sources in the XMM Large Scale Structure Survey (XMM-LSS) field. We match low-frequency-selected (610-MHz) radio sources in the XMM-LSS field with near-infrared K-band observations over the same field from the UKIRT Infrared Deep Sky Survey. We use both the mean V/V<SUB>max</SUB> statistic and the radio luminosity function of these matched sources to quantify the evolution of the comoving space density of the low-luminosity radio sources in our sample. Our results indicate that the low-luminosity sources evolve differently from their high-luminosity counterparts out to a redshift of z∼ 0.8. The derived luminosity function is consistent with an increase in the comoving space density of low-luminosity sources by a factor of ∼1.5 at z= 0.8. We show that the use of the K-z diagram for the radio source population, although coarser than a full photometric redshift analysis, produces consistent results with previous studies using approximately &gt;10 band photometry. This offers a promising method for conducting similar analyses over the whole sky with future near- and mid-infrared surveys.

12225115

2011PhRvD..83e5011K

1012

1012.0435_arXiv.txt

\label{sec1} The standard big bang nucleosynthesis (BBN) model predicts primordial light element abundances which are more or less consistent with abundances inferred from observations of old distant astronomical objects. Deviations from the standard BBN (SBBN) model are, therefore, constrained if predicted abundances in theoretical models change from those in the SBBN. Constraints on the existence of long-lived exotic particles which interact with nuclei by strong force~\cite{Dicus1980,Plaga1995,Mohapatra:1998nd,Kusakabe:2009jt} or Coulomb force~\cite{Pospelov:2006sc,Kohri:2006cn,Cyburt:2006uv,Hamaguchi:2007mp,Bird:2007ge,Kusakabe:2007fu,Kusakabe:2007fv,Kusakabe:2010cb,Jedamzik:2007cp,Jedamzik:2007qk,Kamimura:2008fx,Pospelov:2007js,Kawasaki:2007xb,Jittoh:2007fr,Jittoh:2008eq,Jittoh:2010wh,Pospelov:2008ta,Khlopov:2007ic} have been derived as well as those on the decay of long-lived exotic particles into standard model particles which have electromagnetic or hadronic interactions~\cite{Ellis:1984er,Cyburt:2002uv,Ellis:2005ii,Terasawa:1988my,Kawasaki:1993gz,Kawasaki:1994af,Holtmann:1996cq,Kawasaki:2000qr,Kawasaki:2004yh,Kawasaki:2004qu,Kanzaki:2006hm,Kawasaki:2008qe,Cumberbatch:2007me,Reno:1987qw,Dimopoulos:1987fz,Dimopoulos:1988zz,Dimopoulos:1988ue,Khlopov:1993ye,Sedelnikov:1987ef,Jedamzik:1999di,Jedamzik:2004er,Jedamzik:2004ip,Jedamzik:2005dh,Jedamzik:2006xz,Kusakabe:2006hc,Kusakabe:2008kf,Pospelov:2010cw,Pospelov:2010kq}. A prominent problem relating to the abundances predicted in the SBBN model and inferred from observations is lithium problem~\cite{Melendez:2004ni,Asplund:2005yt}. Primordial lithium abundances are inferred from measurements in metal-poor halo stars (MPHSs). Observed abundances are roughly constant as a function of metallicity~\cite{Spite:1982dd,Ryan2000,Melendez:2004ni,Asplund:2005yt,bon2007,Shi:2006zz,Aoki:2009ce} at $^7$Li/H$=(1-2) \times 10^{-10}$. The theoretical prediction by the SBBN model is, however, a factor of 2--4 higher, e.g., $^7$Li/H=$(5.24^{+0.71}_{-0.67})\times 10^{-10}$~\cite{Cyburt:2008kw}, when its only parameter, the baryon-to-photon ratio, is deduced from the observation with Wilkinson Microwave Anisotropy Probe (WMAP) of the cosmic microwave background (CMB) radiation~\cite{Larson:2010gs}. This discrepancy indicates some mechanism of $^7$Li reduction having operated in some epoch from the BBN to this day. One possible astrophysical process to reduce $^7$Li abundances in stellar surfaces is the gravitational settling in the model including a combination of the atomic and turbulent diffusion~\cite{Richard:2004pj,Korn:2006tv}. The precise trend of Li abundance as a function of effective temperature of stars in the metal-poor globular cluster NGC 6397 is, however, not reproduced theoretically~\cite{Lind:2009ta}. $^6$Li/$^7$Li isotopic ratios of MPHSs have also been measured spectroscopically. The $^6$Li abundance as high as $^6$Li/H$\sim 6\times10^{-12}$ was suggested~\cite{Asplund:2005yt}, which is about 1000 times higher than the SBBN prediction~\footnote{Recently, a new measurement of the cross section of radiative $\alpha$ capture by deuteron and the $^6$Li abundance predicted based upon the result have been reported.~\cite{Hammache:2010bt}}. Convective motions in the atmospheres of MPHSs could cause systematic asymmetries in the observed line profiles and mimic the presence of $^6$Li~\cite{Cayrel:2007te}. A few or several MPHSs, however, have high $^6$Li abundances larger than levels caused by this effect~\cite{Steffen:2009yr}. This high $^6$Li abundance is a problem since the standard Galactic cosmic ray nucleosynthesis models predict negligible amounts of $^6$Li yields compared to the observed level in the epoch corresponding to the metallicity of the stars, i.e., [Fe/H] $<-2$~\cite{Prantzos2006}~\footnote{[A/H]=log(A/H)$-$log(A/H)$_\odot$ is the number ratio of nuclide $A$ to H measured in a logarithmic scale normalized to the solar value.}. The possibility that the $^7$Li and $^6$Li problems stem from uncertainties in nuclear reactions used in theoretical BBN calculation is unlikely~\cite{Boyd:2010kj} unless there remain to be observed new resonant states contributing to $^7$Be destruction~\cite{Cyburt:2009cf,Chakraborty:2010zj}. The $^7$Li reduction needs a destruction mechanism of $^7$Be during or after BBN and before stellar activities since the $^7$Li nuclei observed in MPHSs are thought to have originated from the electron capture process of $^7$Be which is produced in the BBN. Some particle models beyond the standard model include heavy ($m \gg 1$~GeV) long-lived colored particles. The scenarios, i.e., split supersymmetry~\cite{ArkaniHamed:2004fb,ArkaniHamed:2004yi}, weak scale supersymmetry with a long-lived gluino~\cite{Raby:1997bpa,Shirai:2010rr,Covi:2010au} or squark~\cite{Sarid:1999zx} as the next-to-lightest supersymmetric particles, and extended theories with new kinds of colored particles~\cite{Hisano:2010bx,Nakayama:2010vs}, may be tested in experiments such as the Large Hadron Collider. The heavy colored particles would be confined at temperatures below the deconfinement temperature $T_C\sim 180$~MeV inside exotic heavy hadrons, i.e., strongly interacting massive particles (SIMPs) which we call $X$ particles~\cite{Kang:2006yd}. Their thermal relic abundances after the freeze-out of annihilations depend on the annihilation cross sections, and theoretical estimates predict various values which extend over more than several order of magnitude at the heavy mass limit~\cite{Baer:1998pg}. If the annihilation cross section is not different from a typical value for strong interaction, i.e., $\sigma\sim \mathcal{O}({\rm GeV}^{-1})^2$, however, the final abundance of $X$ particles can be derived under the assumption that their abundances are fixed when the annihilation rate becomes smaller than the Hubble expansion rate of the universe~\cite{Kang:2006yd}. The relic abundance can then be written \begin{equation} \frac{N_X}{s} \sim \sqrt[]{\mathstrut \frac{15}{\pi}} \frac{g_\ast^{1/2}}{g_{\ast s}} \frac{m^{1/2}}{\sigma T_B^{3/2} m_{\rm Pl}}~~, \label{eq1} \end{equation} where $N_X$ is the number density of the $X$ particle, $s=2\pi^2 g_{\ast s} T^3/45$ is the entropy density with $g_{\ast s}\sim 10$ the total number of effective massless degrees of freedom in terms of entropy~\cite{kolb1990} just below the QCD phase transition, $g_\ast$ is the total number of effective massless degrees of freedom in terms of number~\cite{kolb1990}, $m$ is the mass ($m\gg 1$~GeV) of the heavy long-lived colored particles, $\sigma$ is the annihilation cross section of the $X$ particle, $T_B$ is the temperature of the universe at which the $X$-particles are formed, and $m_{\rm Pl}$ is the Planck mass. The number abundance of the $X$s with respect to that of baryons is then \begin{equation} \frac{N_X}{n_b} \sim 0.5\times 10^{-8} \left(\frac{g_\ast}{10.75}\right)^{1/2} \left(\frac{m}{\rm TeV}\right)^{1/2} \left(\frac{T_B}{180 {\rm MeV}}\right)^{-3/2} \left(\frac{\sigma}{m_\pi^{-2}}\right)^{-1}~~, \label{eq2} \end{equation} where $n_b$ is the number density of baryons, and $m_\pi\sim 140$~MeV is the mass of pion. The thermal relic abundance is inversely proportional to the annihilation cross section which depends on the particle theory. In addition, there might be a nonthermal production of long-lived colored particles which is not directly related to the thermal production. The final abundance of the $X$ is, therefore, uncertain. So we consider the $X$ abundance as a free parameter in this paper. Observational constraints on hypothetical SIMPs have been studied~\cite{Wolfram:1978gp,Dover:1979sn,Starkman:1990nj}. Effects of exotic neutral stable hadrons on BBN were studied in Ref.~\cite{Dicus1980}. The authors assumed that the strong force between a nucleon and a exotic hadron ($X$) is similar to that between a nucleon $N$ and a $\Lambda$ hyperon. In addition, new hadrons are assumed to be captured in a bound state of $^4$He plus $X$ after BBN. Based upon an analytic estimation, they suggested that beryllium has the largest number fraction $A_X/A$ of bound states with the hadrons among the light elements produced in BBN, where the $A$ and $A_X$ are a nuclide $A$ and a bound state of $A$ with a hadron $X$. Mohapatra and Teplitz~\cite{Mohapatra:1998nd} estimated the cross section of $X$ capture by $^4$He, and suggested that a large fraction of free $X$ particles would not become bound into light nuclides and remain free contrary to the previous suggestion~\cite{Dicus1980}. In deriving the result of those two studies, it has been assumed that exotic hadrons interact with normal nuclei by typical strengths of strong interaction and implicitly assumed that its mass is about that of $\Lambda$ hyperon, i.e., $m_X \sim 1.116$~GeV~\cite{Particle2010}. Effects on BBN of long-lived exotic hadrons of $m\gg 1$~GeV have been studied recently~\cite{Kusakabe:2009jt}. The authors have assumed that the interaction strength between an $X$ particle and a nucleon is similar to that between nucleons. Rates of many reactions associated with the $X$ particle were estimated, and a network calculation of the nucleosynthesis including effects of the $X$ was performed. The constraint on the decay lifetime of such $X$ particles, i.e., $\tau_X \lesssim 200$~s was derived from a comparison of calculated abundances with observational abundance constraints of light elements. Two interesting predictions of the model~\cite{Kusakabe:2009jt} is signatures of the $X$ particles on primordial abundances which should be seen in future astronomical observations: 1) $^9$Be and B can be produced in amounts more than predicted in the SBBN. Future observations of Be and B abundances in MPHSs may show primordial constant values originating from the BBN catalyzed by the $X$ particle. 2) The isotopic ratio $^{10}$B/$^{11}$B tends to be very high. This is different from predictions of other models for boron production, i.e., the cosmic ray nucleosynthesis ($^{10}$B/$^{11}$B$\sim 0.4$~\cite{Prantzos1993,Ramaty1997,Kusakabe2008}) or the supernova neutrino process ($^{10}$B/$^{11}$B$\ll 1$~\cite{Woosley:1995ip,Yoshida:2005uy}). They concluded that the $^6$Li or $^7$Li problems is not solved under their assumption. Since interactions between long-lived exotic hadrons $X$ and a nucleon are not known as well as their masses, we are investigating effects of such particles in various cases of interaction strengths and masses. We found on the way a new possibility that reactions associated with the $X$ particle reduce $^7$Be abundance and that the $^7$Li problem is solved. In this paper we report details of the destruction mechanism of $^7$Be in the presence of the $X$ particle. We carry out a network calculation of BBN in the presence of a long-lived SIMP $X^0$ of a zero charge taking the mass and the strength of interaction with a nucleon as characterizing parameters. In Sec.~\ref{sec2} assumptions on the $X^0$ particle, estimations for binding energies between nuclei and an $X^0$, and rates of important reactions are described. Effects of the $X^0$ decay inside $X$-nuclei are not considered in our model. They should be addressed in the future. In Sec.~\ref{sec3} the destruction processes of $^7$Be and $^7$Li are identified. With results of the network calculations of BBN, we delineate the parameter region in which the $^7$Be and $^7$Li destructions possibly operate. If the $X^0$ particle interacts with nuclei strongly enough to drive $^7$Be destruction but not strongly enough to form a bound state with $^4$He of relative angular momentum $L=1$, then it might solve the $^7$Li problem of standard BBN. In Sec~\ref{sec4} conclusions of this work are summarized, and this model for $^7$Li reduction is compared with other models.

\label{sec4} We have investigated effects on BBN of a long-lived strongly interacting massive particle (SIMP) $X^0$ for different masses $m_X$ and strengths of $XN$ interaction, i.e., $\delta$. Binding energies of bound states of nuclei and an $X^0$ particle, i.e., $X$-nuclei, are calculated for two types of $XN$ potentials, i.e., Gaussian and well types. It is shown that calculated binding energies are not largely dependent upon the potential shapes, and are determined by the interaction strength at a given mass of $X^0$. Evolutions of light element abundances are calculated as a function of the temperature for two specific cases of relatively weak interaction strengths. We found that $^7$Be and $^7$Li can be destroyed by the nuclear capture reactions of free $X^0$ particles. The reactions identified as destruction processes are $X$($^7$Be$,^3$He)$^4$He$_X$ and $X$($^7$Li$,t$)$^4$He$_X$. We show that the lack of an excited state of $^4$He$_X$ with a relative angular momentum $L=1$ is essential for some fraction of the $X^0$ particles to escape capture by $^4$He. We suggest that the $^7$Li problem could be solved based upon a net work calculation of catalyzed BBN, and found the parameter region in the ($m_X$, $\delta$) plane where the $^7$Li problem can be fixed. We note that the results have been derived under the assumption that the $X^0$ particles do not change nuclear structure apart from sticking unaltered nuclei to the particles. This rough approximation is unlikely to be true especially in the case of relatively large strength of interaction since the $XN$ potential is not much weaker than the $NN$ potential and can not be neglected. More realistic estimations of wave functions and binding energies of $X$-nuclei need to include such changes in nuclear structures with the use of three or more-body models. In the case of strongly bound $^4$He, the binding energy of the ground state is 28 MeV which is larger than binding energy of $^4$He$_X$ with respect to separated $^4$He and $X^0$. The effect of change in nuclear density caused by the existence of the $X^0$ then tends to be small for $^4$He. The calculated binding energy of $^4$He$_X$ and the cross section for radiative capture of $X^0$ on $^4$He are, therefore, likely to be accurate. If the $XN$ and $XA$ potentials adopted in this paper describe well the real interaction, the main uncertainty in the BBN calculation would be in the estimations of nuclear reactions involving the $X^0$ particle. In this paper, the cross section of the most important reaction, i.e., $X$($^7$Be,$^3$He)$^4$He$_X$ was estimated only by analogy with $^6$Li($n,\alpha$)$^3$H, and many other cross sections were also estimated using standard nuclear reaction rates. Radiative reaction cross sections were estimated approximately within the framework of two-body models. The rates should, however, be calculated in more rigorous quantum many body models, not by the rough Born approximation, in order to derive realistic values. Although there might be errors in the calculated abundances in the $X$-catalyzed BBN of one order of magnitude or so, we argue at this moment that there is a possibility of $^7$Li reduction in the BBN model including a long-lived sub-SIMP. Effects of possible direct interactions of decay products of $X^0$ with the remaining nuclei $A$ at the decay of $X^0$ in an $X$-nucleus $A_X$ are not taken into account yet. They should be studied in the future in order to better estimate final abundances of the light elements. In addition, nonthermal nucleosynthesis triggered by the decay process of the $X^0$ particle might change the abundances of normal nuclei if the energy injection by the decay were large enough. Recent studies suggest that the radiative decay could lead to the production of $^6$Li to the level at most $\sim 10$ times larger than that observed in MPHSs when the decay life is of the order of $\sim 10^8 - 10^{12}$~s which is associated with $^3$He production~\cite{Kusakabe:2006hc}. The hadronic decay, on the other hand, can be a solution of both the lithium problems although that case gives somewhat elevated deuterium abundances~\cite{Jedamzik:2004er,Cumberbatch:2007me}. For example, we assume that the mass and the initial abundance of the $X$ are $m_X=100$~GeV and $Y_X=1.5\times 10^{-14}$, respectively. The energy injection at the decay of the $X$ into hadronic jets is constrained to be $\lesssim O$(1--100 GeV) if the lifetime is $\tau_X \gtrsim 10^2$~s through abundances of D, $^4$He, $^6$Li and $^3$He depending upon the lifetime (figure 38 of Ref.~\cite{Kawasaki:2004qu}). The energy injection into electromagnetic particles is, on the other hand, constrained to be $\lesssim O$(10 GeV) if the lifetime is $\tau_X\gtrsim 10^7$~s (figure 1 of Ref.~\cite{Kusakabe:2006hc}). This amount of energy injection tends to attain $^6$Li production up to the observed level in MPHSs. We summarize a present status of several models which have effects and thus leave observational signatures on primordial light element abundances. In the BBN model catalyzed by a long-lived sub-SIMP studied in this paper, the abundance of $^7$Li can be reduced below the level of SBBN prediction. In the BBN model catalyzed by a long-lived SIMP, the abundances of $^9$Be or B can be high~\cite{Kusakabe:2009jt}. Moreover, the isotopic abundance ratio, i.e., $^{10}$B/$^{11}$B can be high, which is never predicted in other scenarios for boron production~\cite{Kusakabe:2009jt}. In the BBN model catalyzed by a negatively charged massive particle (CHAMP), the $^6$Li abundance can be high~\cite{Pospelov:2006sc}. Only if the abundance of the CHAMP is more than $(0.04-1)$ times as large as that of baryon~\cite{Kusakabe:2010cb,Kusakabe2010inpc}, the $^7$Li reduction can be possible~\cite{Bird:2007ge}. A signature of CHAMP on $^9$Be abundance has been estimated to be negligible~\cite{Kusakabe:2010cb,Kusakabe2010inpc} in the light of a rigorous quantum mechanical investigation~\cite{Kamimura2010}. The cosmological cosmic ray nucleosynthesis triggered by supernova explosion in an early epoch of the structure formation can produce $^6$Li~\cite{Rollinde:2006zx} as well as $^9$Be and $^{10,11}$B~\cite{Kusakabe2008,Rollinde2008}. In baryon-inhomogeneous BBN models, the abundance of $^9$Be can be higher than in the SBBN~\cite{Boyd1989,Kajino1990a,Kajino1990b,Coc1993,Orito1997}. \appendix*

1012.0435

We identify reactions which destroy Be7 and Li7 during big bang nucleosynthesis (BBN) in the scenario of BBN catalyzed by a long-lived sub-strongly-interacting massive particle (sub-SIMP or X particle). The destruction associated with nonradiative X captures of the nuclei can be realized only if the interaction strength between an X particle and a nucleon is properly weaker than that between two nucleons to a degree depending on the mass of X. Binding energies of nuclei to an X particle are estimated taking the mass and the interaction strength to nuclei of the X as input parameters. Nuclear reaction rates associated with the X are estimated naively and adopted in calculating evolutions of nuclear abundances. We suggest that the Li7 problem, which might be associated with as-yet-unrecognized particle processes operating during BBN, can be solved if the X particle interacts with nuclei strongly enough to drive Be7 destruction but not strongly enough to form a bound state with He4 of relative angular momentum L=1. Justifications of this scenario by rigorous calculations of reaction rates using quantum mechanical many-body models are highly desirable since this result involves many significant uncertainties.

12135967

2010EAS....43..167N

1012

1012.0603_arXiv.txt

Many fields of contemporary astrophysics require a highly-precise determination of basic stellar parameters. A prime example is the investigation of stellar structure and evolution, where asteroseismology has opened up a window to study the governing physical processes in detail (e.g.~Aerts et al.~\cite{aerts08}). Other examples are observational constraints on planet-formation theories by analyses of planet host stars (e.g.~Neves et al.~\cite{neves09}). Predictions of super-/hyper--nova models may be tested via analyses of secondary stars in low-mass X-ray binaries (e.g.~Gonz\'alez Hern\'andez et al.~\cite{GRI08}) and (hyper-)runaway stars (Przybilla et al.~\cite{PNHB08}). Tight constraints on Galactochemical evolution theories may be derived using stars as tracers of the variation of chemical composition with time and location in the Galactic disk (e.g.~Fuhrmann~\cite{fuhrmann08}; Przybilla~\cite{P08}). Precise stellar parameter determinations are decisive in describing the stellar atmospheric structure correctly, i.e.~the temperature gradient and density stratification. This is crucial for the regions where the continuum and spectral lines of interest are formed. A fundamental problem is that a direct determination of stellar parameters is usually not possible. Quantitative spectral analysis has to rely on the comparison of synthetic spectra with observation, which in turn guides computations of improved models. In principle, this requires an iterative approach. In practice, however, approximate methods for the parameter determination are favoured on most occasions to shorten the process, by e.g.~use of photometric indicators or comparison with prescribed (and thus necessarily limited) grids of models. Unfortunately, the gain in efficiency is often accompanied by a loss in accuracy. Quantitative spectroscopy is highly prone to systematic errors. The analysis methodology, the quality of the observed spectra and the details of the assumptions made in the modelling, when taken together, determine the accuracy of a work. In particular, relaxing the assumption of local thermodynamic equilibrium (LTE) for the spectrum synthesis is crucial for improving the accuracy of stellar analyses. In what follows we describe the basic requirements for a highly accurate determination of stellar atmospheric parameters by utilising all available lines of hydrogen and helium, plus multiple metal ionization equilibria, in non-LTE. We concentrate on OB-type stars near the main sequence and on their evolved progeny, BA-type supergiants. The stellar mass range between $\sim$8 and 25\,M$_\odot$ is covered for effective temperatures from $\sim$8000 to 35\,000\,K. The models and analysis technique have been developed and thoroughly tested during the past years (Przybilla et al.~\cite{PBBK06}, PBBK06 henceforth; Nieva \& Przybilla~\cite{NP07,NP08}, NP07/NP08; Przybilla, Nieva \& Butler~\cite{PNB08}, PNB08). One should keep in mind that the same philosophy can, in principle, be applied to any type of star once the underlying stellar atmospheric physics is reproduced consistently. We start with an overview of requirements on atmospheric models, the spectrum synthesis and the quality of observed spectra for a reliable quantitative analysis (Sects.~\ref{modelrequirements} \& \ref{observations}). Then, a self-consistent spectrum analysis technique is described (Sect.~\ref{analysis}). Concrete examples of stellar parameter determinations via non-LTE spectral analyses are discussed (Sect.~\ref{examples}). Finally, common sources of systematic error are briefly addressed (Sect.~\ref{errors}).

Highly-precise and accurate stellar parameters {\it can} be spectroscopically determined, with limiting uncertainties as low as $\sim$1\% in $T_\mathrm{eff}$, $\sim$0.05--0.10\,dex in $\log\,g$ and $\sim$10-20\% in elemental abundances (rms scatter). A self-consistent spectral analysis methodology using non-LTE line formation was presented that allows this to be achieved when typical systematics are avoided. Of crucial importance is to simultaneously bring multiple spectroscopic indicators into agreement, which requires an iterative approach. The method is much more time-consuming than standard approaches for the stellar parameter determination, but it is worth the effort whenever highly-accurate observational constraints are required for astrophysical applications.

1012.0603

We describe a self-consistent spectrum analysis technique employing non-LTE line formation, which allows precise atmospheric parameters of massive stars to be derived: 1σ-uncertainties as low as ~1% in effective temperature and ~0.05-0.10 dex in surface gravity can be achieved. Special emphasis is given to the minimisation of the main sources of systematic errors in the atmospheric model computation, the observed spectra and the quantitative spectral analysis. Examples of applications are discussed for OB-type stars near the main sequence and their evolved progeny, the BA-type supergiants, covering masses of ~8 to 25 M<SUB>⊙</SUB> and a range in effective temperature from ~8000 to 35000 K. Relaxing the assumption of local thermodynamic equilibrium in stellar spectral synthesis has been shown to be decisive for improving the accuracy of quantitative analyses. Despite the present examples, which concentrate on hot, massive stars, the same philosophy can be applied to line-formation calculations for all types of stars, including cooler objects like the Sun, once the underlying stellar atmospheric physics is reproduced consistently.

12214136

2011MNRAS.415.3951F

1012

1012.0573_arXiv.txt

Pulsars in binary systems provide accurate clocks that can be used to infer the properties of the binary orbit and its stellar components with precision. In addition to yielding important constraints on stellar evolution, binary pulsars have been used to test the validity of general relativity and to put stringent constraints on alternative theories of gravity \citep[see, e.g., the review article by][]{2004Sci...304..547S}. The original binary pulsar PSR B1913+16 \citep[][]{1975ApJ...195L..51H}, consisting of a pulsar in orbit around another unseen neutron star, permitted the first tests of general relativity in the strong field regime and the most compelling (albeit indirect) evidence for the gravitational wave emission, via the measurement of its orbital period derivative. Today, the best strong-field tests of general relativity are provided by the ``double pulsar'' PSR J0737-3039 \citep[][]{2003Natur.426..531B}, the only known example of two pulsars in orbit around each other \citep[][]{2008ARA&A..46..541K}. While theoretical considerations suggest that a certain fraction of pulsars should have black hole (BH) companions \citep[e.g.,][]{1991ApJ...379L..17N, 1991ApJ...380L..17P, 1998A&A...332..173P, 1999ApJ...517..318B, 2002ApJ...572..962S, 2002ApJ...572..407B, 2005ApJ...628..343P} and pulsar-black hole (PSR-BH) binaries are a holy grail of pulsar astronomy, none has been found to date. The discovery of such a system would represent a major step forward both from the astrophysical point of view and for the unique gravity tests that it might allow. Clues as to where PSR-BH systems might be found are therefore highly desirable. Most theoretical studies have focused on PSR-BHs formed from primordial binaries in the Galactic field. As dense systems conducive to interactions between stars, globular clusters have also been proposed as sites where PSR-BH binaries may reside \citep[][]{2003ASPC..302..391S}. Indeed, the high rates of stellar encounters in globular clusters explain the overdensity of recycled millisecond pulsars\footnote{A recycled pulsar is one that has been spun up after birth by accretion of material from a binary companion.} (MSPs) and X-ray binaries in globular clusters, by orders of magnitude, relative to the field \citep[e.g.,][]{1987ApJ...312L..23V, 1991A&A...241..137H, 2000ApJ...532L..47R, 2003ApJ...591L.131P, 2006ApJ...646L.143P}. It is natural to expect that analogous processes might lead to the formation of PSR-BH systems. An important limitation of this argument, however, is that BHs may be excessively rare in globular clusters. To date, there is no consensus on whether globular clusters host massive central BHs similar to those commonly found at the centers of galaxies \citep[e.g.,][]{2005ApJ...634.1093G, 2008ApJ...676.1008N, 2010ApJ...710.1063V}. Furthermore, stellar-mass BHs are also expected to be few in globular clusters as a result of dynamical interactions that eject them \citep[][]{1993Natur.364..423S, 1993Natur.364..421K, 2004ApJ...601L.171K, 2009ApJ...690.1370M}. In this paper, we consider the formation of PSR-BH binaries in another kind of dense stellar environment, the center of our Milky Way. Indeed, the nuclear star cluster at the center of our Galaxy should allow many of the stellar interaction processes taking place in globular clusters to operate, and in addition may be particularly rich in BHs. In fact, the surrounding bulge provides a large reservoir of stellar remnants. Because stellar ($\sim10$ M$_{\odot}$) black holes are the most massive components of old stellar populations, they sink to the center owing to dynamical friction \citep[][]{1993ApJ...408..496M, 2000ApJ...545..847M, 2006ApJ...649...91F}. \cite{2000ApJ...545..847M} showed that about 25,000 stellar are expected to have migrated into the central parsec in this way. Here, we show that these BHs are likely to form binaries with pulsars via exchange interactions, and that several MSP-BH binaries should survive in the central parsec today. This makes the Galactic center (GC) a promising experimental target not only for pulsars orbiting the central supermassive black hole \citep[Sgr A$^{*}$;][]{1979IAUS...84..401P, 1997ApJ...475..557C, 2004ApJ...615..253P}, but also stellar-mass BHs. Similar processes should operate in the nuclei of other galaxies, and may actually be more efficient in low-mass galaxies with no or low-mass central black holes \citep[e.g.,][]{2009ApJ...692..917M}. Because of its proximity, the center of the Milky Way is however the most interesting observationally at present. No pulsar has been found in the central parsec so far \citep[][]{2009ApJ...702L.177D}. This dearth of observed pulsars in the GC is understood to arise from the strong scattering of radio waves by the dense, turbulent, ionized plasma in that region \citep[e.g.,][]{1998ApJ...505..715L, 2006MNRAS.373L...6J}, rather than by an intrinsic absence of pulsars. If the scattering of radio pulses broadens them by a time scale comparable to or longer than the pulse period, the signal becomes effectively constant and the pulsar disappears. Since the scattering broadening time scales with frequency as $\tau_{\rm scat} \propto \nu^{-4}$ \citep[e.g.,][]{1997ApJ...475..557C}, radio pulsars can nevertheless in principle be detected in the GC by observing at sufficiently high frequency. Recently, \cite{2010ApJ...715..939M} reported on a 15 GHz search for radio pulsars in the central parsec with the Green Bank Telescope (GBT), the highest frequency search for pulsars toward the GC to date. Observations at this frequency are sensitive to pulsars with periods $\gtrsim50$ ms. Although it was not confirmed in subsequent data taken in 2008, \cite{2010ApJ...715..939M} did detect a 607 ms pulsar candidate in their 2006 observations. It must be noted that the 2008 non-detection does not rule out that the candidate is a real pulsar, as a source in a short-period orbit ($\lesssim$100 yr) orbit around the GC could have its emission beam rotated away from our line of sight on a 2-yr time scale. Alternatively, the pulsar could have moved sufficiently in the inhomogeneous GC plasma to induce significant variations in its pulsed flux, owing either to scattering or scintillation. Regardless of the reality of this pulsar candidate, the upper limit of 90 ordinary pulsars in the central parsec inferred by \cite{2010ApJ...715..939M} is comparable to the expected number based on the simple expectation that there should be about as many ordinary pulsars as progenitor massive stars of mass $>8$ M$_{\odot}$ \citep[e.g.,][]{2010ApJ...708..834B}, since their lifetimes are comparable. Further searches are therefore highly warranted, as even current technology stands a good chance of detecting radio pulsars in the central parsec. For our purpose, it is especially noteworthy that even the \cite{2010ApJ...715..939M} 15 GHz periodicity search, and by extension all previous searches at lower frequencies, were completely insensitive to MSPs with pulse periods $P\approx1-$few ms. Existing observations thus would not have detected the MSP-BH binaries that we predict. Nevertheless, pulsars have been detected at a variety of higher frequencies, from 32 GHz \citep[e.g.,][]{2008A&A...480..623L}, to 43 GHz \citep[e.g.,][]{1997ApJ...488..364K}, to 87 GHz \citep[e.g.,][]{1997A&A...322L..17M}, to 144 GHz \citep[e.g.,][]{2007ApJ...669..561C}. Furthermore, the radio spectra of MSPs are similar to those of ordinary pulsars \citep[e.g.,][]{1998ApJ...501..270K, 1998ApJ...506..863T}, suggesting that radio observations at frequencies high enough to beat down interstellar scattering in the GC for millisecond periods are possible. In theory, GC MSPs could also be observed in higher-energy bands, including in the X$-$rays and $\gamma-$rays. In the rest of this paper, we describe our new formation scenario for MSP-BH binaries in the Galactic center and present simple estimates of the formation rate and survival of these systems (\S \ref{model}). Because of the technical challenges of simulating the dynamics of the GC, including the critical binary processes and the effects of different stellar masses \citep[e.g.,][]{2006ApJ...649...91F, 2009ApJ...700.1933H}, and since substantial uncertainties exist, we focus here on basic analytic considerations. Our here is to outline the relevant physical processes and to combine them in order-of-magnitude estimates for the number of MSP-BH binaries that should survive in the GC today, as a motivation for observational efforts as well as further theoretical studies. In \S \ref{discussion}, we discuss the prospects for detecting the predicted systems, the unique signatures of the proposed formation channel, and the potential implications for physics and astrophysics.

\label{discussion} \subsection{Observational Prospects} \label{observational prospects} Our results suggest that our Galactic center should harbor several MSP-BH binaries, providing a strong motivations for focused searches in this direction. As no pulsar has so far been detected in the central parsec, it is important to consider whether finding MSP-BHs there is feasible. The time scale by which a radio pulse is broadened owing to interstellar scattering is \begin{equation} \tau_{\rm scat} = 116~{\rm ms} \left( \frac{D_{\rm scat}}{100~{\rm pc}} \right)^{-1} \left( \frac{\nu}{10~{\rm GHz}} \right)^{-4}, \end{equation} where $D_{\rm scat}$ is the distance of the effective scattering screen from the GC. Combining all known tracers of ionized gas, \cite{1998ApJ...505..715L} found $D_{\rm scat}=133^{+200}_{-80}$ pc; we assume $D_{\rm scat}=100$ pc. Sensitivity to a pulsed signal of period $P$ requires $\tau_{\rm scat}\lesssim P$. For a period $P=5$ ms at the GC, the minimum observing frequency is therefore $22$ GHz, while $33$ GHz is sufficient to detect a 1 ms source. Let us consider a hypothetical 30 GHz search, which would suffice to detect most MSPs if sufficiently deep. The minimum flux density for a pulsar detection is \begin{equation} \label{radiometer equation} S_{\rm min}=\delta_{\rm beam} \frac{\beta_{\rm sys} \sigma (T_{\rm rec}+T_{\rm sky})}{G\sqrt{N_{p}\Delta \nu t_{\rm int}}}\sqrt{\frac{W_{e}}{P-W_{e}}}, \end{equation} where $\delta_{\rm beam}$ is a factor accounting for the reduction in sensitivity to pulsars located away from the center of the telescope beam, $T_{\rm rec}$ is the receiver temperature on a cold sky, $T_{\rm sky}$ is the sky background temperature, $G$ is the antenna gain, $N_{p}$ is the number of polarizations summed, $\Delta \nu$ is the receiver bandwidth, $t_{\rm int}$ is the integration time, $P$ is the pulse period, $W_{e}$ is the effective pulse width, $\sigma$ is the signal-to-noise (S/N) detection threshold, and $\beta$ is a constant accounting for various system losses \citep[e.g.,][]{dss+84}. \cite{2010ApJ...715..939M} calculated the S/N expected as a function of observing frequency for a 10-hour observations with the GBT with a bandwidth $\Delta \nu = 800$ MHz, for different pulsar periods and spectral properties. We can use equation (\ref{radiometer equation}) to scale their results and explore the potential capabilities of future searches. \cite{2010ApJ...715..939M} predict a $S/N=1$ for a 10-hour, 30 GHz GBT integration with $\Delta \nu=800$ MHz, for a 5-ms pulsar with a typical spectral index $\alpha=-1.7$ ($S_{\nu}\propto \nu^{\alpha}$) and flux density of 1 mJy at 1 GHz (corresponding to $L_{1000}=S_{1000}d^{2}=72.25$ mJy kpc$^{2}$ at the GC, $d=8.5$ kpc), and $S/N=10$ if the spectral index is instead -1. For reference, the observed spectral index distribution is well modeled by a normal distribution with mean $-1.7$ and dispersion $0.35$ \citep[e.g.,][]{2009A&A...505..919S}. To get a sense of the limits of a dedicated GBT campaign with upgraded instrumentation, we consider a 100-hr integration with a $\Delta \nu=8$ GHz bandwidth instead\footnote{A spectrometer capable of such high-frequency, large bandwidth observations is already being developed for the GBT (http://www.gb.nrao.edu/gbsapp/).}, in which case the same S/Ns are achieved for a pulsar with flux density lower by a factor of 10. A futuristic Square Kilometer Array\footnote{http://www.skatelescope.org} (SKA) telescope would have an effective area larger by a factor of $\sim100$, and hence would be sensitive to pulsars $\sim100\times$ fainter still. Thus, an upgraded GBT campaign could in principle detect a $L_{1000}\approx7$ mJy kpc$^{2}$, $\alpha=-1$, 5-ms MSP at the GC with $S/N=10$, while the SKA could in principle probe as faint as $L_{1000}\approx0.7$ mJy kpc$^{2}$ under the same conditions. For a more typical spectral index $\alpha=1.7$, the same S/N is reached for pulsars more luminous at 1 GHz by a factor of 10. For comparison, the $least~luminous$ MSPs currently detected in deep globular cluster observations have $L_{1000}\sim1$ mJy kpc$^{2}$ and the luminosity function extends beyond at least $L_{1000}\sim100$ mJy kpc$^{2}$ \citep[e.g.,][]{2007ApJ...670..363H}. A dedicated GBT search should therefore be capable of detecting MSPs at the GC, while the SKA could potentially probe a substantial fraction of the population. \emph{Fermi} has demonstrated that MSPs are $\gamma-$ray sources \citep[e.g.,][]{2009Sci...325..848A, 2010JCAP...01..005F}. The $\gamma-$ray energy band has the advantage that, unlike the radio, it is not affected by interstellar scattering. However, even \emph{Fermi} can only detect the very brightest, tip-of-the-iceberg $\gamma-$ray pulsars at the GC, so it is unlikely that a significant number of individual $\gamma-$ray MSPs will detected in the central parsec for the foreseeable future. \subsection{Signatures of the Proposed Formation Channel} \label{signatures} Our calculations make definite predictions for the properties of the PSR-BHs formed via the proposed channel: \begin{enumerate} \item The system would be found within $\sim1$ pc of Sgr A$^{\star}$, where the cluster of stellar BHs is assumed to be located. \item The pulsar would be recycled, with a period from $\sim1$ to a few tens of milliseconds, in order to have an original WD companion and to have a sufficiently long spin down time to remain emitting in the radio to the present day. Accordingly, the pulsar should have a low magnetic field, $B\lesssim10^{10}$ G. \item The MSP-BH binary would be relatively wide, as the semi-major axis is multiplied by a factor $\sim M_{\rm BH}/M_{\rm WD}\gtrsim10$ during the exchange interaction. For an original MSP-WD semi-major axis distribution ranging from $\sim0.01$ AU to $\sim0.3$ AU, the MSP-BHs should have semi-major axes ranging from $\sim0.1$ to $\gtrsim3$ AU. \item The MSP-WD would be highly eccentric, $1-e'\sim(M_{\rm WD}/M_{\rm BH})(M_{\rm BH}/M_{\rm NS})^{1/3}$ (\S \ref{survival time}). For our fiducial masses, $e'\sim0.8$, but the eccentricity distribution will be broad. \end{enumerate} These predictions are based on the properties of the MSP-BHs immediately after their formation in exchange interactions. In principle, their internal properties could evolve by the time they are detected. Simple estimates however suggest that they should not change qualitatively. As discussed in \S \ref{survival time}, the spin down time scales of recycled pulsars are typically of order of a few Gyr. In the soft regime, the binary evaporation time $t_{\rm evap}\propto a_{\rm b}'^{-1}$ (eq. (\ref{evaporation time scale})), and so the binaries spend most of their lifetime with a semi-major axis of the order of its value at formation. In the hard regime, the binary semi-major axis is progressively reduced by hardening, but the time scale for this process is comparable to the exchange interaction time scale \citep[e.g.,][]{1993ApJS...85..347H}, which for MSP-BHs at the GC is $\gtrsim10$ Gyr (Fig. \ref{3body time scales}). These simple considerations are in good agreement with the Fokker-Planck modeling of binaries at the GC by \cite{2009ApJ...700.1933H}, who found little evolution of the internal binary properties. Thus, ``Heggie's law'' (according to which soft binaries become softer and hard binaries become harder) holds but the time scales are too long for it to have a large impact. We use the cross section derived by \cite{1996MNRAS.282.1064H} to estimate the time it takes for encounters with other stars to change the eccentricity by an amount $>\delta e_{0}$, starting with a value $e$. Averaging over a Maxwellian velocity distribution, \begin{align} t_{e}(>\delta e_{0}) \approx 0.29 \frac{\sigma_{3}}{G n_{3} M_{123} a_{b}'} & \left( \frac{M_{12}M_{123}}{m_{3}^{2}} \right)^{1/3} \notag \\ & \times \frac{\delta e_{0}^{2/3}}{e^{2/3}(1 - e^{2})^{1/3}}, \end{align} where $M_{12}\equiv m_{1}+m_{2}$ and $M_{123}\equiv m_{1}+m_{2}+m_{3}$. Comparing with the ionization time scale given by equation (\ref{ionization time scale}), we find \begin{equation} \frac{t_{e}(>\delta e_{0})}{t_{\rm ion}} \approx 9.69 \frac{m_{3}^{4/3}}{M_{123}^{2/3}M_{12}^{2/3}} \frac{\delta e_{0}^{2/3}}{e^{2/3}(1 - e^{2})^{1/3}}. \end{equation} For our fiducial MSP-BH mass choices, $e=0.8$ and $\delta e_{0}=0.1$, $t_{e}(>\delta e_{0})/t_{\rm ion}\approx1.9$, indicating that the eccentricity changes should also be modest over the lifetimes of the binaries. It is also noteworthy that even if eccentricity perturbations occurred on a shorter time scale, the systems would be asymptotically driven to a thermal distribution, $f_{\rm th}(e)=2e$ with mean $\langle e \rangle_{\rm th} = 0.67$, so that the eccentricities would remain large on average. \subsection{Applications} \label{applications} The discovery of a MSP-BH system at the GC would have profound implications. At this time, it would be the first known PSR-BH binary and would therefore critically inform our theories of stellar evolution, especially in dense stellar environments. If it has the properties outlined above, it would convincingly originate in an exchange scenario similar to the one we have explored. It would therefore provide strong evidence for the existence of the predicted cluster of stellar BHs around Sgr A$^{\star}$, for which there is currently no observational confirmation. This would in turn have important implications for a number of other phenomena, including HVSs \citep[e.g.,][]{2008MNRAS.383...86O}, gravitational wave sources \citep[particularly for the extreme mass ratio events to be detected by \emph{LISA}\footnote{http://lisa.nasa.gov} but also for \emph{LIGO}\footnote{http://www.ligo.caltech.edu}; e.g.,][]{2007PhRvD..75d2003B, 2009MNRAS.395.2127O, 2010arXiv1010.5781A}, microlensing near Sgr A$^{\star}$ \citep[e.g.,][]{2001ApJ...563..793C, 2001ApJ...551..223A}, and the orbital capture of stars by Sgr A$^{\star}$ \citep[e.g.,][]{2004ApJ...606L..21A}. It would also for the first time provide an accurate clock orbiting a BH and therefore offer a unique probe of the spacetime in a BH potential \citep[e.g.,][]{2003LRR.....6....5S, 2007mru..confE..20K}. While the MSP-BHs formed in this scenario would not be as tightly bound as the relativistic double NS systems from which the best constraints on gravity theories are currently derived \cite[e.g.,][]{2006Sci...314...97K, 2010ApJ...722.1030W}, certain relativistic effects including the periastron shift and the Shapiro time delay can be measured even in relatively loose systems \citep[e.g.,][]{1988ApJ...332..770T, 1991ApJ...371..739R, 1991ApJ...379L..17N}. Measuring these effects would provide much more accurate measurements of the masses of stellar BHs than currently possible in X-ray binaries, in which assumptions regarding the inclination of the binary must be made. Another interesting possibility for nearly edge-on systems would be to measure the gravitational lensing effects of the BH using pulse delays, from which the orientation and spin of the BH could be inferred \citep[e.g.,][]{1991ApJ...379L..17N, 2006ApJ...636L.109B}. \subsection{Caveats} \label{caveats} While our calculations provide straightforward predictions for where and how to find MSP-BH binaries, of their expected properties, and suggest that they should be observationally accessible, they are limited in some respects. As a first investigation of the proposed formation channel, we have focused on simple analytic estimates of the relevant physical processes, and modeled the GC as a steady-state background in which MSP-BHs are created and destroyed. In reality, the GC is a dynamical system in which complicated stellar evolution processes are ongoing. It would therefore be desirable, as the capabilities become available in the future, to perform more detailed dynamical simulations that include the critical binary and stellar evolution effects. Some of the assumptions we have made have also not yet been empirically confirmed. In particular, our results rely on the presence of a central cluster of $\sim25,000$ stellar BHs induced by mass segregation, but this cluster may not be present if relaxation is too inefficient or if a binary massive black hole recently destroyed it \citep[e.g.,][]{2010ApJ...718..739M}. While the disk morphology of the Milky Way argues against a major merger in the last $\sim10$ Gyr, intermediate mass black holes brought in by dwarf galaxies or globular clusters cannibalized by the bulge could also eject stellar BHs from the central cluster. Resonant relaxation, which we have mostly neglected, could furthermore be more efficient at depleting the stellar BH cluster or binaries near Sgr A$^{\star}$ than we have assumed. It will be important to consider these issues in more detail in the future, although detecting MSP-BHs in the central parsec would actually inform these open questions. \subsection{Conclusion} \label{conclusion} We have shown that if dynamical processes analogous to those operating in dense globular clusters occur in the central parsec of our Milky Way, and if this region hosts a cluster of $\sim25,000$ stellar-mass BHs as predicted by mass segregation arguments, then MSP-BH binaries should be formed there in exchange interactions. Taking into account the much higher retention fraction of neutron stars in the Galactic center relative to globular clusters, as a result of the deeper potential well, several of these systems should survive to the present day. Simple scaling arguments indicate that some of these MSP-BHs might be detectable with existing radio telescopes, and that a substantial fraction of the population should be accessible to the Square Kilometer Array. Our predictions therefore suggest a definite roadmap to the detection of the first pulsar-black hole binary, by singling out a small region of the Galaxy where many might reside. In light of the remarkable potential payoffs of such a discovery, it is therefore clear that focused observational searches and further theoretical studies are warranted.

1012.0573

Binaries consisting of a pulsar and a black hole (BH) are a holy grail of astrophysics, both for their significance for stellar evolution and for their potential application as probes of strong gravity. In spite of extensive surveys of our Galaxy and its system of globular clusters, no pulsar-black hole (PSR-BH) binary has been found to date. Clues as to where such systems might exist are therefore important. We show that if the central parsec around Sgr A<SUP>★</SUP> harbours a cluster of ∼25 000 stellar BHs (as predicted by mass-segregation arguments) and if it is also rich in recycled pulsar binaries (by analogy with globular clusters), then three-body exchange interactions should produce PSR-BHs in the Galactic Centre. Simple estimates of the formation rate and survival time of these binaries suggest that a few PSR-BHs should be present in the central parsec today. The proposed formation mechanism makes unique predictions for the PSR-BH properties: (1) the binary would reside within ∼1 pc of Sgr A<SUP>★</SUP>; (2) the pulsar would be recycled, with a period of ∼1 to a few tens of milliseconds, and a low magnetic field B≲ 10<SUP>10</SUP> G; (3) the binary would have high eccentricity, e∼ 0.8, but with a large scatter and (4) the binary would be relatively wide, with semimajor axis a<SUB>b</SUB>∼ 0.1 -≳3 au. The potential discovery of a PSR-BH binary therefore provides a strong motivation for deep, high-frequency radio searches for recycled pulsars towards the Galactic Centre.

12213219

2011MNRAS.413L..51K

1012

1012.0059_arXiv.txt

How galaxies get their gas is a long-standing issue. For decades, the standard theoretical picture of galaxy formation has stipulated that all gas accreted into dark matter haloes is shock heated before it radiatively cools and settles into a galactic disk \citep[][although \citealt{binney77} first suggested this need not be the case]{silk77,rees77}. This picture has recently been revisited both by analytic studies \citep{birnboim03,dekel06} and hydrodynamical simulations (both in 1D: \citealt{birnboim03} and in 3D within an explicit cosmological context, \citealt{keres05,keres09,ocvirk08}; hereafter OPT08). These studies have established that in haloes below a critical mass shocks are unstable and cannot propagate outwards, so that cold diffuse gas and/or cold filaments can penetrate deep into the halo without experiencing shock-heating. In contrast, at the other end of the mass spectrum, very massive haloes easily sustain a virial shock that is stable against gas cooling so that diffuse/filamentary gas is shock heated to the virial temperature of the halo as it enters. Finally, at intermediate halo masses, either gravitationally shock-heated hot gas and/or a hot galactic wind coexists with cold inflowing filaments (OPT08): some dense cold filaments are stable against the pressure force exerted by the hot gaseous material. In other words, the vast majority of galaxy-size host halos, especially at high redshift ($z>2$), are predicted to be threaded by cold gas filaments in a $\Lambda$CDM model of structure formation. Therefore, the question which naturally arises is whether or not the existence of these cold gas filaments can be observationally confirmed, and by which technique. Since the advent of high-resolution spectroscopy has enabled astronomers to study the kinematics of the intergalactic medium at high redshifts in a wealth of detail \citep[e.g.][]{pettini02,adelberger03,shapley03}, and given the fact that cold gas is thought to flow into massive haloes ($10^{12} \msun$) at $z=2.5$ along filaments with velocities of $\gtrsim 200\,{\rm km s^{-1}}$ \citep{dekel09}, it is sensible to think that the spectra of Lyman-break galaxies \citep{steidel96} might reveal these filaments as redshifted absorption features. Interestingly, a recent study reported that few Lyman-break galaxies (LBGs) show redshifted metal absorption lines, suggesting that inflowing gas in haloes with $4\times10^{11}<M_{\rm vir}<10^{12}M_{\odot}$ is rare \citep{steidel10}. Taken at face value, this appears to contradict the theoretical prediction that cold filaments are prominent in the vast majority of high-z halos. However, there exist a variety of reasons as to why these filaments should be very difficult to detect. The first of these is the covering fraction of the filaments. This is estimated in \citep{dekel09} to be around ($\sim 25\%$) for four massive halos with $M_{\rm vir}\sim10^{12}M_{\odot}$ in the \mn\ simulation, counting only relatively dense ($N_{{\rm H}}>10^{20}{\rm cm}^{-2}$) and cold ($T<10^{5}K$) gas within a radial distance $20<r<100$ kpc from the central galaxy hosted by these halos. Whereas $25\%$ is indeed a non-negligible covering fraction, its exact dependence on redshift and halo mass remains to be determined. For instance, \citet{faucher10a} also measured the covering fraction for a Milky way-type progenitor LBG in a smoothed particle hydrodynamic simulation but found a significantly smaller value ($\sim 2\%$). Secondly, low-ionisation lines can be produced not only by cold filamentary gas but also by a galaxy's interstellar medium, so that when using a single galaxy to probe the circumgalactic medium, distinguishing absorption produced by filaments from that produced by the ISM is key to prove/disprove the presence of cold filaments. Thirdly, there is the geometry of the accretion: in order to produce a strong absorption signal, a filament needs to be well aligned with the line of sight to maximise its column density. Finally, there is the issue of metallicity: a metal-poor filament will be transparent to metal line observations. The aim of this letter is to quantify the aforementioned effects to better assess the detectability of the absorption signal produced by cold filaments. We show that the actual probability of detecting such flows with metal lines is much smaller than the high covering fraction derived in \citet{dekel09} would suggest due to i) the low density and ii) low metallicity of the filaments {\em compared with the densities and metallicites of the interstellar medium of the host galaxy. }

Cosmological simulations predict that high-z galaxies grow by acquiring gas from cold streams \citep{dekel09}, but no observational confirmation has been obtained yet. Based on a statistical sample from the \mn\ simulation, we argue that low-ionisation metal absorption features, such as CII $\lambda1334$, arising from intervening cold filaments are extremely hard to distinguish from absorption by the ISM of high-z star-forming galaxies. This is primarily because the optical depth for the low-ionisation transition from cold filamentary gas is minuscule, compared with that of the ISM of the host galaxy. Moreover, the filamentary absorption is not redshifted enough with respect to the ISM absorption, so that the residual ISM absorption in the red wing of the line is still prominent. This small optical depth of filaments mainly finds its source in the intrinsically low densities and metallicities of the cold gas when compared to those of the galaxy's ISM. Another factor is the geometry of the flow which lowers the probability of detecting filaments as their column density will rarely be maximised by being aligned with the line of sight. As an alternative to using a single galaxy, one could probe circumgalactic regions by using a paired background galaxy. This method alleviates the importance of ISM absorption since when probing circumgalactic regions in this way, the absorption by the ISM of the foreground galaxy will occur at a different spatial position from that produced by filaments. Unfortunately, the rare occurrence of suitable foreground-background galaxy pairs makes it difficult to probe more than one line of sight per foreground galaxy. As a result, high resolution individual spectra are hard to obtain and one has to resort to stacking the spectra of multiple galaxies \citep{steidel10}. Contrary to what might be expected, stacking will wash out the cold filament absorption signal since absorption by inflowing gas does not neatly separate from that caused by outflows as was the case when probing the circumgalactic region using single galaxies. Indeed, cold filament absorption against the background galaxy light will not only be redshifted, but also blueshifted as one expects that on average as many cold streams will be detected moving towards as away from the observer. Absorption by outflows will also suffer the same fate. This will make it all the more difficult to argue whether the observed absorption features are driven by infalling gas or by outflows from high-z star-forming galaxies. Moreoever, the metal column density of cold filaments is minuscule, as already mentioned. The hydrogen column densities of the cold filaments are distributed around $10^{20} {\rm cm^{-2}}$ at these redshifts, yielding corresponding carbon column densities around $10^{16} {\rm cm^{-2}}$ if an average $Z=0.001 Z_{\odot}$ is used. Even if all carbon atoms are assumed to be eligible for the CII transition, the optical depth is only around $10^{-2}$ in the line. On the other hand, outflows are expected to be very metal rich, so that even though higher transitions like CIV are expected to dominate the absorption signal, the amount of CII absorption from these outflows might still swamp that produced by the cold filaments. Finally, we have assessed possible numerical resolution issues on column density and metallicity of the filaments using the very high resolution numerical simulation \nut\ suite and found that our conclusions remain by and large unchanged. Based on these considerations, we conclude that the presence of the cold filament is difficult to disprove/prove with low-ionisation metal line absorption. Instead, the Lyman $\alpha$ emission route seems more promising to detect cold filaments, but the line profile is more sensitive to the kinematics of the intervening gas \citep{verhamme06,verhamme08}. Therefore this will require full blown radiative transfer calculations \citep[e.g.][]{faucher10b}.

1012.0059

Cosmological simulations have shown that dark matter haloes are connected to each other by large-scale filamentary structures. Cold gas flowing within this ‘cosmic web’ is believed to be an important source of fuel for star formation at high redshift. However, the presence of such filamentary gas has never been observationally confirmed despite the fact that its covering fraction within massive haloes at high redshift is predicted to be significant (∼25 per cent). In this Letter, we investigate in detail whether such cold gas is detectable using low-ionization metal absorption lines, such as C IIλ1334, as this technique has a proven observational record for detecting gaseous structures. Using a large statistical sample of galaxies from the MARENOSTRUM N-body+ adaptive mesh refinement (AMR) cosmological simulation, we find that the typical covering fraction of the dense, cold gas in 10<SUP>12</SUP> M<SUB>⊙</SUB> haloes at z∼ 2.5 is lower than expected (∼5 per cent). In addition, the absorption signal by the interstellar medium of the galaxy itself turns out to be so deep and so broad in velocity space that it completely drowns that of the filamentary gas. A detectable signal might be obtained from a cold filament exactly aligned with the line of sight, but this configuration is so unlikely that it would require surveying an overwhelmingly large number of candidate galaxies to tease it out. Finally, the predicted metallicity of the cold gas in filaments is extremely low (≤10<SUP>-3</SUP> Z<SUB>⊙</SUB>). If this result persists when higher resolution runs are performed, it would significantly increase the difficulty of detecting filamentary gas inflows using metal lines. However, even if we assume that filaments are enriched to Z<SUB>⊙</SUB>, the absorption signal that we compute is still weak. We are therefore led to conclude that it is extremely difficult to observationally prove or disprove the presence of cold filaments as the favourite accretion mode of galaxies using low-ionization metal absorption lines. The Lyα emission route looks more promising but due to the resonant nature of the line, radiative transfer simulations are required to fully characterize the observed signal.

4533798

2011ApJ...726...31G

1012

1012.0306.txt

\label{sec: Introduction} The launch of the \emph{Hubble Space Telescope} (\emph{HST}) two decades ago made it possible to study the innermost regions of galaxies at spatial resolutions that were previously unattainable at optical wavelengths. The first \emph{HST} imaging surveys of bright early-type galaxies agreed in finding a luminosity-dependent structural dichotomy in the central brightness profiles --- within the innermost few hundred parsecs, galaxies brighter than $M_B\sim-20.5$~mag showed surface brightness profiles that increased very gently towards the center (``core" galaxies) while fainter galaxies exhibited steeper surface brightness cusps (``power-law" galaxies; e.g., \citealt{Fer1994, Lau1995, Fab1997}). The paucity of galaxies with intermediate slopes was striking --- in plots of radially-scaled luminosity density profiles, core and power-law galaxies were seen to define two distinct, and virtually non-overlapping, populations (see, e.g., Figure 3 of \citealt{Geb1996}; hereafter \citetalias{Geb1996}). However, as subsequent studies targeted larger and better defined samples, some galaxies having intermediate slopes were discovered. The distinction between core and power-law galaxies became either less pronounced (e.g., \citealt{Res2001, Rav2001, Lau2007}, hereafter \citetalias{Lau2007}) or disappeared entirely (\citealt{Fer2006astroph, Fer2006, Cot2007}; hereafter \citetalias{Cot2007}. Note that these later studies parameterized the surface brightness profiles using modified S\'{e}rsic profiles rather than the so-called ``Nuker'' profiles used by earlier authors. See \S3.3.1 of \citealt{Fer2006} for a more detailed discussion.) In particular, \citetalias{Cot2007} utilized high-quality \emph{HST} imaging from the Advanced Camera for Surveys Virgo and Fornax Cluster Surveys (ACSVCS, \citealt{Cot2004}; ACSFCS, \citealt{Jor2007}). Taken together, these two surveys represent the largest and most homogeneous imaging database currently available for a well characterized sample of early-type galaxies located in low-mass galaxies clusters in the local universe (i.e., at distances $d \lesssim 20$ Mpc). The distribution of surface brightness profiles for the $\sim140$ ACSVCS/FCS galaxies was found to be a smoothly varying function of galaxy magnitude: galaxies brighter than $M_B \sim -20$~mag showed central luminosity ``deficits" (typically within $\sim40-200$ pc) with respect to the inward extrapolation of the S\'{e}rsic model that best fit the outer parts of the profiles, gradually transitioning toward the fainter galaxies that showed central luminosity ``excesses" with respect to the S\'ersic law (\citealt{Cot2006, Fer2006}). \citetalias{Cot2007} further showed that a bimodality in the central slopes could be introduced by using a biased sample: in particular, Monte-Carlo simulations showed that the bimodal {\it luminosity} distribution of galaxies observed by \citetalias{Lau2007} would lead naturally to a bimodal {\it slope} distribution, even when the intrinsic slope distribution was continuous along the galaxy luminosity function. Since \citetalias{Cot2007} was published, Kormendy~et~al.~(2009; hereafter \citetalias{Kor2009}) have commented on the core/power-law dichotomy issue as well, although they did not compute inner profile slopes. They extracted surface brightness profiles from 40 of the 100 ACSVCS galaxies and combined them with profiles from other space- and ground-based photometry, in some cases adding somewhat to the radial extent of the data. They also included profiles from space- and ground-based imaging of three additional galaxies, NGC~4261, NGC~4636, and M32. Their fits to the surface brightness profiles were determined in a very similar manner to \citetalias{Cot2007}, i.e., fitting modified S\'{e}rsic models (see \S2.1 of \citetalias{Cot2007} and Appendix~A of \citetalias{Kor2009}) and, as such, there were no systematic differences in the fits for individual galaxies, as shown in Figure~75 of \citetalias{Kor2009}. In fact, \citetalias{Kor2009} confirmed the trend from central light deficit to excess along the luminosity function of this sample that was noted by \citet{Fer2006} and \citet{Cot2006, Cot2007}. However, \citetalias{Kor2009} excluded 60\% of the ACSVCS sample -- in particular, the vast majority of galaxies in the $-21.5 \la M_B \la -18.5$ range -- and, unlike \citetalias{Cot2007}, included none of the Fornax cluster galaxies. They consequently found a qualitative gap in the inner slopes of their surface brightness profiles (see their Figure~ 40) and interpreted this gap as confirming the existence of the core/power-law dichotomy. P.~C{\^o}t{\'e}~et~al.~(2011, \emph{in prep}) will provide a much more thorough comparison of the ACSVCS/FCS results with \citetalias{Kor2009}. Several previous authors (e.g., \citetalias{Geb1996}; \citetalias{Lau2007}) who claimed a dichotomy in central surface brightness slopes, extended their work by examining the slopes of three-dimensional (i.e., deprojected) luminosity density profiles. These studies again found that a dichotomy exists, a result that cannot be immediately assumed given how rapidly shallow projected inner profiles deproject to relatively steeper inner profiles (see, e.g., \citealt{Deh1993, Mer1996}; \citetalias{Geb1996}; Figure~\ref{fig:gzcompare}\emph{a-c} of this paper). To address this issue, we show here that the distribution of slopes noted by \citet{Fer2006astroph, Fer2006} and \citetalias{Cot2007} remains continuous once the profiles are deprojected into three-dimensional luminosity density profiles. The deprojections --- which are based on a numerical inversion of the parameterized surface brightness profiles under the assumption of sphericity --- produce individual inner slopes that are consistent with those obtained using the non-parametric methodologies of \citetalias{Geb1996} and \citetalias{Lau2007}. This finding provides additional support for the conclusion that the apparent division of galaxies into core and power-law types is a consequence of the galaxy selection function used in previous studies, which was greatly overabundant in luminous core galaxies, while galaxies in the magnitude range corresponding to the transition between core and power-law types were under-represented. (See Figure 4 of \citetalias{Cot2007}.) At the same time, we note that the characterization of galaxies by the slopes of the central brightness profiles is rather sensitive to a number of factors (including the choice of measurement radius, resolution, and model parameterization). We introduce a parameter, $\Delta_{\rm 3D}$, that quantifies the central deviation of the luminosity density profile from the inward extrapolation of the S\'{e}rsic model fitted to the main body of the galaxy. We show that, when parameterized in this way, early-type galaxies show a systematic progression from central luminosity deficits to excesses (nuclei) along the luminosity function.

\label{sec: Discussion} The main result of this paper is that the deprojected profiles of the ACS Virgo and Fornax Cluster Survey galaxies do not support the existence of a core/power-law dichotomy around $M_B\sim-20.5$~mag. Rather, we find that the inner luminosity density profiles fan out over a continuum of slopes, as \citetalias{Cot2007} found when analyzing the \emph{projected} profiles of the same sample. This result holds whether the compact nuclear components (i.e., nuclei, present in the vast majority of galaxies fainter than $M_B \sim -19$ mag) are included or excluded in the deprojection. This finding is in contrast to the results of \citetalias{Lau2007} who also analyzed the projected and deprojected inner slopes of a sample of galaxies for which ``Nuker"-law fits to the surface brightness profiles were available in the literature. As discussed in \S \ref{sec: Caveat}, the actual details of the deprojection technique are unlikely to be responsible for the difference in our findings. \citetalias{Cot2007} point out that the \citetalias{Lau2007} results are, in fact, biased by their sample selection, which is described by a luminosity function (shown in Figure~\ref{fig:lumfunc}) that is itself bimodal, as \citetalias{Lau2007} themselves also note. Given the observed trend between galaxy magnitudes and inner profile slopes, \citetalias{Cot2007} show that inner slopes drawn from a continuous distribution (such as the one observed for the ACSVCS/FCS galaxies) for galaxies following a bimodal luminosity function (such as the one characterizing the sample of \citetalias{Lau2007}) will produce a bimodal slope distribution that closely resembles the one observed by \citetalias{Lau2007}. The analysis presented in this paper adds further weight to this explanation by showing that, when using a representative galaxy sample such as the ACSVCS/FCS (see Figure~\ref{fig:lumfunc}), both the two-dimensional surface brightness profiles and the three-dimensional deprojected profiles define a nearly continuous sequence as a function of galaxy magnitude. Indeed, as shown in Figure~\ref{fig:profiles}, the apparent contrast in central brightness profiles between the brightest (shallow slopes) and the fainter galaxies (steeper slopes) is less striking in the deprojected profiles, as noted by previous investigators (e.g., \citetalias{Geb1996}, \citetalias{Lau2007}). The absence of a dichotomy between ``core'' and ``power-law'' galaxies should perhaps not be surprising. It is generally believed that the core galaxies that populate the upper end of the luminosity function are nearly spherically symmetric, pressure supported, slowly rotating, boxy systems, while the opposite is true for power-law galaxies. In fact, although isophotal shape, kinematics and stellar populations do show systematic trends as a function of galaxy magnitude \citep{Ben1989, Kor1989, Cao1993, Fer2006}, there is {\it not} a one-to-one correspondence between the core/power-law classification and the above mentioned properties. Most notably, the SAURON team found no clear correspondence between ``core/power-law" galaxies and their ``slow/fast rotators" \citep{Ems2007}. Similarly, P.~C\^ot\'e et al. (2011, \emph{in prep}) find that stellar content and global structural parameters of ACSVCS and ACSFCS galaxies vary systematically along the luminosity function, but do not show any sign of a {\it discontinuity} across the alleged ``core/power-law" divide. From a theoretical standpoint, a dichotomy is also difficult to reconcile with a hierarchical merging scenario for the formation of early-type galaxies. Based on hydrodynamic simulations, \citet{Hop2008, Hop2009} find that the ``core" and ``power-law" galaxies actually form a continuous family, very much in agreement with the results presented in this paper, as well as in \citet{Fer2006} and \citetalias{Cot2007}. A continuity is, in fact, expected given that the key processes involved in the formation of spheroids (e.g., the degree of dissipation) depend on factors such as the gas fractions and masses of the progenitors, which are themselves not believed to be discontinuous. This is not at all to argue that the faintest galaxies in our sample are simply scaled-down versions of the giant ellipticals. There are a number of processes that likely affect these galaxies, including mergers (with various gas fractions), stripping, harassment, cold gas accretion, etc. However, these processes should have differing -- but not \emph{discontinuous} -- levels of importance as we travel down the luminosity function. In our view, parsing complete, or nearly complete, galaxy samples into small subgroups where certain physical mechanisms are expected to dominate and then concluding that these populations are fundamentally distinct will lead to an overly simplified, not to mention logically cyclical, view of the galaxy formation process. As an aside, note that, moving down the luminosity function, our results also show no evidence of a discontinuity across the so called ``giant/dwarf'' transition, traditionally thought to occur at $M_B\approx-17.5$~mag (e.g., \citealt{Kor1985}). One of the arguments often cited in support of the notion that dwarf galaxies are physically distinct from regular ellipticals is that the former have exponential surface brightness profiles (i.e., S\'{e}rsic index $n\sim1$). Indeed, as Figure~\ref{fig:profiles}\emph{h-j} demonstrates, as one travels down to fainter galaxies (from VCC~828 to VCC~1075 in this case) the underlying galaxy at inner radii tends to become very flat in projection. This trend may make it tempting to conclude that galaxies with flat underling inner surface brightness profiles are dwarf ellipticals while those with steeper inner slopes are giants. However, as Figure~\ref{fig:gzcompare} illustrates, whereas $\gamma_{\rm 3D} \approx \gamma_{\rm 2D}+1$ for $\gamma_{\rm 2D} \gtrsim 1$, since $\gamma_{\rm 2D} \lesssim 1 $ tends to deproject to a range of values between $\gamma_{\rm 2D}$ and $\gamma_{\rm 2D}+1$, many of those ``dwarf'' galaxies that appear flat in projection fan out to create a continuous trend with magnitude when deprojected, bridging the gap to the so-called ``giants''. Referring back to Figure~\ref{fig:gammaMBnonuc}, it is apparent that $\gamma_{\rm 3D}$ of the underlying galaxy is continuous with magnitude and does not suddenly flatten around $M_B\approx-17.5$. At this stage, it is worth emphasizing that although S\'{e}rsic profiles were used to parameterize the observed surface brightness profiles globally simply because they provided the best empirical match, the very fact that this family of models so accurately fits the surface brightness profiles (for $R \gtrsim 2\%R_e$) of both ``core" and ``power-law" galaxies, and of high- (``giant") and low-luminosity (``dwarf") galaxies must be a fundamental clue to the physics underlying the hierarchical assembly of baryons within merging dark matter halos. So is there any evidence for a physically motivated origin for S\'{e}rsic model? \citet{Hjo1995} were the first to show that the deviations of the brightness profiles of real galaxies from a de Vaucouleurs $R^{1/4}$ law --- the same deviations that can be accounted for explicitly by a S\'{e}rsic law --- can be reproduced using a simple distribution function constructed on the basis of the statistical mechanics of violent relaxation. In a series of papers by \citet{Ger1997}, \citet{Lim1999}, \citet{Mar2000}, and \citet{Mar2001} it has been shown that elliptical galaxies lie on the intersection between two manifolds --- one a scaling relation between potential energy and mass and the other representing quasi-constant specific entropy. These investigators were able to express this intersection in terms of the three S\'{e}rsic parameters and verify that actual elliptical galaxies fall within the S\'{e}rsic parameter space predicted theoretically. Finally, in a review article, \citet{Lon2006} discusses how N-body numerical simulations of collisionless gravitating systems can reproduce S\'ersic profiles given appropriate initial conditions. Finally, in terms of the origin of galaxies with central light deficits versus those with light excess, it is widely accepted that the light deficits in ``core" galaxies are the result of central scouring by coalescing black hole binaries following predominantly dissipationless galaxy mergers (\citealt{Ebi1991, Fab1997, Gua2008}) whereas nuclei are thought to be mainly the result of gas inflows into the core (\citealt{Mih1994, Cot2006}; \citetalias{Cot2007}; \citealt{Ems2008, Kor2009}), with possible contributions from other processes, such as the infall of star clusters via dynamical friction (\citealt{Tre1975, Tre1976}). It has been suggested that gas inflows played a role in the centers of light-deficit galaxies as well (see, e.g., \citealt{Hop2009b}), although \citet{Ems2008} find that the tidal forces in a galaxy are compressive for S\'{e}rsic indices $n \lesssim 3.5$ such that available gas could collapse and form a cluster of stars in the center while, for $n \gtrsim 3.5$, the tidal forces become disruptive. It may be that, although gas compression becomes increasingly important as one moves down the luminosity function, inflows still occur at higher masses, albeit at reduced levels.

1012.0306

Although early observations with the Hubble Space Telescope (HST) pointed to a sharp dichotomy among early-type galaxies in terms of the logarithmic slope γ' of their central surface brightness profiles, several studies in the past few years have called this finding into question. In particular, recent imaging surveys of 143 early-type galaxies belonging to the Virgo and Fornax Clusters using the Advanced Camera for Surveys (ACS) on board HST have not found a dichotomy in γ', but instead a systematic progression from central luminosity deficit to excess relative to the inward extrapolation of the best-fitting global Sérsic model. Given that earlier studies also found that the dichotomy persisted when analyzing the deprojected density profile slopes, we investigate the distribution of the three-dimensional luminosity density profiles of the ACS Virgo and Fornax Cluster Survey galaxies. Having fitted the surface brightness profiles with modified Sérsic models, we then deproject the galaxies using an Abel integral and measure the inner slopes γ<SUB>3D</SUB> of the resulting luminosity density profiles at various fractions of the effective radius R<SUB>e</SUB> . We find no evidence of a dichotomy, but rather, a continuous variation in the central luminosity profiles as a function of galaxy magnitude. We introduce a parameter, Δ<SUB>3D</SUB>, that measures the central deviation of the deprojected luminosity profiles from the global Sérsic fit, showing that this parameter varies smoothly and systematically along the luminosity function.

12273946

2012CQGra..29h5002C

1012

1012.4811_arXiv.txt

} As is well known, the inflationary paradigm provides solutions to many cosmological problems such as the flatness problem, the horizon or causality problem, and also dilutes unwanted (and unobserved) relics \cite{Guth, Linde, Steinhardt}. It also provides a natural mechanism of producing primordial perturbations that seed the inhomogeneities of the universe \cite{Mukhanov pert,sasaki pert}. The basic idea is that the quantum fluctuations of a classically homogeneous scalar field, the inflaton, source quantum fluctuations of the spacetime metric (the inflaton will create density perturbations which will source the scalar fluctuations of the metric). During the process of inflation, this quantum fluctuation is amplified to become a classical fluctuation, and, at the end of inflation, the fluctuation in the metric induces the density fluctuations of matter that were produced during reheating . This primordial perturbation generated during inflation then is what gives rise to the formation of structure in the universe. In this story, the crucial quantities to be determined are the amplitude of the primordial density and tensor perturbations, as the growth of structure is dependent on their size. The subsequent evolution of the primordial perturbations can be inferred from careful observations of the history of the growth of structure. From this we expect that the density perturbation produced by inflation to be $\mathcal{O}(10^{-5})$. Therefore, it is important to be able to accurately compute these perturbations in order to either preserve or rule out a given inflationary model. For the usual models of inflation, Ricci scalar plus a single canonically normalized scalar field, there are two components that will determine this amplitude. As will be described in more detail below, the first is the form of a function, $z^{\prime\prime}/z$, which arises in the Mukhanov-Sasaki equation. This equation is satisfied by a mode function, $v_{k}$, which arises by a redefinition of the co-moving curvature perturbation in momentum space, $\mathcal{R}_{k}$, upon having written the original action in terms of $\mathcal{R}_{k}$. Since this equation is crucial to finding the curvature perturbation (the same equation is also obeyed by the tensor modes), it is essential that one accurately specifies its form. The second component is the input from vacuum selection, which is equivalent to a boundary condition for the Mukhanov-Sasaki equation. For example, one avenue of study has been to alter the initial state to lie away from the standard Bunch-Davies vacuum\cite{BunchDavies,Danielsson trans Planck,New Physics CMB,Easther:2005a,Chen:2006a,Holman:2007a,Meerburg:2009a,Padmanabhan,Ashoorioon:2012}. Such alternative boundary conditions are typically chosen by conditions set at a given cut-off scale in either momentum or time. These choices will then manifest themselves in physical observables (for example as new features in the power spectrum, or enhanced non-Gaussianity) which can allow one to gain knowledge of the initial state from observation. It should be mentioned that there exist arguments \cite{Susskind} that the Bunch-Davies vacuum might be the most probable vacuum to produce the correct power spectrum from the perspective of technical naturalness, although this does not eliminate the possibility of deviations from the Bunch-Davies vacuum. It is customary to characterize inflation as a period of time where the scale factor grew almost exponentially, a period called de Sitter or quasi-de Sitter inflation (exact exponential growth of the scale factor, $a\propto e^{Ht}$, with $H$ constant, is technically de Sitter inflation, and nearly exponential growth is termed quasi-de Sitter). If the universe inflates as a power law manner $a\varpropto\tau^{p}$ where $\tau$ is the conformal time and $p$ is a constant, than the solution of the Mukahanov-Sasaki equation is known. Notice that de Sitter inflation correspond to $p=-1.$ The solution for general $p$ is given by a linear combination of Bessel functions. The calculation for the perturbation amplitude has been well established for such a case.\cite{Lyth and Stewart} However, for most models of inflation, power law expansion happens only in a short period of time either at the beginning or the end of the inflation. Thus this commonly used approximation may not apply to all inflationary models. If one insists on using the equations derived from the power law limit, one runs the risk of possibly ruling out phenomenologically viable models, or of preserving models that are ruled out by observational data. In the present work we would like to address the possibility that the function $z^{\prime\prime}/z$ inside the Mukhanov equation deviates from the de Sitter limit, and how that may affect one's choice of boundary conditions. It may be that before some point $\tau_{p}$ that using the de Sitter limit is not consistent, and therefore placing boundary conditions at $\tau_{p}$ is more natural. We have thus expanded the standard method of computing the amplitude of the primordial perturbation for such a case. Our method applies to those inflationary models that do not behave with power law (at least partially) expansion with some specific constraints in the background evolution. The paper is organized as follows. In Sec.\ref{standardmethod} we review the standard calculation of primordial density and tensor perturbations where the de Sitter limit is taken. We also explain the physical reasons behind the commonly chosen Bunch-Davies vacuum. In Sec.\ref{vacuumselection} we introduce a new method of vacuum selection by applying this method to a specific inflation model. The principles of generalizing this method to other models is also given. In Sec.\ref{conclusions} we give our conclusions.

\label{conclusions} In order to make predictions testable by observations, inflation needs not only a model, but suitable boundary conditions. Some models of inflation do not seem to fall within the realm where the standard boundary conditions may be naturally applied. With this in mind, we have discussed the introduction of an alternative method which generalizes the standard approach of computing the scalar and tensor power spectra. It is suggests that for those models whose background is analytically solvable, one should re-examine their power spectrum using our method and find how does their spectral index compare with the results from standard method. In general, this procedure will introduce additional parameters into the model, thus allowing more accurate phenomenology, with the usual drawback that introduction of more parameters decreases predictivity. This new method is implemented on a model-by-model basis, hence the generic effects of this approach have yet to be determined. For the specific example discussed in this paper, we explored an inflationary model which has analytically solvable background dynamics. We introduced a quadratic fit (of course, other models may require more complicated fitting functions)for the function $z^{\prime \prime}/z$ which appears in the Mukhanov-Sasaki equation for the scalar modes, and imposed boundary conditions at finite conformal time, $\tau$. It was found that near $\tau=7.05$ the spectral index and its running, both fall into a phenomenologically acceptable range. This calculation gives an example for the implementation of our approach, although the model in question is not fully realized in the sense that it lacks a proper accounting for the cessation of inflation in order to produce the requisite amount of e-fold expansion. The capacity for altering the calculation, and thus the values, of observables predicted by inflation via this new approach is clear. It may therefore be possible that models which were hitherto discarded may need to be re-investigated in the framework of this method. \bigskip

1012.4811

A new approach is given for the implementation of boundary conditions used in solving the Mukhanov-Sasaki equation in the context of inflation. The familiar quantization procedure is reviewed, along with a discussion of where one might expect deviations from the standard approach to arise. The proposed method introduces a (model-dependent) fitting function for the z″/z and a″/a terms in the Mukhanov-Sasaki equation for scalar and tensor modes, as well as imposes the boundary conditions at a finite conformal time. As an example, we employ a fitting function and compute the spectral index, along with its running, for a specific inflationary model, which possesses background equations that are analytically solvable. The observational upper bound on the tensor to scalar ratio is used to constrain the parameters of the boundary conditions in the tensor sector as well. An overview on the generalization of this method is also discussed.

12207099

2011JCAP...11..004R

1012

1012.3939_arXiv.txt

The nature of the dark matter (DM) is still unclear, because its existence is so far only evident via its gravitational interaction with normal matter. There are many possible ways to find DM, other than through is gravitational interactions. Direct detection experiments try to observe nuclear recoils from weak interactions with DM particles. Accelerator experiments search for hints of physics beyond the \textit{standard model} (SM) of particle physics, which may provide clues as to the identity of dark matter. Indirect detection experiments try to identify secondary products of DM annihilation or decay, such as photons, neutrinos and anti-particles. Experiments typically search for characteristic spectral signatures of DM in the cosmic fluxes of such particles, allowing them to (hopefully) differentiate the DM signal from the myriad of astrophysical backgrounds they face. There are many candidates for DM in extensions of the SM. The most popular is the neutralino, a linear combination of the superpartners of the neutral Higgs and electroweak gauge bosons seen in supersymmetric (SUSY) extensions of the SM. If $R$-parity is conserved, and the lightest neutralino is also the lightest supersymmetric particle (LSP), it can -- depending on the underlying SUSY model parameters -- deliver a relic density in the favoured range $0.094 \leq \Omega h^{2} \leq 0.129$. Neutralinos are also Majorana particles, so would self-annihilate. If SUSY is to constitute a valid solution to the well-known \textit{hierarchy problem} of the SM, it must be broken at $\sim$1\,TeV, giving sparticles such as the lightest neutralino masses of between $\sim$10\,GeV and $\sim$10\,TeV. In the annihilation process, very high energy (VHE) $\gamma$-ray photons are produced with energies up to the neutralino mass. The emissivity of annihilating DM is proportional to $\varrho^{2}$, the square of the DM density. It is thus useful to search for VHE $\gamma$-radiation from regions where a high density of DM is expected. One such region is the centre of our own galaxy, the Galactic Centre (GC). Limits on DM annihilation are generally based on assumptions about the form of the annihilation spectrum, ignoring the individual spectra of actual SUSY (or any other) models. This was perhaps reasonable until it was found that internal bremsstrahlung (IB), consisting of both final-state radiation (FSR) and virtual IB (VIB), can make large contributions to the photon spectrum \cite{bib_IB}. In this case, gamma-ray spectra from different supersymmetric models can be very different, even when the neutralino mass is kept fixed. With this new development, it is necessary to compare the observed and predicted energy spectra from annihilation processes on an individual, model-by-model basis. This was first performed in a full SUSY scan using \textit{Fermi}-LAT data on the dwarf galaxy Segue 1 \cite{bib_Pat}. The GC region has also been observed by the High Energy Stereoscopic System (H.E.S.S.), and high-energy gamma radiation has been detected \cite{bib_hessGC1, bib_hessGC5}. Because the observations seem to be incompatible with the total observed flux coming exclusively from neutralino annihilation, the hypotheses that DM annihilation makes a subdominant contribution has been investigated, resulting in limits on the DM self-annihilation cross-section \cite{bib_hessGC2, bib_hessGC4}. In this article we show the results of two full model scans in the parameter space of the Constrained Minimal Supersymmetric SM (CMSSM), comparing model predictions with H.E.S.S. data from the Sagittarius (SgrD), Carina and Sculptor dwarf galaxies, as well as the Galactic halo and Galactic Centre. First we show a simple random scan, producing a set of CMSSM models compatible with constrains on the relic density and accelerator bounds included in DarkSUSY 5.0.4 \cite{bib_DS, bib_DSweb}. Later we show more advanced statistical scans, using the SuperBayeS package \cite{bib_SB1, bib_SB4, bib_SBweb, bib_SB3, bib_SB5, bib_SB2, bib_SB6}. In section \ref{sec_hess} we introduce the H.E.S.S. experiment and the data that we use for this work. Section \ref{sec_theo} is about the theoretical framework of supersymmetric DM, and Section \ref{sec_ana} describes our analysis of the H.E.S.S. data. Section \ref{subsec_GC1} gives our results for the random scan using a spectrum from the GC source. Section \ref{subsec_GC2} introduces the CMSSM parameter scan with SuperBayeS, considering the same GC spectrum. Section \ref{subsec_SgD} describes a SuperBayeS scan taking into account the H.E.S.S. observations on the SgrD, whilst Secs.~\ref{sec_2dwarfs} and \ref{sec_halo} introduce further constraints from the Carina and Sculptor dwarfs, and the Galactic halo, respectively. Section \ref{sec_sum} finishes with a summary and outlook.

\label{sec_sum} We have performed a scan over the CMSSM parameter space, taking into account a large range of experimental data at the composite likelihood level, and using nested sampling. We have done this in order to check what constraints are placed on CMSSM models by the combination of H.E.S.S. observations of dwarf spheroidal galaxies, the Galactic halo and the Galactic Centre. Due to the strong astrophysical $\gamma$-ray source in or very near the GC, the search for DM there is strongly handicapped, so the data are not very constraining. With unrealistic assumptions about the DM density profile around the GC, we showed some example constraints on the coannihilation region and focus-point neutralinos with large masses. These examples show how the scanning technique will be useful for future observations with the next generation of $\gamma$-ray experiments, such as CTA. For dwarf galaxies and the Galactic halo we also obtained constraints on the coannihilation region and high-mass parts of the focus point, even with realistic density profiles. These constraints result from the combination of the energy reach of H.E.S.S. and a full treatment of IB. Our results give the tightest constraints to date upon the coannihilation region of the MSSM. There are however still large uncertainties in the DM density profile of the SgrD, due to strong tidal forces \cite{bib_hessSGD,bib_Viana}. This is unfortunate, as the SgrD potentially provides the strongest constraint on CMSSM coannihilation models. Future scans and limits based on the SgrD should become more robust as they eventually come to include the the dark matter halo parameters as nuisances, and observational constraints upon those parameters improve. Ultimately however, we see that including observations of Carina and Sculptor along with those of the SgrD, and assuming median values of all $J$ factors, results in almost as strong a constraint on the coannihilation region as taking just the SgrD on its own, and using a maximal $J$ factor. This speaks strongly to the robustness of the results we have presented in this paper. The recently presented results from LHC \cite{bib_Aad:2011hh} are not directly comparable with our results, since constraints are presented for fixed $\tan \beta=3$ and $A_{0}=0$, which is actually not part of the most favoured $68\%$ region that we find. However, ATLAS constrains gaugino masses below about $310 \, \text{GeV}$ and scalar masses below about $740 \, \text{GeV}$. Most of the favoured region that we find is at either larger gaugino or larger scalar masses. Thus, the present ATLAS constraint can be expected to have a minor effect on the results presented here, see e.g. \cite{bib_Bertone:2011nj} for a more detailed discussion.

1012.3939

In order to place limits on dark matter (DM) properties using γ-ray observations, previous analyses have often assumed a very simple parametrisation of the γ-ray annihilation yield; typically, it has been assumed that annihilation proceeds through a single channel only. In realistic DM models, annihilation may occur into many different final states, making this quite a rough ansatz. With additional processes like virtual internal bremsstrahlung and final state radiation, this ansatz becomes even more incorrect, and the need for scans of explicit model parameter spaces becomes clear. Here we present scans of the parameter space of the Constrained Minimal Supersymmetric Standard Model (CMSSM), considering γ-ray spectra from three dwarf galaxies, the Galactic Centre region and the broader Galactic halo, as obtained with the High-Energy Stereoscopic System (H.E.S.S.). We present a series of likelihood scans combining the H.E.S.S. data with other experimental results. We show that including combined observations of the Sagittarius, Carina and Sculptor dwarf galaxies strongly disfavour the coannihilation region of the CMSSM and models with large annihilation cross-sections. Without the Sagittarius dwarf, which admittedly has a rather uncertain dark matter profile, the results are similar, but weaker. The Galactic Centre search is complicated by a strong (unknown) γ-ray source, and we see no significant effect on the CMSSM parameter space when assuming a realistic Galactic Centre DM density profile.

12299388

2012PhRvD..85f3510D

1012

1012.4458_arXiv.txt

For many decades we have been aware that dark matter exists through its gravitational effects on galaxies and the Universe as a whole.\footnote{ See, e.g., Ref.~\cite{D'Amico:2009df} for a pedagogical review.} Over this time, compelling candidates for dark matter have been proposed that often try to link the origin of dark matter to other outstanding problems in physics and cosmology. Amongst others these include, for example, weakly interacting massive particles (WIMPs) (e.g., \cite{WIMPs}), axions \cite{axions}, sterile neutrinos \cite{sterile}, SuperWIMPs \cite{SuperWIMPs}, and asymmetric dark matter (ADM) models that link the origin of dark-matter and baryon number (e.g., \cite{ADMBaryon}). Many of these models are already well constrained from a remarkable effort in experimental dark matter physics including direct and indirect searches for dark matter and constraints from cosmology, and improved measurements and data from the Large Hadron Collidor (LHC) have the potential to further shed light on dark matter physics (see, e.g., Ref.~\cite{Feng:2010gw}). Despite these intriguing possibilities for dark matter physics it remains entirely possible that dark matter belongs to hidden sector and lacks standard model gauge charge. In such a case dark matter is expected to have at most extra-weak interactions with the standard model, and in the extreme case interacts only via gravitational forces. In such a case is dark matter physics wholly inaccessible to us? We take heart in the knowledge that the only reason we know about dark matter is via its gravitational effects in cosmology and astrophysics. The most scrutinized candidate for dark matter is arguably the WIMP because the so-called ``WIMP miracle" ensures that the simple process of its weak-scale interactions freezing out from a thermal plasma give it nearly the right relic abundance to be all of the dark matter. But as this result only relies on the size of the total annihilation cross section being close to typical electroweak value $\sigma_{\rm EW} \sim 10^{-8}~{\rm GeV}^2$ (and not that the annihilation products are standard model particles) the WIMP miracle can be extended to a ``WIMPless miracle'' in hidden-sector models that also have a similar value for the cross section but with a range of masses extending down to $O(10)$\,keV \cite{Feng:2008mu}. In Ref.~\cite{Feng:2008mu} a gauge mediated model of supersymmetry breaking is adopted that naturally gives $\sigma \simeq \sigma_{\rm EW}$ and WIMP-like dark matter that decouples from the hidden plasma when it is nonrelativistic. Also see Ref.~\cite{Cheung:2010gj} for different possible origins of hidden sector dark matter. In Ref.~\cite{Sigurdson:2009uz} hidden-sector models that decoupled when ultrarelativistic were considered. In this work we present a unified treatment of the freeze-out of thermal relics in hidden sectors for arbitrary cross section $\sigma$, only requiring that constraints from cosmology are satisfied, and find viable dark matter that freezes out when it is relativistic, semirelativistic, or nonrelativistic. Already, there are stringent bounds on hidden sector dark matter from cosmology if the hidden-sector reheated to a temperature at or above the temperature of the visible sector. If the hidden sector has more than a few light species for dark matter to annihilate into then constraints on the number of relativistic degrees of freedom from big bang nucleosynthesis (BBN) rule out such models \cite{Ackerman:2008gi,Simha:2008zj,Izotov:2010ca}. However, in hidden sector models the reheat temperature of the hidden sector can easily be different from that of the visible (standard model) sector depending on the details of the reheating process or if the visible sector and hidden sector contain a different particle spectrum (different number of degrees of freedom as a function of temperature) and thus cool differently. Note that implicit in saying the visible and hidden sectors are at different temperatures is the assumption that the two sectors are not in thermal contact (or at least they lost thermal contact considerably before BBN) and processes that could thermalize them in a Hubble time are inefficient. As discussed above, we take a model-independent approach to hidden-sector dark matter in this paper and let the cross section $\sigma$ be a free parameter. In addition we allow the mass $m_{\chi}$ of our dark matter particle $\chi$ and the hidden-to-visible temperature ratio $\xi=T^h_f/T_f$ (at dark matter freeze-out) be additional parameters and map of the viable region of this three-dimensional parameter space. Depending on the values of $m_{\chi}$, $\xi$, and $\sigma$ the freeze-out process can be relativistic, semirelativistic or nonrelativistic. Our paper is organized as follows: In Sec.~\ref{sec:ra} we discuss the standard Boltzmann equations involved in the freeze-out process and how they must be extended for the general hidden-sector case with $\xi < 1$. In Sec.~\ref{sec:ta} we discuss our approach to modelling the temperature dependence of the annihilation cross section in detail, how we treat the problematic semirelativistic regime, and a numerical approach to relativistic and nonrelativistic decoupling within the same framework. In Sec.~\ref{sec:fs} we calculate bounds arising from cosmological limits on the free-streaming length of dark matter, while in Sec.~\ref{sec:tg} we calculate limits from the generalized Tremaine-Gunn bound on the phase space density of dark matter. We discuss our results in Sec.~\ref{sec:results}, the strength of hidden-visible interactions in Sec.~\ref{sec:detection}, and conclude in Sec.~\ref{sec:conclusion}.

\label{sec:conclusion} Dark matter might belong to a hidden sector with new particle physics totally unknown to us. Even if this is the case, we can put constraints on its particle properties from its gravitational effects on the Universe. The constraints we discuss here are valid for a wide class of hidden-sector models of dark matter where the dark matter is a thermal relic particle that was once in equilibrium with a hidden-sector plasma. We have allowed the dark-matter annihilation cross section to be a free parameter and the hidden sector temperature to be different from the visible sector. Due to these extra freedoms we have shown, in contrast to a visible-sector WIMP, we must include cases where the dark matter is relativistic, semirelativisitc, and nonrelativistic at freeze-out. By solving the general freeze-out scenario numerically we have treated all cases in a unified way, and have found the region of hidden dark matter parameter space that can account for the observed density of dark matter while remaining consistent with current constraints on the free-streaming length and phase-space density of dark matter. A thermal relic in a hidden sector with dark matter mass $m_{\chi} \leq 1.5~{\rm keV}$ is incompatible with current constraints. This lower bound on the mass of hidden-sector is insensitive to the details of hidden-sector particle physics and relies on gravitationally-mediated signatures of dark matter from cosmology. \\

1012.4458

We explore the model-independent constraints from cosmology on a dark-matter particle with no prominent standard model interactions that interacts and thermalizes with other particles in a hidden sector. Without specifying detailed hidden-sector particle physics, we characterize the relevant physics by the annihilation cross-section, mass, and temperature ratio of the hidden-to-visible sectors. While encompassing the standard cold weakly interacting massive particle (WIMP) scenario, we do not require the freeze-out process to be nonrelativistic. Rather, freeze-out may also occur when dark-matter particles are semirelativistic or relativistic. We solve the Boltzmann equation to find the conditions that hidden-sector dark matter accounts for the observed dark-matter density, satisfies the Tremaine-Gunn bound on dark-matter phase-space density, and has a free-streaming length consistent with cosmological constraints on the matter power spectrum. We show that for masses ≲1.5keV, no region of parameter space satisfies all these constraints. This is a gravitationally mediated lower bound on the dark-matter mass for any model in which the primary component of dark matter once had efficient interactions—even if it has never been in equilibrium with the standard model.

12226112

2011PhRvD..84d2002A

1012

1012.1561_arXiv.txt

The last solar cycle 24 started at the beginning of 2008 \cite{svalgaard09}. Even so, in this second semester of 2010, we are still at a low level of activity, in an anomalous extended period of minimal solar activity. But sunspots are starting to appear again, and spots are the manifestation of the magnetic field poking through the surface of the Sun. This Sun cycle anomaly is the first in the spatial era, i.e., where the Sun is monitored by spacecraft detectors, but there are such registered from 107 years ago, of a similar pattern. That happened during the transition between cycles 13 and 14. Solar flares are among the most powerful astrophysical phenomena in the solar system. Initially the observation and the study of these solar flares used detectors on the surface of the Earth, mainly by neutron monitor experiments. A lot has been obtained with these observations, such as the anti-correlation between solar activity and the flow of galactic cosmic rays. The existence of a prompt and late emission in flares and their correlations with Coronal Mass Emission (CME), Forbuch events, a fall in the cosmic ray intensity, due to a solar disturbance crossing the Earth, and so forth. . However, in most cases, the observations are restricted to flares with high intensity, those with an X-ray flux at 1 UA is classified by GOES as X-Class and M-Class flares, flux above $10^{-4}$ and $10^{-5}$ W m$^2$ respectively. Recently the Tupi experiment has reported experimental evidence of muon excess in association with high energy particles (protons and ions) with energies above the pion production threshold (because they produce muons in the atmosphere), emitted by flares of small scale, those with an X-ray flux below $10^{-5}$\ W m$^2$ or C-Class flares \cite{navia05,augusto05}. Even in events observed at ground level in association with large solar flares, the acceleration mechanism producing particles (protons, ions) up to several tens of GeV is not well understood. The situation becomes still more critical in the case of ground events associated with solar flares of small scale. Here we argue for the possibility of a ``scale-free'' energy distribution of particles accelerated by solar flares. This mean that high and low fluxes of solar particles, associated to big and small flares, have the same energy spectrum, up to energies exceeding the pion production. The difference between them is only the intensity. This hypothesis is corroborated by new observations of two solar flares of small scale observed by Tupi telescopes, in association with energetic gamma-rays detected by Fermi GBM (designed to observe gamma-ray bursts) and with X-ray flux detected by GOES 14; both are spacecraft detectors orbiting the Earth. \paragraph{Outline} This paper is organized as follows: In Section 2 a brief description of the Tupi experiment is presented and includes a comment on the location of the Tupi telescopes. In Section 3, we argue why the telescopes have a high sensitivity. Section 4 contains a brief description of micro and mini solar flares. A brief study on muon excess and solar flare association is presented in Section 5 Section 6 is devoted to showing the results of the association between muon excesses at ground level and two solar mini flares, registered by GOES-14 (X-ray flux) and Fermi GBM (gamma-ray counting rate). Section 7 contains an analysis of the pitch angle and muon excess intensity, and Section 8 contains conclusions and remarks.

\label{sec8} Solar flares release energy in many forms---electro-magnetic (gamma-rays and X-rays), particles (protons, ions and electrons), and mass flows. Flares are characterized by their brightness in X-ray flux. The biggest flares are X-Class flares (flux between 10$^{-4}$ and 10$^{-3}$\ Watt m$^2$). M-Class flares have one tenth the energy and C-Class flares have one tenth of the X-ray flux seen in M-Class flares. In most cases only flares of X-class and M-class are observed by ground detectors. However, we report here two C-class flares observed as muon excesses by the vertical Tupi telescope located at sea level and within the South Atlantic Anomaly (SAA) region. We argue here that the high sensitivity attained by the Tupi telescopes is a consequence of its location, as in the SAA region the shielding effect of the magnetosphere is not perfect and shows a `dip'. The SAA is an area of anomalously weak geomagnetic field strength. This characteristic offers muon telescopes (inside the SAA region) the possibility of achieving a low rigidity of response to primary and secondary charged particles ($\geq 0.1$ GV). We have shown that both events (muon excess) are in excellent correlation with the X-ray emission observed by GOES 14, as well as in excellent correlation with the gamma-ray emission observed by Fermi GBM. In addition, from an analysis on the basis of a Monte Carlo study, rise time in the muon time profiles, and their pitch angle, we conclude that the second event on 03 November 2010 is constituted at least by two extremely coherent muon pulses, which means that the particles producing muons in the Earth atmosphere were emitted by the Sun in the same direction as an IMF line crossing the Earth and their transport was practically non-diffuse. Already for the first event analysed, and registered on 4 November 2010, the transport conditions between the Sun and the Earth were not very favourable. Even so it has a good correlation with both the Goes 14 X-ray flux and the Fermi GBM gamma-ray counting rate. If muons are detected at sea level with energies above 0.1 GeV, this means primary (protons) with energies above the threshold of pion production ($\sim 1 GeV$). Even so, what is the mechanism in the flare that accelerates protons to these energies? Especially for flares of small scale, such as those of C-Class, is this a question. Here we have argued for the possibility of a `scale-free' energy distribution of particles accelerated by solar flares. Large and small scale flares have the same energy spectrum, up to energies exceeding the pion production, they differ only in their intensity. If this hypothesis is correct, the vertical telescope is registering muons produced by protons (ions) whose energy corresponds to the tail of the spectrum. Consequently, the energy distribution of the emitted protons has to be a power law spectrum, as power law distributions are characterized as scale free distributions. Finally, we would like to point out that solar flare events also have been searched for in the first year of Fermi LAT data (August 2008-August 2009). Up until now there has been no evidence of high energy emission from solar flares detected by the LAT \cite{iafrate09}.

1012.1561

This paper presents results of an ongoing survey on the association between muon excesses at ground level, registered by the Tupi telescopes, and transient solar events, whose gamma-ray and x-ray emissions were reported by the Fermi Gamma Burst Monitor and the Geostationary Operational Environmental Satellite 14, respectively. We show that solar flares of small scale, those with prompt x-ray emission classified by the Geostationary Operational Environmental Satellite as C-Class with power 10<SUP>-6</SUP> to 10<SUP>-5</SUP>Wattsm<SUP>-2</SUP> at 1 AU, may give rise to muon excess probably associated with solar protons and ions emitted by the flare and arriving at the Earth as a coherent particle pulse. The Tupi telescopes are within the central region of the South Atlantic Anomaly, where the geomagnetic field intensity is the lowest on the Earth. Here we argue for the possibility of a “scale-free” power-law energy spectrum of particles accelerated by solar flares. For energies around and exceeding the pion production, large and small scale flares have the same power-law energy spectrum. The difference is only in the intensity. The Tupi events give support to this conjecture.

12203408

2011EAS....52..201V

1012

1012.4943_arXiv.txt

In normal galaxies like our own, most stars have about the mass of the Sun, and only $\sim$1\% of stars is more massive than 10\,\msol. Despite this rarity, high-mass stars are a major source of radiative and mechanical energy input to the interstellar medium, through ionizing UV radiation, strong stellar winds, and supernova explosions. High-mass stars are thus important for the evolution of their host galaxies, and form a link to starburst galaxies and the early Universe, where our view of star formation is dominated by the high-mass range. The formation of high-mass stars is less well understood than the low-mass case, due to large distances, short timescales, and heavy extinction. Scenarios of monolithic accretion via a disk have had some success after suitable modification, in particular increased temperatures and stronger turbulence in the parent cloud. Alternatively, pre-stellar cores or protostellar envelopes may merge and/or accrete from the same reservoir. In any case, feedback is important due to the clustered nature of high-mass star formation. While the angular resolution of ALMA will be essential to understand high-mass star formation on the `disk' scale, Herschel-HIFI will clarify the picture on the `envelope' and 'cloud' scales in at least three ways. First, HIFI gives full access to the \hho\ molecule, which is a key probe of interstellar physics and chemistry. Second, observations of hydride molecules are valueable probes of gas processing by radiation and shocks. Third, broad-band spectral surveys give a full view of the chemical composition of the gas, which is sensitive to parameters which are not directly observable, such as energetic (X-ray / cosmic-ray) irradiation and time. This review concentrates on \hho\ observations; the papers by Benz and Ceccarelli in this volume treat hydrides and spectral surveys, respectively. The \hho\ molecule is a sensitive probe of physical conditions in interstellar gas clouds, in particular kinetic temperature and volume density. Being a major carrier of oxygen, the third most abundant element in the Universe, it influences the abundances of many molecular species. Previous space telescopes (ISO, Spitzer, SWAS, Odin) have observed \hho\ lines, but had limited spectral or angular resolution, or limited line coverage. The HIFI instrument on ESA's Herschel space observatory gives the first high-resolution view of the bulk of interstellar \hho, and \hho\ is the subject of a dedicated Guaranteed Time program, `Water in Star-forming regions with Herschel' (WISH; \citealt{ewine:wish}). This paper describes the results of the first observations within the high-mass subprogram of WISH. For a precise overview of all planned observations in this subprogram, and the results for low-mass protostars, see the contributions by Herpin and Kristensen in this volume.

The first results from the high-mass WISH program show that around high-mass protostars, \hho\ is present in three distinct physical components: envelopes, outflows, and foreground clouds. The abundance of \hho\ is low in the envelopes and the foreground clouds, and higher in protostellar outflows. These results are similar to those for low-mass protostars, where outflows dominate the \hho\ line profiles \citep{kristensen:ngc1333}. Envelopes and foreground clouds are barely visible (if at all) in the data for low-mass protostars, which is very likely just due to limited sensitivity. Several projects are planned to follow up on these initial results. The \hho\ 557\,GHz line will be studied toward a few positions in infrared dark clouds. These data will clarify the role of water during the earliest stages of high-mass star formation, and also form a link with low-mass pre-stellar cores \citep{caselli:b68}. A survey of 19 high-mass protostellar objects in the ground state line of p-\hho\ may reveal basic trends of the \hho\ abundance with luminosity, mass, and other physical properties of the sources such as the presence of a hot core and an ultracompact H{\sc II} region. At the same time, multi-line studies will be valueable to construct radial abundance profiles of \hho, which constrain possible formation and destruction routes for \hho. Two-dimensional models may be used to derive detailed models of the source structure, and to evaluate the uncertainties of derived \hho\ abundances. Large-scale (arcmin-sized) maps of \hho\ emission and absorption towards star-forming complexes will constrain the spatial extent of dense gas and its role in clustered star formation. These maps may also be used to search for \hhop\ emission, which so far only has been seen in the nucleus of Mrk 231 \citep{vdWerf:mrk231}. Finally, the PACS instrument will be used for imaging of lines of CO, OH, \hho\ and other molecules at far-infrared wavelengths, which will constrain the distribution of high-excitation molecular gas. Together, these observations will be a significant step towards understanding the physical and chemical processes during high-mass star formation.

1012.4943

This paper reviews the first results of observations of H<SUB>2</SUB>O line emission with Herschel-HIFI towards high-mass star-forming regions, obtained within the WISH guaranteed time program. The data reveal three kinds of gas-phase H<SUB>2</SUB>O: "cloud water" in cold tenuous foreground clouds, "envelope water" in dense protostellar envelopes, and "outflow water" in protostellar outflows. The low H<SUB>2</SUB>O abundance (10<SUP>-10</SUP>-10<SUP>-9</SUP>) in foreground clouds and protostellar envelopes is due to rapid photodissociation and freeze-out on dust grains, respectively. The outflows show higher H<SUB>2</SUB>O abundances (10<SUP>-7</SUP>-10<SUP>-6</SUP>) due to grain mantle evaporation and (probably) neutral-neutral reactions.

12224834

2011PhRvD..83g5002B

1012

1012.3960.txt

Directional detection of Galactic dark matter has been first proposed by D.~N.~Spergel \cite{spergel} highlighting the fact that even low angular resolution directional detectors could be used to show a clear asymmetry in the forward/backward distribution of weakly interacting massive particle (WIMP) events with respect to the direction of the Cygnus constellation.\\ Beyond the simple asymmetry feature, it has recently been shown that dedicated statistical data analysis of forthcoming directional detectors~\cite{white,mimac,drift,dmtpc,newage} could lead either to a competitive exclusion \cite{billard.exclusion} or to a conclusive discovery \cite{billard.disco,green.disco}, depending on the value of the WIMP-nucleon cross section. In the latter case, by using a map based likelihood analysis and even in the presence of a sizeable background, it is possible to show that the main incoming direction does correspond to the direction of the Cygnus constellation ($\ell_{\odot}, b_{\odot}$). This is indeed the discovery proof of this detection strategy and it has been shown that a 10 kg $\rm CF_4$ detector (MIMAC) operated during 3 years, would allow for a high significance discovery down to $\sigma^{SD} \simeq 10^{-4} \ {\rm pb}$ \cite{billard.idm2010}. In this paper, we strive to go one step beyond by trying to constrain the properties of Galactic dark matter with directional detection.\\ Indeed, constraining WIMP parameters (mass $m_\chi$ and cross section $\sigma_n$) with upcoming dark matter experiments is a main concern of current phenomenological studies, using either indirect detection \cite{bernal,bernal2}, direct detection on its own \cite{bernal2,green.masse1,green.masse2,Drees:2007hr,drees.masse,Shan:2010hr}, in combination with collider data~\cite{Bertone:2010rv} or with the measurements of halo star kinematics \cite{Strigari:2009zb}. The quest for a model-independent formalism is a difficult task as the signal expected in direct detection depends on the properties of both the WIMP particle (mass and cross section) and the Galactic dark matter halo (three-dimensional local WIMP velocity distribution and density). This approach is of particular interest in the context of competitive upcoming experiments which might be able to give positive WIMP detection instead of background rejection. M.~Drees and C.~L.~Shan have proposed a model-independent reconstruction of the WIMP velocity distribution as well as its various moments (mean velocity, dispersions, ...), providing the WIMP mass is {\it a priori} known \cite{Drees:2007hr} or deduced from positive signals from at least two direct detectors with different target nuclei \cite{drees.masse}. The complementary approach is to constrain the WIMP properties with the help of a high dimensional multivariate analysis and within the framework of a general halo model, with a large number of parameters. Thus, the main strength of this study, and hence of directional detection, is the possibility of constraining the properties of both the dark matter particle and the dark matter halo with a single experiment. The choice of the fitting model must be well motivated {\it e.g.} by N-body simulations, as it remains as an {\it ansatz}.\\ Directional detection presents a high identification potential thanks to the use of the double-differential spectrum ${\mathrm{d}^2R}/{\mathrm{d}E_R\mathrm{d}\Omega_R}$, also called the directional event rate, in a given recoil energy range. Indeed, its shape depends both on the WIMP mass and WIMP velocity distribution, while the magnitude mainly depends on the product of the local WIMP density and the WIMP-nucleon cross section. Within the framework of a multivariate recoil event analysis using a Markov chain Monte Carlo (MCMC), we show for the first time the possibility to constrain, with a single directional experiment, the unknown WIMP parameters, both from particle physics ($m_\chi, \sigma$) and Galactic halo (velocity dispersion along the three axis), leading to an identification of non-baryonic dark matter. It is, of course, possible to include external data, {\it e.g.} halo star kinematics as in \cite{Strigari:2009zb}, and to relax some astrophysical inputs, as $\rho_0$ for instance. However, in this work, we focus on the contribution of directional detection on its own, highlighting the need for future large directional detectors.\\ The paper is organized as follows. In Sec. II, the dark matter halo modeling is introduced while the directional detection framework is presented in Sec. III. Then, the Markov chain Monte Carlo analysis is detailed in Sec. IV, highlighting the performance of such a method in the context of directional detection. Sec. V presents the results of this 8 parameter analysis for a directional detector with a sizeable background contamination and in the case of a benchmark dark matter model. Departures from this input model, by changing the WIMP mass, the velocity anisotropy and the background assumptions are presented in Sec. VI.

We have shown that identification of dark matter might be achieved with a 10 kg $\rm CF_4$ directional detector, allowing 3D track reconstruction with sense recognition down to 5 keV and operated during three years. To fully exploit upcoming data, we propose a new high dimensional multivariate analysis method based on a Markov chain Monte Carlo analysis of recoil events, allowing to constrain, in a single directional experiment, the WIMP parameters, both from particle physics (mass and cross section) and Galactic halo (velocity dispersion along the three axis) and within the framework of a given ansatz. Indeed, the combination of information from the angular and energy distributions leads to robust allowed regions in the ($m_{\chi},\log_{10}(\sigma_n)$) plane, since the halo model is also being constrained with the MCMC analysis from the same dataset of a single directional detection experiment. Moreover, the velocity anisotropy parameter $\beta$, related to the three velocity dispersions, could be sufficiently constrained to discriminate between various halo models with future directional detectors such as the one proposed by the MIMAC collaboration \cite{mimac}.

1012.3960

Directional detection is a promising dark matter search strategy. Indeed, weakly interacting massive particle (WIMP)-induced recoils would present a direction dependence toward the Cygnus constellation, while background-induced recoils exhibit an isotropic distribution in the Galactic rest frame. Taking advantage of these characteristic features, and even in the presence of a sizeable background, it has recently been shown that data from forthcoming directional detectors could lead either to a competitive exclusion or to a conclusive discovery, depending on the value of the WIMP-nucleon cross section. However, it is possible to further exploit these upcoming data by using the strong dependence of the WIMP signal with: the WIMP mass and the local WIMP velocity distribution. Using a Markov chain Monte Carlo analysis of recoil events, we show for the first time the possibility to constrain the unknown WIMP parameters, both from particle physics (mass and cross section) and Galactic halo (velocity dispersion along the three axis), leading to an identification of non-baryonic dark matter.

12213831

2011MNRAS.413.1395O

1012

1012.3752_arXiv.txt

The knowledge of distances to galaxies is important to deduce intrinsic galaxy properties (absolute magnitude, size, etc.) from the observed properties (colours, sizes, angles, flux, apparent size, etc.).\\ Since pioneering works measuring redshift of bright galaxies (e.g., (\citet{shapley1932}, \citet{humason1956}), many efforts have been invested in mapping the light and matter in the Universe. Statistical analysis of galaxy properties and systems can be invaluable tools to study large-scale structure and evolution in the universe.\\ In recent years, multi-band photometry has been performed for several millions of galaxies, whereas spectroscopic redshifts have been measured only for a small fraction of the photometric data. The Sloan Digital Sky survey has obtained multi-band images for approximately one hundred billion galaxies (and the next surveys foresee increasing the number of objects up to billions), while spectroscopic measurements have been obtained for nearly one million galaxies. A solution to the difficulty of obtaining spectroscopic redshifts relies on the use of photometric redshift techniques. Although the redshifts calculated through these techniques are far less accurate than the spectroscopic measurements, these approximate distance estimates allow for useful analysis in fields such as extragalactic astronomy and observational cosmology. Two basic family of methods are commonly employed to calculate photometric redshifts. In the template matching approach, a set of spectral energy distribution (SED) templates is fitted to the observations (e.g., colours). In the empirical approach, on the other side, photometric redshifts are obtained from a large and representative training set of galaxies with both photometry and precise redshift estimations. The advantage of the first method is that it can also provide additional information, like the spectral type, k-corrections and absolute magnitudes. The accuracy of these estimations is limited by the SED models. The empirical model overcomes this limitation through the use of a training set which, however, should be large and representative enough to provide accurate redshift estimations.\\ The observed spectral energy distribution of distant galaxies is redshifted with respect to that in the galaxy rest frame. The k-correction term \citep{oke,hogg01} applied to the apparent magnitude measured in a given photometric band takes into account this effect, allowing to compare photometric properties of galaxies at different redshifts. The estimation of k-corrections, then, is a requirement for many studies of distant galaxies. It is possible to model k-corrections as a function of redshift and galaxy morphological type \citep{fuku95, Mannucci01}. \citet{lahav95} and \citet{banerji10} use ANNs to obtain morphological classification of galaxies. However, these techniques employ as training sets, objects that had been classified by human eye subject to some degree of ambiguity and uncertainty, particularly at large redshifts. A more direct way to obtain k-corrections is by modelling galaxy SEDs as a function of wavelength. Usually template fitting of observed galaxy fluxes is employed to reconstruct the SED of the galaxy (\citet{blanton03}, \citet{blanton07}). The feasibility and accuracy of this method relies in the quality of the models. For objects with spectral data, k-corrections can also be obtained directly. \citet{roche} used this technique to calculate k-corrections for early-type galaxies from the SDSS-DR6, providing individual estimates for each galaxy. However, this technique is restricted to a limited number of galaxies with spectroscopy. Recent works (\citet{chilinga}, \citet{westra}) have approximated k-corrections with analytical functions of redshift, parametrized with some property characterizing galaxy type. \citet{chilinga} used different observed colour indices to approximate k-corrections for nine filters ($ugrizYJHK$). \citet{westra} used spectra from the Smithsonian Hectospec Lensing Survey to obtain direct measurements of k-corrections by parametrization with the ratio of the average flux red and bluewards of the 4000\AA{} break ($D_n4000$). These kinds of parametrization simplify the computation of k-corrections and improve their accuracy.\\ In this paper, we present a galaxy photometric redshift ($\zphot$) catalogue and a method for calculating k-corrections, for the seventh Data Release (DR7) of the Sloan Digital Sky Survey (SDSS) imaging catalogue (\citet{blanton03}, \citet{lrgs}, \citet{gunn98}, \citet{mgs}, \citet{sdss}). To compute photometric redshifts we used the ANNz software package (\citet{annz}), which have been shown to be a reliable tool. There are also two sets of photometric redshifts in the SDSS database: \citet{dr7} employ empirical, template-based and hybrid-techniques approaches to photometric redshift estimation, whereas \citet{zp} adopt a neural network method and provide two different estimations, D1 and CC2. Here we also compare our redshift estimations with those in the CC2 catalogue from \citet{zp}, which uses colours and concentration indices to infer redshifts. For the computation of k-corrections we propose a joint parametrization in terms of redshift and the $(g-r)$ colour in a certain reference frame for all SDSS bands, as well as an algorithm to determine them from the photometric data. We compare our results with those found in literature. This paper is organized as follows. In Section 2 we describe the data used in our analysis. Section 3 presents our approach to calculate photometric redshifts in SDSS-DR7, analysing their advantages and limitations. Our estimation of k-corrections is presented in Section 4. Finally, Section 5 summarizes the results obtained in this work. Throughout this paper, we adopt a cosmological model characterized by the parameters $\Omega_m=0.3$, $\Omega_{\Lambda}=0.7$ and $H_0=75~h~ {\rm km~s^{-1}~Mpc^{-1}}$. \begin{table*} \begin{minipage}{175mm} \caption{Description of the samples used in this work.} \begin{center}\begin{tabular}{@{}ccl@{}} \hline \hline sample name & number of objects & Description \\ \hline Sz1 & $\sim 550000$ & Selected from SDSS-DR7 spectroscopic data.\\ & & Used as training and validation sets in the computation of photometric redshifts \\ \hline Sz2 & $\sim 70000$ & Selected from SDSS-DR7 MGS (excluding training set galaxies). \\ & & Used as testing set for photometric redshifts.\\ \hline Sz3 & $\sim 82000$ & Selected from SDSS-DR6 photometric data.\\ & & Used to compare photometric redshift estimation with \citet{zp}.\\ \hline Sk1 & $\sim 122000$ & Selected from SDSS-DR7 MGS taking into account apparent magnitude and redshift limits (see text).\\ & & Used for k-correction calibration.\\ \hline Sk2 & $\sim 575000$ & Selected from SDSS-DR7 photometric data with our photometric redshifts estimation. \\ & & Used to compute k-correction at higher redshift and compare with literature.\\ \hline \hline \label{t1} \end{tabular} \end{center} \end{minipage} \end{table*}

In this work we present a new set of photometric redshift ($\zphot$) and k-correction estimations for the SDSS-DR7 photometric catalogue available on the World Wide Web. In order to calculate $\zphot$, artificial neural networks were applied using the public code ANNz. The improvements in the SDSS-DR7 photometric redshift estimation are:\\ 1) We added the concentration index and the Petrossian radii in $g$ and $r$ bands to the usual five magnitudes used in previous similar works. These additional inputs improve $\zphot$ estimations because the concentration index provides information regarding the slope of the galaxy brightness profile, helping us to break the degeneracies in the redshift-colour relation due to morphology. The Petrossian radius is a robust measure of how shallow the brightness profile is and contain information about the angular size, that is related to the distance.\\ 2) The choice of different galaxy samples for the training set (MGS, LRG and AGN sample) provides a wide sampling of different galaxy types at various redshifts, allowing to improve $\zphot$ estimates. Our $\zphot$ estimates have a $rms\simeq 0.0227$, and the resulting galaxy distribution shows a good agreement with the theoretical distribution derived from the SDSS galaxy luminosity functions. We have used the \texttt{k-correct\_v4.2} code (\citet{blanton07}) for the MGS and we have performed a linear fit between reference frame $(g-r)$ colour and redshift, extrapolating this relation at high redshifts. We propose an iterative procedure to estimate k-corrections from the observed photometry in the $g$ and $r$ bands. Using initial values that depend on the concentration index and the observed colour, we obtain the k-correction for the other bands. Our results show that the use of this simple linear relation between the reference frame $(g-r)$ colour and redshift is as accurate as those obtained in previous work. A clear plus of our approach is the low computational time. Our k-correction estimations do not use templates, avoiding statistical errors in the lack of homogeneity in spectral features, and minimizing systematical errors caused by an assumed spectral energy distribution (SED). This can be noticed in the smooth behavior of the distribution of k-corrections, even for intermediate redshifts. The analysis of the distribution of k-corrected absolute magnitudes show that the shape of these distributions has a bell-shaped skewed to fainter magnitudes and the mean of the distributions move towards fainter magnitudes as the redshift range decreases. We have computed the luminosity function in $g$ and $r$ bands trhough $1/Vmax$ method taking into account the incompleteness with a $V/Vmax$ test (\citet{schmidt}). We notice that the curves derived from this work and that of \citet{zucca} are consistent taking into account the $(g-B)$ value of a typical Sbc galaxy at $z \sim 0.5$ \citet{fuku95}. The analysis of the $(g-r)$ colour distribution for galaxies brighter than $M_r=-21.5$ shows that the galaxies in our samples are shifted bluewards as redshift increases. This trend leads to the emergence of bimodality in the full redshift range. From the colour-colour diagrams, we can conclude that the behavior of colours at low redshift is in good agreement with the trends of the spectroscopic sample. As redshift increases we see a broadening of the contours and an increase in the blue galaxy population. However, the distribution of galaxies in the colour-colour diagram remains centered with respect to the spectroscopic data.

1012.3752

We present a catalogue of galaxy photometric redshifts and k-corrections for the Sloan Digital Sky Survey Data Release 7 (SDSS-DR7), available on the World Wide Web. The photometric redshifts were estimated with an artificial neural network using five ugriz bands, concentration indices and Petrosian radii in the g and r bands. We have explored our redshift estimates with different training sets, thus concluding that the best choice for improving redshift accuracy comprises the main galaxy sample (MGS), the luminous red galaxies and the galaxies of active galactic nuclei covering the redshift range 0 &lt; z≤ 0.3. For the MGS, the photometric redshift estimates agree with the spectroscopic values within rms = 0.0227. The distribution of photometric redshifts derived in the range 0 &lt; z<SUB>phot</SUB>≤ 0.6 agrees well with the model predictions. <P />k-corrections were derived by calibration of the K-CORRECT_V4.2 code results for the MGS with the reference-frame (z= 0.1) (g-r) colours. We adopt a linear dependence of k-corrections on redshift and (g-r) colours that provide suitable distributions of luminosity and colours for galaxies up to redshift z<SUB>phot</SUB>= 0.6 comparable to the results in the literature. Thus, our k-correction estimate procedure is a powerful, low computational time algorithm capable of reproducing suitable results that can be used for testing galaxy properties at intermediate redshifts using the large SDSS data base.

12101885

2010AIPC.1273...31D

1012

1012.5342_arXiv.txt

Stars with masses, M$\simless$5M$_{\odot}$ (ie. $>$95\% of the stellar population by number) will ultimately become white dwarfs, with electron degenerate cores composed of He or C and O. Those much rarer stars with M$\simgreat$10M$_{\odot}$ will burn their non-degenerate cores through to Fe and then die as core-collapse Type II supernovae. However, the final evolutionary fate of single stars in the intervening mass range remains less certain. Early stellar evolutionary models predicted that when the partially degenerate CO core of a heavy-weight intermediate mass early-asymptotic giant branch (E-AGB) star achieved M$\sim$1.1M$_{\odot}$ it ignited and burned to Ne and O before collapsing. The star was then expected to explode as an electron capture Type II supernova (e.g. \cite{nomoto84}). However, more modern calculations which include improved physics suggest that a sizeable proportion of the stars in this initial mass range may instead pass through a super-AGB phase before ending their lives as ultra-massive, M$\sim$1.05-1.3M$_{\odot}$, ONe white dwarfs (e.g. \cite{garcia97}). Refining our knowledge of the fate of stars in this initial mass range is important for understanding the chemo-dynamical evolution of the Galaxy. Supernovae inject substantial amounts of kinematic energy into the ISM (e.g. \cite{cox74}). Moreover, due to the highly compact nature of the core during the super-AGB phase, the convective envelopes of these stars are subject to intense hot bottom burning and consequently they synthesise substantial amounts of $^{14}$N and $^{13}$C which are returned to the ISM (e.g. \cite{siess10}). Here we report preliminary results from our search of the moderately rich and relatively nearby open clusters NGC2287 and NGC3532 for the white dwarf remnants of stars which are believed to have had initial masses within this range of interest. \begin{figure} \includegraphics[height=.6\textheight,angle=270]{dobbie01} \caption{$V$, $B$-$V$ colour-magnitude diagrams for NGC3532 (left) and NGC2287 (right). Theoretical evolutionary tracks for 0.7M$_{\odot}$ and 1.1M$_{\odot}$ CO core \cite{holberg06} and 1.16 ONe core \cite{althaus07} white dwarfs, adjusted to account for the foreground reddening towards and the distance of each population, are overplotted. Previously known white dwarf members of the clusters recovered here are also highlighted (large triangles). } \end{figure}

\begin{figure} \includegraphics[height=.6\textheight,angle=270]{dobbie03} \caption{Location of the seven white dwarf members of NGC3532 in initial mass-final mass space for an assumed cluster age of 300Myr. A theoretical IFMR \cite{kovetz09} is overplotted (thick grey line).} \end{figure} $B$ and $V$ band surveys of NGC2287 and NGC3532 have unearthed four new white dwarfs which are probable members of these two open clusters. Somewhat surprisingly, despite a range in effective temperature of $\sim$20\%, corresponding to a range in age of $\sim$100Myr, the masses of the four heaviest white dwarf members of NGC3532 all lie between M$_{WD}$$\approx$0.9-1.0M$_{\odot}$. This suggests that there is little change in remnant mass for initial masses, 4.5M$_{\odot}$ $\simless$M$_{init}$$\simless$6.5M$_{\odot}$, ie. the IFMR is relatively flatter here than at immediately lower masses as expected from theory. WD J0646-203 is possibly the most massive cluster white dwarf identified to date and with M$_{WD}$=1.12/1.08$\pm$0.04M$_{\odot}$, based on CO/ONe models, seems likely to be composed of ONe. It has a cooling time consistent with it having evolved from a single star. We have in hand deep imaging for a number of other open clusters which we aim to use to further hone understanding of the form of the IFMR. NGC752, NGC6940 and NGC2477 will allow us to probe the relatively unexplored region 2M$_{\odot}$M$_{init}$$\simless$3M$_{\odot}$ which includes the dividing line between stars that do and do not experience the helium flash. Excitingly, we have $\sim$80 candidate members of NGC2477 with V$\simless$24.2 which could provide the chance to probe the relative form of the IFMR from M$_{init}$=2M$_{\odot}$ upto the electron capture SNe limiting mass in unprecedented detail. Data for the young and extremely rich NGC6791 will allow us to further scrutinise the upper limit on the mass of a white dwarf formed via single star evolution.

1012.5342

We present the preliminary results of a survey of the open clusters NGC3532 and NGC2287 for new white dwarf members which can help improve understanding of the form of the upper end of the stellar initial mass-final mass relation. We identify four objects with cooling times, distances and proper motions consistent with membership of these clusters. We find that despite a range in age of ~100 Myrs the masses of the four heaviest white dwarfs in NGC3532 span the narrow mass interval M<SUB>WD</SUB>~0.9-1.0M<SUB>solar</SUB> suggesting that the initial mass-final mass relation is relatively flatter over 4.5M<SUB>solar</SUB>&lt;~M<SUB>init</SUB>&lt;~6.5M<SUB>solar</SUB> than at immediately lower masses. Additionally, we have unearthed WD J0646-203 which is possibly the most massive cluster white dwarf identified to date. With M<SUB>WD</SUB>~1.1M<SUB>solar</SUB> it seems likely to be composed of ONe and has a cooling time consistent with it having evolved from a single star.

12168719

2011ApJ...731...66R

1012

1012.0617_arXiv.txt

\subsection{Convection, double-diffusive convection (semi-convection) and fingering convection (thermohaline convection)} One of the longest-standing problems in stellar and planetary astrophysics is that of modeling the transport of heat and chemical species within turbulent regions. The best-studied and most ubiquitously relevant case is that of overturning convection through a chemically homogeneous gas layer. There, the well-known Schwarzchild criterion is used to determine the extent of the convective region, while the transport properties through the layer are commonly modeled using mixing-length theory \citep{Biermann1932}. The success of well-calibrated mixing-length models in explaining many observable properties of stars is quite remarkable. However, much less is known about convection in the presence of additional factors such as strong rotation, strong magnetic fields and strong compositional gradients \citep{spiegel1972}. In all cases, the linear stability of the system is well-understood \citep{chandrasekhar1961}, but characterizing its fully-nonlinear transport properties remains the subject of ongoing research. In this work, we focus on the case of convection in the presence of a strong stabilizing compositional gradient, but in the absence of rotation or magnetic field. This regime is often called ``semi-convection'' in astrophysics \citep{schwartzchildharm1958}, although we prefer to use the terminology ``double-diffusive convection'' commonly used in oceanography to clarify the true nature of the instability responsible for the turbulence. It has long been known that the relevant criterion for instability to overturning convection in the presence of a compositional gradient is {\it not} the Schwarzchild criterion, \begin{equation} \nabla - \nabla_{\rm ad} = \left( \frac{\partial \ln T}{\partial \ln p} \right ) - \left( \frac{\partial \ln T}{\partial \ln p} \right)_{\rm ad} > 0 \mbox{ , } \end{equation} but the Ledoux criterion \citep{ledoux1947}: \begin{eqnarray} \nabla - \nabla_{\rm ad} &>& \nabla_\mu \nonumber \\ \Leftrightarrow \left( \frac{\partial \ln T}{\partial \ln p} \right) - \left( \frac{\partial \ln T}{\partial \ln p} \right)_{\rm ad} &>& \left( \frac{\partial \ln \mu}{\partial \ln p} \right) \mbox{ , } \end{eqnarray} where $T$ is the temperature, $\mu$ the mean molecular weight, $p$ the pressure, and where the subscript ``ad'' expresses a derivative at constant specific entropy. In fact, both of these criteria merely express the same property when written in terms of the density stratification: \begin{equation} \left( \frac{\partial \rho}{\partial p} \right)_{\rm ad} > \left( \frac{\partial \rho}{\partial p} \right) \mbox{ . } \end{equation} A system is unstable to overturning convection if the density of a parcel of fluid, raised adiabatically and in pressure equilibrium from its original position, is lower than that of its new surroundings. The question of what happens to regions which are stable according to the Ledoux criterion but unstable according to the Schwarzchild criterion was first raised by \citet{schwartzchildharm1958}. It was later found that this regime is in fact {\it also} linearly unstable \citep{walin1964,kato1966}, but through a {\it double-diffusive} instability, i.e. an instability which cannot occur unless the thermal diffusivity of the fluid, $\kappa_T$, is larger than its compositional diffusivity $\kappa_\mu$. This condition is however automatically satisfied in stellar and planetary interiors where the diffusivity ratio $\tau = \kappa_\mu / \kappa_T$ (often called the inverse Lewis number), can be as low as $10^{-7}$. As a result, a wide range of situations arise in which double-diffusive convection occurs and controls transport within the object. A commonly studied case is that of semi-convection at the edge of core-convective stars \citep{ledoux1947,tayler1954,schwartzchildharm1958,merryfield1995}. In moderately massive stars for example, a mean molecular weight gradient develops over time at the edge of the core, and eventually begins to affect convection. When it is strong enough to stabilize the fluid, the fully-convective region shrinks in size, leaving behind a ``semi-convective'' region in which transport is controlled by double-diffusive processes instead. Other related examples are reviewed by \citet{merryfield1995}. The possible role of double-diffusive convection in regulating thermal and compositional transport in the interior of giant planets was recognized later, and has been discussed in the context of convective planetary envelopes where the stabilizing component is helium \citep{stevenson1977}, higher-metallicity material at the edge of a rocky core (see \citet{stevenson1985} and in particular Figure 2 of his paper), or methane \citep{gierash1987}. Double-diffusive convection has also been proposed to explain the abnormally large radii of some transiting exoplanets \citep{chabrier2007htg}. Finally, it has recently been invoked as a new mechanism for driving pulsations in white dwarfs \citep{shibahashi2007,kurtz2008}. Before we move on to describe existing work on double-diffusive convection, we note that it should not be confused with that arising from the related double-diffusive ``fingering'' instability \citep{stern1960sfa}. The latter also occurs in Ledoux-stable systems, but in the opposite situation when the more rapidly diffusing thermal field is stably stratified while the slowly-diffusing compositional field is unstably stratified. Its turbulent manifestation is often referred to as ``thermohaline convection'' in astrophysics, by analogy with the oceanic case in which the compositional gradient is due to salt. We prefer a more general terminology and use the alternative name ``fingering convection''. Figure \ref{fig:regimes} illustrates for clarity the various regimes of convective instability. See \citet{traxler2010b} for a study of fingering convection in the astrophysical regime. \subsection{Previous work on double-diffusive convection} Very little is known about mixing by double-diffusive convection, despite its obvious importance in stellar and planetary astrophysics. Linear stability reveals that the unstable modes take the form of overstable gravity waves \citep{walin1964,kato1966,baines1969}. What governs the saturation of the instability in the astrophysically-relevant parameter regime, however, remains essentially unknown. Note that by ``astrophysically-relevant'' we imply a low diffusivity ratio, $\tau \ll 1$, and a low Prandtl number, ${\rm Pr} = \nu/\kappa_T \ll 1$, where $\nu$ is the viscosity. Both numbers are typically of the order of $10^{-5} - 10^{-7}$ in stellar and planetary interiors. To add to the complexity of the problem, double-diffusive convection is known in some cases to lead to thermo-compositional layering, i.e. to the development of stacks of well-mixed fully-convective layers separated by strongly stratified interfaces. This double-diffusive layering is commonly observed for example in the arctic ocean \citep{neal1969,toole2006} where cool fresh water lies on top of warmer, saltier water. It has been studied extensively in laboratory experiments \citep{turner1968,huppert1979,huppert1980,turner1985mc}. An important result of these studies is that turbulent mixing in the presence of layers is significantly enhanced compared with that of a system which has the same {\it overall} contrast in temperature and composition, but where the stratification is everywhere much smoother. Whether layer formation occurs in double-diffusive convection at low Prandtl number actually remains to be determined -- it is usually {\it assumed} \citep{spruit1992,chabrier2007htg}, by analogy with the high-Prandtl number oceanographic case. It is important to realize, however, that such analogies can be misleading. This was recently demonstrated by \citet{traxler2010b} in the case of fingering convection. Similar thermohaline staircases are ubiquitously observed in {\it fingering-unstable} regions of the ocean \citep{schmitt2005edm}, and have been shown to form spontaneously through a secondary linear instability of homogeneous fingering convection \citep{radko2003mlf,traxler2010,stellmach2010} (see \S\ref{sec:layers} too). However, \citet{traxler2010b} demonstrated that this secondary instability does not happen in the astrophysical context and concluded that thermo-compositional layers are not expected in that case. In other words, given that the analogy with the heat-salt system doesn't hold in the fingering regime, one should be extremely cautious about using it {\it a priori} in the double-diffusive regime. To summarize, it is known that ``homogeneous'' and ``layered'' double-diffusive convection have rather different mixing properties. A good mixing parametrization therefore needs to incorporate both cases, and must include a criterion to decide whether the system considered lies in one or the other. Existing parametrizations, by contrast, have so far either ignored the possible effect of layering \citep{schwartzchildharm1958,langer1983}, or relied on it \citep{spruit1992}. A few numerical simulations of double-diffusive convection have been performed to date to address the problem. The first of this kind (to our knowledge), were presented by \citet{merryfield1995}. He ran a series of two-dimensional (2D) anelastic simulations, horizontally-periodic, and bounded in the vertical direction by two plates. Double-diffusive convection was forced through the imposed boundary conditions, which maintained a given overall temperature and compositional contrast across the domain. \citet{merryfield1995} focused on understanding how the mechanism responsible for the saturation of the initial double-diffusive instability depends on the governing parameters, and in particular the strength of the thermal driving, the Prandtl number and the diffusivity ratio. He also compared the outcome of his simulations with existing parametrizations of double-diffusive convection both in the absence of layers, and in the presence of initially forced layers. One of the main difficulties encountered was the development of numerical instabilities in cases with strong thermal driving, which prevented him from drawing definite conclusions on the long-term statistical properties of the turbulence. In addition, in many of the runs the flows were eventually influenced by the presence of domain boundaries. In subsequent years, two additional attempts at modeling double-diffusive convection were made. \citet{biello2001}, as part of his PhD thesis, ran a series of 2D fully-compressible simulations which complement those of \citet{merryfield1995}. There, the system was also confined between two plates, but the boundary conditions were ``fixed flux'' conditions. \citet{biello2001} was interested in studying more specifically the layer formation process, and his experimental setup was similar to that of heat-salt laboratory experiments \citep[e.g.][]{huppert1979}. For this purpose, the simulations were initialized with {\it stable} uniform gradients in temperature and composition, but destabilized at $t=0$ by increasing the heat flux at the bottom boundary. He found that the first bottom layer easily forms, but did not observe any subsequent layer formation. He analyzed the dynamics of the interface, and concluded that interfacial transport was dominated by wave-breaking rather than by diffusive processes as is often assumed. However, boundary effects also began to influence the results of his simulations after some time. \citet{bascoul2007}, also as part of his PhD thesis, studied a similar 2D time-dependent system, in which the initial background state had a homogeneous composition and neutrally stable temperature gradient, and where the system was destabilized by an imposed heat and mean-molecular weight flux through the bottom boundary. He also observed the growth of a convective layer near the bottom boundary and studied its development, for high Prandtl number (heat-salt regime) and for low Prandtl number (``astrophysical'' regime). He was unable, however, to run his simulations long enough to achieve statistical equilibrium. In this paper, we present a new series of three-dimensional (3D) numerical experiments to study mixing by double-diffusive convection. We approach the problem from a different but complementary angle, and try to address some of the inherent shortcomings of the experimental setup used in previous studies: we focus {\it specifically} on measuring the quasi-steady statistical properties of double-diffusive turbulence, and use a numerical setup which minimizes the effects of domain boundaries. We discuss the model setup and briefly summarize its linear stability properties in \S\ref{sec:model}. The numerical algorithm and the selection of the experimental parameters are described in \S\ref{sec:num}. In each case, as described in \S\ref{sec:homogen}, we extract values for the transport coefficients while the system is in a state of homogeneous turbulence. However, we find that for more unstable systems a secondary instability leads to the formation of thermo-compositional layers. The layers continue evolving and successively merge, and each merger is accompanied with a significant increase in the transport coefficients. These results are discussed in \S\ref{sec:layers}, and compared with a recent theory of layer formation in the oceanographic context \citep{radko2003mlf}. We find good agreement between this theory and our numerical results, which enables us to deduce a general criterion for the spontaneous formation of layers in double-diffusive convection in the astrophysical context. We discuss our results in \S\ref{sec:ccl}.

\label{sec:ccl} \subsection{Summary of the results} In this work, we have studied a set of numerical simulations of double-diffusive convection, in a triply-periodic domain, for Prandtl number ${\rm Pr} = \nu/\kappa_T = 1/3$ and diffusivity ratio $ \tau = \kappa_\mu/\kappa_T = 1/3$. We have explored the entire instability range, varying the inverse density ratio $R_0^{-1}$ between 1 (the onset of direct overturning convection) and $({\rm Pr} + 1)/({\rm Pr} + \tau)$ (the marginal stability limit). Our simulations were performed in a ``small'' domain spanning, in the horizontal direction, about five wavelengths of the fastest-growing double-diffusive mode (i.e. $L_x = L_y = 100d$ where $d$ is a thermal diffusion lengthscale), and in the vertical direction, $L_z = 100d$ or $L_z = 178d$ depending on the runs. In all cases we initialized a double-diffusively unstable system with infinitesimal perturbations, and found that these first grow exponentially according to linear theory, then saturate into a state of homogeneous double-diffusive convection. In that state, the turbulent contribution to thermal and compositional transport is significant but much smaller than that expected from standard convection, ranging from 5-10 times the diffusive rate near the onset of direct convective instability, and rapidly dropping towards zero as $R_0^{-1}$ increases towards marginal stability (see \S\ref{sec:homogen}). For small $R_0^{-1} $, however, the system does not remain in this homogeneously convecting state. Instead, thermo-compositional layers rapidly appear, and transport through the system strongly increases. We showed that the layer formation process is governed by Radko's $\gamma-$instability theory \citep{radko2003mlf,stellmach2010,traxler2010b}, both qualitatively and quantitatively. In particular, it explains why our simulations with $R_0^{-1} <1.35$ transition into layers while those with $R_0^{-1} > 1.35$ do not. The key factor is the variation of the total buoyancy flux ratio $\gamma_{\rm tot}$ with density ratio (see \S\ref{sec:gamma-instab}): layers can only form when $\gamma_{\rm tot}$ decreases with $R_0$. In the layered phase, we found that the flux through the staircase depends sensitively on the mean layer height $H_L$. Given the large variability of the measured fluxes during the layered phase, our results are roughly consistent both with Spruit's theory \citep{spruit1992} and with heat transport between two solid plates (as given by equation (\ref{eq:convection-from-wall})). Further simulations will be needed to help distinguish between these two possibilities -- or perhaps suggest an alternative one. Finally, note that in our small-domain simulations, the mergers always proceed until a single layer is left. In that sense, the dynamical evolution of the system is always eventually influenced by the domain size. \subsection{Discussion of the applicability of our results to real systems} Our initial goals were threefold: (a) to characterize transport by homogeneous double-diffusive convection (i.e. in the absence of layers), (b) to determine if, under which conditions, and through which process thermo-compositional layers may form and (c) to characterize transport by layered double-diffusive convection when appropriate. To answer part (a) in detail, a much larger number of simulations will be needed, using progressively smaller Prandtl numbers and diffusivity ratios. These are the subject of an ongoing investigation. We hope to find similar scaling laws for ${\rm Nu}_T$ and ${\rm Nu}_\mu$ as functions of {\rm Pr} and $\tau$ as the one found by \citet{traxler2010} for fingering convection. By contrast with the case of fingering convection, however, we now know that thermo-compositional staircases {\it can} form spontaneously from double-diffusive convection. Radko's criterion for layer formation, namely that $\gamma_{\rm tot}$ should decrease with $R_0$, is the answer to part (b) of our goals, but does require knowledge of the function $\gamma_{\rm tot}(R_\rho)$ to be applied in practice. The latter must be determined separately for each parameter set $({\rm Pr}, \tau)$. Finally, our findings have provided some insight into part (c). Since transport through a staircase depends on the layer height only (for given fluid parameters and overall stratification), the problem shifts to estimating actual layer heights in astrophysical objects. The layers we observe in our simulations have a strong tendency to merge, which suggests two possible outcomes: the mergers continue indefinitely, until the scale of the equilibrium layers is commensurate with the system size; or the mergers stop for other reasons (see below), with an equilibrium layer height significantly smaller than the system size. It is of course crucial to know which of these two scenarios is correct, as they imply vastly different transport rates through the staircase. Unfortunately, our simulations were not able to provide a definitive answer to this question. In the two cases studied, the final layer height was equal to the domain height, but this should not be interpreted as a result in favor of the first scenario since this could simply mean that our domain was too small to ``contain'' the intrinsic equilibrium layer height of the second scenario. \citet{radko2005dtl} proposed a theory supporting the idea that mergers stop before layers reach the system size, and deduced a means of estimating the equilibrium layer height. Starting from an initial staircase with uniform ``jumps'' in temperature and chemical composition across the interfaces, he studied how the staircase evolves if it is perturbed slightly, by making some of the jumps larger and some of the jumps weaker. He concluded that the staircase is unstable to mergers if the total buoyancy flux ratio through the interfaces is a decreasing function of the density ratio across the interfaces -- a criterion very similar to the $\gamma-$instability criterion. We have tried to test Radko's merger theory against our simulations, but this has unfortunately proven to be difficult. The statistical fluctuations in the measured turbulent fluxes (see Figure \ref{fig:nuevol-layers} for example) are too large to detect a significant variation of the interfacial flux ratio as the mergers proceed. We would need a much larger domain to improve the signal to ``noise'' ratio to a point where our results could be compared with his theory. We would also need a much taller domain (at least a few times taller than the equilibrium layer height) to see if the merger process indeed stops as expected. Finally, we would need to integrate the simulation long enough to establish convincingly that the mergers have indeed stopped. Unfortunately, running equivalent simulations in a much larger domain, and for long enough to observe the layer formation and merger process, is impossible within current numerical limitations. \subsection{Future prospects} The preliminary findings presented in this paper still enable us to lay out a clear path towards obtaining better parametrizations of mixing by double-diffusive convection in the near future, using currently available computational resources: \begin{itemize} \item Firstly, we must gain a better understanding of the instability saturation mechanism at low Prandtl number and low diffusivity ratio, in order to determine the flux laws ${\rm Nu}_T({\rm Pr}, \tau, R_0)$ and ${\rm Nu}_\mu({\rm Pr}, \tau, R_0)$ for homogeneous double-diffusive convection. These flux laws are needed to determine when layers are expected to form, and can be used ``as is'' to parametrize mixing otherwise. They {\it can} be measured using ``small-domain'' simulations similar to the ones we have presented here, at least for values of ${\rm Pr}$ and $\tau$ as low as about 0.01 or so. Semi-analytical weakly-nonlinear models will then be helpful to guide extrapolations to the much lower parameter values appropriate of the astrophysical regime. \item Secondly, we must gain a general understanding of mixing in the layered case, at low Prandtl number and low diffusivity ratio. In order to do this, we need to determine how interfacial transport depends on the fluid parameters $({\rm Pr},\tau)$ and on the interfacial density ratio (ie. a density ratio based on the difference in temperature and composition across the layers). We must also understand how transport scales {\it within} the convective layers, as a function of the same parameters but also as a function of the layer height. This can be done today using simulations in which a single layer is pre-seeded, to bypass the rather lengthy layer formation and merger phases. Using this information, we will be able to test the basic flux laws which are central to Radko's merger theory more quantitatively \citep{radko2005dtl}. If this theory holds, then one can straightforwardly deduce the equilibrium layer height for a given parameter set, and ultimately quantify the staircase transport properties. \end{itemize}

1012.0617

Double-diffusive convection, often referred to as semi-convection in astrophysics, occurs in thermally and compositionally stratified systems which are stable according to the Ledoux criterion but unstable according to the Schwarzschild criterion. This process has been given relatively little attention so far, and its properties remain poorly constrained. In this paper, we present and analyze a set of three-dimensional simulations of this phenomenon in a Cartesian domain under the Boussinesq approximation. We find that in some cases the double-diffusive convection saturates into a state of homogeneous turbulence, but with turbulent fluxes several orders of magnitude smaller than those expected from direct overturning convection. In other cases, the system rapidly and spontaneously develops closely packed thermo-compositional layers, which later successively merge until a single layer is left. We compare the output of our simulations with an existing theory of layer formation in the oceanographic context and find very good agreement between the model and our results. The thermal and compositional mixing rates increase significantly during layer formation and increase even further with each merger. We find that the heat flux through the staircase is a simple function of the layer height. We conclude by proposing a new approach to studying transport by double-diffusive convection in astrophysics.

12133647

2010arXiv1012.5204K

1012

1012.5204_arXiv.txt

The cosmic microwave background (CMB) is radiation that comes from the early universe. In the early universe, ordinary matter was in the form of hot hydrogen and helium plasma which was almost homogeneously distributed in space. Almost all electrons were free. Because of scattering from these electrons, the mean free path of photons was short compared to cosmological distance scales: the universe was opaque. As the universe expanded, this plasma cooled, and first the helium ions, then also the hydrogen ions captured the free electrons: the plasma was converted into gas and the universe became transparent. After that the photons of the thermal radiation of this primordial plasma have travelled through the universe and we observe them today as the CMB. The CMB is close to isotropic, i.e., the microwave sky appears almost equally bright in every direction. With sensitive instruments we can, however, see small variations, the CMB anisotropy. \begin{figure}[ht] \begin{center} \includegraphics[width=14cm]{Planck_FIRST_LIGHT_SURVEY_lo.eps} \caption{The sky at optical and microwave wavelengths: A sky map of the first two week of observations by the Planck satellite at the 70 GHz frequency, superimposed on an optical image of the sky. From Ref.~\cite{Psite}. \emph{Credit: ESA, LFI \& HFI Consortia (Planck), Optical image: Axel Mellinger.}} \label{fig:fls} \end{center} \end{figure} There is also ``foreground'' microwave radiation that comes from astrophysical sources, our own galaxy and other galaxies. In Fig.~\ref{fig:fls} we see the radiation from the Milky Way as a horizontal red band in the microwave image, whereas further away from the galactic plane we see variations in the intensity of the CMB. The foreground can be separated from the CMB by measuring at several frequencies, since it has a different electromagnetic spectrum. The formation of helium and hydrogen atoms is called recombination, although in this context it is a misnomer, since this was the first time the ions and electrons formed atoms. The related increase of the photon free mean path beyond cosmological distance scales is called photon decoupling. This happened when the age of the universe was about 380 000 years old. At this time there were small density variations, about one part in ten thousand in the primordial plasma/gas. After photon decoupling, the over-densities began to grow by gravitational attraction and eventually led to the formation of galaxies and stars hundreds of millions of years later. \begin{figure}[ht] \begin{center} \includegraphics[width=16cm]{101080_7yrFullSky_WMAP_1280W_lo.eps} \caption{Temperature anisotropy of the CMB according to 7 years of measurements by the WMAP satellite. This is a false-color image, where yellow and red indicate hotter than average, and blue colder than average. From Ref.~\cite{WMAPsite}. \emph{Credit: NASA / WMAP Science Team.}} \label{fig:wmap7} \end{center} \end{figure} When looking at the CMB we are thus looking at the 380 000 year old early universe. We see those distant parts of the universe from which it has taken the whole remaining part of the history of the universe for the light to travel from there to here. The observed small variations in the CMB reflect the small density variations at that time. See Fig.~\ref{fig:wmap7}. Because of the finite speed of light, everything we see lies on our past light cone (see Figure~\ref{fig:intr_and_jour}). The intersection of our past light cone with the time of photon decoupling forms a sphere, which we call the sphere of last scattering. It is this sphere that we observe when we observe the CMB: we see each photon coming from the location where it last scattered from an electron. When the photons travel from the last scattering sphere to here they are redshifted by the expansion of the universe. The universe has expanded by a factor of 1100 since last scattering, and therefore the photon wavelengths have been stretched by that factor. Photons decoupled when the temperature of the primordial plasma/gas was about 3000 K, and therefore the photons had then a blackbody spectrum with that temperature. When all wavelengths of a blackbody spectrum are stretched by the same factor, the spectrum remains blackbody, but its temperature falls with the same factor. The observed mean temperature of the CMB is $T_0 = 2.725\pm0.001$ K today \cite{COBEtemp}. However, because of the inhomogeneity of the universe, photons coming from different directions have suffered slightly different redshifts, which is another contribution to the observed CMB anisotropy. Thus the variation $\delta T(\theta,\phi)$ of the observed temperature $T(\theta,\phi) = 2.725 K + \delta T(\theta,\phi)$ can be divided into two contributions, $\delta T_\mathrm{intr}$ that is due to inhomogeneous conditions at the last scattering surface, and $\delta T_\mathrm{jour}$ that arises as the photons travel from the last scattering sphere to here. \begin{figure}[ht] \begin{center} \includegraphics[width=12cm]{intr_and_jour.eps} \caption{A spacetime diagram of our past light cone.} \label{fig:intr_and_jour} \end{center} \end{figure} An important thing of the anisotropy $\delta T(\theta,\phi)$ is that it is small. The root-mean-square variation is about $100 \mu$K, or \beq \frac{\delta T}{T_0} \sim 4\times10^{-5} \,. \eeq While this makes observing this anisotropy very difficult, it simplifies understanding and calculating the physics that causes this anisotropy: The primordial density perturbations $\delta\rho$ that are responsible for this anisotropy must have also been very small, and we can calculate their evolution using first-order perturbation theory around a homogeneous and isotropic model of the universe, the so-called background model. The deviations from this background model are small, so we can ignore any products of two or more such small quantities. This makes the equations linear, so that they can be easily Fourier transformed, and lead to a direct relation between initial and final values. CMB was discovered by Penzias and Wilson \cite{PeWi} in 1964, using a microwave antenna at Bell Laboratories in Holmdel, New Jersey. The CMB anisotropy was first measured by the COBE satellite \cite{COBEanis} in 1992, and much more accurate measurements have later been taken by the WMAP satellite and are currently being taken by the Planck satellite. In this lecture I sketch our present understanding how the CMB anisotropy arises. The relevant physics involves quantum field theory in curved spacetime (for the generation of primordial perturbation) and general relativistic perturbation theory (for their evolution and effect on the CMB), and it is not possible to give a full account in this short lecture. However, many parts of the relevant physics are relatively easy to understand, and I try to present those here; for the other parts I just have to give results without derivation, in an attempt to present a continuous story. I also give a short overview of the ongoing Planck satellite mission to observe the CMB.

1012.5204

This lecture is a sketch of the physics of the cosmic microwave background. The observed anisotropy can be divided into four main contributions: variations in the temperature and gravitational potential of the primordial plasma, Doppler effect from its motion, and a net red/blueshift the photons accumulate from traveling through evolving gravitational potentials on their way from the primordial plasma to here. These variations are due to primordial perturbations, probably caused by quantum fluctuations in the very early universe. The ongoing Planck satellite mission to observe the cosmic microwave background is also described.

2173923

2011MNRAS.412.2391A

1012

1012.3034_arXiv.txt

The observation of the diffuse, X--ray emitting medium (a.k.a. intra-cluster medium, or ICM) of galaxy clusters provides quantities like its mass, temperature ($T$) and X--ray luminosity ($L_{X}$). The analysis of the scaling relation between these physical quantities gives considerable insight into the physical processes in the ICM (e.g. Rosati et al. 2002 and reference therein). On the other hand, the evolution of these scaling relations is difficult to predict theoretically (e.g. Norman 2010). The simplest model (Kaiser 1986), in which the ICM evolution is governed only by gravity, predicts an $L_X-T$ relation shallower than observed (Markevich 1998). This suggests that non-gravitational energy inputs, such as merger shocks or feedback from active galactic nuclei (AGNs) and star formation, need to be considered. More sophisticated models sensitively depend on the assumed physics of the baryons, and their predictions can be tuned to be in good agreement with observed scaling relations (Kravtov et al. 2005; Nagai, Kratsov \& Vikhlinin 2007; Bode, Ostriker \& Vikhlinin 2009) measured in the nearby Universe, if one accept an overprediction of the baryon fraction in stars by an order of magnitude (Gonzalez, Zarisky \& Zabludoff 2007, Andreon 2010). The most direct way to probe ICM evolution is to measure the scaling relations over a wide range of redshifts. Here a difficulty arises: many cluster samples with known $L_X$ and $T$ are either X-ray selected, or are heterogenous collections of objects without a simple and accountable selection function. In both cases, neglecting the selection function may bias the $L_X-T$ relation (Stanek et al. 2006; Pacaud et al. 2007; Nord et al. 2008), because at a given temperature clusters more luminous enter more easely in the sample (they can be seen on a larger volume, have smaller temperature errors, and are more frequently in archive and samples). Therefore, the mean $L_X$ at a given $T$ can be systematically over-estimated, unless one accounts for the selection function (e.g. Gelman et al. 2004, Pacaud et al. 2007, Andreon \& Hurn 2010). The requirement of a known selection function restricts the choice of the available samples and the redshift baseline making hard to detect deviations from a self--similar evolution for lack of extension at high redshift. For example, $z\le 1.05$ for Pacaud et al. (2007), and $z<0.2$ for Pratt et al. (2009). Only a handful of clusters are known at high $z$ (four at $z>1.4$). In this paper we use the only two suitable for this study, namely JKCS\,041, % probably the most distant cluster known to date, and ISCS\,J1438+3414 (at $z=1.41$, Stanford et al., 2005), the second most distant cluster that can be used for studying scaling relations. Note that the redshift of JKCS\,041, conservatively estimated at $z=1.9$ in Andreon et al. (2009) and has now a red-sequence estimated redshift of $z=2.20\pm0.11$ (Andreon \& Huertas-Company 2010). Both are optically-NIR selected, i.e. are detected through their galaxies, and have been subsequently followed up in X rays (see Andreon et al. 2009 for JKCS\,041 and this paper for ISCS\,J1438+3414) to derive $L_X$ and T for the gas. Though small, this sample is free from the biases that affect X-ray selected samples, since these clusters are considered independently from their X--ray luminosity. By using them, we extend the redshift baseline to $z\sim 2$, where the self-similar model predicts a brightening 1.7 times larger than at $z=1$. \begin{figure} \centerline{% \psfig{figure=HST_Xraycont_100_fix.ps,width=8.5truecm,clip=}% } \caption[h]{Contours from an adaptively smoothed Chandra image in the [0.3-2] keV energy band superposed onto an Hubble Space Telescope (F850LP) image of ISCS\,J1438+3414.} \end{figure} We adopt the following cosmological parameters: $\Omega_\Lambda=0.7$, $\Omega_m=0.3$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$. The scale, at $z=1.41$, is 8.4 kpc arcsec$^{-1}$. As point estimate and error measurements, we quote posterior mean and standard deviation when a Bayesian approch is esplicitely mentioned, or, otherwise the usual profile likelihood-based estimates (e.g. XSPEC error, $-2 \Delta \ln \mathcal{L} = 1$).

The large redshift leverage considered in this paper has provided a direct, though not yet compelling, evidence that clusters do not evolve self-similarly in the last $10.6$~Gyr, about three quarter of the current age of the Universe. We remark that our result relies on a large redshift leverage, rather than on a detailed analysis of small effects on large samples at lower redshift. If confirmed, the trend we have found implies that non-gravitational effects, such as baryon physics, % began long ago to shape the clusters' scaling relations. In particular, the observed evolution is in line with the predictions of simulation that include high-redshift pre-heating and radiative cooling in addition to shock heating, such as those in Short et al. (2010). They predict that our clusters should be a factor 3 to 4 fainter than self--similar evolution while we observe a factor 5. Instead, their models that include feedback directly tied to galaxy formation or that incorporate gravitional heating only strongly disagree with our observations. This conclusion should not over-emphasized, because we are still a long away from having the numerical resolution required to really implement these mechanisms (e.g. Norman 2010), for example to follow the formation of stars, whose feedback is deemed important for the evolution of the gas properties. It is of the utmost importance to extend the sample of non X-ray selected clusters to $z> 1.4$, to confirm the modulation provided by non-gravitational phenomena in the cluster evolution. We emphasize the need of non-X-ray selected samples: X-ray selected samples should be treated with caution when used in this context, because the probability that an object is in the sample is not random in $L_X$ at a given $T$. Optically/near-infrared selected samples should instead be used since their selection is not due to their X-ray properties, unless we were able to predict their individual X-ray luminosity relative to the average X-ray luminosity at a given $T$ in absence of X-ray data and we were to make use of this information to select the objects. If confirmed, the breakdown of the self--similar evolution, would have important consequences for the cosmological studies. Indeed, the evolution with redshift of the scaling relations is very sensitive to cosmological parameters (e.g. Allen et al. 2004; Albrecht et al. 2006, Report of the Dark Energy Task Force, and references therein). A proper assessment of the intrinsic processes shaping the scaling relations is fundamental for the use of galaxy cluster surveys, such as the planned WFXT (Conconi et al. 2010) and JDEM (Sholl et al. 2009), as probes of the cosmological parameters.

1012.3034

The evolution of the properties of the hot gas that fills the potential well of galaxy clusters is poorly known, since models are unable to give robust predictions, and observations lack a sufficient redshift leverage and are affected by selection effects. Here, with just two high-redshift, z≈ 1.8, clusters avoiding selection biases, we obtain a significant extension of the redshift range and we begin to constrain the possible evolution of the X-ray luminosity versus temperature relation. The two clusters, JKCS 041 at z= 2.2 and ISCS J1438+3414 at z= 1.41, are, respectively, the most distant and the second most distant clusters, overall, that can be used for studying scaling relations. Their location in the X-ray luminosity versus temperature plane, with an X-ray luminosity five times lower than expected, suggests at the 95 per cent confidence level that the evolution of the intracluster medium has not been self-similar in the last three-quarters of the age of the Universe. Our conclusion is reinforced by data on a third, X-ray-selected, high-redshift cluster, too faint for its temperature when compared to a sample of similarly selected objects. Our data suggest that non-gravitational effects, such as the baryon physics, influence the evolution of galaxy clusters. Precise knowledge of evolution is central for using galaxy clusters as cosmological probes in planned X-ray surveys, such as WFXT or JDEM.

12225907

2011PhRvD..83h4002H

1012

1012.4939_arXiv.txt

\label{sec:introduction} There has been excellent progress towards gravitational wave astronomy over recent years. The first generation of large scale gravitational wave interferometers reached unprecedented sensitivities and have undertaken extended science runs. The U.S. \ac{LIGO} \cite{Abbott:2009li}, the French--Italian Virgo \cite{Acernese:2006bj} and the German--British GEO600 \cite{Willke:2007zz} detectors now form a collaborative network of interferometers. The data from these detectors has been analyzed for gravitational waves from compact binary coalescence \cite{Abadie:2010yba}, stochastic background \cite{Abbott:2009ws}, unmodelled burst \cite{Abadie:2010mt} and pulsar \cite{Collaboration:2009rfa} sources. \ac{S6} and \ac{VSR23} ended in October 2010 and yielded the most sensitive data yet taken; the analysis of this data is ongoing. In the meantime, the detectors are being upgraded to their advanced configurations \cite{Harry:2010zz,avlligowebsite, advvirgowebsite}, with the expectation of a ten fold improvement in sensitivity. With these sensitivities, it is expected that gravitational waves will be observed regularly \cite{Temp:2010cfa}. Furthermore, with a proposed advanced detector in Japan \cite{lcgtwebsite}, a possible detector in Australia \cite{aigowebsite}, and 3rd generation detectors on the horizon \cite{Hild:2008ng}, future prospects are promising. As the gravitational wave community matures it is essential that a relationship is built between \ac{GW} and \ac{EM} astronomers. The \ac{GW} emission from a source is likely to provide complementary information to emission in various \ac{EM} bands, and a joint observation is significantly more likely to answer outstanding astrophysical questions. Already this relationship is beginning to mature. A number of \ac{EM} transients have already been followed up in \ac{GW} data \cite{Abbott:2007rh, Abadie:2010uf, Collaboration:2009kk}. Additionally, infrastructure is also being put in place to allow for \ac{EM} follow-up of \ac{GW} observations \cite{Kanner:2008zh}. \ac{CBC} are one of the most promising sources of gravitational waves, and also an ideal candidate for joint \ac{GW}-\ac{EM} astronomy. During the late stages of inspiral and merger, a compact binary emits a distinctive, ``chirping'' gravitational wave signal. Furthermore, \ac{CBC}s containing at least one \ac{NS} are expected to emit electromagnetically. Specifically, \ac{BNS} and \ac{NSBH} mergers are the preferred progenitor model for the short \ac{GRB} \cite{nakar:2007,Shibata:2007zm}. It is also possible that these mergers will be observable electromagnetically as orphan afterglows \cite{nakar:2007}, optical \cite{Metzger:2010sy} or radio transients \cite{Predoi:2009af}. Since \ac{GRB}s are well localized both in time and on the sky by \ac{EM} observations, the corresponding \ac{GW} search can be simplified by reducing the volume of parameter space relative to an all-sky, all-time search. Targeted searches for \ac{CBC} waveforms associated to short \ac{GRB}s were performed using data from \ac{S5} and \ac{VSR1} \cite{Abbott:2007rh, Abadie:2010uf}. In this paper, we introduce a targeted, coherent search algorithm for detecting \ac{GW} from \ac{CBC}. This targeted search is designed as a follow-up to \ac{EM} transients, and in particular \ac{GRB}s. Previous searches for this source have made use of a coincidence requirement --- namely that a signal with consistent parameters is observed in two or more detectors in the network. The analysis introduced here makes use of the data from all operational detectors, and combines the data in a coherent manner before matched filtering against \ac{CBC} template waveforms. In a coherent analysis, it is straightforward to restrict the signal model to only two independent polarizations. This allows for the rejection of incoherent background noise, and consequently increases the sensitivity of the search if more than two detectors are operating. The data output by gravitational wave interferometers is neither stationary nor Gaussian, but is contaminated by noise transients of instrumental and environmental origin. This makes the task of doing data analysis a complex one, and matched filtering alone is not sufficient to distinguish signal from noise. Regardless of whether a coincident or coherent search is performed, the most significant events by \ac{SNR} would always be dominated by non-Gaussian transients, or ``glitches'', in the data. A significant effort goes into understanding the cause of these glitches \cite{Blackburn:2008ah} and removing times of poor data quality from the analysis. While these efforts greatly reduce the number of glitches they cannot remove them entirely. Therefore the analysis must also employ methods to distinguish signal from noise transients. In previous \ac{CBC} searches, signal consistency tests \cite{Allen:2004gu, Hanna:2008} have proved very effective at removing the non-Gaussian background. We extend these tests to the coherent analysis described in this paper and demonstrate their continued effectiveness. In addition, coherent analyses naturally lend themselves to multi-detector consistency tests, such as the null stream \cite{Guersel:1989th}. We describe a number of such consistency tests for this templated \ac{CBC} search and again demonstrate their efficacy. Finally, we are able to show that the various signal consistency tests are sufficient to remove the majority of non-Gaussian transients and render the search almost as sensitive as if the data were Gaussian and stationary. The layout of this paper is as follows. In section \ref{sec:coh_matched_filter}, we describe the formulation of a targeted coherent triggered search for \ac{CBC} signals. In section \ref{sec:snr_cont_consistency} we discuss an implementation of the null stream formalism and other multi-detector consistency tests. In section \ref{sec:chi2tests} we describe a number of $\chi^{2}$ tests than can be applied in a coherent search to try to separate and veto glitches. Finally, in \ref{sec:s4data} we outline an implementation of a targetted, coherent search for \ac{CBC} and present results on both simulated, Gaussian data and real detector data taken from \ac{S4}.

We have presented a formulation of a targeted coherent search for compact binary coalescences. For Gaussian noise, the coherent \ac{SNR} would be ideal for separating signals from the noise background. However, since data from gravitational wave interferometers is neither Gaussian nor stationary, we have also discussed a number of methods of separating the non-stationary noise background from the signal population. These tests include various $\chi^{2}$ tests, which were originally designed for use in single detectors. We have extended them to the network analysis and demonstrated their continued efficacy. Additionally, the coherent analysis allows for some additional tests which are not readily available in the coincidence case. The most significant of these is the null \ac{SNR} which can be used to reject events which are not consistent with two gravitational wave polarizations. Additionally, we explored consistency tests between the recovered amplitudes of the gravitational wave and found that a simple \ac{SNR} threshold on the two most sensitive detectors gave excellent results. The analysis described in this paper has been implemented and in the final section we showed results of a test run. This made use of the \ac{S4} data from the LIGO detectors. Although the data was far from Gaussian, after the application of all of the signal consistency tests the results were remarkably close to what would be expected in Gaussian noise. This analysis is available to be used in searches for GW inspiral signals associated with GRBs in more recent LIGO and Virgo data, such as S6 and VSR2 and VSR3. There are a number of ways in which this analysis could be enhanced to broaden its use and increase its sensitivity. First, a number of \ac{GRB}s, particularly those observed by Fermi \cite{fermiwebsite} and IPN \cite{ipnwebsite} are not localized sufficiently accurately that the error box can be treated as a point on the sky. Thus, it would be nice to extend this analysis to allow for a region of the sky to be covered. This would require looping over the relevant sky points; incorporating the correct detector sensitivities $F_{+,\times}$ and time delays. In principle, this would not greatly slow down the analysis as the majority of time is taken in performing the single detector filters and these would \textit{not} need to be re-calculated. As well as looking at a patch on the sky, the analysis could be extended to cover the whole sky, as appropriate for an un-triggered search. This brings in a host of new complications which have been met and dealt with by other coherent search methods \cite{1367-2630-12-5-053034, Klimenko:2008fu}. In order to obtain a good estimate of the background for an all sky, un-triggered search we would need to implement background estimation and time shifting the data would likely be the best way to do this. Since \ac{GRB}s are thought to be rather tightly beamed, it is reasonable to take them as being face on, or close to. In this case, the gravitational waves are circularly polarized and there is, in effect, only a single polarization. This opens the possibility of limiting the signal space to just this one polarization and adding an extra ``null'' test. Alternatively, it should be possible to perform a Bayesian marginalization over the astrophysically expected distributions of the various parameters. The progenitors of short \ac{GRB}s are thought to be \ac{BNS} or \ac{NSBH}. The search we have described is ideal for the \ac{BNS} case as the spins of the neutron stars are unlikely to have a significant effect on the waveform. However, when one of the components of the binary is a black hole, the spin could be large. Furthermore, the mass ratio is likely to be relatively large. In this case, the spin of the black hole can have a significant effect on the observed waveform \cite{PBCV04}. Consequently, we would like to extend this search to incorporate spin effects. The infrastructure described in this paper can already accept spinning waveforms, but the implementation of signal based vetoes proves somewhat more complex. Work is underway on this \cite{HFspin}.

1012.4939

We introduce a method for conducting a targeted, coherent search for compact binary coalescences. The search is tailored to be used as a follow-up to electromagnetic transients such as gamma-ray bursts. We derive the coherent search statistic for Gaussian detector noise and discuss the benefits of a coherent, multidetector search over coincidence methods. To mitigate the effects of nonstationary data, we introduce a number of signal consistency tests, including the null signal-to-noise ratio, amplitude consistency, and several χ<SUP>2</SUP> tests. We demonstrate the search performance on Gaussian noise and on data from LIGO’s fourth science run and verify that the signal consistency tests are capable of removing the majority of noise transients, giving the search an efficiency comparable to that achieved in Gaussian noise.

12218071

2011PhLB..699..320C

1012

1012.3943_arXiv.txt

Nowadays, there is great interest in modified $f(R)$ gravity (see~\cite{Francaviglia} for recent reviews). Such theories are interesting in that they generalize Einstein's general relativity and provide insights into the consequences of quantum corrections to its equations in the high energy regime ($R\rightarrow \infty$). In cosmology, the interest in these theories comes from the fact that they can naturally drive an accelerating cosmic expansion without introducing dark energy, as happens for instance in the standard $\Lambda$CDM cosmology. However, the freedom in the choice of different functional forms of $f(R)$ gives rise to the problem of how to constrain the many possible $f(R)$ gravity theories. In this regard, much efforts have been developed so far, mainly from the theoretical viewpoint~\cite{Ferraris}. General principles such as the so-called energy conditions~\cite{energy_conditions}, nonlocal causal structure~\cite{RSantos}, have also been taken into account in order to clarify its subtleties. More recently, observational constraints from several cosmological data sets have been explored for testing the viability of these theories~\cite{Amarzguioui}. An important aspect that is worth emphasizing concerns the two different variational approaches that may be followed when one works with $f(R)$ gravity theories, namely, the metric and the Palatini formalisms (see, e.g., \cite{Francaviglia}). In the metric formalism the connections are assumed to be the Christoffel symbols and the variation of the action is taken with respect to the metric, whereas in the Palatini variational approach the metric and the affine connections are treated as independent fields and the variation is taken with respect to both. Because in the Palatini approach the connections depend on the particular $f(R)$, while in metric formalism the connections are defined {\it a priori} as the Christoffel symbols, the same $f(R)$ seems to lead to different space-time structures. In fact, these approaches are certainly equivalents in the context of general relativity (GR), i.e., in the case of linear Hilbert action; for a general $f(R)$ term in the action, they seem to provide completely different theories, with very distinct equations of motion. The Palatini variational approach, for instance, leads to 2nd order differential field equations, while the resulting field equations in the metric approach are 4th order coupled differential equations, which presents quite unpleasant behavior. These differences also extend to the observational aspects. For instance, we note that cosmological models based on a power-law functional form in the metric formulation fail in reproducing the standard matter-dominated era followed by an acceleration phase~\cite{Amendola} (see, however, \cite{cap}), whereas in the Palatini approach, analysis of a dynamical autonomous systems for the same Lagrangian density have shown that such theories admit the three post inflationary phases of the standard cosmolog ical model~\cite{Tavakol}. Although being mathematically more simple and successful in passing cosmological tests, we do not yet have a clear comprehension of the properties of the Palatini formulation of $f(R)$ gravity in other scenarios, and issues such as solar system experiments~\cite{Faraoni}, the Newtonian limit~\cite{Meng-Wang} and the Cauchy problem~\cite{Lanahan} are still contentious (regarding this last issue, see however Refs.~\cite{Cauchy-Problem}). In this paper, we explore cosmological consequences of a class of exponential $f(R)$-gravity models in the Palatini formalism. In order to test the observational viability of these scenarios, we use one of the latest type Ia Supernovae (SNe Ia) sample, the so-called Union2 compilation~\cite{Amanullah} along with 11 determinations of the expansion rate $H(z)$~\cite{newh,svj}. In what concerns the past evolution of the Universe, we show that for some intervals of model parameters a matter-dominated era is followed by a late time accelerating phase, differently from some results in the metric approach. Another interesting feature of this class of models is the possibility of a transient cosmic acceleration, which can lead the Universe to a new matter-dominated era in the future. This particular result seems to be in agreement with current requirements from String/M theory.

\label{Conclusion} Cosmological models based on $f(R)$-gravity may exhibit a natural acceleration mechanism without introducing a dark energy component. In this paper, we have investigated cosmological consequences of a class of exponential $f(R)$-gravity in the Palatini formalism, as given by Eq. (\ref{expo-gravity}). We have performed consistency checks and tested the observational viability of these scenarios by using one of the latest SNe Ia data, the so-called Union2 sample with 557 data points and 11 measurements of the expansion rate $H(z)$ at intermediary and high-$z$. We have found a good agreement between these observations and the theoretical predictions of the model, with the reduced $\chi^2_{min}/\nu \simeq 1$ for the three tests performed. Differently from the dynamical behavior of other $f(R)$ scenarios discussed in the literature (either in metric or Palatini formalisms), we have found solutions of transient cosmic acceleration in which the large-scale modification of gravity will drive the Universe to a new matter-dominated era in the future. As mentioned earlier, this kind of solution is in full agreement with theoretical requeriments from String/M theories, as first pointed out in Ref.~\cite{fischler}. Finally, we have also shown that, differently from the results of Ref.~\cite{Amendola} for power-law $f(R)$ gravity in the metric formalism, exponential $f(R)$ models corresponding to the best-fit solutions from SNe Ia and SNe Ia + $H(z)$ $\chi^2$ minimization have the usual matter-dominated phase followed by a late time cosmic acceleration (see also \cite{cap} for a discussion).

1012.3943

We investigate cosmological consequences of a class of exponential f(R) gravity in the Palatini formalism. By using the current largest type Ia Supernova sample along with determinations of the cosmic expansion at intermediary and high-z we impose tight constraints on the model parameters. Differently from other f(R) models, we find solutions of transient acceleration, in which the large-scale modification of gravity will drive the Universe to a new decelerated era in the future. We also show that a viable cosmological history with the usual matter-dominated era followed by an accelerating phase is predicted for some intervals of model parameters.

1118418

2011PhRvD..83h3505P

1012

1012.3458_arXiv.txt

\par Many experiments are currently searching for Dark Matter (DM) in the form of Weakly Interacting Massive Particles (WIMPs), by looking for rare scattering events off nuclei in the detectors, and many others are planned for the next decade \cite{book,Lewin,Bergstrom,Munoz,Bertone,CerdenoGreen}. This direct DM detection strategy has brought over the last year several interesting observations and upper limits. The results of the DAMA/LIBRA \cite{dama} and, more recently, the CoGeNT \cite{cogent} collaborations have been tentatively interpreted as due to DM particles. It appears however that these results cannot be fully reconciled with other experimental findings, in particular with the null searches from XENON100 \cite{xe100,gelmini1,Bezrukov} or CDMS \cite{cdms10}, and are also in tension with ZEPLIN-III \cite{zeplin}. In this context, the next generation of low-background, underground detectors is eagerly awaited and will hopefully confirm or rule out a DM interpretation. \par If convincing evidence is obtained for DM particles with direct detection experiments, the obvious next step will be to attempt a reconstruction of the physical parameters of the DM particle, namely its mass and scattering cross-section (see e.g.~Refs.~\cite{Goudelis,Green1,ST}). This is a non-trivial task, hindered by the different uncertainties associated with the computation of WIMP-induced recoil spectra. In particular, Galactic model uncertainties -- i.e.~uncertainties pertaining to the density and velocity distribution of WIMPs in our neighbourhood -- play a crucial role. In attempting reconstruction, the simplest assumption to make is a fixed local DM density $\rho_0=0.3 \textrm{ GeV/cm}^3$ and a ``standard halo model'', i.e.~an isotropic isothermal sphere density profile and a Maxwell-Boltzmann distribution of velocities with a given galactic escape velocity $v_{esc}$ and one-dimensional dispersion $\sigma^2\equiv v_0^2/2 = v_{lsr}^2/2$ ($v_0$ being the most probable velocity and $v_{lsr}$ the local circular velocity, see below). However, the Galactic model parameters are only estimated to varying degrees of accuracy, so that the true local population of DM likely deviates from the highly idealised standard halo model. \par Several attempts have been made to improve on the standard approach \cite{ST,Ling,McCabe,Green2}. In the case of a detected signal at one experiment, recent analyses have studied how complementary detectors can extract dark matter properties, independent of our knowledge of the Galactic model~\cite{fox1}. Certain properties of dark matter may also be extracted under assumptions about the nature of the nuclear recoil events~\cite{fox2}. Furthermore, eventual multiple signals at different targets have been shown to be useful in constraining both dark matter and astrophysical properties \cite{APeter} and in extracting spin-dependent and spin-independent couplings \cite{Vergados,BertoneCerdeno}. Here, using a Bayesian approach, we study how uncertainties on Galactic model parameters affect the determination of the DM mass $m_{\chi}$ and spin-independent WIMP-proton scattering cross-section $\sigma_{SI}^p$. In particular we focus on realistic experimental capabilities for the future generation of ton-scale detectors -- to be reached within the next 10 years -- with noble liquids (argon, xenon) and cryogenic (germanium) technologies. \par The main focus of this paper is the complementarity between different detection targets. It is well-known (see e.g.~\cite{Lewin}) that different targets are sensitive to different directions in the $m_{\chi}-\sigma_{SI}^p$ plane, which is very useful to achieve improved reconstruction capabilities -- or more stringent bounds in the case of null results. This problem has often been addressed without taking proper account of Galactic model uncertainties. Using xenon (Xe), argon (Ar) and germanium (Ge) as case-studies, we ascertain to what extent unknowns in Galactic model parameters limit target complementarity. A thorough understanding of complementarity will be crucial in the near future since it provides us with a sound handle to compare experiments and, if needed, decide upon the best target to bet on future detectors. Our results also have important consequences for the combination of collider observables and direct detection results (for a recent work see \cite{Fornasa}). \par Besides degrading the extraction of physical properties like $m_{\chi}$ and $\sigma_{SI}^p$, uncertainties in the Galactic model will challenge our ability to distinguish between different particle physics frameworks in case of a positive signal. Other relevant unknowns are hadronic uncertainties, related essentially to the content of nucleons \cite{Ellis}. Here, we undertake a model-independent approach without specifying an underlying WIMP theory and using $m_{\chi}$ and $\sigma_{SI}^p$ as our phenomenological parameters -- for this reason we shall not address hadronic uncertainties (hidden in $\sigma_{SI}^p$). A comprehensive work complementary to ours and done in the supersymmetric framework has been presented recently \cite{Akrami1,Akrami2}. \par The paper is organised as follows. In the next section, we give some basic formulae for WIMP-nucleus recoil rates in direct detection experiments. In Section \ref{exp} the upcoming experimental capabilities are detailed, while Section \ref{statsection} describes our Bayesian approach. We outline the relevant Galactic model uncertainties and our modelling of the velocity distribution function in Section \ref{astrounc} and present our results in Section \ref{results} before concluding in Section \ref{secconc}.

\label{secconc} \par We have discussed the reconstruction of the key phenomenological parameters of WIMPs, namely mass and scattering cross-section off nuclei, in case of positive detection with one or more direct DM experiments planned for the next decade. We have in particular studied the complementarity of ton scale experiments with Xe, Ar and Ge targets, adopting experimental configurations that may realistically become available over this time scale. To quantify the degree of complementarity of different targets we have introduced a figure of merit measuring the inverse of the area enclosed by the 95\% marginalised contours in the plane $\log_{10}(m_\chi)-\log_{10}(\sigma_{SI}^p)$. There is a high degree of complementarity of different targets: for our benchmark with $m_\chi=50$ GeV and our fiducial set of Galactic model parameters, the relative error on the reconstructed mass goes from 8.1\% for an analysis based on a xenon experiment only, to 5.2\% for a combined analysis with germanium, to 4.5\% adding also argon. Allowing the parameters to vary within the observational uncertainties significantly degrades the reconstruction of the mass, increasing the relative error by up to a factor of $\sim$4 for xenon and germanium, especially due to the uncertainty on $\rho_0$ and $v_0$. However, we found that combining data from Ar, Ge and Xe should allow to reconstruct a 50 GeV WIMP mass to 11.8\% accuracy even under weaker astrophysical constraints than currently available. Although the mass reconstruction accuracy may appear modest, any improvement of this reconstruction is important, in particular in view of the possible measurement of the same quantity at the Large Hadron Collider at CERN. The existence of a particle with a mass compatible, within the respective uncertainties, with that deduced from direct detection experiments would provide a convincing proof that the particles produced in accelerators are stable over cosmological time scales. Although this is not sufficient to claim discovery of DM \cite{Fornasa}, it would certainly be reassuring. Despite the strong dependence of direct detection experiments on the Galactic model degrades the reconstruction of DM properties, it does open up the possibility to potentially constrain the local distribution of DM, in case of detection with multiple targets. For example in the case of a low mass 50 GeV WIMP, we have shown that the local circular velocity can be determined from direct detection data alone more accurately than it is presently measured using the local distribution of stars and gas clouds. Additionally, directly detecting DM provides the most realistic way of measuring the local DM velocity distribution. This will in principle provide invaluable information on the structure and formation of the Milky Way halo. \vspace{0.5cm} \par {\it Acknowledgements:} G.B., R.T.~and M.P.~would like to thank the organisers of the workshop ``Dark Matter all around'' for a stimulating meeting. We wish to thank the authors of the paper \cite{Akrami1} for providing their preliminary results, as well as Henrique Araujo and Alastair Currie for useful discussions. We also acknowledge support from the SNF grant 20AS21-29329 and the University of Zurich. M.P.~is supported by Funda\c{c}\~{a}o para a Ci\^encia e Tecnologia (Minist\'erio da Ci\^encia, Tecnologia e Ensino Superior).

1012.3458

We investigate the reconstruction capabilities of the dark matter mass and spin-independent cross section from future ton-scale direct detection experiments using germanium, xenon, or argon as targets. Adopting realistic values for the exposure, energy threshold, and resolution of dark matter experiments which will come online within 5 to 10 years, the degree of complementarity between different targets is quantified. We investigate how the uncertainty in the astrophysical parameters controlling the local dark matter density and velocity distribution affects the reconstruction. For a 50 GeV WIMP, astrophysical uncertainties degrade the accuracy in the mass reconstruction by up to a factor of ∼4 for xenon and germanium, compared to the case when astrophysical quantities are fixed. However, the combination of argon, germanium, and xenon data increases the constraining power by a factor of ∼2 compared to germanium or xenon alone. We show that future direct detection experiments can achieve self-calibration of some astrophysical parameters, and they will be able to constrain the WIMP mass with only very weak external astrophysical constraints.

12164341

2011AIPC.1356...35T

1012

1012.2803_arXiv.txt

\label{sec:intro} Symbiotic stars are interpreted as interacting binaries consisting of a cool visual primary and a hot compact secondary component accreting matter from the atmosphere of its companion. Their spectral variability is determined from the orbital motion and the outburst events of the hot component. These events are often accompanied by intensive loss of mass in the form of optically thick shells, stellar wind outflow and bipolar collimated jets. The stellar wind is possible to be bipolar and highly collimated too. The ejected material forms complicated flow structure in the space between the components, whose elements determine the observed properties of the star. The interacting binary Z And is considered as a prototype of the classical symbiotic stars. It consists of a normal cool giant of spectral type M4.5 \citep{MS}, a hot compact component with a temperature of about $1.5\times10^5$ K \citep{Sok06} and an extended circumbinary nebula formed by the winds of the stellar components and partly photoionized by the compact object. Z And has undergone several active phases consisting of repeated optical brightenings with amplitudes up to $2$--$3$ mag and characterized by intensive loss of mass \citep{SS,Boyarchuk,FC95,Bis06,Sk06,Sok06,TTB08,Sk09,TTB10}. The last active phase of Z And began at the end of August 2000 \citep{Sk00} and continued up to that time including six optical brightenings. The maxima of the light during the active phase were in December 2000 ($ V\sim 8$\fm8), November 2002 ($V \sim 9$\fm8), September 2004 ($V \sim 9$\fm1), July 2006 ($V \sim 8$\fm6), January 2008 ($V\sim 9$\fm5) and January 2010 ($V \sim 8$\fm4) \citep{Sk09} (we also used data of the American Association of Variable Star Observers, AAVSO, and via private communication from A.~Skopal). High resolution optical data indicating highly different physical conditions (velocity, density, temperature) in the line emission regions were obtained during this active phase \citep{Sk06,Sok06,TTB07,TTB08,Sk09,TTB10}. The current study is based on observations acquired during the first and fourth outbursts. The main features of the flow structure in the system during the active phase indicated by our data can be summarized as follows: \begin{enumerate} \item We suppose that an accretion disk is present in the system. Only in this way we can explain the presence of the \mbox{He\,{\sc ii}} $\lambda$4686 line in the spectrum and its behavior during the first outburst \citep{TTB08}. \item Stellar wind was observed in the system during all the outbursts. Two wind components with different velocity regimes were observed simultaneously---wind propagating at a moderate velocity, $0$--$100$ \kms \citep{Sok06,Sk06,TTB08}, and wind with a high velocity, in excess of $500$ \kms \citep{Sk06,TTB08}. \item Bipolar collimated outflow appeared during the fourth outburst. In the beginning of July 2006 the H$_\alpha$ line had an absorption component shifted by $-1400$ \kms\, from its center. It went into emission and later, in July--December 2006, the spectrum contained additional emission components on either side of the central peak corresponding to velocities of $1200$--$1500$ \kms\, \citep{BL,TTB07,Sk09}. The H$_\beta$ line had the same features \citep{TTB07,Sk09}. Both the emission and absorption high-velocity components are assumed to be formed in a bipolar collimated outflow \citep{TTB07}. \end{enumerate} The main aim of the current study is to suggest a model for the flow structure in the system to explain all of these features. We believe that it will be helpful in explaining not only the data acquired during this active phase, but also all spectroscopic data taken during earlier phases, those following 1939, 1960 and 1984 and possibly the activity of other classical symbiotic stars. \section[]{The model of the flow structure in Z~And} \label{sec:model} According to the theoretical models \citep{Bis02,Mitsumoto05,Bis06} an accretion disk is formed around the compact object in the Z~And system when the wind of its giant has quiescent parameters. The radius of the disk is about 50$R_{\sun}$\, for a wind velocity of 20--25 \kms and thus the outer part of the disk is optically thin. The mass of the disk is estimated as product of one quarter of the mass-loss rate of the giant $\sim\!\!2\times10^{-7}M_{\sun}$ yr$^{-1}$ \citep{FC88}, and the typical time interval between the active phases---about 10 years. Based on the size and the mass of the disk of $\sim\!\!5 \times 10^{-7}M_{\sun}$\, we can assume that its inner region is optically thick in the quiescent state of the system. To change the system from quiescent to active state a sufficiently large increase of the accretion rate is needed. Accretion of a considerable fraction of the disk's mass is required in order for an outburst to develop even in the combined model where the increased nuclear burning rate is taken into account \citep{Bis06}. Maximal increase of the accretion rate is possible in the framework of the mechanism proposed by \citet{Bis02} and \citet{Mitsumoto05}. According to this mechanism, even a small increase of the velocity of the wind of the donor is sufficient to change the accretion regime. During the transition from disk accretion to accretion from the flow, the disk is partially disrupted and the increased velocity of the wind causes falling of the material of the disk onto the accretor's surface. However, even in this case a considerable amount of mass (up to $50$--$80$ per cent) stays in the disk. Then massive accretion disk exists in the system during the active phase too. \begin{figure} \includegraphics[width=0.47\textwidth]{fig01_a.eps} \includegraphics[width=0.47\textwidth]{fig01_b.eps} \caption{Left panel: Schematic model of the region around the hot component during the first outburst. Right panel: The same, but in the plane perpendicular to the orbital plane where the emission regions are shown. (From spectroscopy presented in \citet{TTB08}.)} \end{figure} During the outburst the high velocity wind of the hot compact component collides with the accretion disk, its velocity decreases near the orbital plane but does not change at higher stellar latitude. A consequence of this is that two different kinds of stellar wind are observed---P~Cyg wind with a low velocity and an optically thin high velocity wind. As a result of the collision the density of the wind close to the orbital plane increases and the level of the observed photosphere locates at a larger distance from the star. Thus, an optically thick disk-like shell forms, which occults the compact object. Since its effective temperature is lower it is responsible for the continuum energy redistribution and the increase of the optical flux of the star. The interaction of the wind with the disk is equivalent to interaction of two winds and thus can form shocked region of collisional ionization with temperature of up to $10^{6}$ K (Fig.~1) \citep{NW,Bis06}. During the active phase the wind from the hot component ``strips'' the accretion disk carrying some part of its material in the circumbinary envelope. After each outburst some part of this material does not leave the potential well and after the cessation of the wind begins to accrete again. Because of the initial amount of angular momentum a disk-like envelope extending to larger distance from the orbital plane than the accretion disk itself forms (Fig.~2). The existence of a centrifugal barrier creates hollow cones around the axis of rotation with small opening angle \citep{I,BB}. \begin{figure}[!htb] \includegraphics[width=0.47\textwidth]{fig02_a.eps} \includegraphics[width=0.47\textwidth]{fig02_b.eps} \caption{The same like in Fig.~1 but during recurrent strong outburst.} \end{figure} The disk-like envelope does not exist during the first outburst of the active phase. During the recurrent outbursts the extended disk-like envelope can collimate the wind which escapes the compact object predominantly via the two hollow cones (Fig.~2). The bipolar outflow can be observed as high-velocity satellite components situated on either side of the mean peak of the line. Their presence in the spectrum depends on the density of the disk-like envelope (i.e. on the system's activity during its previous phases) and on the wind intensity. They can appear only if the density of the disk-like envelope is sufficient to provide collimation and the mass-loss rate of the compact component is high enough. According to our model these spectral features can be observed during outbursts accompanied by loss of mass at high rate and preceded by similar strong outbursts. According to the current theory the existence of bipolar collimated jets is supposed to be due to presence in the system of magnetized disk which transforms the potential energy of the accreting material into kinetic energy of the outflowing jets. This means that accretion must be realized to provide existence of the jets. In the spectrum of Z~And during its 2006 outburst, however, different indications of stellar wind presented along with satellite emission components. If we believe that ``traditional'' jets were observed, it is necessary to explain accretion onto the star at the time when a strong mass outflow from its surface was observed. On the other hand, the proposed model of collimated stellar wind makes it possible to explain all features of the observed spectrum without any important contradictions. Another question related to the nature of the collimated outflow is that about the orbit inclination. The velocity of the radiatively accelerated wind cannot exceed $3000$ \kms. The highest observed outflow velocity during the 2006 outburst was $1500$ \kms. A preliminary analysis of our spectra obtained during late December 2009 showed velocities of $1700$--$1800$ \kms. Taking into account these velocities we obtain an upper limit of the orbit inclination of about $55$\degr. That is why the view about the nature of the outflow can shed some light on the question about the orbit inclination. We suppose that the model suggested by us can be used to explain the behavior of the line spectrum of other classical symbiotic stars during their active phases too since the classical symbiotic stars have the next general characteristics: \begin{enumerate} \item The mass transfer in the majority of the classical symbiotic stars is realized by means of the stellar wind of the giant. The theoretical computations of \citet{Bis02} and \citet{Mitsumoto05} show that for systems with parameters close to those of Z~And an accretion disk from wind accretion forms around the compact object. The accretion disk prevents the outflowing material during the active phase and outflow with two-component velocity regime forms. \item The systems with parameters close to Z~And have similar accretion rate of their compact object. The accretion rate determines the regime of hydrogen burning. The view that hydrogen burns in a steady state in the classical symbiotic stars is commonly accepted. When the accretion rate goes above the upper limit of the steady burning range the white dwarf expands which is observed as optical outburst with typical duration of one year. The first outburst opening the active phase can be followed by repeated outbursts. During the repeated outbursts an extended envelope surrounding the accretion disk exists in the system and it can be responsible for the collimation of the stellar wind. \end{enumerate}

1012.2803

Results of the study of the symbiotic binary Z And during its recent active phase 2000-2010 when it experienced a series of six optical outbursts are presented. High-resolution spectra obtained during the first and fourth outburst, which was the strongest one, have been analyzed. These data are compared with results of theoretical computations. The comparison provides information about the behavior of the system during the entire active phase rather than during an individual outburst. In particular it was found fundamental difference between the first outburst which opened the active phase and the recurrent outbursts--namely, the presence of bipolar collimated optical outflow during some of the recurrent outbursts. A scenario that can explain all the spectroscopic phenomena observed during this active phase as well as previous active phases of Z And is proposed. The possibility to use this scenario for explanation of the line spectrum of other classical symbiotic stars during their active phases is motivated.

12306454

2012PThPh.127..355I

1012

1012.0070_arXiv.txt

Gravitational lensing has become an important subject in modern astronomy and cosmology \cite{Schneider,Weinberg}. It has many applications as gravitational telescopes in various fields ranging from extra-solar planets to dark matter and dark energy at cosmological scales \cite{Refregier}. For instance, it is successful in detecting extra-solar planetary systems \cite{SW,MP,GL,Bond,Beaulieu}. Gaudi et al.\cite{Gaudi} have found an analogy of the Sun-Jupiter-Saturn system through lensing. Recently gravitational lensing has been used to constrain modified gravity at cosmological scale \cite{Reyes}. This paper considers the gravitational lensing by point-mass systems on multiple planes, where the number of planes is arbitrary. Such a multi-plane treatment is important. In microlensing studies, we usually assume a binary lens on a single lens plane. In order to discuss its validity, we can consider two lens planes and later take the limit that two lens planes merge. In this way, it will become possible to estimate the effect caused by a separation between the double lens planes. Another importance is for gravitational lensing in cosmology. Clearly, galaxies at different redshifts and dark matter inhomogeneities must be described by not a single-plane but multi-plane method. It has long been a challenging problem to express the image positions as functions of lens and source parameters \cite{Asada02a,Asada03}. For this purpose, we present a method of Taylor-series expansion to solve the multi-plane lens equation in terms of mass ratios by extending the previous single-plane work \cite{Asada09}. In particular, we carefully investigate, as a non-trivial task, the denominators of the lens equation with singular points. The multi-plane lensed-image counting theorem states that the lower bound on the image number is $2^N$ for $N$ planes with a single point mass on each plane (page 458 in Petters, Levine and Wambsganss \cite{PLW} and references therein). However, the counting theorem tells nothing about the image positions. Therefore, it is important to discuss {\it how} such image positions are realized in an analytical method. Under three assumptions of weak gravitational fields, thin lenses and small deflection angles, gravitational lensing is usually described as a mapping from the lens plane onto the source plane \cite{SW}. Bourassa and Kantowski \cite{BKN,BK} introduced a complex notation to describe gravitational lensing. Their notation was used to describe lenses with elliptical or spheroidal symmetry \cite{Borgeest,Bray,Schramm}. For $N$ point lenses, Witt \cite{Witt90} succeeded in recasting the lens equation into a single-complex-variable polynomial. This is in an elegant form and thus has been often used in investigations of point-mass lenses. The single-variable polynomial due to $N$ point lenses on a single plane has the degree of $N^2+1$, though the maximum number of images is known as $5(N-1)$ \cite{Rhie01,Rhie03,HN06,HN08}. This means that unphysical roots are included in the polynomial (for detailed discussions on the disappearance and appearance of images near fold and cusp caustics for general lens systems, see also Petters, Levine and Wambsganss \cite{PLW} and references therein). Following Asada \cite{Asada09}, we consider the lens equation in dual complex variables, so that we can avoid inclusions of unphysical roots. This paper is organized as follows. In Section 2, the formulation of multi-plane lens systems with complex variables is briefly summarized. The lens equation is iteratively solved. In section 3, we present iterative solutions for a two-plane case and give an algorithm for computing image positions for an arbitrary number of lens planes in terms of mass ratios. In section 4, we discuss how lensed-image positions are realized in the present method. Section 5 presents numerical tests. Section 6 is devoted to the conclusion.

We made a systematic attempt to determine, as a function of lens and source parameters, the positions of images by multi-plane gravitational lenses. We presented a method of Taylor-series expansion to solve the multi-plane lens equation in terms of mass ratios except for the neighborhood of the caustics. In concordance with the multi-plane lensed-image counting theorem that the lower bound on the image number is $2^N$ for $N$ planes with a single point mass on each plane, our iterative results directly show how $2^N$ images are realized except for the neighborhood of the caustics. It is left as a future work to compare the present result with state-of-art numerical simulations.

1012.0070

Continuing work initiated in an earlier publication (H. Asada, Mon. Not. R. Astron. Soc. 394 (2009), 818), we make a systematic attempt to determine, as a function of lens and source parameters, the positions of images by multi-plane gravitational lenses. By extending the previous single-plane work, we present a method of Taylor-series expansion to solve the multi-plane lens equation in terms of mass ratios except for the neighborhood of the caustics. The advantage of this method is that it allows a systematic iterative analysis and clarifies the dependence on lens and source parameters. In concordance with the multi-plane lensed-image counting theorem that the lower bound on the image number is 2^N for N planes with a single point mass on each plane, our iterative results show how 2^N images are realized. Numerical tests are done to investigate if the Taylor expansion method is robust. The method with a small mass ratio works well for changing a plane separation, whereas it breaks down in the inner domain near the caustics.

12234660

2012APh....35..537B

1012

1012.0041.txt

High energy Cosmic Ray (CR) positrons are promising targets for indirect searches of Galactic particle Dark Matter (DM) \cite{DMreview}. The recent results reported by the PAMELA \cite{PAMELA} and Fermi collaborations \cite{Abdo:2009zk,Ackermann:2010ij,Ackermann:2011rq} on the positron fraction $e^{+}/(e^{+}+ e^{-})$ and on the $(e^{+}+ e^{-})$ (CRE) spectra in the GeV $\div$ TeV energy range show large discrepancies with standard astrophysical model predictions and have indeed raised a large number of interpretations in terms of DM. In particular, it has been shown that a good fit of the PAMELA/Fermi data can be achieved with $e^+e^-$ produced by a DM particle of $\sim$TeV mass annihilating or decaying predominantly via leptonic channels, and several models realizing this scenario have been proposed \cite{Bergstrom:2008gr,Cholis:2008hb,Cirelli:2008pk,ArkaniHamed:2008qn,Bergstrom:2009fa}. However, interpretations based on discrete astrophysical extra sources (like e.g.~pulsars, or stochastic local sources) \cite{Blasi:2009hv,Blasi:2010de, Serpico:2008te,CRE_interpretation1,DiBernardo:2010is,Blasi:2011fm,Profumo:2008ms,Mertsch:2010fn} have been shown to provide equally good fits to the data (see e.g. \cite{Serpico:2011wg} for an extended critical review of the subject). It is, however, very unlikely to distinguish the two scenarios using as observables only the CRE fluxes, even with the larger statistics expected in the future \cite{Pato:2010im}. It is thus mandatory to find other observables accessible to experiments, that are as much model independent as possible and can provide a clear discrimination between a DM dominated scenario and an ``astrophysically'' dominated one. The intrinsic degree of dipole anisotropy in the arrival directions of high energy CREs expected from a DM scenario, $\delta_{DM}$, is indeed insensitive to many uncertainties, and constitutes, to a good approximation, a universal characteristics of galactic DM, \red{as we will show in the following}. The reason why the dipole anisotropy has a very weak dependence on the many unknowns involved in the problem is, on the one hand, the very short ($\sim 1~\kpc$) electron path above $\sim100~\GeV$ which makes this quantity very local in origin, on the other hand, the fact that it is a flux ratio (see Eq.~\ref{eq:delta}) so that most of the uncertainties cancel each other in the ratio. Furthermore, the anisotropy signal from DM is intrinsically very different from the one due to local discrete sources. Pulsars are rare (their number in the Galaxy is estimated to be around $10^5$ \cite{FaucherGiguere:2009df}) and powerful objects and can induce very large anisotropies typically dominated by a single or a few nearby objects. On the other hand, the number of galactic DM substructures is ${\cal O}(10^{17})$ and they produce a ``collective'' anisotropy which is never dominated by a single clump. The flux from a very nearby clump would be always accompanied by the large, dominant and almost isotropic flux from the whole population of clumps, which washes out the single clump anisotropy. Therefore, the dipole anisotropy offers a viable criterion to discriminate among CREs produced by DM or in local sources. \red{Anisotropies in the DM component have been also studied in gamma-rays (see for example \cite{Cuoco:2010jb,SiegalGaskins:2009ux,Fornasa:2009qh}) which have clearly the advantage of being independent of the choice of a diffusion models. On the other hand, as we will show, CREs anisotropies are also fairly independent of the propagation setup over a wide range of possible diffusion models and thus offer an interesting complementary anisotropy probe of DM.} Besides the anisotropy from DM and \emph{local discrete} astrophysical sources, there is also a third source of anisotropy which needs to be considered in order to have a complete picture, i.e. the anisotropy from the large scale distribution of the astrophysical sources considered as a whole. This component, as we will show, gives generally a smaller anisotropy with respect to the first two components above. On the experimental side, Fermi-LAT recently placed the first upper limits on the integrated dipole anisotropy of the arrival directions of CRE with $E> 60~\GeV$ \cite{Ackermann:2010ip}, and there are prospects for its actual observation after a few years of data taking, if local pulsars contribute significantly to the CRE fluxes above $\sim100~\GeV$ \cite{DiBernardo:2010is}. Also AMS-02 \cite{AMSweb} is now taking data, but its sensitivity to CRE anisotropy is much lower than the one of Fermi \cite{Pato:2010im}. We will show that if Fermi-LAT or future experiments will find an anisotropy larger than the maximum DM anisotropy we derive here, then a dominant DM contribution to the CRE anisotropy can be excluded in a basically model independent way, pointing instead to local discrete astrophysical CRE sources as the main source of anisotropy. Therefore, the observation of anisotropy at the level within reach by Fermi-LAT in the next years will be able to constrain significantly the flux of CREs possibly contributed by DM annihilations in the Galaxy. On the other hand, to identify pulsars as responsible for a possible anisotropy would require at least a careful analysis of their spectral characteristics and of the direction and intensity of the anisotropy. This paper is structured as follows: in Section \ref{sec:DMint} we describe how we compute the interstellar CRE density due to astrophysical sources and due to DM annihilations in the smooth halo and in substructures. In section \ref{sec:DManiso} we detail how we simulate the distribution of galactic DM substructures. \red{In section \ref{sec:universal} we compute the total intrinsic DM anisotropy (i.e. the anisotropy when, ideally, only DM contributes to the total CRE emission) while in section \ref{sec:mixed} we discuss the anisotropy for a mixed scenario in which both DM and standard (non discrete) astrophysical sources contribute to the CRE flux}. Section \ref{sec:discussion} is finally devoted to our final comments and conclusions.

\label{sec:discussion} % \begin{figure*}[tbp] \centering \includegraphics[width=0.45\textwidth]{fig6a.pdf} \hspace{2pc} \includegraphics[width=0.45\textwidth]{fig6b.pdf} \caption{\emph{Left}: Black solid contours: anisotropy at 500 GeV due to a single clump as function of its mass and distance from the Earth when substructures of mass down to $10^{-6}~M_{\odot}$ are included. Grey dashed contours: expected number of clumps of a given mass closer than a given distance to the Earth. Numbers represent the $\log_{10}$ of the related quantities. The plot is drawn for 3 TeV DM fully annihilating into muon pairs, assuming a NFW profile and KRA propagation setup. \emph{Right}: same as left panel, but when no substructures other than the clump itself (and no main Halo) are included. The white dashed curves represent the log$_{10}$ of the boost factor (with respect to thermal cross section) required for the CRE fluxes to be in agreement with the observed fluxes at 500 GeV.} \label{fig:singleclump} \end{figure*} % Our findings result from a MonteCarlo computation of the local distribution of DM substructures and a possible bias of this approach is that we might have missed configurations whose probability is less than 1\%, in which, e.g., a large mass clump emerges isolated and very close to the Earth. This could in principle produce an anisotropy larger than what we quote as a ``maximum''. We checked, however, that this configuration cannot produce a high degree of anisotropy. Indeed, even in the unlikely case of a $10^{8}~M_{\odot}$ clump at 100~pc from Earth (whose probability is $< 0.1\%$ \cite{Brun:2009aj}) the anisotropy is strongly suppressed by the nearly isotropic flux of the much more abundant smaller substructures and it is thus always diluted below $\max(\delta_{DM})$. %Diffusion below 100~GeV smoothes out completely the single clump anisotropy. At higher energies, instead, the nearly %isotropic flux of the much more abundant smaller mass substructures present within 100~pc dilutes the anisotropy produced %by the close-by, high mass clump. % %even if it is possible to find a clump with $M > %10^{8}~M_{\odot}$ at $d \simeq 100~\pc$ from the Earth --and the %probability for this is $< 0.1\%$ \cite{Brun:2009aj}-- we should %take into account the much more (by several orders of magnitude) %abundant substructures with smaller mass which are present within %$100~\pc$. Their flux would dilute the anisotropy produced by the %close-by, high mass clump. % % This feature makes the DM signal intrinsically different from the one expected from pulsars. Indeed, while there might be a close-by, isolated pulsar, that can possibly lead to a large anisotropy \cite{DiBernardo:2010is}, it is not possible to reproduce this configuration with DM. The situation is different also from $\gamma$-rays, where this clump would be a quite bright point source. %%Finally, close-by small mass substructures, such as %%$10^{-6}~M_{\odot}$ clumps, cannot be responsible for a large anisotropy either. %%%, as it can be inferred extrapolating from Fig.~\ref{fig:500GeV}. Another possible caveat is that low mass clumps ($m_{cl}\sim 10^{-6}~M_{\odot}$) are so abundant that in principle they can be found within 1 pc from Earth, hence CREs could reach the Earth before diffusing significantly. %which is the distance where transition from diffusive to ballistic propagation for TeV CREs occurs. Based on their number density, we expect to find only a few substructures with mass $10^{-6}~M_{\odot}$ within 1 pc from Earth. These clumps would look more like point-like sources of $e^+e^-$ rather than like a dipole. Even in this case, however, their point-like flux both in $e^+e^-$ and $\gamma$-rays would be several orders of magnitude below the Fermi sensitivity. These points are illustrated in the left panel of Fig.~\ref{fig:singleclump} where the contours show the anisotropy at 500 GeV of a single clump as function of its mass and distance from the Earth. The plot represents the same case assumed in drawing Fig.~\ref{fig:triangle}: 3 TeV DM fully annihilating into muon pairs and assuming NFW profile and KRA propagation setup. The anisotropy is calculated as in Eq.~\ref{eq:DManiso} with at the numerator the gradient of the flux of the clump and at the denominator the total flux from all the clumps and the smooth Halo so to represent the effective contribution of the clump to the total anisotropy. As it can be seen, the anisotropy has a plateau at about $10^{-2}$ (corresponding to the case where the observer is well inside a very massive clump) which does not exceed the maximum possible anisotropy at the same energy. A possible exception to the above scenario is the extreme case in which a massive DM clump is the ``only" relevant CRE source. In fact, any other contribution from other clumps or the smooth halo would dilute the anisotropy of this single clump. We note here that in this scenario, in order for the halo contribution to be negligible with respect to the one of the single clump, one needs to invoke a strong suppression of the annihilation cross section of the DM in the halo much below the thermal value. Moreover, clumps of at least $10^7~M_{\odot}$ are known experimentally to exist in the form of dwarf galaxies \cite{Strigari:2008ib} and they would share the same boost factor as the hypothetical nearby clump, so that their presence would dilute the anisotropy of the single clump anyway. This case is discussed here only because of its extreme geometry and is analyzed in the right panel of Fig.~\ref{fig:singleclump}. The anisotropy is again calculated as in Eq.~\ref{eq:DManiso} but this time with at the numerator the gradient of the flux of the clump and at the denominator the flux of the clump itself. The needed boost factor with respect to the thermal cross section in order for the clump to produce all the observed CRE flux at 500 GeV is shown as the white, dotted curves. This gives the maximal possible anisotropy and, as we expected, it can exceed the upper limits of Fig.~\ref{fig:env}, being possibly as large as $10^{-1}$. However, the only clump configurations which would give such a high anisotropy lie in a region of parameter space (small mass or large distance) where unrealistically high boost factors (of the order of $10^6$ or higher) are required. If we restrict the allowed region to the more reasonable case of a nearby massive clump (lower right corner of the plot), again the maximal anisotropy does not exceed $\sim 10^{-2}$ since the anisotropy is somewhat reduced by the fact that the observer is well inside the clump and sees the structure of the clump. We notice that large boost factors (e.g.~via Sommerfeld enhancement) are strongly constrained by analyses of the CMB distortion during the recombination epoch \cite{Galli:2009zc,Slatyer:2009yq,Galli:2011rz}. Further constraints on the boost factor also come for DM annihilation in the core of the Earth \cite{Albuquerque:2011ma}. Also, in the case of a small clump, besides unrealistically high boost factors being required, it would be not justified anymore to not include the contribution of the other equally or more massive clumps, falling back to a configuration of low anisotropy similar to the one of Fig.~\ref{fig:singleclump} (left). For both cases it should be further stressed that, although above $\sim200$ GeV the positron fraction and thus the possible DM fraction is unconstrained, it is unlikely that DM constitutes 100\% of the CRE flux. Considering the contribution of the astrophysical component will lower the overall anisotropy of the clump although the precise decrease will be dependent on the particular AP model employed. %A possible exception to the above scenario is the extreme case in which a very massive nearby clump is the only DM substructure present in all the Galaxy, so that no contribution from other clumps would dilute the anisotropy of this single clump. %We analyze this scenario in the right panel of Fig.~\ref{fig:singleclump}. The anisotropy is again calculated as in Eq.~\ref{eq:DManiso} but this time with at the numerator the gradient of the flux of the clump and at the denominator the flux of the clump itself. This represents the extreme scenario in which the clump makes all the observed CRE flux and gives the maximal possible anisotropy. It can be seen indeed that in this case anisotropies of the order of $10^{-1}$ would be possible, exceeding the upper limits of Fig.~\ref{fig:env}. However, the clump configurations which would give such a high anisotropy lie in a region of parameter space (small mass or large distance) where unrealistically high boost factors (of the order of $10^{5}$ or higher) would be required. If we restrict the allowed region to the more reasonable case of a nearby massive clump (lower right corner of the plot), again the maximal anisotropy does not exceed $\sim 10^{-2}$ since the anisotropy is somewhat reduced by the fact that the observer is well inside the clump and sees the structure of the clump. For both cases it should be further stressed that, although above $\sim200$ GeV the positron fraction and thus the possible DM fraction is unconstrained, it is unlikely that DM constitutes 100\% of the CRE flux. Considering the contribution of the astrophysical component will lower the overall anisotropy of the clump although the precise decrease will be dependent on the particular AP model employed. % Another remark concerns the density profiles we considered. While $N$-body simulations suggest spiked halo and subhalo matter density profiles, astrophysical observations of many dwarf spiral galaxies point to a shallower, Burkert-like density profile \cite{Salucci:2000ps}. Our results are stable under the relevant change from a spiked to a cored profile. Indeed, high energy CREs arriving at Earth do not carry information on the DM distribution in the galactic center, as they propagate only a few kpc in the interstellar medium. The anisotropy is not sensitive to the internal concentration of the subhaloes as well, because diffusion over kpc scales smooths out the effect of a possible cusped over-density region. For the same reason, in the case of decaying DM we find similar results as in the case of annihilating DM. We also neglected the effects of a possible proper motion of substructures. Indeed, as it was pointed out in \cite{Regis:2009qt} for the case of an isolated substructure, a dynamical treatment would lead to a slightly enhanced dipole anisotropy only for sources moving towards the Solar System. However, while this effect can be relevant for a single clump, it is expected to average away for a population of clumps as considered here. Finally, a word of caution must be said about our choice of diffusion models. It might in fact be that the local anisotropy observed on Earth can be affected by local magnetic turbulence, which would break the assumption of isotropic and spatially uniform diffusion we used in this work. Hints in this direction may come from the observations of dipole anisotropies and on anisotropies on angular scales of the order of $10^{\circ}\div30^{\circ}$ in the \red{CR hadronic component} at energies $\gtrsim 10~\TeV$ \cite{Abdo:2008aw,Abdo:2008kr,Vernetto:2009xm,Abbasi:2010mf,Toscano:2011dc,Abbasi:2011zk}, \red{as discussed in several works} \cite{Battaner:2009zf,Salvati:2008dx,Drury:2008ns,Malkov:2010yq,Lazarian:2010sq,Giacinti:2011mz}. \red{Anisotropies typically increase as a function of energy on account of the increasing gyro-radius, and thus the intensity of the anisotropies observed above $\sim 10$ TeV will be correspondingly decreased when rescaled to our energy range ($\sim$100 GeV - 1 TeV). On the other hand leptons in this energy range have a much smaller horizon with respect to hadrons (due to their shorter propagation length) and this typically can increase the anisotropies, since local effects are more important. Thus, it would be difficult to understand which is the dominant effect and to assess precisely the effects of local turbulence on our results. We remark, however, that the observed hadronic dipolar anisotropy seems generally in agreement with the hypothesis of isotropic and homogeneous diffusion \cite{Blasi:2011fm}, while local magnetic turbulence seems to be required mainly to explain the anisotropy at higher multipoles \cite{Giacinti:2011mz}. Since we are considering in this work only dipolar anisotropies, our results are likely less affected by the above effects. } %We remark, however, that we are considering in this work only dipolar anisotropies, which suffer less from the effects of local magnetic turbulence on scales on the order of the particle gyro-radius (which are instead needed to explain smaller scale anisotropies, see \cite{Giacinti:2011mz}), and in particular in the DM case the rather large dipole anisotropy is mainly caused by the clumpy distribution of sources (large fluctuations in the presence of discrete source distribution were also found and discussed in the case of standard astrophysical sources in \cite{Blasi:2011fm}). Moreover, at the energies that we consider in this work, local turbulence effects on the dipole anisotropy are expected to be further reduced, because of the smaller gyro-radius of the involved CRs. %Although the models we chose were the results of detailed statistical %analyses which excluded models with larger $\alpha$, we remark that %such models would indeed raise the DM anisotropy. However, we checked %that even in the case $\alpha=0.7$ the predicted anisotropy would not %overshoot current upper limits. % %We finally remark that no boost factor has been included in our calculation. Indeed any global boost factor of the annihilation cross section would simplify in an exact way in the definition of $\delta_{DM}$, while it is unlikely that the energy-dependent boost factor due to the clumpy DM distribution could be larger than $\cal O$(1) \cite{Pieri:2009je,Lavalle:1900wn}. \red{In summary, barring the above caveats, we demonstrate that our results on DM anisotropy are robust with respect to several choices of propagation setup and of DM spatial distribution and particle model, and we thus propose to use them as a criterion to reject or at least disfavor a DM dominated scenario in the case of detection of a large anisotropy in high energy CREs. }

1012.0041

We study the dipole anisotropy in the arrival directions of high energy CR electrons and positrons (CRE) of dark matter (DM) origin. We show that this quantity is very weakly model dependent and offers a viable criterion to discriminate among CRE from DM or from local discrete sources, like e.g. pulsars. In particular, we find that the maximum anisotropy which DM can provide is to a very good approximation a universal quantity and, as a consequence, if a larger anisotropy is detected, this would constitute a strong evidence for the presence of astrophysical local discrete CRE sources, whose anisotropy, instead, can be naturally larger than the DM upper limit. We further find that the main source of anisotropy from DM is given by the fluctuation in the number density of DM sub-structures in the vicinity of the observer and we thus devote special attention to the study of the variance in the sub-structures realization implementing a dedicated Montecarlo simulation. Such scenarios will be probed in the next years by Fermi-LAT, providing new hints, or constraints, about the nature of DM.

2101094

2011Icar..212..416M

1012

1012.0727.txt

The efficiency of the formation of rocky planets and the cores of gas giant planets depends sensitively on the amount of solids present in each location of the solar nebula. In most models of planet formation the solar nebula has been assumed to resemble the structures proposed by \citet{1977Ap&SS..51..153W} and \citet{1981IAUS...93..113H}. In this model, and subsequent incarnations of it, the surface density of the gas is approximately given by \begin{equation} \label{eq:gas mmsn} \Sigma^{\mathrm{MMSN}}_{\mathrm{gas}}(r) = 1700 \left(\frac{R}{1\;\mathrm{AU}}\right)^{-3/2}\quad \mathrm{gram/cm}^2 \end{equation} \citep[see][]{2006plfo.book..129T} where the acronym MMSN stands for ``minimum mass solar nebula'', a name by which this model is often identified. In the inner regions (small values of $R$) of the nebula the dust grains are refractory in nature, while beyond some radius $R_{\mathrm{ice}}$ water ice is present, most likely in the form of ice mantels surrounding the refractory grains. In the Hayashi MMSN model this location is assumed to be $R_{\mathrm{ice}}=2.7$ AU, which is consistent for temperatures expected in an optically thin nebula. Often this water ice sublimation boundary radius is called the `snow line'. Since ices may contain a considerable mass compared to the mass in refractory grains, the total surface density was assumed to jump at $R=R_{\mathrm{ice}}$: \begin{equation} \label{eq:mmsn} \Sigma^{\mathrm{MMSN}}_{\mathrm{solid}}(r) = 7.1 F_{\mathrm{ice}} \left(\frac{R}{1\;\mathrm{AU}}\right)^{-3/2}\quad \mathrm{gram/cm}^2 \end{equation} with \begin{equation} F_{\mathrm{ice}} =\left\{ \begin{matrix} 1, & \hbox{for} & R< R_{\mathrm{ice}} \\ 4.2, & \hbox{for} & R> R_{\mathrm{ice}} \end{matrix}\right. \end{equation} \citep{2006plfo.book..129T} This model serves in many planet formation models to compute the midplane gas and solid density and the isolation mass of planetary embryos. However, it has been shown that the outcomes of planet formation models depends critically on the value of $R_{\mathrm{ice}}$ and the strength of the ice jump $F_{\mathrm{ice}}$ \citep[see e.g.][]{2009A&A...501.1139M}. Also dust coagulation and planet formation and migration models are very sensitive to these parameters \citep[e.g.][]{2008A&A...480..859B, 2008ApJ...685..584I}. Simply taking the values to be $R_{\mathrm{ice}}=2.7$\,AU and $F_{\mathrm{ice}}=4.2$ (for $R>R_{\mathrm{ice}}$) is probably too simplistic and may lead to major errors in predictions of the progress of the planet formation process. Besides the amount of solid mass available also the strength of the bonds between dust grains is of importance to their aggregation behavior. The presence of ice coatings on grains can significantly increase their 'stickiness', allowing grains to coagulate more easily. Perhaps even more importantly the bonds formed are stronger preventing aggregates that formed to be destroyed again by collisions. Over the last decade or so, there has been enormous progress in astronomy in the understanding of the structure and evolution of `protoplanetary disks', i.e.~the dust+gas disks surrounding many very young stars. These disks very likely give a good impression of what our own solar nebula looked like 4.567 billion years ago, and the knowledge gained in that field can help getting a better handle on the structure of the solar nebula and the location of the ice condensation front. Indeed, modeling tools that were originally meant for modeling protoplanetary disks and comparing model predictions to observations have been used to make models of the solar nebula and the snow line. For instance, the 1+1D disk structure model of \citet{2000A&A...358..378H}, in which the equations are first solved vertically (1D) and then connected radially (1+1D), has been used by \citet{2001ApJ...554..391H} to study water in the solar nebula, in particular concerning the issue of the D/H ratio in meteorites. \citet{2005ApJ...627L.153D} made a model of the surface density of the solar nebula disk as a function of radial coordinate and accretion rate. He used this model to study the ice condensation front \citep{2005ApJ...620..994D}. This model includes both the effect of irradiation of the disk by the young sun and the heating of the disk near the midplane through viscous dissipation of potential energy (i.e.\ accretional heating). Like in the Hersant study, these models are 1+1D models, i.e.\ they are 1-D vertical models of density and temperature as a function of $z$. The Davis models include radiative transfer using a variable eddington factor method. \citet{2006ApJ...640.1115L} also modeled this, using an updated version of the \citet{1997ApJ...490..368C} two-layer disk model, i.e.\ slightly simpler than the 1+1D modeling of Davis. Recently, \citet{2009Icar..200..672D} combined complex models of ice formation and viscous evolution of the disk with a simple radiative transfer approximation to compute the evolution of the solid surface density. While these modeling efforts have improved the understanding of the distribution of ices in the solar nebula substantially, they still rely on rather dramatic simplifications of the treatment of radiative transfer. The reason for this is very understandable: it is quite a numerical challenge to model multi-dimensional radiative transfer in protoplanetary disks that are very optically thick and actively accreting (and hence producing heat near the midplane). However, in recent years multi-dimensional radiative transfer tools, in particular those based on the Monte Carlo method, have improved dramatically in speed. Extreme optical depths are now not necessarily a problem anymore. We have developed a code, MCMax, that can now handle extreme optical depths efficiently, even if many of the photon packages originate from the optically thickest regions of the disk \citep{2009A&A...497..155M}. In addition, we have adjusted the code to fully take into account the density dependend sublimation state of both the refractory as well as the icy material \citep{2009A&A...506.1199K}. It is therefore a right time to revisit the problem of the snow line with the full force of multi-dimensional radiative transfer and investigate (a) where is the ice line as a function of time and (b) what is the gas and solid surface density distribution in the disk as a function of time. This is the goal of this paper.

We successfully computed the solid density structure and snow line in the early Solar nebula for various parameters of the mass accretion rate and grain composition. This was done using full 3-D radiative transfer taking into account both the energy released by viscous heating as well as radiation from the Sun. The density structure of both the gas and the solids are computed self consistently using an $\alpha$ description for the viscosity and detailed evaporation physics for the dust and ice. We find that: \begin{itemize} \item The location of the snow line is determined by the opacity of the grains and the mass accretion rate. We show that approximate radiative transfer methods to compute the location of the snow line are quite accurate. \item At high mass accretion rates the midplane temperature in the inner few AU can exceed the silicate evaporation temperature. This causes a balance between condensation and evaporation of matter acting like a thermostat keeping the midplane at the silicate evaporation temperature over a significant region. The solid mass in the inner regions is dramatically reduced by this effect and models computing the solid state surface density of the early Solar system should take this into account. \item For our standard set of parameters the surface density in the solid state in the region of the terrestrial planets is below the MMSN for all values of the mass accretion rate. This poses serious constraints on the scenarios of planet formation. It could point to a significantly lower opacity value at the time of planet formation, which can be achieved by increasing the average grain size or significantly change the chemical composition. \item At high accretion rates the increased midplane temperature of the disk due to viscous heating causes the midplane temperatures to be $>1000\,$K in a relatively large part of the disk. Under these conditions thermal processing of dust material as well as sintering of aggregates becomes important. \item We find that for the high mass accretion rates the effects of convective cooling of the midplane are important to take into account. \end{itemize}

1012.0727

The precise location of the water ice condensation front (‘snow line’) in the protosolar nebula has been a debate for a long time. Its importance stems from the expected substantial jump in the abundance of solids beyond the snow line, which is conducive to planet formation, and from the higher ‘stickiness’ in collisions of ice-coated dust grains, which may help the process of coagulation of dust and the formation of planetesimals. In an optically thin nebula, the location of the snow line is easily calculated to be around 3 AU, subject to brightness variations of the young Sun. However, in its first 5-10 myr, the solar nebula was optically thick, implying a smaller snowline radius due to shielding from direct sunlight, but also a larger radius because of viscous heating. Several models have attempted to treat these opposing effects. However, until recently treatments beyond an approximate 1 + 1D radiative transfer were unfeasible. We revisit the problem with a fully self-consistent 3D treatment in an axisymmetric disk model, including a density-dependent treatment of the dust and ice sublimation. We find that the location of the snow line is very sensitive to the opacities of the dust grains and the mass accretion rate of the disk. We show that previous approximate treatments are quite efficient at determining the location of the snow line if the energy budget is locally dominated by viscous accretion. Using this result we derive an analytic estimate of the location of the snow line that compares very well with results from this and previous studies. Using solar abundances of the elements we compute the abundance of dust and ice and find that the expected jump in solid surface density at the snow line is smaller than previously assumed. We further show that in the inner few AU the refractory species are also partly evaporated, leading to a significantly smaller solid state surface density in the regions where the rocky planets were formed.

2774981

2011MNRAS.414.1545F

1012

1012.2305_arXiv.txt

\label{sct:introduction} It is now commonly accepted that the formation of structures in the Universe originated from seed density fluctuations in the dark matter fluid that were laid down during inflation \citep*{HA83.1}. Gravitational instability has the effect of amplifying density perturbations, that at a given point enter the non-linear regime and collapse to form clumps, voids, and more in general the complex Large Scale Structure (LSS henceforth) that we observe today. For a given distribution of the initial conditions, the statistical properties of the LSS are determined uniquely by the subsequent expansion history of the Universe that is, ultimately, on its matter and energy content. In the past decade much effort has been directed toward understanding the effect of dark energy on the formation of structures (see \citealt*{CU09.1,GR09.2,SA09.1}; \citealt{DE10.2}; \citealt*{MO10.1} for recent examples) in order to gain insights on its nature, specifically whether it truly is a cosmological constant or it does have some kind of evolution with cosmic time. After several pioneering works \citep{ME90.1,MO91.1,WE92.1}, only recently has the question of the initial conditions gained renewed attention, specifically about the shape of the primordial density fluctuations distribution. While the simplest models of inflation (single slow-rolling scalar field) predicts this distribution to be virtually indistinguishable from a Gaussian, a plethora of more complex models have been proposed that, in addition to solving the standard cosmological problems, allow for significant and possibly scale dependent deviations from Gaussianity. The study of non-Gaussian cosmologies and the effect they have on the formation and evolution of cosmic structures is thus extremely important in order to rule out inflationary models, and hence to have a better handle on the physics of the early Universe. Moreover, studies on possible detectability of primordial non-Gaussianity are very timely, given the recent claims for a positive skewness of the primordial density fluctuations distribution coming from the Cosmic Microwave Background (CMB) and from the angular correlation function of radio-selected quasars \citep{XI10.1}. The problem of constraining deviations from primordial Gaussianity by means different from the CMB intrinsic anisotropies has recently attracted much efforts in the literature, with studies directed towards the abundance of non-linear structures (\citealt*{MA00.2}; \citealt{VE00.1}; \citealt*{MA04.1}; \citealt{GR07.1,GR09.1,MA10.2}), halo biasing (\citealt{DA08.1,MC08.1}; \citealt*{FE09.1}), galaxy bispectrum \citep{SE07.2,JE09.1}, mass density distribution \citep{GR08.2} and topology \citep{MA03.2,HI08.2}, cosmic shear (\citealt*{FE10.1}, Pace et al. 2010), integrated Sachs-Wolfe effect \citep*{AF08.1,CA08.1}, Ly$\alpha$ flux from low-density intergalactic medium \citep{VI09.1}, $21-$cm fluctuations \citep*{CO06.2,PI07.1} and reionization \citep{CR09.1}. In this work we focused attention on the spatial distribution of galaxies and galaxy clusters as tracers of the LSS. We were particularly interested in comparing the performances of the power spectra of the two individual tracers in constraining primordial non-Gaussianity, and evaluate the improvements in forecasted constraints given by the addition of the cluster-galaxy cross power spectrum. Throughout the paper we assumed a fiducial future all-sky optical/near-infrared survey on the model of \emph{Euclid} \citep{LA09.1,BE10.1}. In order to fully exploit the potentials of both the imaging and the spectroscopy part of \emph{Euclid} we considered galaxies as spectroscopically selected according to their H$\alpha$ flux, and galaxy clusters as selected to be the high signal-to-noise ratio (S/N) peaks in full-sky cosmic shear maps. This approach has the advantage of allowing treatment of the galaxy and cluster samples as independent. The results obtained here are relevant for other planned missions with a similar concept to \emph{Euclid}, such as the NASA Wide Field Infrared Survey (WFIRST). The rest of the paper is organized as follows. In Sections \ref{sct:ng} and \ref{sct:mf} we summarize the non-Gaussian models that we have explored in this work, as well as their effect on the mass function and large scale bias of dark matter halos. In Section \ref{sct:halomodel} we describe the halo model, the physically motivated framework that we adopted for modeling the power spectrum of clusters and galaxies as well as the cluster-galaxy cross spectrum. In Section \ref{sct:results} we summarize our results and in Section \ref{sct:discussion} we discuss them, with particular emphasis on alternative survey configurations. Finally, in Section \ref{sct:conclusions} we present our conclusions. Throughout this work we refer to the fiducial cosmological model as the one defined by the parameter set derived by the latest analysis of the WMAP data \citep{KO10.1}. This means that density parameters for matter, cosmological constant, and baryons are equal to $\Omega_{\mathrm{m},0}=0.272$, $\Omega_{\Lambda,0} = 0.728$, and $\Omega_{\mathrm{b},0}=0.046$, respectively, the Hubble constant equals $h\equiv H_0/(100 \;\mathrm{km\;s}^{-1}\;\mathrm{Mpc}^{-1})=0.704$, and the normalization of the matter power spectrum is set by $\sigma_8=0.809$.

\label{sct:conclusions} We adopted the physically motivated halo model in order to compute the effect of different kinds of primordial non-Gaussianity on the power spectrum of galaxies and galaxy clusters, as well as on the cluster-galaxy cross spectrum. Specifically, we considered galaxies selected spectroscopically according to their H$\alpha$ flux, and galaxy clusters selected as the largest S/N peaks in cosmic shear maps, having in mind future wide field optical/near-infrared surveys such as \emph{Euclid} and WFIRST. We additionally performed a Fisher matrix analysis in order to forecast the errors on the joint determination of the level of non-Gaussianity $f_\mathrm{NL}$ and the amplitude of the matter power spectrum $\sigma_8$. The main findings of this work can be summarized as follows. \begin{itemize} \item The non-Gaussian corrections to the power spectrum of tracers of the LSS resembles the modifications to the large-scale bias of dark matter halos, as one might naively expect. The largest effect is seen for the local bispectrum shape, while the smallest appears for the enfolded shape. The power spectrum of massive galaxy clusters is more heavily modified than the spectrum of galaxies, because the former are originally more biased than the latter. The modification to the cluster-galaxy cross spectrum lies somewhere in between. \item Galaxies have a much higher constraining power on both $f_\mathrm{NL}$ and $\sigma_8$ as compared to galaxy clusters, due to the much lower abundance of the latter that is not adequately compensated by the larger effect on the relative power spectrum. Assuming a \emph{Euclid}-like survey, while spectroscopically selected galaxies can constrain $f_\mathrm{NL}$ to the level of a few and $\sigma_8$ to the level of $\sim 3 \times 10^{-4}$, errors on parameters derived by clusters alone are at the level of $\Delta f_\mathrm{NL} \sim 10$ and $\Delta\sigma_8 \sim 8\times 10^{-3}$ (with some dependence on the primordial bispectrum shape). \item When considering the cluster-galaxy cross spectrum alone, the forecasted constraints on $\sigma_8$ are comparable to the cluster-only constraints, while the constraints on $\sigma_8$ are improved by more than one order of magnitude, reaching a predicted error comparable with that coming from galaxies alone, $\Delta f_\mathrm{NL} \sim$ a few. This result highlights the high complementarity of the cluster power spectrum and the cluster-galaxy cross spectrum as cosmological probes. \item While the constraints on $\sigma_8$ coming from the galaxy power spectrum alone are almost unchanged when combined with the cluster-galaxy cross correlation, the constraints on $f_\mathrm{NL}$ are improved. This is true only to a slight level for the local and equilateral bispectrum shapes, however the error on $f_\mathrm{NL}$ can be reduced by a factor of $\sim 2$ for the enfolded and orthogonal cases. The addition of the cluster power spectrum carries only slight change to this conclusion. \item As expected the parameters $f_\mathrm{NL}$ and $\sigma_8$ are degenerate with respect to the power spectra of LSS tracers. This degeneracy is more marked for the cluster power spectrum, with correlation coefficients reaching up to $0.3-0.4$. The degeneracy on the other hand is almost insignificant for the galaxy power spectrum, in the sense that the constraining power weights much more on $\sigma_8$ than on $f_\mathrm{NL}$. The degeneracy with respect to the cluster-galaxy cross spectrum lies in between the former two, with the exception of the orthogonal bispectrum shape. \item We considered several survey configurations alternative to the fiducial \emph{Euclid}-like one, finding that a multi-slit spectroscopic instrument would allow more stringent constraints both on $f_\mathrm{NL}$ and $\sigma_8$. This improvement is due to a combination of the fact that with this other arrangement the selected galaxy population is different, and that the data analysis can be pushed down to a lower minimum redshift. This result is interesting with respect to future wide field survey concepts alternative to \emph{Euclid}, such as WFIRST. \end{itemize} The improvement in the constraints on both $f_\mathrm{NL}$ and $\sigma_8$ when combining the cluster and/or galaxy power spectra with the cluster-galaxy cross spectrum is reminescent of the fact that the latter is sensitive to different scales in a different way with respect to the former. The fact that this improvement is more important for the enfolded and orthogonal bispectrum shapes is extremely interesting, since these models are still relatively poorly studied compared to the equilateral and, almost ubiquitous, local shapes. Additionally, the amazing constraining power of the galaxy power spectrum, even when considered alone, stresses the importance of the spectroscopy part for future \emph{Euclid}-like missions when it comes to restrict cosmology on the basis of the spatial distribution of objects. The present work reinforces the conclusion that the spatial distribution of tracers of the LSS, especially galaxies, is an incredibly powerful probe for primordial non-Gaussianity, thanks to the very strong impact that the latter has on the large scale bias of dark matter halos \citep*{CA08.1,CA10.1}. The combination of cluster and/or galaxy power spectra with the cross spectrum of clusters and galaxies significantly improves the forecasted constraints, especially for the least studied non-Gaussian shapes. Merging all the information from future wide field surveys such as \emph{Euclid} and WFIRST promise to bring constraints on $f_\mathrm{NL}$ to the unity level, and constraints on $\sigma_8$ to the level of $\sim 10^{-4}$, in both cases superior to future CMB experiments \citep{SE09.1,VE09.1}.

1012.2305

We investigate the constraints on primordial non-Gaussianity with varied bispectrum shapes that can be derived from the power spectrum of galaxies and clusters of galaxies detected in future wide field optical/near-infrared surveys. Having in mind the proposed ESA space mission Euclid as a specific example, we combine the spatial distribution of spectroscopically selected galaxies with that of weak lensing selected clusters. We use the physically motivated halo model in order to represent the correlation function of arbitrary tracers of the large-scale structure in the Universe. As naively expected, we find that galaxies are much more effective in jointly constrain the level of primordial non-Gaussianity f<SUB>NL</SUB> and the amplitude of the matter power spectrum σ<SUB>8</SUB> than clusters of galaxies, due to the much lower abundance of the latter that is not adequately compensated by the larger effect on the power spectrum. Nevertheless, combination of the galaxy power spectrum with the cluster-galaxy cross-spectrum can decrease the error on the determination of f<SUB>NL</SUB> by up to a factor of ∼2. This decrement is particularly evident for the less studied non-Gaussian bispectrum shapes, the so-called enfolded and the orthogonal ones. Setting constraints on these models can shed new light on various aspects of the physics of the early Universe, and hence it is of extreme importance. By combining the power spectra of clusters and galaxies with the cluster-galaxy cross-spectrum we find constraints on primordial non-Gaussianity of the order Δf<SUB>NL</SUB>∼ a few, competitive and possibly superior to future cosmic microwave background experiments.

12225164

2011PhRvD..84d4011S

1012

1012.5239_arXiv.txt

The observations of type Ia supernovae interpreted within isotropic and homogeneous cosmological models imply the accelerated expansion of the Universe. These observations \cite{Riess:1998cb,Perlmutter:1998np} were made in late nineties and ignited a ``megabit bomb''\footnote{Stanis\l aw Lem, {\it Summa technologiae}, 1964.} of an enormous number of publications. Among many hypothesis that were proposed, the most conservative one assumed that inhomogeneities in the energy distribution may mimic the accelerated expansion. Although such a hypothesis does not solve the old cosmological constant problem (why this constant is so small), it gives a hope for understanding why the vacuum energy density in the concordance model is of the same order as the present matter energy density \cite{rasanen-2004-0402}. Inhomogeneities may have a twofold effect. Firstly, the averaging procedure in general relativity is not well understood yet. Hence, assuming homogeneity and then solving the Einstein equations could not lead to the proper metric \cite{0264-9381-4-6-025}. Secondly, the light propagates differently in inhomogeneous spacetimes. This may modify the luminosity distance -- redshift relation that is crucial for an interpretation of the type Ia supernovae data. In this article, we follow \cite{einstraus,1969ApJ...155...89K,2007PhRvD..76l3004M,btt,Biswas:2007gi,Vanderveld:2008vi,2009MNRAS.400.2185C,Valkenburg:2009iw} and study exact solutions to the Einstein equations, the so-called Swiss-Cheese (SC) models. In such models, inhomogeneities do not influence the global dynamics by construction. Therefore, the averaging problem will not be investigated here. The SC models provide convenient settings for studies of light propagation in inhomogeneous spacetimes. The SC models are constructed out of the Einstein-de Sitter (EdS) solution with spherical regions removed. The Lema\^itre-Tolman (LT) solutions are matched to the spacetime in the excised regions.\footnote{For another possibility see \cite{PhysRevD.82.103510}.} Since inhomogeneities in the real Universe are not spherically symmetric, it is not obvious how to choose density profiles of the inhomogeneous regions. Therefore, we treat the SC model as a toy model of the Universe and we search for a reasonable ``extremal'' setting to determine the maximal effect of inhomogeneities on the luminosity distance -- redshift curve. If shear is neglected, the upper bound on the luminosity distance for a given redshift is determined by the so-called {\it empty beam} formula \cite{1973ApJ...180L..31D}.\footnote{For light bundles which have not passed through a caustic \cite{1992grle.book.....S}.} The numerical code that we have developed give us large freedom in the construction of models. The light may travel non-radially trough arbitrary size inhomogeneous regions whose centres do not have to lie in a plane or on a regular lattice. This allows us to investigate more general settings than these presented so far in the literature. We solve numerically the fully relativistic system of equations.

In this article, the effect of inhomogeneities on light propagation was investigated in the framework of the SC models. We have examined the angular diameter distance -- redshift relation. This type of distance is related to the luminosity distance by the reciprocity theorem. Therefore, the theoretical angular diameter distance -- redshift relation is crucial for an interpretation of the type Ia supernovae data. Our analysis confirms that inhomogeneities may partially mimic the accelerated expansion of the Universe provided the light propagates through regions with lower than the average density. The effect is small and it becomes negligible if the average density along the beam does not deviate from the corresponding EdS value. In light of our work and the weak field gravitational lensing analysis \cite{Vanderveld:2008vi}, the result \cite{2007PhRvD..76l3004M} that suggest more significant influence of inhomogeneities is due to a peculiar setting of the underlying model. The precise size of the effect depends on the details of the model that was studied. Since the SC models are toy-models of the real Universe, it is speculative to base on them the final conclusions. Nevertheless, what our analysis shows is that, within the models we have studied, the effect of inhomogeneities remains too small\footnote{In the aligned SI model (our ``extremal'' setting with randomized photon trajectories), the effect of inhomogeneities on the angular diameter distance was one order smaller than what is needed to eliminate dark energy.} to explain the type Ia supernovae observations without dark energy. Our analysis of the fully relativistic and non-linear models did not reveal any stronger effect on the angular diameter distance than that predicted by the {\it partially filled beam} approximation. Randomization reduces the effect considerably in accordance with \cite{Vanderveld:2008vi}, \cite{2009MNRAS.400.2185C}, \cite{btt2} (and the others). We have directly evaluated the effect of shear on the angular diameter distance. Within our models, the effect of shear was negligible, but the models that do not take into account the rotation of the principal axes of shear (e.g.\ \cite{Valkenburg:2009iw}) may overestimate its effect around two orders of magnitude. In these models, the overestimated shear plays a minor role and it may lead to the small underestimation of the effect of inhomogeneities. Our results suggest that the size of inhomogeneities is not crucial for the angular diameter distance, provided that non-radial models are sufficiently randomized. We have found that the {\it partially filled beam} formula \eqref{eq:emb} gives good approximation to the angular diameter distance if the average density along the beam is much lower than the corresponding EdS value. Finally, we stress the analysis presented in this article does not touch directly the ``averaging problem'' in general relativity. The SC models behave on large scales by construction as the RW models, thus the influence of inhomogeneities on the global expansion rate of the Universe cannot be studied within this framework. \vspace{0.5cm} \noindent{\sc Acknowledgements} The main part of this research was conducted during my stay at the University of Geneva and was founded by the Foundation for Polish Science through its {\sc Kolumb} program. I would like to express my gratitude to Ruth Durrer for hospitality and useful discussions. I thank Krzysztof Bolejko, Piotr Chru\'sciel, Bart\l{}omiej Kos, Syksy R\"as\"anen for comments and Zdzis\l{}aw Golda, Andrzej Woszczyna for suggesting this topic to me. I acknowledge private communication with Christina Sormani and Wessel Valkenburg. The numerical calculations were carried out with the supercomputer ``Deszno'' purchased thanks to the financial support of the European Regional Development Fund in the framework of the Polish Innovation Economy Operational Program (contract no.\ POIG.02.01.00-12-023/08). I acknowledge the use of {\sc GNU Scientific Library} \cite{GSL} and {\sc Mathematica} together with the {\sc xAct} \cite{xAct} package. \clearpage \appendix

1012.5239

We study the effect of inhomogeneities on light propagation. The Sachs equations are solved numerically in the Swiss-cheese models with inhomogeneities modeled by the Lemaître-Tolman solutions. Our results imply that, within the models we study, inhomogeneities may partially mimic the accelerated expansion of the Universe provided the light propagates through regions with lower than the average density. The effect of inhomogeneities is small and full randomization of the photons’ trajectories reduces it to an insignificant level.

12171560

2011Ap&SS.336..257R

1012

1012.1306_arXiv.txt

Remnants of historical supernovae had been known since Lundmark's identification of the Crab Nebula with SN 1054 AD \citep{hubble28}. The first radio source to be identified as a previously unknown supernova remnant (SNR) was Cassiopeia A, proposed by Shklovskii in 1953, who also suggested that synchrotron radiation was the mechanism for producing radio emission \citep{shklovskii53}, based on the observed power-law spectrum $S_\nu \propto \nu^{-\alpha}$, with $\alpha \sim 0.8$ for Cas A. Minkowski in 1957 confirmed the identification from optical observations \citep{minkowski57}. Thus from the early days of radio astronomy, it was recognized that energetic particles were present in SNRs. Basic synchrotron physics tells us that emission at frequency $\nu$ is primarily by electrons with energy $E = 15 ( \nu({\rm GHz})/B(\mu{\rm G}) )^{1/2}$ GeV, so radio emission at frequencies of a few hundred MHz, typical at the time, immediately implied the presence of electrons with Lorentz factors of $10^3 - 10^4$. \begin{figure} \includegraphics[width=2.5truein]{greensnrs11-09histbw.eps} \caption{Spectral-index distribution of SNRs in Green's catalogue with adequately measured radio spectra, in bins of 0.1 in $\alpha.$ } \label{spix} \end{figure} Radio remains the spectral region in which SNRs are most consistently identified. Green's famous catalogue of Galactic SNRs (Green 2009; available online at http://www.mrao.cam.ac.uk/surveys/snrs/) lists 274 objects, essentially all with known radio properties. The mean spectral index is about 0.5, but with a significant spread of order 0.2 (Figure~\ref{spix}). This spread is a significant problem for theories of particle acceleration described below. Also significant is a tendency for the historical remnants (less than 2000 years old) to have steeper indices ($\alpha \gapprox 0.6$), a trend continued by radio supernovae, not represented here, which can have indices as steep as 0.9 -- 1.0 \citep{weiler09}. A few remnants such as Cas A have very well-sampled radio spectra (see references in Green 2009); Cas A's spectrum is well described by a single power-law with spectral index of 0.77 between 100 MHz and 100 GHz. However, most remnants are represented by only a few data points with substantial error bars. Other historical remnants have spectra with suggestions of concave-up curvature \citep{reynolds92}, naturally explained by efficient shock acceleration (see below). Magnetic field strengths are difficult to measure in SNRs; since the intensity of synchrotron radiation from a power-law distribution of electrons $N(E) = K E^{-s}$ electrons cm$^{-3}$ erg$^{-1}$ scales as $K B^{1+\alpha}$, radio synchrotron fluxes only give roughly the product of the energy densities in magnetic field and electrons. X-ray evidence described below indicates that magnetic fields are substantially amplified over typical interstellar values of a few microGauss, but the only avenue for estimating energetics from radio data is the use of minimum-energy equipartition arguments. These arguments indicate that the minimum energy in electrons and magnetic field in typical SNRs is far below the $\sim 10^{51}$ erg explosion energies; SNRs are not efficient synchrotron radiators. The equipartition magnetic fields so derived tend to be low, but as there is no strong physical argument that equipartition should hold, or even among which particles (should one include ions?), the equipartition magnetic field strength is really just a proxy for mean surface brightness. Magnetic-field orientations can be usefully derived from radio polarization directions. A uniform synchrotron source with spectral index $\alpha = 0.5$ has a polarized fraction of about 70\%, but very few remnants show values above 40\%. Young SNRs have much lower values, typically of order 10\% -- 15\% (see references in Reynolds \& Gilmore 1993), implying that their magnetic fields are primarily disordered. The ordered components, however, tend to be radial. In older remnants, magnetic-field orientations are typically confused, but it is fairly common to see a tangential orientation, which one would expect if a high-compression radiative shock compresses upstream magnetic field, increasing the tangential component by a factor of the compression ratio. Even though radio observations of SNRs do not require a large fraction of the SN energy, it is possible to argue that young SNRs require acceleration of new electrons -- simply borrowing and compressing relativistic electrons from the cosmic-ray population in the ISM is inadequate both because of high observed surface brightnesses and spectra much different from those of low-energy cosmic-ray electrons \citep{reynolds08a}. Radio observations of SNRs have left several unexplained puzzles, most several decades old now. What is responsible for spectral indices flatter than 0.5? (While a few pulsar-wind nebulae contaminate the low-$\alpha$ end of the distribution of Figure 1, most of the remnants with $\alpha < 0.5$ are shells.) In a few cases, contamination with flat-spectrum thermal emission may be responsible. Steeper spectra than 0.5 can be obtained with very low Mach-number shocks, but this requires ${\cal M} \lapprox 10$ and is unlikely to be the case for as many remnants as Figure~\ref{spix} requires. Finally, the question of the radial orientation of the ordered component of magnetic field in young remnants remains unexplained, though the operation of fluid instabilities at the contact interface between shocked ISM and shocked ejecta is often invoked (e.g., Jun, Jones, \& Norman 1996). \begin{figure} \includegraphics[width=2truein]{sn1006xhalfbw.eps} \includegraphics[width=2truein]{g1.9xsmonewbw.eps} \caption{{\sl Chandra} images of (top) SN 1006 (NASA/CfA) and (bottom) G1.9+0.3 \citep{reynolds08b} between 1.5 and 7 keV, smoothed with platelets as described in \cite{willett07}.} \label{g1.9} \end{figure}

1012.1306

It has been known for over 50 years that the radio emission from shell supernova remnants (SNRs) indicates the presence of electrons with energies in the GeV range emitting synchrotron radiation. The discovery of nonthermal X-ray emission from supernova remnants is now 30 years old, and its interpretation as the extension of the radio synchrotron spectrum requires electrons with energies of up to 100 TeV. SNRs are now detected at GeV and TeV photon energies as well. Strong suggestions of the presence of energetic ions exist, but conclusive evidence remains elusive. Several arguments suggest that magnetic fields in SNRs are amplified by orders of magnitude from their values in the ambient interstellar medium. Supernova remnants are thus an excellent laboratory in which to study processes taking place in very high Mach-number shocks. I review the observations of high-energy emission from SNRs, and the theoretical framework in which those observations are interpreted.

3941613

2011AdSpR..47.2118G

1012

1012.1901_arXiv.txt

\label{} Coronal mass ejections (CMEs) are spectacular ejections from the solar atmosphere of coronal material containing plasma threaded by magnetic fields, but despite over thirty years of study, the basic physics that expels these plasma clouds into the solar system is still not well understood \citep{kunow2006}. Precise measurements of CME properties, such as their rate of occurance, basic morphology and kinematics, first became possible with the launch of the Large Angle Coronagraph--Spectrograph \citep[LASCO;][]{brueckner1995} on SOHO in December, 1995. Further progress in understanding CMEs, and their three-dimensional properties in particular, was made possible with the Sun-Earth Connection Coronal \& Heliospheric Investigation \citep[SECCHI;][]{howard2008}, flown aboard NASA's recently launched Solar Terrestrial Relations Observatory (STEREO). The SECCHI instrument combines solar disk, coronagraph, and heliospheric observations from two distinct perspectives and is well suited to exploring the physics of CMEs, both at their source on the solar disk, and during their propagation to Earth \citep[e.g.][]{maloney2009}. CME kinematics have typically been determined manually using simple point-and-click methodologies such as those used in creating the LASCO CME Catalog \citep{gopal2009}. That is, using a mouse, the scientist clicks along a particular CME feature, such as the apex or front, in order to determine their positions in an image. The position of a given feature can then be plotted as a function of time to create height-time and velocity-time curves. Many authors have used point-and-click techniques to reconstruct the height-time profile of CMEs, and consequentially to derive CME velocities and accelerations projected onto the plane-of-sky \citep{gallagher2003,schrijver2008,temmer2009}. The draw-backs of this method are that it is slow, open to observer bias, and only allows a small number of points on the CME front to be tracked. In fact, the apex of the CME is the only feature usually tracked. This method is therefore not suitable for tracking the complex morphologies and kinematics of CMEs. Furthermore, they are certainly not practical for realtime operations, or the requirements of space weather monitoring and forecasting (e.g., at the NOAA Space Weather Prediction Center). We envisage that automated image processing techniques could be used to automatically identify and track CME leading-edges in images from coronagraphs such as SOHO/LASCO or STEREO/COR1/2. Image processing techniques enable us to automatically measure CME properties such as angular distribution about the occulting disk, the position angles over which it was launched, and its velocity up to approximately 30 solar radii. These CME properties are known to be associated with its probability of impacting the Earth, its arrival time at 1~AU, and how geoeffective it may be. For example, Halo CMEs and events with a large western angular extent are more likely to be directed towards the Earth, while fast CMEs are likely to be more geoeffective \citep{moon2005}. Due to the large quantity of data from LASCO and SECCHI, image processing techniques are essential for accurately identifying and characterising CME properties. \citet{robbrecht2004} developed a system that autonomously detects CMEs in image sequences from LASCO. Their software, Computer Aided CME Tracking (CACTus\footnote{http://sidc.oma.be/cactus/}), relies on the detection of bright ridges in CME height-time maps using the Hough transform. The main drawback of this method is that the Hough transform imposes a linear height-time evolution, thereby forcing constant velocity profiles for each bright feature. This method is therefore inappropriate for studying CME kinematics in the low corona, where non-constant acceleration may be at play \citep{byrne2009}. A complimentary system, the Solar Eruptive Event Detection System \citep[SEEDS;][]{olmedo2008}, is an automatic detection based on LASCO/C2 running difference images again unwrapped into polar coordinates. The algorithm uses a simple intensity threshold to segment the images and hence determines the CME's height, velocity and acceleration profiles. A major disadvantage of these systems is that neither report information on the morphological properties of CMEs or their predominant direction of propagation. This is a disadvantage for space weather purposes, as CME width and direction are known to be of importance to predicting geoeffectiveness at 1~AU \citep[e.g.,][]{michalek2008,kim2008}. A quite different image processing method was described by \citet{colaninno2006} to detect and track CMEs. Their algorithm, based on optical flow techniques, offers a number of attractive features, such as the ability to monitor the velocity field of a CME across its entire volume. There are limitations to the optical flow techniques, though, such as the assumption that the intensity of CME features do not change from frame-to-frame. Not only do CME features change as a CME erupts, but the brightness (from Thompson scattered photospheric photons) decreases systematically with distance from the Sun. This makes detection of CMEs using optical flow techniques challenging at large distances from the Sun. The algorithm has also not been implemented in a realtime manner for use by scientists and space weather forecasters. The reader is referred to \citet{robbrecht2005} for a comrehensive review of traditional CME detection techniques. CMEs are intrinsically multiscale features, making their detection using wavelets and other multiscale techniques an attractive proposition. \citet{stenborg2003} were the first to apply a wavelet-based technique to study the multiscale nature of coronal structures in SOHO images. Their method employed a multi-level decomposition scheme using the {\`a} trous wavelet transform. However, their technique only enhances coronal structures. It does not define, characterise or extract image features. This is a drawback for real-time applications or when attempting to study the detailed kinematics of multiple CME features. A number of authors have further explored multscale techniques to enhnace the visibility of a CME's front. \citet{young2008} used a derivative-of-a-Gaussian approximation of a wavelet to first decompose LASCO images into a variety of spatial scales. The gradient of each scale was then obtained and the CME front or leading-edge then isolated by identifying local maxima at each wavelet scale. An advantage of this method is that multiscale techniques can be used in conjunction with bootstrapping techniques to estimate the uncertainty in CME properties. This is particularly important for studies of CME kinematics, where the acceleration is estimated using inherently noisy numerical differencing schemes. \citet{byrne2009} extended these multiscale methods to take advantage of both the magnitude and angle of the gradient of the multiscale decomposition. Here, the CME front was found to have a well-defined signature in wavelet magnitude, and in angle. The resulting multiscale vector map as a function of time could then be used to design a multiscale spatio-temporal filter for CME front segmentation. \citet{byrne2009} showed that the results of the CME detection methods discussed above can introduce large errors in the kinematics of CMEs. They showed that for certain events, the results of CACTus, CDAW and SEEDS can differ significantly from multiscale methods. Existing on-line systems fit either a linear model to the height-time of the CME apex, implying constant velocity and zero acceleration (e.g. CACTus) or a second order polynomial, producing a linear velocity and constant acceleration (e.g. SEEDS). The multiscale decomposition discussed here and by \citet{byrne2009} minimises the uncertaintly in the CME height measurements; the resulting errors in velocity and acceleration are consequently only determined by the numerical errors associated with the differencing scheme used. Given that the estimated time of arrival of a CME at 1~AU can be approximated by \begin{equation} t_{1AU} = t_{Sun} + \int_{1R_{Sun}}^{1AU} dr/v(r), \end{equation} an accurate estimation of a CME's velocity profile in the low corona, $v(r)$, is essential to producing reliable forecasts of arrival times at Earth and other positions in the solar system. If autonomous CME tracking is to be used to more accurately predict CME arrival times at Earth, the velocity of the CME must be know to a high degree of accuracy; an uncertainty of only $\pm$10~km~s$^{-1}$ in CME velocity corresponds to an uncertaintly in the predicted arrival time at Earth of more than $\pm$3-hours. It should also be noted that steam interactions in the heliosphere can modify interplanetary CME velocities \citep[e.g.,][]{maloney2009}. An excellent overview of the physics of space weather and its effects are given in \citet{bothmerNdaglis2007}. In this paper a new set of multiscale transforms, namely wavelets, ridgelets, and curvelets, are used to identify the kinematic and morphological properties of CMEs to a high degree of accuracy. In Section 2, the wavelet, ridgelet and curvelet transforms are described, while their application to CME front enhancement and detection are discussed in Section 3. Our conclusions and prospects for future work are then given in Section 4.

Scientists and space weather forecasters primarily monitor solar activity using simple but robust image and signal processing techniques \citep[e.g., ][]{bothmerNdaglis2007}. Straightforward autonomous systems, such as SolarSoft Latest Events\footnote{http://www.lmsal.com/solarsoft/last\_events} and SolarMonitor\footnote{http://www.SolarMontitor.org} \citep{gallagher2002} have proven to be very popular for a number of reasons. They continuously deliver data, such as images, event movies, and flare times and positions, in a near-realtime manner and with a consitent level of accuracy. The latter is particularly important for large-scale statistical studies of, for example, active region or solar flare properties over an entire solar cycle. The data are thus not subject to observer bias or periods that were not well observed or characterised manually. For the purposes of operational space weather monitoring, such as at NOAA's Space Weather Prediction Center, solar data products must be continuously available and be produced in a self-consistent manner. Systems such as SolarMonitor have gone some way to achieving these goals. In this paper, a number of multiscale transforms have been evaluated in terms of their appropriateness for detecting the morphology and kinematics of CMEs. The wavelet-based technique was found to offer a fast and robust method for decomposing coronagraph images into a set of predefined length-scales. This was effective at removing noise and isolating particular CME features. Due to their isotropic symmetry and localization in space, wavelets are very sensitive to noise, cosmic ray hits and background stars. The wavelet transform can also be sensitive to small-scale symmetric features in CMEs, such as knot-like structures along a CME front, and so must be used with some care. An additional source of error can be introduced when deciding which wavelet scales to use when isolating the CME front, as the front width and visibility relative to the background varies with scale. In our work, we have selected scales that maximise the contrast of the front, although other criteria could equally well be allied. A particular difficulty with wavelets is their unsuitability for detecting the curvilinear structure of many CME features. One finding regarding the wavelet transform was that CME fronts can become disjointed features in scale-space. Ridgelets and curvelets avoid a number of these issues. The ridgelet transform in particular was found to be well suited to identifying the curved structures observed in CMEs. Not only can curvelets be used to enhance the contrast of CMEs, but the coefficiants of the transform give important morphological properties of a CME front, such as the positions of all points along the front, its curvature, its inclination, etc. As with all transforms, ridglets and curvelets also have their draw-backs. For example, these transforms require an expert user to manually select the particular filter properties that best match the shapes that one desires from the images. Furthermore, the application of transforms that are optimised to detect particular features (e.g., localised or curved shapes), may suppress the innate complexity of the emitting structures that we see in images of CMEs. These advanced image processing techniques are likely to be of use for real-time space weather operations, and in addition, to the analysis of large data-volumes from future ground- and space-based imagers. For space weather applications, advanced image processing techniques can be used to identify and track a plethora of features and events related to solar activity. For example, Lockheed's Latest Events service identifies flare locations and occurrence times using EUV and X-ray imaging data, and then automatically determines the source active region. For space weather applications, accurate flare occurrence times and locations is important for making forecasts of their possible impact at Earth. Of particular importance to forecasting space weather effects at Earth is identifying the properties of CMEs. The methods discussed in this paper are capable of measuring CME properties that are known to be related to adverse space weather at Earth. These CME properties include leading-edge position, angular width and launch position angle. See for example \citet{moon2005}, \citet{michalek2008}, and \citep{kim2008} for details of how these properties relate to space weather effects at Earth. While data rates from spacecraft such as SOHO are low enough to make human analysis of images feasible ($\leq$1~GB per day), missions such as the Solar Dynamics Observatoy (SDO), will have significantly highter data rates. SDO for example has projected data rates that make an interactive analysis impossible ($\sim$2~TB per day). The multiscale methods discussed in this paper naturally lend themselves to the automatic identification and characterisation of features related to solar activity, and therefore provide a stepping-stone to future autonomous monitoring of solar activity for space weather purposes.

1012.1901

Coronal mass ejections (CMEs) are large-scale eruptions of plasma and magnetic field that can produce adverse space weather at Earth and other locations in the Heliosphere. Due to the intrinsic multiscale nature of features in coronagraph images, wavelet and multiscale image processing techniques are well suited to enhancing the visibility of CMEs and suppressing noise. However, wavelets are better suited to identifying point-like features, such as noise or background stars, than to enhancing the visibility of the curved form of a typical CME front. Higher order multiscale techniques, such as ridgelets and curvelets, were therefore explored to characterise the morphology (width, curvature) and kinematics (position, velocity, acceleration) of CMEs. Curvelets in particular were found to be well suited to characterising CME properties in a self-consistent manner. Curvelets are thus likely to be of benefit to autonomous monitoring of CME properties for space weather applications.

12168249

2011ApJ...730...26F

1012

1012.4009_arXiv.txt

\label{sec:intro} Progress in our understanding of short-duration gamma-ray bursts (GRBs) and their progenitors relies on detailed studies of their afterglows and host galaxy environments. Of particular interest are bursts with precise sub-arcsecond positions, which can provide unambiguous host galaxy associations, redshifts, and burst properties (energy, density). Such localizations require the detection of ultraviolet, optical, near-infrared, and/or radio afterglows; or alternatively an X-ray detection with the {\it Chandra} X-ray Observatory. As of December 2010, only 20 short bursts have been precisely localized in this manner. Of these, 14 have clearly identified hosts\footnotemark\footnotetext{These bursts are 050709: \citep{ffp+05,hwf+05}; 050724: \citep{bpc+05}; 051221A: \citep{sbk+06}; 060121 \citep{ltf+06,pcg+06}; 060313: \citep{rvp+06}; 061006: \citep{amc+09}; 070707: \citep{pac+08}; 070714B: \citep{gfl+09}; 070724: \citep{bcf+09}; 071227: \citep{amc+09}; 080905: \citep{rwl+10}; 090426: \citep{aap+09,lbb+10}; 090510: \citep{mkr+10}; and 100117A: this paper.} (10 with spectroscopic redshifts), 5 do not have unambiguous host associations\footnotemark\footnotetext{GRBs 061201: \citep{sap+07}, 070809: \citep{gcn6739}, 080503: \citep{pmg+09}, 090305: \citep{gcn8934,gcn8933}, and 090515: \citep{rot+10}.} \citep{ber10a}, and 1 has not been reported in the literature so far\footnotemark\footnotetext{GRB 091109b: \citep{gcn10156}.}. For a recent summary see \citet{ber10b}. Only in a single case out of the 10 hosts with spectroscopic identifications is the galaxy known to be early-type with no evidence for on-going star formation activity (GRB\,050724: \citealt{bpc+05}); the remaining hosts are star forming galaxies, albeit at a level that is on average significantly lower than in long GRB hosts \citep{ber09}, particularly when accounting for their higher luminosities and stellar masses \citep{ber09,lb10}. The hosts with measured redshifts span a range of $z\approx 0.1-1$ (e.g., \citealt{ber09}), with the exception of GRB\,090426 at $z=2.609$ \citep{aap+09,lbb+10}; in the three remaining cases the hosts are too faint for a spectroscopic redshift determination, but are likely to be located at $z\gtrsim 1$ \citep{bfp+07}. At the same time, there is tentative evidence for early-type hosts in the sample of short bursts with optical positions and no coincident hosts based on chance coincidence probabilities (GRBs 070809 and 090515; \citealt{ber10a}). The host demographics and redshift distribution provide important constraints on the nature of the progenitors. For example, an abundance of low redshifts and early-type hosts would point to a population that is skewed to old ages, $\gtrsim {\rm few}$ Gyr \citep{zr07}. However, from the existing host galaxy demographics and redshift distributions it appears that the progenitors span a broad range of ages, $\sim 0.1-{\rm few}$ Gyr, and may indeed be over-represented in late-type galaxies with intermediate-age populations ($\sim 0.3$ Gyr; \citealt{bfp+07,lb10}). Afterglow detections are also important for determining the GRB and circumburst medium properties. To date, however, little detailed information about these properties has been drawn from the existing (though sparse) broad-band afterglow detections (e.g., \citealt{bpc+05,pan06,rvp+06,sbk+06,ber10a}), mainly due to the faintness of short GRB afterglows. Early time optical observations also probe possible emission from radioactive material synthesized and ejected in a binary compact object merger, a so-called Li-Paczynski mini-supernova \citep{lp98,mmd+10}. No such emission has been conclusively detected to date (e.g., \citealt{hsg+05,bpp+06,bcf+09}). Here we report the discovery of the optical afterglow and host galaxy of the short \grb. From spectroscopy and optical/near-IR imaging we find that the host is an early-type galaxy at $z=0.915$, making it only the second unambiguous early-type host association for a short GRB with a significantly higher redshift than the previous event, GRB 050724 at $z=0.257$ \citep{bpc+05}. The precise position also allows us to measure the burst offset, while the optical flux provides constraints on the circumburst density. We present the afterglow and host discovery in \S\ref{sec:obs}. In \S\ref{sec:host} we study the host redshift and stellar population properties, while in \S\ref{sec:ag} we analyze the afterglow properties. Finally, we draw conclusions about the nature of this burst and implications for the short GRB sample in \S\ref{sec:discussion}. Throughout the paper we use the standard cosmological parameters, $H_{0}=71$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{m}=0.27$, and $\Omega_{\Lambda}=0.73$.

\label{sec:discussion} In the sample of 14 short GRBs with optical afterglows and coincident hosts, \grb\ is only the second event unambiguously associated with an early-type galaxy (the other being GRB\,050724; \citealt{bpc+05}). Additional cases of early-type hosts have been proposed. In particular, GRB\,050509b is likely associated with an early-type cluster galaxy but this is based on only an X-ray position \citep{bpp+06}. Two additional bursts (070809 and 090515) lack coincident hosts despite optical afterglow detections, but in both cases the galaxies with the lowest probability of chance coincidence are early-type galaxies \citep{ber10a}. Even if we accept these additional early-type host associations as genuine, the host of \grb\ is located at a significantly higher redshift than the previous events, with $z\approx 0.23-0.47$. \grb\ also has the highest isotropic-equivalent gamma-ray energy of these events by a factor of a few, with $E_{\rm\gamma,iso}=2.1 \times 10^{50}$ erg ($15-150$ keV). These results suggest that some of the optically-faint host galaxies identified to date (e.g., \citealt{bfp+07}) may be bright near-IR sources due to a dominant old population. It also indicates that the presence of short GRBs in early-type galaxies does not necessarily point to progenitor ages of $\sim 10$ Gyr. Instead, the typical ages of short GRB progenitors in early-type hosts appear to be $\sim 1-4$ Gyr \citep{lb10}, which may lead to early-type hosts even at $z\approx 3$. At the inferred redshift of $z=0.915$ the projected physical offset of \grb\ is only $470\pm 310$ pc. Our previous analysis of short GRB offsets revealed a median projected offset of about 5 kpc \citep{fbf10}. In this context, \grb\ has the smallest offset measured to date. We note that the only other burst with a secure early-type host (GRB\,050724) also has a small offset of about $2.7$ kpc. Given the age of the stellar population of $\sim 1$ Gyr in both cases (see also \citealt{lb10}), these small offsets indicate that \grb\ and GRB\,050724 did not originate from progenitors with a substantial kick (unless the kick direction in both cases is nearly aligned with our line of sight). Given the lack of any recent star formation activity, we can also rule out the possibility of a highly kicked progenitor system with a short merger time. On the other hand, the proposed associations of short GRBs 050509b, 070809, and 090515 with early-type hosts at offsets of tens of kpc \citep{bpp+06,ber10a} indicates that some progenitors may experience large kicks. The cases of GRBs 070809 and 090515 is particularly intriguing since both had optical afterglows of comparable brightness to \grb\ (Figure~\ref{fig:optag}), suggestive of a similar circumburst density despite a potential large difference in offsets. Only a few short GRBs have circumburst density measurements, reflecting a general lack of multi-wavelength afterglow detections. GRB\,051221A had an estimated density of $n\sim 10^{-3}$ cm$^{-3}$ \citep{sbk+06}, GRB\,050724 had $n\approx 0.01-0.1$ cm$^{-3}$ \citep{bpc+05}, and GRB\,050709 had\footnotemark\footnotetext{Only an upper bound is available due to the lack of a radio detection.} $n\lesssim 0.1$ cm$^{-3}$ \citep{pan06}. For \grb\ we estimate $n\sim 10^{-4}-10$ cm$^{-3}$, continuing the trend of relatively low circumburst densities for short GRBs. This is particularly striking in comparison to the circumburst densities inferred for long GRBs, with a median of $\langle n\rangle\approx 1-10$ cm$^{-3}$ (e.g., \citealt{sbk+06}). Our discovery of the afterglow and $z=0.915$ early-type host of \grb\ continues to support the conclusion that short GRBs exist at $z\sim 1$ and beyond \citep{bfp+07}. However, unlike all previous short GRB hosts at these redshifts \citep{bfp+07,gfl+09,aap+09,lbb+10}, the host of \grb\ exhibits no evidence for star formation activity and is instead dominated by a $\sim 1$ Gyr old stellar population. With its faint optical afterglow it is possible that previous such events have been missed due to shallow optical afterglows searches, thereby potentially biasing the known host population against $z\gtrsim 1$ early-type hosts.

1012.4009

We present the discovery of the optical afterglow and early-type host galaxy of the short-duration GRB 100117A. The faint afterglow is detected 8.3 hr after the burst with r <SUB>AB</SUB> = 25.46 ± 0.20 mag. Follow-up optical and near-infrared observations uncover a coincident compact red galaxy, identified as an early-type galaxy at a spectroscopic redshift of z ≈ 0.915 with a mass of ~3 × 10<SUP>10</SUP> M <SUB>sun</SUB>, an age of ~1 Gyr, and a luminosity of L <SUB>B</SUB> ~= 0.5 L <SUB>*</SUB>. From a possible weak detection of [O II]λ3727 emission at z = 0.915 we infer an upper bound on the star formation rate of ~0.1 M <SUB>sun</SUB> yr<SUP>-1</SUP>, leading to a specific star formation rate of lsim0.004 Gyr<SUP>-1</SUP>. Thus, GRB 100117A is only the second short burst to date with a secure early-type host (the other being GRB 050724 at z = 0.257) and it has one of the highest short gamma-ray burst (GRB) redshifts. The offset between the host center and the burst position, 470 ± 310 pc, is the smallest to date. Combined with the old stellar population age, this indicates that the burst likely originated from a progenitor with no significant kick velocity. However, from the brightness of the optical afterglow we infer a relatively low density of n ≈ 3 × 10<SUP>-4</SUP> epsilon<SUP>-3</SUP> <SUB> e,-1</SUB>epsilon<SUP>-1.75</SUP> <SUB> B,-1</SUB> cm<SUP>-3</SUP>. The combination of an optically faint afterglow and host suggests that previous such events may have been missed, thereby potentially biasing the known short GRB host population against z &gt;~ 1 early-type hosts.

12224490

2011PhRvD..83f3515H

1012

1012.3904_arXiv.txt

There has been convincing evidence indicating that our universe is composed of nearly $25\%$ cold dark matter (DM) plus a small fraction of baryonic matter and around $70\%$ dark energy (DE)\cite{1}. One leading candidate of such a DE is the cosmological constant, representing a vacuum energy density with constant equation of state $w=-1$. However, it is difficult to understand such a cosmological constant in terms of fundamental physics. Its observed value is far below that estimated in quantum field theory, what is referred to as the cosmological constant problem. Moreover, using the cosmological constant to explain the DE, there is no natural understanding why the constant vacuum energy and matter energy densities are precisely of the same order today. This is the so-called coincidence problem. Considering that DE and DM are dominant sources of the content of the universe, it is natural, in the framework of field theory, to consider the inevitable interaction between them \cite{secondref}. An appropriate interaction between DE and DM can provide a mechanism to alleviate the coincidence problem \cite{10}-\cite{14}. A non-minimal coupling in dark sectors can affect significantly the expansion history of the universe and the density perturbation evolution, changing the growth history of cosmological structures. The possibility that DE and DM interact with each other has been widely discussed recently \cite{10}-\cite{AbdallaPLB09}. A number of studies have been devoted to analyze the constraints on the dark sectors mutual interaction from the probes of the cosmic expansion history by using the WMAP, SNIa, BAO and SDSS data etc \cite{71}-\cite{pp}. Interestingly it was disclosed that the late ISW effect has the unique ability to provide insight into the coupling between dark sectors \cite{hePRD09}. Furthermore, complementary probes of the coupling within the dark sectors have been carried out in the study of the growth of cosmic structure \cite{31}-\cite{AbdallaPLB09}. It was found that a non-zero interaction between dark sectors leaves a clear change in the growth index \cite{31,Caldera}. In addition, it was suggested that the dynamical equilibrium of collapsed structures such as clusters would acquire a modification due to the coupling between DE and DM \cite{pt,AbdallaPLB09}. Comparing the naive virial masses of a large sample of clusters with their masses estimated by X-ray and by weak lensing data, a small positive coupling has been tightly constrained \cite{AbdallaPLB09}, which agrees with the results given in \cite{hePRD09} from CMB. The small positive coupling indicates that there is energy transfer from DE to DM, which can help to alleviate the coincidence problem \cite{hePRD09,heJCAP08}. Both DE and DM are currently only detected via their gravitational effects and any change in the DE density is conventionally attributed to its equation of state $w$. This status leads to an inevitable degeneracy while extracting the signature of the interaction between dark sectors and other cosmological parameters. In this work, we will first discuss the degeneracy between the DE and DM coupling and the equation of state (EoS) of DE in the background dynamics. Furthermore, we will extend our discussion to the perturbed spacetime by considering the perturbation evolution of DE and DM. We review the formalism of the perturbation theory when there is an interaction between dark sectors. Based upon this formalism we explore the possibility of breaking the degeneracy between the coupling and other cosmological parameters, such as the EoS parameter $w$ of DE as well as the DM abundance. This can help us extract a tighter constraint on the interaction between dark sectors from observations.

In this paper we have reviewed the formalism of the perturbation theory when there is an interaction between DE and DM. We have proposed a way to construct the coupling vector in a self consistent manner both in the perturbed form and in the background. Based upon the perturbation formalism we have studied the signature of the interaction between dark sectors from CMB angular power spectrum. Theoretically we found that there are possible ways to break the degeneracy between the interaction, DE EoS and DM abundance. This can help to get tight constraint on the interaction between DE and DM. We have performed the global fitting by using the CMB power spectrum data from WMAP7Y results together with latest SNIa, BAO and $\rm{H_0}$ data to constrain the interaction between DE and DM. When the interaction between DE and DM takes the form proportional to the energy density of DM and the total dark sectors, in $1\sigma$ range the coupling is found to be positive. The tight positive coupling indicates that there is energy flow from DE to DM, which can help us to alleviate the cosmological coincidence problem. The question of how to improve the model is now much related to find a field theory based model for the interaction and how to relate the model to the standard model of particle interactions. This is currently under study. \textbf{Acknowledgement:} This work has been supported partially by NNSF of China No. 10878001 and the National Basic Research Program of China under grant 2010CB833000. EA wishes to thank FAPESP and CNPq (Brazil) for financial support.

1012.3904

Cosmological analyses based on currently available observations are unable to rule out a sizeable coupling between dark energy and dark matter. However, the signature of the coupling is not easy to grasp, since the coupling is degenerate with other cosmological parameters, such as the dark energy equation of state and the dark matter abundance. We discuss possible ways to break such degeneracy. Based on the perturbation formalism, we carry out the global fitting by using the latest observational data and get a tight constraint on the interaction between dark sectors. We find that the appropriate interaction can alleviate the coincidence problem.

1048694

2011MNRAS.413.1479S

1012

1012.3073_arXiv.txt

Supermassive black holes (BHs) are believed to dwell in almost all galaxy spheroids (hereafter bulges) as a result of past quasar activity (Soltan~1982, Marconi et al.~2004, Shankar~2009 and references therein). A direct link between bulge formation and the growth of the central BHs has been inferred from tight relations between BH mass (\mbh) and the bulge structural parameters, such as velocity dispersion ($\sigma$, Ferrarese and Merrit~2000, Gebhardt et al.~2000), luminosity (\lbul, Kormendy \& Richstone~1995, Marconi \& Hunt~2003 - hereafter, MH03) and mass (Magorrian et al.~1998, MH03, H{\"a}ring \& Rix~2004, hereafter HR04).\\ The underlying physics is encoded not only in the slopes of the correlations, but also in the intrinsic or cosmic scatter ($rms$), i.e.~the scatter not accounted for by measurement uncertainties. Particular care has been taken with its measurement and, for instance, over the last 15 years the intrinsic scatter of \mbh- \lbul\ decreased from an initial value of $rms\sim0.5$~dex (in log~\mbh, Kormendy \& Richstone 1995) to $\sim 0.3-0.4$~dex (MH03, G\"ultekin et al.~2009, Hu~2009 hereafter G09 and H09 respectively). \mbh-$\sigma$ is considered one of the tightest relations ($rms\sim0.25-0.3$~dex, Tremaine et al.~2002, T02), suggesting that bulge dynamics or mass rather than luminosity are driving the correlations. Indeed the \mbh-\mdyn\ relation has a scatter comparable with \mbh-$\sigma$ ($\sim0.3$~dex, MH03, HR04, H09 and references therein). Nonetheless, slopes and $rms$ of scaling relations are still a subject of debate; an accurate assessment of the actual $rms$ is indeed crucial to draw conclusions on the physical origin of the correlations and turns out to be fundamental to both constrain BH-galaxy evolutionary models and compute the space density of BHs (e.g.~G09, Hopkins et al.~2007). Since a successful theoretical model should be able to account for \emph{all} empirical BH-bulge relations, discerning which are the truly fundamental ones, there has also been a growing interest in searching for all possible connections of the BH with other structural parameters. These include the effective radius \re\ (MH03) and its various combinations with $\sigma$ such as virial mass, gravitational binding energy of the host spheroid (e.g.~MH03, Aller \& Richstone 2007, H09, Feoli et al.~2010), thus leading to a possible fundamental plane for BHs (Barway \& Kembhavi~2007, Hopkins et al.~2007). To investigate the interplay between \mbh, bulge dynamics and luminosity, it is crucial (\emph{i}) to perform a proper photometric decomposition of the galaxy components (e.g.~disk, bulge, bar etc.) and (\emph{ii}) to obtain good estimators of stellar mass (\mstar), dynamical mass (\mdyn) and the galaxy mass-to-light ratio (M/L). An accurate bulge-disk decomposition can be performed by simultaneously fitting multiple bi-dimensional components to galaxy images. To compute the bulge mass, one can apply the virial theorem which turns out to be quite accurate (Cappellari et al.~2006) but requires a spectroscopical estimate of the velocity dispersion in addition to the bulge effective radius \re. Alternatively, one can estimate \mstar\ by combining \lbul\ with a M/L ratio based on calibrated relations available in the literature (see, e.g., Bell \& de Jong~ 2001, Bell et al.~2003, Cappellari et al.~2006). In this case, \mstar\ can be affected by dust extinction, depending on wavebands, and by the assumption of a constant M/L over the entire galaxy. In any case, it is believed that \mstar\ is more easily estimated than \mdyn\ especially at higher redshifts where it is the only host property not affected by strong biases (Merloni et al.~2010, Trakhtenbrot \& Netzer~2010, Lamastra et al.~2010). Recently, attention has been drawn to the coexistence of pseudobulges and BHs (Graham 2008b, Greene, Ho \& Barth~2008, H09, Nowak et al.~2010, Erwin~2010). Pseudobulges are bulges which are photometrically and morphologically disk-like (possibly containing typical disk features such as bars, rings or ovals) and which present 'cool' kinematics, dominated by rotation (Kormendy \& Kennicutt~2004). Apparently, pseudobulges follow their own relation with BHs, hosting less massive BHs than classical bulges (Greene, Ho, \& Barth~2008, H09). Indeed, in these kind of structures, \mbh\ seems to better correlate with the small classical bulge component only (Nowak et al.~2010, Erwin~2010). However, these results are only based on a few pseudobulges ($\sim5$) which have been analyzed as a separate class, and the poor statistics preclude a firm conclusion. In this paper we present a mid-infrared (MIR) view of the \mbh-bulge scaling relations based on a 2D photometric decomposition of \spitzer/IRAC images at \wav. The superb performance of \spitzer/IRAC (Fazio et al.~2004) provides images of unprecedented quality in the MIR for the 57 galaxies analyzed here, and IRAC sensitivity permits the clear identification of morphological features. Therefore we identify classical bulges and pseudobulges with our analysis and compare their properties with \mbh\ for one of the largest samples ever considered. \wav\ observations are an ideal tracer of \mstar, and are less affected by extinction than shorter wavelengths. Our aims are to:\\ - determine the bulge structural parameters with high accuracy;\\ - investigate the \mbh-bulge scaling relations and their intrinsic scatter, taking into account the nature of the bulge (classical vs pseudobulges);\\ - calibrate the M/L ratio thus supplying the reference \mbh-\mstar\ in the local Universe for studies of the \mbh-galaxy scaling relations at higher $z$, based on \mstar\ derived from luminosity measurements (by, e.g., Merloni et al.~2010, Trakhtenbrot \& Netzer~2010). In Section~2 we describe sample selection, data reduction and the grid-method adopted to decompose the IRAC \wav\ images. We then present the fundamental plane (FP) for spheroids defined by the objects in our sample. \mbh-bulge scaling relations and the linear regression methods adopted to estimate slopes and intrinsic dispersion of the correlations are described in Section~3. Section~4 is dedicated to a discussion about (a) the utility of \mbh-\mdyn\ as a benchmark to probe BH vs host galaxy co-evolution, and (b) pseudobulges location in the scaling relations. Throughout this paper we assume the standard cosmology with H$_0=70$~km~s$^{-1}$~Mpc$^{-1}$, $\Omega_M=0.3$, $\Omega_\Lambda=0.7$. \begin{table*}{OBSERVATIONS} \begin{center} \centerline{\begin{tabular}{lccccc} \hline Source Name & PID & Date & Frame Time & Frames & Position\\ (1) & (2) & (3) & (4) & (5) & (6) \\ \hline Circinus & 40936& 2007-09-09 & 12 & 1 & 4 \\ IC1459 & 20371& 2005-11-28 & 2 & 1 & 22 \\ IC2560 & 40936& 2007-07-05 & 12 & 1 & 4 \\ IC4296 & 20371& 2006-02-13 & 2 & 1 & 22 \\ NGC221 & 00069& 2004-07-19 & 12 & 1 & 5 \\ NGC524 & 50630& 2008-09-18 & 30 & 1 & 5 \\ NGC821 & 20371& 2005-08-21 & 2 & 1 & 22 \\ NGC1023 & 00069& 2004-02-11 & 30 & 1 & 5 \\ NGC1068 & 00032& 2004-01-16 & 12 & 2 & 5 \\ NGC1300 & 61065& 2009-09-06 & 30 & 1 & 11 \\ NGC1316 & 00159& 2004-02-23 & 30 & 1 & 9 \\ NGC2549 & 50630& 2008-12-23 & 30 & 1 & 5 \\ NGC2748 & 61063& 2009-12-02 & 30 & 1 & 4 \\ NGC2778 & 30318& 2006-11-24 & 100 & 1 & 5 \\ NGC2787 & 03674& 2004-10-30 & 30 & 1 & 3 \\ NGC2974 & 30318& 2006-12-28 & 100 & 1 & 5 \\ NGC3031 & 00159& 2004-05-01 & 30 & 1 & 12 \\ NGC3079 & 00059& 2004-04-05 & 12 & 1 & 4 \\ NGC3115 & 00069& 2004-04-29 & 12 & 1 & 4 \\ NGC3227 & 03269& 2004-12-21 & 12 & 1 & 1 \\ NGC3245 & 03674& 2004-11-28 & 30 & 1 & 3 \\ NGC3368 & 00069& 2004-05-19 & 30 & 1 & 5 \\ NGC3377 & 00069& 2004-05-27 & 30 & 1 & 5 \\ NGC3379 & 00069& 2004-12-15 & 12 & 1 & 5 \\ NGC3384 & 30318& 2006-12-27 & 100 & 1 & 5 \\ NGC3414 & 50630& 2009-01-29 & 30 & 1 & 5 \\ NGC3489 & 00069& 2004-05-19 & 12 & 1 & 5 \\ NGC3585 & 30318& 2006-07-09 & 100 & 1 & 5 \\ NGC3607 & 00069& 2004-05-19 & 12 & 1 & 5 \\ NGC3608 & 30218& 2006-12-27 & 100 &1 & 5 \\ NGC3998 & 00069& 2004-04-21 & 12 & 1 & 5 \\ NGC4026 & 30318& 2006-12-26 & 100 &1 &5 \\ NGC4151 & 03269& 2004-12-17 & 12 & 1 & 5 \\ NGC4258 & 20801& 2005-12-25 & 30 & 2 & 7 \\ NGC4261 & 00069& 2004-05-27 & 12 & 1 & 5 \\ NGC4374 & 00069& 2004-05-27 & 12 & 1 & 5 \\ NGC4459 & 03649& 2005-01-22 & 12 & 1 &5 \\ NGC4473 & 03649& 2005-01-22 & 12 & 1 & 5 \\ M87$^a$ & 03228 & 2005-06-11 & 30 & 1 & 5 \\ & 03228 & 2005-06-11 & 30 & 1 & 5 \\ NGC4486A & 03228 & 2005-06-11 & 30 & 1 & 5 \\ NGC4552 & 00159 & 2004-05-27 & 30 & 1 & 7 \\ NGC4564 & 20371 & 2006-02-09 & 2 & 1 & 22 \\ NGC4594 & 00159 & 2004-06-10 & 30 & 1 & 6 \\ NGC4596 & 03674 & 2005-06-10 & 30 & 1 & 3 \\ NGC4621 & 03649 & 2005-06-10 & 12 & 1 & 5 \\ NGC4649 & 00069 & 2004-06-10 & 12 & 1 & 5 \\ NGC4697 & 03403 & 2005-06-17 & 30 & 1 & 5 \\ NGC5077 & 00069 & 2005-05-11 & 12 & 1 & 5 \\ CenA & 00101 & 2004-02-11 & 12 & 1 & 5 \\ NGC5576 & 03403 & 2005-07-15 & 30 & 1 & 5 \\ NGC5813 & 00069 & 2004-02-17 & 12 & 1 & 5 \\ NGC5845 & 20371 & 2005-08-23 & 2 & 1 & 22 \\ NGC5846 & 00069 & 2004-03-09 & 12 & 1 & 5 \\ NGC6251 & 02418 & 2004-12-16 & 30 & 2 & 6 \\ NGC7052 & 30877 & 2006-11-26 & 30 & 1 & 5 \\ NGC7457 & 30318 & 2006-07-12 & 100 & 1 & 5 \\ NGC7582 & 03269 & 2004-11-27 & 12 & 1 & 1 \\ \hline \end{tabular}} \end{center} \caption{Observations Log. Columns: (1) Source name. $^a$ mean of 2 observations (see Section~2.2 for details). (2) Proposal identification number (PID). (3) Observing date. (4) Exposure time for each frame. (5) Numbers of frames. (6) Number of positions to realize the source maps.} \label{tb:obs} \end{table*}

In this work we have investigated the scaling relations observed in the local Universe between \mbh\ and the structural parameters of the host bulges. The analysis is based on a 2D decomposition of \wav\ Spitzer/IRAC images of 57 early- and late- type galaxies with \mbh\ measurements. Given the well known degeneracy between the bulge S\'{e}rsic index $n$ and effective radius \re, we have adopted a grid-method to fit IRAC images and determine $n$ rather than let it freely vary. The accuracy of our analysis is verified by the agreement of our bulges with the mid-infrared fundamental plane determination by JI08. Our galaxies extend their \wav\ FP by doubling the \re\ range towards smaller radii and none of our galaxies deviate significantly from the JI08 FP. This allows us to reliably (\emph{i}) calibrate M/L at \wav, (\emph{ii}) study the \mbh- scaling relation and (\emph{iii}) identify pseudobulges in our sample. We obtained a tight ($rms=0.10$~dex) \mdyn-\lbs\ relation: \[ \log\mdyn/\msun= 11.04+1.18\times[\log(\lbs/L_{3.6,\odot})-11], \] which allows us to estimate stellar masses based on luminosities at \wav\ with high accuracy. The \wav\ luminosity appears to be the best tracer of \mstar\ yet found. The relations between \mbh, luminosity, masses, and effective radius, fitted with a Bayesian approach to linear regression are: \[\log\mbh/\msun=8.19+0.93\times[\log(\lbs/L_{3.6,\odot})-11],\] \[\log\mbh/\msun=8.20+0.79\times[\log(\mdyn/\msun)-11],\] \[\log\mbh/\msun=8.16+0.79\times[\log(\mstar/\msun)-11],\] \[\log\mbh/\msun=8.22+0.9\times[\log(\re/kpc)-0.4],\] \emph{all} with the same intrinsic dispersion of $rms\sim0.35$~dex except for \mbh-\re\ which has $rms\sim0.5$~dex. Our \mbh-\mstar\ relation can be used as the local reference for high redshift studies which probe the cosmic evolution of \mbh-galaxy scaling scaling relations. These \wav\ \mbh-\lbs, \mbh-\mdyn, \mbh-\mstar\ relations turn out to be as tight as \mbh-$\sigma$ and, moreover, they are consistent with previous determinations from the literature at shorter wavelengths. We independently identified as pseudobulges those galaxies with S\'{e}rsic index lower than 2 and found 9 sources that satisfy this criterion. Of these, 4 pseudobulges lie on scaling relations within the observed scatter, while those with \mbh\ lower than $10^7$~\msun\ are significantly displaced. We discussed the different physical and evolutionary origins for such behavior, while considering the presence of nuclear morphological components not reproduced by our two-dimensional decomposition. Finally, we verified the existence of a possible FP for supermassive BHs, relating \mbh\ with two or more bulge properties. We did not find any correlation between the residuals of \mbh-$\sigma$ and the effective radius, showing that our data do not require the existence of any BH fundamental plane.

1012.3073

We present a mid-infrared investigation of the scaling relations between supermassive black hole masses (M<SUB>BH</SUB>) and the structural parameters of the host spheroids in local galaxies. This work is based on 2D bulge-disc decompositions of Spitzer/IRAC 3.6 μm images of 57 galaxies with M<SUB>BH</SUB> estimates. We first verify the accuracy of our decomposition by examining the Fundamental Plane (FP) of spheroids at 3.6 μm. Our estimates of effective radii (R<SUB>e</SUB>) and average surface brightnesses, combined with velocity dispersions from the literature, define a FP relation consistent with previous determinations but doubling the observed range in R<SUB>e</SUB>. None of our galaxies is an outlier of the FP, demonstrating the accuracy of our bulge-disc decomposition which also allows us to independently identify pseudo-bulges in our sample. We calibrate M/L at 3.6 μm by using the tight M<SUB>dyn</SUB>-L<SUB>bul</SUB> relation (∼0.1 dex intrinsic dispersion) and find that no colour corrections are required to estimate the stellar mass. The 3.6 μm luminosity is thus the best tracer of stellar mass yet studied. We then explore the connection between M<SUB>BH</SUB> and bulge structural parameters (luminosity, mass, effective radius). We find tight correlations of M<SUB>BH</SUB> with both 3.6 μm bulge luminosity and dynamical mass (M<SUB>BH</SUB>/M<SUB>dyn</SUB>∼ 1/1000), with intrinsic dispersions of ∼0.35 dex, similar to the M<SUB>BH</SUB>-σ relation. Our results are consistent with previous determinations at shorter wavelengths. By using our calibrated M/L, we rescale M<SUB>BH</SUB>-L<SUB>bul</SUB> to obtain the M<SUB>BH</SUB>-M<SUB>★</SUB> relation, which can be used as the local reference for high-z studies which probe the cosmic evolution of M<SUB>BH</SUB>-galaxy relations and where the stellar mass is inferred directly from luminosity measurements. The analysis of pseudo-bulges shows that four out of nine lie on the scaling relations within the observed scatter, while those with small M<SUB>BH</SUB> are significantly displaced. We explore the different origins for such behaviour while considering the possibility of nuclear morphological components not reproduced by our 2D decomposition.

12168932

2011ApJ...734...11K

1012

1012.4465_arXiv.txt

\label{Intro} Interacting galaxies in the nearby universe are laboratories for direct studies of several physical processes that were important in the formation and early evolution of galaxies. From such studies, we hope to learn, for example, how interactions and mergers affect the cycle in which baryonic matter is converted from diffuse interstellar gas into dense molecular clouds, then into star clusters, and eventually, by disruption, into a relatively smooth stellar field. It is clear that interactions and mergers boost the rate of star and cluster formation. But do they also change the rate at which clusters are disrupted? This is the question we address in this paper. The most intensively studied interacting galaxies are the \object{Antennae} (\object{NGC 4038}/39), at a distance $D \sim 20 \Mpc$ \citep{SchweizerEtAl2008AJ}. They consist of two normal disk galaxies that began to collide a few $\times 10^8 \yr$ ago. The stellar population and interstellar medium (ISM) of the Antennae have been observed over an enormous range of wavelengths, from X-ray to radio \citep[see, e.g.][and references therein]{ZhangFallWhitmore2001ApJ, HibbardEtAl2001AJ, KassinEtAl2003AJ, ZezasEtAl2006ApJS..166..211Z, 2009ApJ...699.1982B, 2010A&A...518L..44K}. The star clusters have been the focus of numerous studies based on observations with the {\it Hubble Space Telescope (HST)}, culminating in well-determined luminosity, mass, age, and space distributions \citep[see][and references therein]{2010AJ....140...75W}. The Antennae have also been the focus of several dynamical simulations, first by \citet[][hereafter TT72]{Toomre&Toomre1972ApJ} and then by \citet{Barnes1988ApJ}. These pioneering studies demonstrated that gravity alone can account for the gross features of the observed morphology and kinematics of the stellar components of the merger. Subsequent simulations have included an interstellar medium and star formation, with the additional goal of matching the observed space distribution of young stars in the Antennae \citep[][hereafter MBR93, TCB10, and K10, respectively]{MihosBothunRichstone1993ApJ, 2010ApJ...720L.149T, 2010ApJ...715L..88K}. There have also been two recent attempts to match the observed age distribution of the clusters, with different assumptions about their disruption histories \citep[][K10]{BastianEtAl2009ApJ...701..607B}. The purpose of this paper is to reexamine the issues raised by the observed age distribution of the clusters in the Antennae. In Section \ref{CFH}, we review the evidence for a quasi-universal age distribution of star clusters in different galaxies, and in Section \ref{simulations}, we assemble the star formation histories from all available {\it N}-body+hydrodynamical simulations of the Antennae. We then combine these and compare the results with observations in Section \ref{results}. We summarize our conclusions and their implications in Section \ref{discussion}. In particular, we show in this paper that there is nothing special about the disruption history of clusters in the interacting Antennae galaxies; it is similar to that in quiescent (non-interacting) galaxies. Before proceeding, we offer a few remarks about nomenclature. We use the term ``cluster'' for any concentrated aggregate of stars with a density much higher than that of the surrounding stellar field, whether or not it also contains gas and whether or not it is gravitationally bound. This is the standard definition in the star formation community \citep[see, e.g.][]{2003ARA&A..41...57L, 2007ARA&A..45..565M}. Some authors use the term ``cluster'' only for gas-free or gravitationally bound objects. Such definitions are not appropriate in the present context for two reasons: (1) A key element in our analysis is the connection between the formation rates of stars and clusters. We would break this connection artificially if we were to exclude the gas-rich clusters in which stars form. (2) It is virtually impossible to tell from observations which clusters satisfy the virial theorem precisely and which do not, especially at the distance of the Antennae. Indeed, {\it N}-body simulations show that an unbound cluster retains the appearance of a bound cluster for remarkably long times, more than 10 crossing times \citep{2007MNRAS.380.1589B}.

\label{discussion} The main conclusion of this paper is that the star clusters in the interacting Antennae galaxies are disrupted in much the same way as those in other galaxies. In most if not all star-forming galaxies, including the Antennae, the observed age distribution of clusters can be approximated by a power law, $dN/d\tau \propto \tau^{\gamma}$, with an exponent in the range $-1.0 \la \gamma \la -0.7$ (see \citealp{2010ApJ...713.1343C} and Section \ref{CFH} here). In general, this must reflect the combined formation and disruption histories of the clusters. However, variations in the formation rate are expected to be a minor influence because $dN/d\tau$ is so similar in different galaxies and declines by such a large factor, typically $\sim 10^2$, over a relatively small range of age, $\tau \la 10^8$--$10^9$~yr (i.e., less than 10\% of the lifetime of the galaxies). Indeed, in several well-studied galaxies (the Milky Way and the Magellanic Clouds), the SFR is known from observations to have been nearly constant for the past $\sim 10^9$~yr, compelling evidence that the decline in $dN/d\tau$ is mainly a consequence of disruption \citep{1998gaas.book.....B, 2004AJ....127.1531H, 2009AJ....138.1243H, 2010ApJ...711.1263C}. The interpretation is less straightforward for the Antennae galaxies, since they are too far away to determine their star formation history from observations. Furthermore, it is natural to wonder whether the interaction could trigger enough recent formation to explain the shape of $dN/d\tau$ without disruption. \citet{FallChandarWhitmore2005ApJ} argued that the formation rate would vary by factors of a few on the orbital timescale ($\sim 10^8$~yr), too gradually to account for most of the decline in $dN/d\tau$ (see also \citealp{WhitmoreChandarFall2007AJ} and \citealp{2009ApJ...704..453F}). The results presented here support this suggestion. We have assembled the star formation histories in all the available $N$-body$+$hydrodynamical simulations of the Antennae. These are based on different numerical methods, different orbits, and different prescriptions for star formation and stellar feedback. The treatment of small-scale processes is still approximate at best, due to the low resolution in the simulations compared to molecular clouds and clumps in the real ISM. The star formation rates differ greatly among the simulations, both in absolute level and in time dependence. Nevertheless, we find that they {\it all} vary slowly, by factors of only $1.3 - 2.5$ in the past $10^8$~yr. When we combine the formation rates in the simulations with a power-law model for the survival fraction of clusters, $f_{\rm surv} \propto \tau^\delta$, we find good agreement with the observed age distribution over the range $10^6$~yr $\lesssim \tau \lesssim 10^9$~yr for $-0.9 \lesssim \delta \lesssim -0.6$. The similarity between $\delta$ and $\gamma$ indicates that $dN/d\tau$ is shaped mainly by the disruption of clusters rather than variations in their formation rate. The only caveat to this conclusion stems from our assumption that the formation rates of clusters and stars are proportional to each other, i.e. $(dN/d\tau)_{\rm form} = c \cdot dN_*/d\tau$ with $c = {\rm constant}$ (cf. Equation (\ref{eq:S-CFH})). This is certainly true if most stars form in clusters, a hypothesis consistent with $\mathrm{H}{\alpha}$ observations of the Antennae \citep{FallChandarWhitmore2005ApJ}. However, even if we were to abandon this assumption entirely, and allow $c$ to vary arbitrarily with $\tau$, a slightly modified version of the analysis presented above would then lead to the alternative result $c(\tau) \cdot f_{\rm surv}(\tau) \propto \tau^{\delta}$ with $-0.9 \lesssim \delta \lesssim -0.6$. We would then be left with the problem of explaining why the product $c \cdot f_{\rm surv}$ but not $f_{\rm surv}$ in the Antennae happens to be the same as $f_{\rm surv}$ alone in other galaxies. The simplest interpretation --- the one advocated here --- is that the formation rates of clusters and stars {\it do} track each other and that the survival fractions and hence the disruption histories {\it are} similar in different galaxies, including the Antennae.

1012.4465

We re-examine the age distribution of star clusters in the Antennae in the context of N-body+hydrodynamical simulations of these interacting galaxies. All of the simulations that account for the observed morphology and other properties of the Antennae have star formation rates that vary relatively slowly with time, by factors of only 1.3-2.5 in the past 10<SUP>8</SUP> yr. In contrast, the observed age distribution of the clusters declines approximately as a power law, dN/dτvpropτ<SUP>γ</SUP> with γ = -1.0, for ages 10<SUP>6</SUP> yr &lt;~ τ &lt;~ 10<SUP>9</SUP> yr. These two facts can only be reconciled if the clusters are disrupted progressively for at least ~10<SUP>8</SUP> yr and possibly ~10<SUP>9</SUP> yr. When we combine the simulated formation rates with a power-law model, f <SUB>surv</SUB>vpropτ<SUP>δ</SUP>, for the fraction of clusters that survive to each age τ, we match the observed age distribution with exponents in the range -0.9 &lt;~ δ &lt;~ -0.6 (with a slightly different δ for each simulation). The similarity between δ and γ indicates that dN/dτ is shaped mainly by the disruption of clusters rather than variations in their formation rate. Thus, the situation in the interacting Antennae resembles that in relatively quiescent galaxies such as the Milky Way and the Magellanic Clouds.

12225049

2011PhRvD..83j3503C

1012

1012.5596_arXiv.txt

} \renewcommand{\thesection}{\arabic{section}} Probing the large-scale geometry of the Universe at cosmological scales is, undoubtedly, one of the most outstanding issues in modern cosmology. The standard assumptions of homogeneity and isotropy (namely, the Cosmological Principle~\cite{Weinberg}) can now be tested via new and very accurate data coming from the study of cosmic microwave background (CMB) radiation, especially from the Wilkinson Microwave Anisotropy Probe (WMAP)~\cite{WMAP7}, and data from type Ia supernovae, such as those collected in the so-called Union~\cite{Kowalski} and Union2~\cite{Union2} compilations. Indeed, concerning the tests of isotropy, an anisotropic model of Universe, known as ``ellipsoidal universe''~\cite{prl,prd,prd2,ADE,parallax,Cea1,Cea2}, may be even favored by an observed anomalous feature of the CMB power spectrum --the lack of power on large angular scales-- while being consistent with other cosmological data. It is not excluded that such an anisotropic model of the Universe could even account for three other large-scale ``anomalies'' of the isotropic standard cosmological model (for a brief but pointed discussion see Ref.~\cite{Antoniou}): the detection of large-scale velocity flows significatively larger than those predicted in standard cosmology~\cite{A1}, a statistically significant alignment and planarity of the CMB quadrupole and octupole modes~\cite{A2}, and the observation of large-scale alignment in quasar polarization vectors~\cite{A3}. It should be stressed, however, that the above large-angle anomalies are still subject to an intense debate, since they could be indeed related to some common systematic. In this paper, we present an analysis of the large-scale isotropy assumption by means of magnitude-redshift data of type Ia supernovae (SNe). In particular, we use data from both Union and Union2 compilations, consisting of 307 and 557 type Ia SNe respectively, to set constraints on the parameters of an anisotropic model of the Universe (for earlier work on the possibility to test the Cosmological Principle with SN data, see Ref.~\cite{Kolatt,Gupta1,Gupta2,Koivisto-Mota}, while for recent works, see Ref.~\cite{Colin,Quartin} and references therein). We assume an anisotropic Bianchi type I cosmological model~\cite{Bianchi}, characterized by a cosmic shear $\Sigma$ in the presence of a dark energy fluid with anisotropic equation of state, characterized by a skewness $\delta$. This fluid, firstly studied by Barrow in Ref.~\cite{Barrow}, could be produced by the dynamics of a cosmic vector field, as shown by Koivisto and Mota in Ref.~\cite{Koivisto-Mota}. Other effects could give rise to an ellipsoidal universe, such as a large-scale cosmic magnetic field~\cite{prl,prd,prd2}, or a dark energy fluid having a nonvanishing velocity with respect to the CMB frame~\cite{Jimenez}. For an incomplete list of such mechanisms of universe anisotropization see, e.g., Ref.~\cite{Antoniou} and references therein. Testing the dark energy anisotropic model with type Ia SNe implies that we can only constrain the anisotropy parameters ($\Sigma$ and $\delta$) at relatively recent times (i.e., at redshift $z \lesssim 1.6$), their earlier evolution being largely unconstrained. We find no evidence in favor of anisotropies of either geometric origin ($\Sigma \neq 0$) or dark-energy origin ($\delta \neq 0$). However, we can put significant upper and lower bounds on the deviations of $\Sigma$ and $\delta$ from zero. The paper is organized as follows. In the next Section we set up the formalism of a cosmological model with anisotropic fluid while, in Section III, we derive the magnitude-redshift relation for such a universe. In Section IV, we use magnitude-redshift data of SNe from the Union and Union2 compilations to constrain all the free parameters of the model, including $\Sigma$ and $\delta$, so as to test the isotropy of the observable universe at $z \lesssim 1.6$. In Section V, we draw our conclusions.

} \renewcommand{\thesection}{\arabic{section}} The analysis of cosmic microwave background radiation reveals that the level of anisotropy in the large-scale geometry of the Universe must be below $10^{-5}$~\cite{parallax} at redshifts $z \sim 10^3$. This constraint do not exclude that a much larger level of anisotropy is allowed at recent times, $z \sim \mathcal{O}(1)$, as first noted in Ref.~\cite{Jimenez} and Ref.~\cite{Koivisto-Mota}. In particular the authors of Ref.~\cite{Koivisto-Mota} analyzed the magnitude-redshift data on type Ia supernovae in the ``GOLD'' data set of Riess et al.~\cite{Riess2006} (consisting of 182 supernovae), in a nonstandard cosmological model with anisotropic dark energy. They found that deviations from isotropy in the equation of state of dark energy were constrained at the level of $|\delta| < few \times 10^{-1}$. In this paper, we have extended the work~\cite{Koivisto-Mota} by constraining also the present level of cosmic anisotropy $\Sigma_0$, namely, the anisotropy in the large-scale geometry of the Universe, and by analyzing the magnitude-redshift data on type Ia supernovae in the Union and Union2 data sets of Kowalski et al.~\cite{Kowalski} (consisting of 307 supernovae) and Amanullah et al.~\cite{Union2} (consisting of 557 supernovae), respectively. In particular, by using Union2 data, we have confirmed the results of~\cite{Koivisto-Mota} about the skewness, finding \begin{equation} -0.16 < \delta < 0.12 \;\; (1\sigma \; \mbox{C.L.})\ , \end{equation} and we have put first limits on the present cosmic shear parameter $\Sigma_0$, \begin{equation} -0.012 < \Sigma_0 < 0.012 \;\; (1\sigma \; \mbox{C.L.})\ . \end{equation} We conclude that a standard isotropic universe is consistent with SN data, any deviation at redshifts $z \lesssim 1.6$ being constrained by the above results.

1012.5596

We analyze the magnitude-redshift data of type Ia supernovae included in the Union and Union2 compilations in the framework of an anisotropic Bianchi type I cosmological model and in the presence of a dark-energy fluid with anisotropic equation of state. We find that the amount of deviation from isotropy of the equation of state of dark energy, the skewness δ, and the present level of anisotropy of the large-scale geometry of the Universe, the actual shear Σ<SUB>0</SUB>, are constrained in the ranges -0.16≲δ≲0.12 and -0.012≲Σ<SUB>0</SUB>≲0.012 (1σ C.L.) by Union2 data. Supernova data are then compatible with a standard isotropic universe (δ=Σ<SUB>0</SUB>=0), but a large level of anisotropy, both in the geometry of the Universe and in the equation of state of dark energy, is allowed.

12205582

2011IJBC...21.2321K

1012

1012.2463_arXiv.txt

The aim of this paper is to study the orbital behavior at the neighborhood of a complex unstable periodic orbit in a 3D autonomous Hamiltonian system of galactic type. Complex instability is a type of instability of periodic orbits that appears in Hamiltonian systems of three or more degrees of freedom. \par In order to study the dynamical behavior at the neighborhood of a complex unstable periodic orbit we use the method of surfaces of section [Poincar\'{e} 1892], which has many applications to Dynamical Astronomy [for a review see e.g. Contopoulos 2002]. A basic problem in Hamiltonian systems of three degrees of freedom is the visualization of the 4D surfaces of section. Let us assume the phase space of an autonomous Hamiltonian system, that has 6 dimensions, e.g. in Cartesian coordinates, $(x,y,z,\dot x,\dot y,\dot z)$. For a given Jacobi constant a trajectory lies on a 5D manifold. In this manifold the surface of section is 4D. This does not allow us to visualize directly the surface of section. \par Patsis and Zachilas [1994] proposed a method to visualize 4D spaces of section. It is based on rotation of the 3D projections of the figures in order to understand the geometry of the formed structures and on color for understanding the distribution of the consequents in the 4th dimension. We use for this application the ``Mathematica'' package [Wolfram 1999]. We work in Cartesian coordinates and we consider the $y=0$ with $\dot y>0$ cross section. A set of three coordinates (e.g. $(x,\dot x, \dot z)$) are used for the 3D projection, while the fourth coordinate (in our example $z$) will determine the color of the consequents. There is a normalization of the color values in the [min($z$), max($z$)] interval, which is mapped to [0,1]. Following the intrinsic ``Mathematica'' subroutines our viewpoint is given in spherical coordinates. The distance $d$ of the center of the figure from the observer is given by ``Mathematica'' in a special scaled coordinate system, in which the longest side of the bounding box has length 1. For all figures we use $d=1$. The method associates the smooth distribution or the mixing of colors, with specific types of dynamical behavior in the 4th dimension [Patsis and Zachilas 1994][see also Katsanikas and Patsis 2011]. \par The calculation of the linear stability of a periodic orbit follows the method of Broucke [1969] and Hadjidemetriou [1975]. We first consider small deviations from its initial conditions and then integrate the orbit again to the next upward intersection. In this way a 4D map (Poincar\'{e} map) is established and relates the initial with the final point. The relation of the final deviations of this neighboring orbit from the periodic one, with the initially introduced deviations, can be written in vector form as $\xi$ = M $\xi_{0}$. Here $\xi$ is the final deviation, $\xi_{0}$ is the initial deviation, and M is a $4\times 4$ matrix, called the monodromy matrix. It can be shown, that the characteristic equation can be written in the form $\lambda^4 + a \lambda^3 + \beta \lambda^2 + a \lambda +1 = 0 $. Its solutions $\lambda_{i},\; i=1,2,3,4$, due to the symplectic identity of the monodromy matrix, that obey the relations $\lambda_{1} \lambda_{3}=1$ and $\lambda_{2} \lambda_{4}=1$ can be written as: \begin{equation} \lambda_i, \frac{1}{\lambda_i} = \frac {- b_i \pm \sqrt{b_i^2 - 4}}{2}, i=1,2 \end{equation} where \begin{equation} b_{1, 2} = \frac {a \pm \sqrt{\Delta}} {2} \end{equation} and \begin{equation} \Delta = a^2 - 4 (\beta - 2) \end{equation} \par The quantities $b_{1}$ and $b_{2}$ are called the stability indices. Following the notation of Contopoulos and Magnenat [1985], if $\Delta > 0,\; |b_1| < 2$ and $|b_2| < 2$, all four eigenvalues are complex on the unit circle and the periodic orbit is called ``stable'' (S). If $\Delta > 0$ and $|b_1| > 2, \; |b_2| < 2$ or $|b_1| < 2, \;|b_2| > 2$, the periodic orbit is called ``simple unstable" (U). In this case two eigenvalues are on the real axis and two are complex on the unit circle. If $\Delta > 0$ and $|b_1| > 2$ and $|b_2| > 2$, the periodic orbit is called ``double unstable" (DU) and the four eigenvalues are on the real axis. Finally, if $\Delta < 0$ the periodic orbit is called ``complex unstable'' ($\Delta$). In this case the four eigenvalues are complex numbers and they are off the unit circle. For the generalization of this kind of instability in Hamiltonian systems of N degrees of freedom the reader may refer to Skokos [2001]. If we have a stable one-parameter (in our system the Jacobi constant) family of periodic orbits, the four eigenvalues of a stable periodic orbit are complex on the unit circle. By varying the parameter we have a pairwise collision of eigenvalues on two conjugate points of the unit circle. From the Krein-Moser theorem [e.g. Contopoulos 2002 p.298] we can decide if after the collision of the eigenvalues they will remain on the unit circle and the periodic orbits of the family will stay stable, or if the eigenvalues will move out from the unit circle into the complex plane forming a complex quadruplet. In this latter case the periodic orbits of the family will become complex unstable and we will have a transition from stability to complex instability (also known as Hamiltonian Hopf Bifurcation). From an analytical point of view, the transition to complex instability has been studied using the Hamiltonian itself [Heggie 1985, Broer et al. 2007, Oll\'{e} et al. 2008] or 4D symplectic maps [Bridges et al. 1995]. In both cases, the approach consists of normal forms techniques [Oll\'{e} et al. 2005a,b] to simplify the Hamiltonian (or the map) and describe the local phase space structure near the critical periodic orbit (or fixed point in the discrete context). Such analysis shows, that the transition to complex instability gives rise to bifurcating invariant 2D tori in the flow context or invariant curves in 4D symplectic maps (as Poincar\'{e} map). The numerical computation of these \begin{figure} \begin{center} \includegraphics[scale=0.5]{diagr1_mz.eps} \caption{Stability diagram for $-5.5 < E_j < -4.5$, that shows the stability of family x1 (with black line) and its bifurcating family of p.o. x1v1 (with red line).} \label{cstab1} \end{center} \end{figure} invariant objects has been done for Hamiltonian systems of three degrees of freedom [Pfenniger 1985b, Oll\'{e} and Pfenniger 1998, Oll\'{e} et al. 2004] and for 4D symplectic maps [Pfenniger 1985a, Jorba and Oll\'{e} 2004]. It is remarkable, that there exists not one but two kinds of Hamiltonian Hopf bifurcations (as it happens in the usual dissipative setting), depending on the coefficients of the normal form [Van der Meer 1985]. These two kinds of bifurcations are usually called direct (supercritical) and inverse (subcritical). From a numerical point of view we can distinguish the two kinds of Hamiltonian Hopf bifurcation if we take firstly initial conditions in the vicinity of a complex unstable periodic orbit near the transition point from stability to complex instability. After that, we plot the consequents of the corresponding orbit in the surface of section. If the consequents are confined, we have a direct Hamiltonian Hopf Bifurcation. Otherwise, if the orbit escapes we have an inverse Hamiltonian Hopf Bifurcation. In this paper we examine the structure of the invariant surfaces in the 4D surface of section in the neighborhood of a complex unstable periodic orbit after a direct Hamiltonian Hopf bifurcation. The inverse Hopf bifurcation is not considered in the present paper, because in such a case all orbits close to a complex unstable periodic orbit escape [Jorba and Oll\'{e} 2004] and no confined structures appear.

In this paper we studied the phase space structure at the neighborhood of a complex unstable periodic orbit after a direct Hamiltonian Hopf bifurcation. Close to a Hamiltonian Hopf bifurcation we observe that: \begin{enumerate} \item We have a spiral structure formed by the consequents at the neighborhood of complex unstable periodic orbits in the 2D projections of the 4D surface of section as in Contopoulos et al [1994] and Papadaki et al [1995]. Here we find that we have this spiral structure in the 3D projections and we observe a smooth color variation along their arms. This means that this spiral structure is a 4D object. \item The consequents near a complex unstable periodic orbit arrive at a maximum distance from the periodic orbit and they move inwards. By repeating this process, a disk structure is formed in the 3D projection of the 4D surface of section, which is called a confined torus [Pfenniger 1985a,b, Jorba \& Oll\'{e} 2004 and Oll\'{e} et al 2004]. On this disk structure we observe a smooth color succession. This means that we have a disk structure (the confined torus) also in the 4D space of section. \item If we apply larger perturbations to the initial conditions we observe, that the disk structures become toroidals with smooth color variation and holes at their centers. For a critical value of the $\Delta x$ perturbation of the initial conditions we observe for a number of intersections a toroidal surface with smooth color variation. However, later, the consequents leave this toroidal surface and move away occupying a larger volume in the phase space. This is a case of stickiness [Contopoulos \& Harsoula 2008] in the case of complex instability and confined tori. This is the first time that we visualize stickiness in the neighborhood of a complex unstable periodic orbit. \item As the value of energy increases we do not find confined tori anymore, but clouds of points in the neighborhood of complex unstable periodic orbits. These clouds have mixing of colors. This means that we have strong chaos in the 4D space of section when we go far from the transition point from stability to complex instability. \item The calculation of the ``finite time'' $LCN$, gives a global value of the variation of the volume filled by the consequents of the orbits. We found that the $LCN(t)$ curve levels off after more than 4500 intersections and is around $2.5 \times 10^{-2}$ for the case of the cloud (Fig.~\ref{stick2}), and around $6 \times 10^{-3}$ for the confined torus (Fig.~\ref{csur3}). On the other hand, the method of color and rotation gives in a much more direct and detailed way the volumes of the phase space occupied by an orbit during its integration. This is of particular importance in Galactic Dynamics. \end{enumerate} \vspace{2cm} \textit

1012.2463

We study the orbital behavior at the neighborhood of complex unstable periodic orbits in a 3D autonomous Hamiltonian system of galactic type. At a transition of a family of periodic orbits from stability to complex instability (also known as Hamiltonian Hopf Bifurcation) the four eigenvalues of the stable periodic orbits move out of the unit circle. Then the periodic orbits become complex unstable. In this paper we first integrate initial conditions close to the ones of a complex unstable periodic orbit, which is close to the transition point. Then, we plot the consequents of the corresponding orbit in a 4D surface of section. To visualize this surface of section we use the method of color and rotation [Patsis and Zachilas 1994]. We find that the consequents are contained in 2D "confined tori". Then, we investigate the structure of the phase space in the neighborhood of complex unstable periodic orbits, which are further away from the transition point. In these cases we observe clouds of points in the 4D surfaces of section. The transition between the two types of orbital behavior is abrupt.

12272916

2012ASSP...28..155R

1012

1012.0185_arXiv.txt

Dwarf elliptical galaxies are the most common galaxy type in the nearby universe, accounting, for about 75\% of all objects in the Virgo cluster (e.g. \cite{trentham:2002}). They display a variety of shapes (from round to significantly flattened) and (sub)structures (disk features, bars or spiral arms). These different properties may suggest that more than just one mechanism was involved in their formation. Even though dEs may provide important clues on the main processes involved in galaxy assembly and evolution, many of their properties are still unknown and no present formation scenario(s) can fully explain observational characteristics of all dEs subclasses. A detailed analysis of the structure, internal kinematics and stellar populations of dEs is, therefore, essential to understanding the evolutionary properties of this class of objects.

We have presented the results of our SAURON study of four bright nucleated dwarf ellipticals from the Virgo Cluster. We have obtained reliable stellar velocity and velocity dispersion maps. We have also measured line-strength indices that allowed us to estimate ages and metallicities and compare the trends with those obtained for giant ellipticals from the SAURON sample. We find one rotating galaxy with misaligned photometric and kinematic major axes, indicating the presence of a bar. Two out of three of our objects that do not appear to rotate are significantly flattened, we speculate that they might be triaxial systems. We calculate line-strength indices and compare them with the predictions from the new MILES models \cite{vazdekis:2010}. The index-index diagrams show the presence of age and metallicity gradients different for each galaxy. The comparison with the Virgo giant ellipticals (Es) sample from the SAURON project \cite{kuntschner:2010} reveals that dEs have lower metallicity and are on average younger than Es. Both types seem to follow the same trends when index values as well as age, Z, and abundances are plotted as a function of velocity dispersion and dynamical mass. We conclude that dEs can, in principle, be a low-mass extension of Es, given that they seem to follow the same trends with mass. However, dEs as progenitors of Es seem less likely since the former have much lower abundance ratios than the latter. This argument is consistent with detailed photometric studies (e.g. \cite{kormendy:2009} and references therein).

1012.0185

Dwarf elliptical galaxies (dEs) are the most common galaxy type in nearby galaxy clusters; even so, many of their basic properties have yet to be quantified. Here we present the results of our study of 4 Virgo dwarf ellipticals obtained with the SAURON integral field unit on the William Herschel Telescope (La Palma, Spain). While traditional long-slit observations are likely to miss more complicated kinematic features, with SAURON we are able to study both kinematics and stellar populations in two dimensions, obtaining a much more detailed view of the mass distribution and star formation histories. What is visible even in such a small sample is that dEs are not a uniform group, not only morphologically, but also as far as their kinematic and stellar population properties are concerned. We find the presence of substructures, varying degrees of flattening and of rotation, as well as differences in age and metallicity gradients. We confirm that two of our galaxies are significantly flattened, yet non-rotating objects, which makes them likely triaxial systems. The comparison between the dwarf and the giant groups shows that dEs could be a low-mass extension of Es in the sense that they do seem to follow the same trends with mass. However, dEs as progenitors of Es seem less likely as we have seen that dEs have much lower abundance ratios.

12214178

2011MNRAS.418..536E

1012

1012.2978_arXiv.txt

Weak lensing by large-scale structure, so-called cosmic shear, was first detected in 2000 independently by several groups \citep{bre00,kwl00,wme00,wtk00}, and has recently progressed to an important tool in cosmology. Latest results \citep[e.g.][]{wmh05,smw06,hmv06,ses06,hss07,mrl07,fsh08,shj10} already indicate its great ability to constrain cosmological parameters which will be enhanced by large upcoming ground based surveys like KIDS\footnote{http://www.astro-wise.org/projects/KIDS/}, DES\footnote{www.darkenergysurvey.org/}, LSST\footnote{http://www.lsst.org/lsst}, and satellite missions Euclid\footnote{sci.esa.int/euclid/} and WFIRST\footnote{http://wfirst.gsfc.nasa.gov/}.\\ In order to meet the requirements for analyzing high precision data sets from these future surveys cosmic shear needs to overcome several challenges. On the observational side accurate shape measurements \citep[see][for latest developments]{hwb06,mhb07,bbb09,ber10} and photometric redshift information \citep{beh10,hac10} probably pose the most challenging problems. On the astrophysical side intrinsic alignments can mimic a shear signal \citep{his04,mhi06,jma10,mbb10} and need to be either removed \citep{kis03,jos08,jos09}, or modeled carefully in the subsequent analysis \citep{brk07,job09,scb10,kbs10}.\\ An important step to quantify and check for these systematics is the decomposition into E- and B-modes, where, to leading order, gravitational lensing only creates E-modes. In principle B-modes can arise from the limited validity of the Born approximation \citep{jsw00,hhw09} or redshift source clustering \citep{svm02}. Predictions coming from numerical simulations differ on the impact of these effects \citep[e.g.][]{hrh00,cnp01,jin02}, however, the observed B-mode amplitude is higher than one would expect from the foregoing explanations. Lensing bias \citep{srd09} is another possible source of B-modes; again the amplitude compared to the E-mode is small \citep{krh10}. Most likely, a B-mode detection indicates remaining systematics in the observation and data analysis, in particular an insufficient PSF-correction. \\ There exist several methods to perform the EB-mode decomposition; as a general classification they can be separated into real- and Fourier-space methods, where the latter are inspired by the pseudo-CL formalism invented to analyze the CMB polarization power spectrum \citep{hgn02,bct05}. A finite survey area and masking effects introduce a mixing of E- and B-modes in the CL's (also called \ti{leakage}), which prohibits a clean separation \citep{lct02}. The effect cannot be removed completely, however for CMB polarization the arising leakage B-mode can be suppressed \citep{lew03,smi06,kin10} to a level that enables for a possible detection of primordial B-modes (depending on the primordial tensor-to-scalar ratio).\\ \cite{hth10} extend the Pseudo-CL formalism to weak lensing, and test it on ray-tracing simulations, finding an B-mode leakage at the percent level and below depending on the Fourier mode $\ell$. However, the strength of this effect depends on masking and survey geometry and it needs to be examined on a case by case basis. Compared to CMB polarization additional difficulties arise in lensing when calculating the errors/covariances of the Pseudo CL's. For CMB polarization this covariance can be expressed in terms of the Pseudo CL's themselves assuming that the underlying field is Gaussian \citep{chc05}. This assumption is justified for the CMB field (if primordial non-Gaussianity is small/zero), however it is not true for the late time shear field where non-linear structure growth leads to non-negligible non-Gaussian effects \citep{whu00,svh07,taj09,sht09}. Here, higher order terms arise in the Pseudo CL shear covariance and the impact of EB-mode leakage on these terms is unknown.\\ Several real space EB-decomposing methods suffer from EB-mode mixing as well, e.g., the aperture mass dispersion, the shear E-mode correlation function, and the shear dispersion \citep{kse06}. These statistics can be calculated in terms of the shear two-point correlation function (2PCF), however all three measures need information on scales outside the interval $[\tmin;\tmax]$ of the measured 2PCF. This information can of course be modeled using a theoretical 2PCF, however this biases the results towards the cosmological model assumed in the 2PCF extension. The ring statistics \citep{sck07,esk10,fuk10} and more recently the COSEBIs \citep[][hereafter SEK10]{sek10} provide a new method to perform an EB-mode decomposition using a 2PCF measured over a finite angular range, thereby avoiding any EB-mode leakage/mixing.\\ Independent of whether cosmic shear data is analyzed in Fourier or real space both methods rely on accurate predictions for the corresponding shear measure in order to constrain cosmological models to the desired precision. \cite{hut05} find that one needs to know the power spectrum of density fluctuations ($P_\delta$) to $<$1\% accuracy over a range of 1-10 h/Mpc in order to obtain cosmic shear predictions that are sufficiently accurate for LSST. The Halofit fitting formula for $P_\delta$ described in \cite{smp03} is only accurate to $\sim$10\% (depending on cosmology and scale). The Coyote Universe emulator \citep{heh09,hww10,lhw10} has significantly improved on this issue; it models $P_\delta$ with percent accuracy in a 5 dimensional parameter space over the range $k \in [0.002;3.4] \, \mr{h/Mpc}$ within $z \in [0;1]$.\\ This paper has two goals: First, we built a predictions pipeline for second-order shear measures based on the Coyote Universe emulator and a modified Halofit. For improved weak lensing predictions from future numerical simulations we examine to which scale $k$ and redshift $z$ the density power spectrum must be known to model various second-order shear statistics to the desired accuracy for DES and LSST. The second goal is to examine the COSEBIs in a 5-dimensional parameter space. As outlined in SEK10 the COSEBIS do not only offer a check for B-mode systematics, moreover the authors argue that the COSEBIs' E-mode is the quantity that should enter a cosmic shear likelihood analysis. The 2PCF itself in particular should not be used for this purpose. The reason for this is that even when finding no B-modes in any EB-mode decomposition tests, the 2PCF can still be affected by B-modes that 1) mimic a constant or linear shear field 2) are localized in the power spectrum but present on scales ($\ell$) larger than the survey area.\\ The (logarithmic) COSEBIs are designed to contain all EB-decomposable cosmological information, and compress it into few data points (modes). SEK10 examined this in the two-dimensional parameter space $\om$ vs. $\sig$; we extend their analysis to 5 parameters which correspond to the cosmological parameter space covered by the Coyote Universe.\\ The paper starts with a short description of our weak lensing prediction pipeline; here and in Sect. \ref{sec:cosebis}, we also determine the requirements for future numerical simulations needed for sufficiently accurate weak lensing predictions for DES and LSST. In Sect. \ref{sec:like_analysis} we examine the performance of the COSEBIs in a 5-dimensional parameter space, in particular we address the questions 1) how many modes of the COSEBIs need to be included before the cosmological information saturates and 2) how does this saturation limit compare to the case of a pure E-mode 2PCF (which we can of course simulate but is never guaranteed in a real data set). We discuss our findings and conclude in Sect. \ref{sec:conclusions}.

\label{sec:conclusions} Accurate predictions of cosmic shear measures play an important role in the analysis of future weak lensing data sets. In the first part of this paper we present our new weak lensing predictions pipeline that is based on the Coyote Universe emulator and extends the emulator using the Halofit code. A shear power spectrum derived from this new pipeline differs from a shear power spectrum calculated from Halofit only by 6-11 \%, depending on the Fourier mode $\ell$.\\ We consider this as a first step in developing a weak lensing predictions pipeline that meets the requirements for DES and later for LSST, Euclid, and WFIRST. For this we examine to which $k$ in h/Mpc and $z$, the density power spectrum must be modeled accurately by numerical simulations in order to meet those requirements. We find that $\kmax=8$ $\mr{h/Mpc}$ causes a bias in the shear power spectrum at $\ell=4000$ that is within the statistical errors (intrinsic shape-noise and cosmic variance) of a DES-like survey, whereas for LSST already $\kmax=15$ $\mr{h/Mpc}$ is needed. \\ A future pipeline for weak lensing predictions that models the density power spectrum to such large $k$, the shear power spectrum to such high $\ell$, respectively, will have to take baryons into account \citep{jzl06,rzk08}. For example, \cite{jzl06} find that the shear power spectrum on scales of $\ell \in [1000,10000]$ are affected by 1-10\%, depending on the level of complexity in the treatment of baryons.\\ In addition to the treatment of baryons additional corrections need to taken into account. For example, shape distortions probe the reduced shear $g=\gamma/(1-\kappa)$ instead of the shear ($\gamma$) itself \citep[e.g.,][]{sha09}, and when calculating the reduced shear from the measured mean ellipticity higher order terms in $g$ need to be taken into account. Corrections to the Born approximation are necessary to account for multiple deflections of light rays \citep[e.g.,][]{hhw09}, and corrections to biases that occur because de(magnification) of galaxies due to lensing correlates with selection criteria of the considered galaxy sample \citep[e.g.,][]{srd09}. These corrections are calculated in \cite{krh10}; the authors find that reduced shear and magnification bias are important already for DES-like surveys, whereas the other effects will only become important for weak lensing data from LSST.\\ In the second part of this paper we extend earlier studies of the COSEBIs, the most recently developed EB-mode decomposing second-order cosmic shear measure. In particular, we are interested in their performance in a high-dimensional cosmological parameter space. Compared to other second-order shear statistics the COSEBIs can be calculated from a shear 2PCF measured on a finite interval $[\tmin;\tmax]$. The 2PCF again is independent of any masking effects or survey geometry and thereby avoids several difficulties present in other second-order shear measurement methods, e.g. the shear power spectrum. \\ The COSEBIs can be imagined as a filtered version of the 2PCF (or the shear power spectrum); instead of having angular scale $\vt$ or Fourier modes $\ell$ as an argument, they are a function of the order of the polynomial which is used as their filter function. Furthermore, they are designed to condense the second-order cosmic shear information into a (small) number of discrete modes. In SEK10 two different filter functions are examined: Filter functions that sample the 2PCF linearly can be expressed very conveniently in terms of Legendre polynomials, however the number of modes that is needed to capture all the cosmological information is rather large in this case. The logarithmic filter functions are much more efficient; SEK10 find that $\sim$ 5 modes are sufficient to capture the cosmological information in the two-dimensional parameter space $\om$ vs. $\sig$.\\ Here, we extend their analysis and examine the performance of the COSEBIs in a 5-dimensional parameter space. We find that the bulk of the cosmological information is contained on angular scale $<$100', whereas scales above this threshold only contribute little to constraints on cosmology. Furthermore, we find that $\sim$8 modes of the COSEBIs are sufficient to capture all the EB-mode decomposable cosmological information. This number is still relatively small, in particular as we show that the COSEBIs can be calculated accurately up to the 9th mode when knowing $\pd$ ``only'' up to $\kmax=10$ h/Mpc. From these results we conclude that the COSEBIs are fairly robust against theoretical uncertainties in modeling the density power spectrum, and represent a good choice as a second-order weak lensing statistics.

1012.2978

Analysing future weak-lensing data sets from KIDS, Dark Energy Survey (DES), LSST, Euclid and WFIRST requires precise predictions for the weak-lensing measures. In this paper, we present a weak-lensing prediction code based on the Coyote Universe emulator. The Coyote Universe emulator predicts the (non-linear) power spectrum of density fluctuations (P<SUB>δ</SUB>) to high accuracy for k∈[0.002; 3.4] h Mpc<SUP>-1</SUP> within the redshift interval z∈[0; 1]; outside this regime, we extend P<SUB>δ</SUB> using a modified HALOFIT code. <P />This pipeline is used to calculate various second-order cosmic shear statistics, e.g., shear power spectrum, shear-shear correlation function, ring statistics and Complete Orthogonal Set of EB-mode Integrals (COSEBIs), and we examine how the upper limit in k (and z), to which P<SUB>δ</SUB> is known, impacts on these statistics. For example, we find that k<SUB>max</SUB>∼ 8 h Mpc<SUP>-1</SUP> causes a bias in the shear power spectrum at ℓ∼ 4000 that is comparable to the statistical errors (intrinsic shape noise and cosmic variance) of a DES-like survey, whereas for LSST-like errors k<SUB>max</SUB>∼ 15 h Mpc<SUP>-1</SUP> is needed to limit the bias at ℓ∼ 4000. <P />For the most recently developed second-order shear statistics, the COSEBIs, we find that nine modes can be calculated accurately knowing P<SUB>δ</SUB> to k<SUB>max</SUB>= 10 h Mpc<SUP>-1</SUP>. The COSEBIs allow for an EB-mode decomposition using a shear-shear correlation function measured over a finite range, thereby avoiding any EB-mode mixing due to finite survey size. We perform a detailed study in a five-dimensional parameter space in order to examine whether all cosmological information is captured by these nine modes with the result that already 7-8 modes are sufficient.

12202213

2011CaJPh..89..395L

1012

1012.5243_arXiv.txt

Accurate oscillator strength values are of central importance for the analysis and interpretation of astrophysical spectra. Detailed radiative transfer modeling of the profiles of absorption and emission lines observed in the spectrum of the Sun and stars requires reliable atomic data to correctly infer the physical properties of the atmospheric line formation region. These data are however currently unavailable for a large fraction of the weak and medium-strong absorption lines observed in optical stellar spectra. In recent years a number of online databases have become available that offer atomic line data compiled from a large variety of sources in the scientific literature, such as {\sc nist}, {\sc vald}, the Kurucz website, Topbase, etc. Thanks to the efforts of the atomic database developers the detailed synthesis of stellar spectra has become much more accessible to astrophysicists without having to fully address the fundamental atomic physics needed to acquire these important atomic input data. On the other hand, the successful application of publicly available atomic databases entirely depends on the quality of the data compiled in them, but which is frequently not (or more appropriately cannot always be) guaranteed. The demand for spectral standard star atlases is steadily increasing with the fast improvements of resolving power and the quality by which stars of all spectral types are observed with modern spectrographs. Printed atlases of stellar spectra often only provide a small list of identified features without an assessment of the reliability of the spectral line identifications. Users often have no means to determine if the spectral line identifications are valid, or if they have been revised since publication. On the other hand, many atomic databases offer line data (that can be text queried online) based in part on theoretical calculations that have not been tested against observed stellar spectra. Users cannot readily verify the quality of the line data, or if it applies to their spectroscopic observations. The data can contain line identifications that do not apply to an observed stellar spectrum because of unknown atmospheric formation conditions or elemental abundance differences with the solar values. Conversely, observed spectral features often cannot be identified because the quality of the provided atomic and molecular line data is limited and requires further improvements. In this paper we present a critical evaluation of 911 oscillator strength values of astrophysical importance based on detailed spectral synthesis calculations of the weak and medium-strong neutral lines spectrum of the Sun between 400 nm and 680 nm, and of two bright spectroscopic reference stars, Procyon and $\epsilon$ Eri, observed with large spectral dispersion, and with S/N ratios exceeding 2,000. The observed and best-fit theoretical spectra are publicly available in the online SpectroWeb\footnote{\tt alobel.freeshell.org/spectrowebl.html} database \cite{lobe1}. Users can directly assess the quality of the absorption line identifications by comparing the observed spectra with state-of-the-art computed spectra. The graphical web interface enables users to select 10 or 25 \AA\, wide spectral regions-of-interest from an interactive list of observed wavelengths, together with the employed atomic line data. We investigate the accuracy of the log(gf)-values obtained with best fits to the observed spectra using an average curve-of growth analysis method. We also discuss a comparison of computed and observed line equivalent widths to investigate remarkable trends in a number of important neutral elements that reveal systematic over-estimations of literature (e.g., that are offered in online atomic databases) log(gf)-values for weak Fe~{\sc i}, Ni~{\sc i}, and Si~{\sc i} lines with central line depths below 15 \%.

We measure accurate log(gf)-values of 911 atomic absorption lines of neutral elements observed in high-resolution stellar spectra. We perform detailed synthesis calculations of the optical spectrum of the Sun, Procyon, and $\epsilon$ Eri observed with very large S/N ratios exceeding 2,000. We find systematic over-estimations in the log(gf)-values adopted from atomic databases for weak lines of iron-peak elements (such as Fe~{\sc i}, Ni~{\sc i}, Cr~{\sc i}, and Ti~{\sc i}), and of Si~{\sc i}. We demonstrate that the computed equivalent line widths strongly correlate to observed values previously published for the lines in the optical spectrum of the Sun. An average curve-of-growth analysis reveals that the errors of the log(gf)-values we measure for weak lines are smaller than the standard mean error. The systematic trends we find for weak lines therefore cannot be attributed to large systematic errors in the detailed spectral line synthesis modeling method of the three stars. We attribute the remarkable systematic trends to the limited accuracy of predicted log(gf)-values of weak absorption lines having central depths below 15 \% currently offered in online atomic databases. The adjusted log(gf)-values we measure are available in the online SpectroWeb database, together with the observed and computed spectra and the central line depths and equivalent width values computed for the three stars. Further updates of the log(gf)-values and other atomic line data based on accurate laboratory measurements and advanced semi-empiric computations are urgently needed for reliable line identifications in the optical spectra of stars of all spectral subtypes.

1012.5243

Herein we develop a new method to determine oscillator strength values of atomic absorption lines with state-of-the-art detailed spectral synthesis calculations of the optical spectrum of the Sun and of standard spectral reference stars. We update the log(gf) values of 911 neutral lines observed in the KPNO-FTS flux spectrum of the Sun and high-resolution echelle spectra (R = 80 000) of Procyon (F5 IV-V) and Eps Eri (K2 V) observed with large signal-to-noise (S/N) ratios of 2000 using the new Mercator-Hermes spectrograph at La Palma Observatory (Spain). We find for 483 Fe I, 85 Ni I, and 51 Si I absorption lines in the sample a systematic overestimation of the literature log(gf) values with central line depths below 15%. We employ a curve-of-growth analysis technique to test the accuracy of the new oscillator strength values and compare calculated equivalent line widths to the Moore, Minnaert, and Houtgast atlas of the Sun. The online SpectroWeb database at http://spectra.freeshell.org interactively displays the observed and synthetic spectra and provides the new log(gf) values together with important atomic line data. The graphical database is under development for stellar reference spectra of every spectral sub-class observed with large spectral resolution and S/N ratios.

2564448

2011ApJ...728...66S

1012

1012.1526_arXiv.txt

\citet{brown2004a} introduced the concept of completeness to study the selection effects introduced by observatory architectures on direct searches for sub-stellar companions. Assuming distributions for semi-major axis and eccentricity of planetary orbits, Brown calculated the probability that a companion would fall outside the telescope's central obscuration during an observation of a star. \citet{brown2005} subsequently expanded this concept to include the selection effects due to the photometric restrictions on observability introduced by telescope optics, and \citet{brown2009} demonstrated how completeness could be evaluated for indirect companion detection methods such as astrometry. Completeness has also been extended to account for multiple observations of one star at different times \citep{brown2010new}, and has been utilized in mission analysis and development for a variety of proposed exoplanet observatories \citep{savransky2010,brown2009}. The direct detection (imaging) completeness is evaluated by assuming that a companion will be observable if its angular separation from the star is greater than the observatory's inner working angle (IWA), and illuminated such that the difference in brightness between star and companion ($\Delta$mag) is below a threshold value, called the limiting $\Delta$mag, or $\Delta$mag$_0$. The IWA represents the minimum angular separation between the telescope center-line and detectable objects on the sky. It is determined by the size of a central obscuration, or the capability of adaptive optics systems to remove light from certain areas of the image plane, or the size and geometry of external occulting optics. $\Delta$mag$_0$ represents the point where systematic errors produce unresolvable confusion between planet signal and background noise. To calculate the completeness, probability distributions (or constant values) are assumed for planetary orbital elements and physical properties (see \S\ref{sec:orbit_defs}). A large, equal number of samples is generated from each distribution, and the star-planet angular separation and $\Delta$mag are calculated for each set of samples. When binned in a two-dimensional histogram, these generate a density function representing the probability that a planet drawn at random from the assumed population will have a given angular separation and $\Delta$mag. Integrating over this density yields a cumulative distribution function (CDF), which can be used to determine the probability that an observatory with a given $\Delta$mag$_0$ and IWA, observing a specific star once, will be able to detect a planet belonging to the assumed population. The procedure described above is a Monte-Carlo sampling of a bivariate distribution function of non-independent arguments (since star-planet separation and $\Delta$mag are functions of the same parameters). This means that to find any one point (or section) of the completeness distribution, it is necessary to sample it completely. Because of the relatively high dimensionality of the initial parameter space and wide range of values certain parameters can take, complete sampling requires a large number of Monte-Carlo trials. The first simulation in \citet{brown2005}, for example, includes 100 million samples, and it can be shown that certain low probability areas of the function are under-sampled. Any alternate method of sampling completeness requires at least some knowledge of its density function. In particular, Markov chain methods such as Metropolis-Hastings \citep{hastings1970monte} perform significantly better if the proposal distribution (a function used to `propose' new samples that are then either accepted or rejected) closely approximates the target distribution. A special case of Metropolis-Hastings, known as Gibbs sampling, can be used to generate a sequence of samples from the joint distribution of two variables if their conditional probabilities are known. Our goal is to derive the distribution functions of the arguments to the completeness functions. We do so by starting with an ensemble of Keplerian orbits whose orientation is uniformly distributed in space, and deriving the distribution functions of these orbits' Keplerian parameters. Rather than constraining the population of orbits, we make no assumptions as to the distribution of orbital semi-major axis and eccentricity. This allows us to derive the completely general distribution functions presented in \S\ref{sec:pdfs}. Building upon this, we further consider parameters related to direct planetary observations in \S\ref{sec:pdfsObs}, and make the discovery that the planetary phase angle (star-planet-observer angle) is independent of any of the orbital parameters. This makes it possible to write distribution functions for quantities directly related to the two parameters of the direct detection completeness function. It is important to note that, while the specific application explored here is direct imaging, the distributions derived in this paper are more broadly applicable to exoplanet studies in general. For example, Keplerian fits are often employed in doppler spectroscopy surveys, making these derivations useful for inferring the true distributions of orbital parameters derived from radial velocity data sets. Similarly, statistical analyses play an important role in other methods of exoplanet study, including transit photometry and microlensing surveys \citep{gould2010frequency}.

Starting with a uniform distribution of orbital orientations and arbitrary PDFs for eccentricity and semi-major axis, we have derived expressions for the distributions of the remaining Keplerian orbital elements, and of parameters used to describe direct exoplanet detections. We have demonstrated the independence of phase angle from the distributions of the Keplerian orbital elements, allowing for the calculation of the distributions of both apparent separation and flux ratio, which can be directly sampled to generate the completeness distribution. At the same time, we have provided fully analytic forms for the distributions of Keplerian orbital elements based on uniform distributions of eccentricity and semi-major axis, as well as logarithmically distributed semi-major axes, assumptions often made in research related to planet-finding mission planning and analysis. These forms should allow researchers to increase the efficiency of their simulations, and to empirically check for global errors generated by under-sampling. The same equations and procedures described here will also be useful for other statistical work associated with planet-finding, including calculating the likelihood of transits or inferring the true distribution of orbital parameters from doppler spectroscopy surveys.

1012.1526

Starting with just the assumption of uniformly distributed orbital orientations, we derive expressions for the distributions of the Keplerian orbital elements as functions of arbitrary distributions of eccentricity and semimajor axis. We present methods for finding the probability density functions of the true anomaly, eccentric anomaly, orbital radius, and other parameters used in describing direct planetary observations. We also demonstrate the independence of the distribution of phase angle, which is highly significant in the study of direct searches, and present examples validating the derived expressions.

12208111

2011JGRA..116.3104S

1012

1012.3523_arXiv.txt

Refractive scattering of radiation by density turbulence in the Sun's corona and solar wind leads to angular broadening of embedded radio sources, and of cosmic sources observed through these media. The process is similar to the twinkling of stars and modified``seeing'' caused by density turbulence in Earth's atmosphere and ionosphere. This scattering process has been investigated for many years using geometrical optics [e.g., Steinberg et al 1971] and the parabolic wave equation [e.g., Lee and Jokipii, 1975; Coles and Harmon 1989; Bastian 1994; Cairns 1998]. Scattering is thought to affect the observed properties of type II and III solar radio bursts in several ways: greatly increasing the angular sizes of the sources [e.g., Riddle 1974], causing the time profiles to have exponential decreases [e.g., Robinson and Cairns, 1998], and causing anomalously low brightness temperatures at decametric wavelengths [e.g., Thejappa \& Kundu 1992]. The primary motivation of this paper is to investigate the constraints imposed on models of density turbulence in the solar corona by recent observations at 327 MHz [Mercier et al., 2006], made by combining visibilities from the Giant Metrewave Radio Telescope (GMRT) in Pune, India, and the Nancay Radioheliograph (NRH) in France. The maps of Mercier et al. [2006] show structures ranging from the smallest observed size of $49^{''}$ to that of the whole Sun, with dynamic ranges as high as a few hundred. These features make them the best meter wavelength snapshot maps of the solar corona to date. Mercier et al. [2006] found the smallest steady angular size of type I solar noise storms to be $49^{''}$ in their high dynamic range, full disk, 17 second snapshots. We therefore adopt the smallest observed source size of $49^{''}$ in these maps to be a canonical number for comparison with our model predictions. The only other two-dimensional map showing small source sizes that we are aware of is that of Zlobec et al [1992], who observed a source as small as $30^{''}$ at 327 MHz. However, the dynamic range of their map was severely limited, and included only a very small range of size scales. It should be noted that the lower limit to the observed size is imposed by the resolution of the instrument if scattering by density turbulence is weak enough. The scattering calculations shown below imply that there is a possibility that smaller solar sources will be detected in the future by instruments with improved angular resolution. In this paper we use a formalism based on the paraxial wave equation and the structure function, together with observationally based models for the density turbulence that scatter the radiation, to predict the size of sources in the solar corona at 327 MHz. Our results can be interpreted as the scatter-broadened image of an ideal point source in the solar corona. We have used an empirical model for the amplitude $C_{N}^{2}(R)$ of coronal turbulence that is based directly a fit to the scattering measure obtained from VLBI observations of cosmic sources broadened by scattering in the outer solar corona and inner solar wind. Here $R$ is the heliocentric distance. We have assumed that this model is valid throughout the corona, specifically at smaller $R$. We also consider the effects of spherical and plane wave propagation, variations of the inner scale $l_{i}(R)$ and power-law index $\alpha$ of the turbulence on the predicted source sizes. In most cases, we find that the models predict sizes that are at least an order of magnitude below the smallest observed size of 49$^{''}$ at 327 MHz. Our formalism and analyses differ primarily from those of Bastian [1994] in the models for $C_{N}^{2}(R)$ and the electron density profile $n_{e}(R)$, while our applications are to metric rather than centimetric and decimetric emissions. Since our predictions are much smaller and Bastian's [1994] predictions much larger than $49''$ at $327$ MHz, the analyses demonstrate the importance of knowing $C_{N}^{2}(R)$, $n_{e}(R)$, and $l_{i}(R)$ much better for future observations and predictions of solar sources. These quantities are also relevant to the heating and outward flow of the coronal plasma, with activity localized to specific ranges of $R$ potentially leading to larger $C_{N}^{2}(R)$ and so enhanced scattering at, say, decimetric frequencies than expected at, say, metric frequencies. The paper is organized as follows. In \S~2 we summarize the scattering formalism and observations of the density turbulence. Coronal density models are described in \S~3. The results are presented in \S~4, including estimates of the predicted angular broadening, and the implications for coronal density turbulence. The conclusions are presented in \S~5.

The highest resolution meter wavelength observations of the solar corona reveal compact sources around 49$^{''}$ in size at 327 MHz. The main aim of this paper is to employ an observationally-motivated model for the turbulence amplitude $C_{N}^{2}$, and see what it implies for the predicted scattering angle for radio sources located in the solar corona. We reference our calculations to the same frequency (viz. 327 MHz) at which the smallest source size is observed. We employ the parabolic wave equation, together with the standard asymptotic forms for the phase structure function, which are valid for situations where the effective baseline is either much larger or much smaller than the inner scale. We define the predicted scattering angle $\theta_{c}$ via Eq. (\ref{eq3b2b}) as the angle where the phase structure function falls to 1/e times its peak value. Effectively, this means that the scattering angles predicted here should be interpreted as the scatter-broadened image of an ideal point source in the solar corona. The real source will have an intrinsic size (i.e., it will not be a point source) and the observable source will be the convolution of the intrinsic source profile with $\theta_{c}$, provided there are no instrumental limitations. The results in this paper should therefore be regarded as lower limits to the observable source size set by scattering. The general consensus now seems to be that there is not much about the instrinsic source size that can be gleaned from scatter-broadened images [Bougeret \& Steinberg, 1977; Melrose, 1980; Bastian, 1994]. However, if the intrinsic source size and $\theta_{c}$ are similar in size, then the observed source size will be larger than $\theta_{c}$ by a factor near $\sqrt{2}$. Note that this factor cannot account for the large discrepancies between the minimum source size observed at 327 MHz [Mercier et al., 2006] and those predicted here. If, on the other hand, the intrinsic source size is much larger, then scatter broadening does not play a significant role. We have included refractive index effects that can be important when the radiation is emitted near the fundamental plasma level, but found them to be relatively unimportant. The inner scale is included via the Coles and Harmon [1986] model, interpreted in terms of cyclotron damping of MHD waves, and so depends primarily on the ambient electron density. We employ a hybrid model for the electron density that yields reasonable heights for meter wavelength emission at the fundamental. In view of the uncertainity in its value in the inner corona, the power law index $\alpha$ characterizing the turbulent spectrum is taken to be a free parameter. We consider both plane wave and spherical wave propagation. For the geometry we consider, where the source is embedded in the scattering medium, the spherical wave description is arguably more appropriate. We have thus explored a wide variety of effects. We observe that there is no significant difference in the predicted scattering angle between fundamental and second harmonic emission. We also find that the removal of the refractive index effect causes a negligible change in the predicted scattering angle. We find that the spherical divergence effect results in a significant lowering of the predicted scattering angle (by around 2 orders of magnitude). We find that removing inner scale effects by artificially setting the inner scale to be equal to a very small value (instead of determining it self-consistently from Eq~\ref{eq2}) results in a significant enhancement of the predicted source size. The enhancement is greatest for flatter spectra, where it can be a factor of around $50$, and it progressively disappears for steeper spectra. As mentioned earlier, the power law index of the turbulent spectrum is a free parameter. There is a formal divergence at $\alpha = 4$ in Eqs (\ref{eqp2v41}) and (\ref{eqp2v42}), and we therefore limit the computations to a maximum value of $\alpha = 3.97$. The maximum value of the predicted scattering angle thus occurs at $\alpha = 3.97$. For plane wave propagation, the predicted source size for Kolmogorov turbulence is around 10$^{''}$ lower than the observed one. For spherical wave propagation, we find that the maximum value of the predicted scattering angle is at least 25 times smaller than the observed one. It is emphasized that plane wave propagation is relevant to the well-accepted empirical formula of Erickson [1964], which predicts large source sizes, since it pertains to observations of celestial sources through the solar wind. Even so, with current estimates of $l_{i}(R)$ implying that $s_{\rm eff} > l_{i}$ (Figure 4), additional scattering is required to bring Erickson's result into quantitative agreement with the calculations here for plane wave scattering. Alternatively, the inner scale $l_{i}(R)$ should be smaller than that predicted by the Coles and Harmon [1989] model assumed here, so that $s_{\rm eff} < l_{i}$. The crucial result of this paper is that the predicted source sizes are considerably smaller than the observed lower limit of $49''$ when spherical wave propagation and inner scale effects are included, as they should be for sources in the solar corona. This broad trend of the models substantially underpredicting the source size can be interpreted in three ways that are not exclusive and can occur in combination. First, it could imply that source sizes much smaller than those that have been observed so far actually exist in the solar corona, and can potentially be observed. All the instances of observations of small sources to date have been limited by the instrument resolution; it is therefore quite likely that smaller sources can be detected when instrument resolutions are improved. Second, this broad trend can be taken to imply that our naive extrapolation of the empirical form for the turbulence amplitude $C_{N}^{2}(R)$ to the inner corona is not justified. The results could be taken to imply that $C_{N}^{2}(R)$ in the inner corona is far higher than suggested by the empirical formula (\ref{eq3}). This could be due to the functional form for $C_{N}^{2}(R)$ increasing more rapidly with decreasing $R$ or due to a larger normalization factor or both effects. Third, the model (\ref{eq2}) may significantly overestimate $l_{i}(R)$, meaning that the turbulent cascade extends to smaller length scales (larger $q$) and leads to more scattering. These proposals all appear reasonable, and we regard all three as viable. Finally, we discuss the connection between our work and that of Bastian [1994]. The methodology is similar, and we investigate similar issues such as the effects of the inner scale, turbulence index $\alpha$, and spherical versus planar wave propagation. Bastian's [1994] findings are contrary to ours: we find that our model predictions are substantially below the minimum observed size of $49''$, while Bastian's [1994] model predictions are substantially above $49''$. Thus, a priori, both models need revision. A major difference is in the choice of a model for $C_{N}^{2}$. Bastian [1994] uses a model for $C_{N}^{2}$ which is proportional to the square of the background electron density and assumes the density model of Riddle (1974), which also involves a constant of proportionality. These two constants of proportionality are absorbed into one and fixed by normalizing the structure function $D_{20}(10 {\rm km})$ for a baseline of 10 km, an observing wavelength of $20$ cm, and an elongation of 5 $R_{\odot}$ . In contrast, as explained earlier, the $C_{N}^{2}$ model we use is determined by an empirical fit to VLBI scattering observations between 10--50 $R_{\odot}$; this was motivated by the need to use a $C_{N}^{2}$ model that is derived as directly as possible from observations. A minor matter is that Bastian [1994] discusses the disk to limb variation in the predicted scattering angle, whereas our treatment is valid only for sources that are reasonably close to disk center. In order to do so, we would need to use the general formalism used here, together with an integration path that incorporates the appropriate extra path length needed for sources that are displaced from the disk center. An appropriate means of comparing the normalizations of the two treatments is thus to compare the normalizations of the structure function. Using (\ref{eq3b2}) -- (\ref{eq3b3}), we write the structure functions as {\begin{eqnarray} \nonumber D_{sf}(s) = \frac{4 \, \pi^{2}\,s^{2}}{\lambda^{2}}\,\theta_{c\,sf}^{2} \, , \\ D_{pf}(s) = \biggl ( \frac{2\, \pi\, s}{\lambda} \biggr )^{\alpha - 2}\, \theta_{c\,pf}^{\alpha - 2}\, , \label{eq3a5} \end{eqnarray}} where $\theta_{c\, sf}(s)$ is the value of $\theta_{c}$ for spherical wave propagation using branch (\ref{eqp2v42}) and $\theta_{c\, pf}(s)$ corresponds to plane wave propagation for branch (\ref{eqp2v41}). Then using our models for $C_{N}^{2}(R)$, $l_{i}(R)$, and $n_{hyb}(R)$ we find that $D_{sf}(s) = 2.8 \times 10^{-3}$ rad$^{2}$ for s = 10 km, $\lambda$ = 91 cm (corresponding to 327 MHz), $\alpha = 11/3$ and a starting height corresponding to 327 MHz fundamental emission. The same prescription and parameters yield $D_{pf}(s) = 7.7 \times 10^{-3}$ rad$^{5/3}$. Since the structure functions we derive are based on integrations over heliocentric distance, we cannot assign a specific elongation to them. In comparison, Bastian normalizes $C_{N}^{2}$ by assuming $D_{20 {\rm cm}}(10 {\rm km}) = $4--12 rad$^{2}$, based on measurements of by Coles and Harmon [1989] and Armstrong et al. [1990] of cosmic sources (implying primarily planar wave effects) at an elongation of 5 $R_{\odot}$. In order to normalize Bastian's [1994] values for the structure function to a wavelength of 91 cm, we concentrate on the structure function for spherical wave propagation. Inspection of Eqs [\ref{eq3b2a}] and [\ref{eq3a5}] reveals that $D_{sf}(s) \propto \lambda^{2}$. Therefore, $D_{20 {\rm cm}}(10 {\rm km})/D_{91 {\rm cm}}(10 {\rm km}) = 21$. Bastian's [1994] range of values for $D_{20 {\rm cm}}(10 {\rm km})$ thus corresponds to $D_{91 {\rm cm}}(10 {\rm km}) = 82 - 250$~rad$^{2}$. The difference in $D_{91 {\rm cm}}(10 {\rm km})$ between the two prescriptions is thus a factor of $\approx (30 - 90)\times 10^{3}$, corresponding to a factor $\approx 170 - 300$ in $\theta_{c}$. This large difference $\approx (30 - 90)\times 10^{3}$ in the normalization of the structure function is primarily indicative of a corresponding difference in the normalization of $C_{N}^{2}$ between the two treatments; this is because neglect of inner scale effects increases $\theta_{c}$ by less than a factor of $10$ for Kolmogorov turbulence in Figure 7. In this connection, we note that Bastian's [1994] normalization for $D_{20 {\rm cm}}(10 {\rm km})$ is based on values of $D(s)$ measured at an elongation of 5 $R_{\odot}$. We also note (e.g., Fig 1 of Coles \& Harmon [1989]) that values of $D(s)$ measured at larger elongations can be considerably lower (by as much as a few orders or magnitude, depending upon the elongation). This is significant, since the model for $C_{N}^{2}$ that we use in this paper is based on observations between 10 and 50 $R_{\odot}$. In summary, the foregoing results demonstrate conclusively that spherical wave propagation effects are vital for solar sources, with plane wave predictions several orders of magnitude larger than the spherical predictions. Similarly, inner scale effects are quantitatively important, while fundamental versus harmonic radiation effects are relatively small. The results and discussion above demonstrate the importance of accurate models for $C_{N}^{2}(R)$ and to a lesser extent models of $l_{i}(R)$ and $\alpha(R)$. This paper's prescription for $C_{N}^{2}$ (Eq~[\ref{eq3}]) is empirical and directly based on observations (but extrapolated to smaller $R$), does not have any normalization constants that need to be determined, and leads to scattered sizes for a point source that are smaller than the minimum source size observed to date ($49''$ by Mercier et al. [2006]). Thus smaller source sizes than $49''$ may be observable. In contrast, another well-known prescription [Bastian, 1994] predicts much stronger scattering with source sizes always larger than $49''$: while this is inconsistent with the minimum source size observed to date at 327 MHz, it may provide the extra scattering required to account for Erickson's empirical angular broadening result for cosmic sources viewed through the solar wind. While this paper's results extend and confirm previous theoretical results pertaining to spherical vs. plane wave effects and provide the first explanation of the small source sizes recently observed, it is also clear that more observational and theoretical work is required on $C_{N}^{2}(R)$ especially, but also on $l_{i}(R)$ and $\alpha(k,R)$. This includes temporal variations over the solar cycle but also spatial variations between radio source regions and other regions of the corona. Increases in $C_{N}^{2}(R)$ and decreases in $l_{i}(R)$ would lead to more scattering. Work on both $n_{e}(R)$ and $\delta n(R) / n_{e}(R)$ may be useful [e.g., Efimov et al., 2008; Cairns et al., 2009]. It is quite possible that scattering observations and theory will provide useful constraints on these five quantities and therefore on the processes heating the solar corona and accelerating the solar wind.

1012.3523

We seek to reconcile observations of small source sizes in the solar corona at 327 MHz with predictions of scattering models that incorporate refractive index effects, inner scale effects, and a spherically diverging wavefront. We use an empirical prescription for the turbulence amplitude C<SUB>N</SUB><SUP>2</SUP>(R) based on very long baseline interferometry observations by Spangler et al. of compact radio sources against the solar wind for heliocentric distances R ≈ 10-50 R<SUB>⊙</SUB>. We use the Coles and Harmon model for the inner scale l<SUB>i</SUB>(R), which is presumed to arise from cyclotron damping. In view of the prevalent uncertainty in the power law index that characterizes solar wind turbulence at various heliocentric distances, we retain this index as a free parameter. We find that the inclusion of spherical divergence effects suppresses the predicted source size substantially. We also find that inner scale effects significantly reduce the predicted source size. An important general finding for solar sources is that the calculations substantially underpredict the observed source size. Three possible, nonexclusive, interpretations of this general result are proposed. First and simplest, future observations with better angular resolution will detect much smaller sources. Consistent with this, previous observations of small sources in the corona at metric wavelengths are limited by the instrument resolution. Second, the spatially varying level of turbulence C<SUB>N</SUB><SUP>2</SUP>(R) is much larger in the inner corona than predicted by straightforward extrapolation sunward of the empirical prescription, which was based on observations between 10 and 50 R<SUB>⊙</SUB>. Either the functional form or the constant of proportionality could be different. Third, perhaps the inner scale is smaller than the model, leading to increased scattering. These results and interpretations are discussed and compared with earlier work.

12132442

2010arXiv1012.3009T

1012

1012.3009_arXiv.txt

\label{sec:introduction} Discoveries of empirical correlations of gamma-ray bursts (GRBs) raised many researches on early universe using GRBs. One of the most well studied correlations is the one between the spectral peak energy ($E_p$) and isotropic equivalent energy (\Eiso) called \amati correlation \citep{amati02,sakamoto04,lamb04,amati06,amati09}. \citet{yonetoku04} found a similar but tighter correlation between $E_p$ and 1-second peak luminosity called the \yonetoku correlation. These correlations are tight but they have large dispersions such as $\sigma_{int}=0.33$ in $\log L_p$ and $\sigma_{int}=0.37$ in $\log$\Eiso which can not be explained as statistical errors of $E_p$ , \Eiso and $L_p$ \citep{yonetoku10}. \citet{ggl04} found that $E_p$ tightly correlates with the collimation-corrected gamma-ray energy ($E_{\gamma}$). \citet{firmani06} proposed that adding the high signal time scale ($T_{0.45}$) to the \yonetoku relation reduces the dispersion of the correlation. This correlation is defined by using only prompt emission properties like \amati, \yonetoku correlations so that it seems to be promising tools to constrain the cosmological parameters. However, this correlation is not confirmed by later studies \citep{rossi08, collazzi08}. More recently, \citet{tsutsui09} found that the luminosity time ($T_L=E_{\rm iso} / L_p$) also improves both the \amati and \yonetoku correlations. These correlations were used to investigate the star formation history \citep{yonetoku04}, the reionaization epoch \citep{murakami05}, and the cosmological expansion history of the early universe \citep{takahashi,oguri,ghirlanda06,schaefer07, kodama08, liang08,cardone09, tsutsui09}. However, in spite of high correlation coefficients, there have been many cautions to use these empirical correlations for cosmology \citep{np04,bp05,butler,sn09}. To establish these correlations in GRBs prompt emissions as tools to determine cosmological parameters, we must investigate the origins of systematic errors and the way to remove them. We note that there are many factors to cause systematic errors besides intrinsic dispersions of their prompt emissions. For example the sensitivity of the detectors , the evolution effects of GRBs, the confusion with other sources, the lack of unknown parameters like the jet opening angle $\theta_{jet}$, etc. All of these effects might arise the additional systematic errors over the intrinsic dispersions of GRBs. Possible selection effects on these correlations are studied by many authors with contrasting results \citep{butler, ghirlanda08, nava08, sn09, amati09, yonetoku10}. However previous studies did not consider the difference of spectral models to determine $E_{p}$ well. As shown in \cite{kaneko06}, it often happens that high energy power-law index $\beta$ for the Band model with four free parameters can not be determined by the data so that the cutoff power-law (CPL) model with three free parameters is used to fit the data. CPL model might be good if the peak energy is close to the high energy end of the detector band width. One can not use Band model but CPL model if the event is so dim that the number of high energy photons is very small. Importantly, simulations in \cite{kaneko06} showed that, if the signal-to-noise ratio is relatively low , a true spectrum with the shape of the Band model can be fitted by CPL model with $E^{\rm obs}_p$ which is larger than the true value of $E^{\rm obs}_p$ up to $\sim 100~{\rm keV}$. Therefore CPL model might overestimate $E_p$. While , if we fit a true CPL spectrum by the Band model, the estimated value of $E_{p}^{obs}$ is almost equal to the true value since the large value of $-\beta$ looks like an exponential function. In reality \citet{sn09} found that $E_{p}$ estimated using the CPL model by \citet{kaneko06} are systematically harder than $E_{p}$ estimated using the Band model by \citet{yonetoku04}. Although the systematic difference between the peak energies fitted by the Band model and the ones fitted by the CPL model are reported, how this difference affect the spectral-brightness correlations of GRBs has been hardly studied so that we shall study this problem in this paper. The purpose of this letter is to investigate the effect of uncertainty in using different spectral models which determine $E_p$ on the \amati , \yonetoku and \tsutsui correlations, using our database developed in \cite{yonetoku10}. We examine this model bias by dividing the samples into two data sets as gold and bronze according to the quality of spectral observation. In this paper, we assume, if signal-to-noise ratio is high enough, all of the spectrum of GRBs are well expressed by the Band function, The structure of this letter is as follows. First we describe our database of 109 GRBs with known redshift and well-determined spectral parameters, 1-second peak luminosity, and isotropic energy in section~\ref{sec:data}. We construct the \tsutsui, \yonetoku, and \amati correlations with only gold data set in section~\ref{sec:correlations}. Finally we will give summary in section~\ref{sec:summary}.

\label{sec:summary} In this paper, using database constructed by \citet{yonetoku10}, we examine the model bias, that is, Band or CPL, on \tsutsui, \yonetoku and \amati correlations. We found that GRBs with the peak energies fitted by the CPL model are distributed in systematically harder and/or dimmer side of the \tsutsui, \yonetoku and \amati correlations than the ones by the Band function. There might be two interpretations about this result. The first is that these correlations have much larger intrinsic dispersion than that of observed one. If we had the more sensitive detector and could observe dimmer GRBs, the dispersion of the relations would become larger\citep{butler,sn09}. Another is that the use of the CPL model to estimate the peak energies causes this systematic difference. As simulated by \citet{kaneko06}, the Band function spectrum is well fitted by the CPL model if the detector does not have enough sensitivity to observe the high-energy photons. However the peak energies fitted by the CPL models are always higher than that of the simulated Band function spectrum \citep[see table.~3 in ][]{kaneko06}. Thus, it seems to be natural to conclude that the latter is more acceptable. In short, using only the peak energies determined by the Band function, we would get tighter correlations. If we could have much more GRBs by which we can uniformly analyze the data with the Band function, GRBs would be more powerful tool to constrain cosmological parameters. We found seven outliers in our gold data set. We classify these outliers in two classes as \begin{enumerate} \item [](dimmer and/or harder) 980425, 980613, 090328, 091003 \item [](brighter and/or softer) 000131, 091020, 091127 \end{enumerate} Although we do not know how and why these outliers are different from ordinary GRBs except for the distribution in the \tsutsui space, the effect of eliminating these GRBs is obvious, that is, the correlation becomes tighter. To find the characteristics which distinguish these outliers from ordinary GRBs is urgent. We here point out possible origins of these outliers. Let us assume that if we observe the jet nearly on axis , \tsutsui correlation would be very tight. However if we observe the jet with a certain viewing angle, we might have some dispersions on the observed \tsutsui correlation. In other words, we might not avoid some dispersion in the \tsutsui correlation from viewing angle, especially when the observer locates near the edge of the jet of GRBs. If we will know how to distinguish these outliers from the ordinary gold GRBs, the \tsutsui correlation might be much tighter and very useful in determining the nature of dark energy in redshift larger than $\sim$3. We should note that even in the Period-Luminosity relation of Cepheid variable there are $\sim$ 10\% outliers \citep{Riess09} so it is not surprising that there are $\sim$ 20\% outliers in the \tsutsui relation. \citet{butler}, using the Bayesian approach to estimate $E_p$, indicated that dim events close to the detector sensitivity would make large scatter on the $E_p$--$E_{\rm iso}$ and $E_p$--$L_p$ relations and that there is a significant threshold effect. Thus, they conclude that the \amati correlation have larger intrinsic dispersion than observed if we do not suffer from a threshold effect. Recently, \citet{sn09} argue that using hardness ratio instead of $E_{p}$ they also find that $E_{p}^{obs}$-Fluence correlation become more wider if we will be able to determine $E_{p}$ of dimmer events. However, there is a possible bias by using different method to estimate $E_p$. Even the difference of $E_{p}$ between the Band and CPL models causes the systematic errors so that using the other method to estimate $E_{p}$ might cause the additional systematic errors. The smaller the intrinsic dispersion of the relation is, the more the correlation suffers from these systematic effects. This might be why the $E_{p}$--$T_{0.45}$--$L_{p}$ relation is not confirmed by later studies \citep{firmani06, rossi08,collazzi08}. Kaneko et al. (2006) suggested the $E_{p}$ value of CPL function becomes systematically higher than the one of the Band function. If the 65 bronze data previously analyzed by the CPL function are reconsidered by the Band function with the fixed $\beta$ as an average value of $-2.25$ (Preece et al. 2000), they might show the distribution around the best fit line of each correlation estimated with 41 gold data set. They have a good potential to become a "silver" data set. To do so, we need help from each instrument team, and this is a future work. Finally, we note that there would be many reasons which cause systematic errors on the correlation in addition to intrinsic property of GRBs. These systematic errors must be carefully estimated and removed from the correlation analysis one by one. If we will finish it, GRBs become more powerful and unique standard candles to investigate the nature of the dark energy at high redshift larger than $\sim$ 3.

1012.3009

We reconsider correlations among the spectral peak energy ($E_p$), 1-second peak luminosity ($L_p$) and isotropic energy (\Eiso), using the database constructed by \citet{yonetoku10} which consists of 109 Gamma-Ray Bursts (GRBs) whose redshifts are known and $E_p$, $L_p$ and \Eiso are well determined. We divide the events into two groups by their data quality. One (gold data set) consists of GRBs with peak energies determined by the Band model with four free parameters. On the other hand, GRBs in the other group (bronze data set) have relatively poor energy spectra so that their peak energies were determined by the Band model with fixed spectral index (i.e. three free parameters) or by the Cut-off power law (CPL) model with three free parameters. Using only the gold data set we found the intrinsic dispersion in $\log L_p$ ($=\sigma_{\rm int}$) is 0.13 and 0.22 for \tsutsui correlation ($T_L \equiv E_{\rm iso}/L_p$) and \yonetoku correlation, respectively. We also find that GRBs in the bronze data set have systematically larger $E_p$ than expected by the correlations constructed with the gold data set. This means that the intrinsic dispersion of correlations among $E_p$, $L_p$, and \Eiso of GRBs depends on the quality of data set. At present, using \tsutsui correlation with gold data set, we would be able to determine the luminosity distance with $\sim 16%$ error, which might be useful to determine the nature of the dark energy at high redshift $z &gt; 3$.

12201862

2011ASPC..448..255R

1012

1012.1183_arXiv.txt

The Sun and cool stars are known to harbor magnetic fields leading to all phenomena summarized under the term \emph{stellar activity}. It may be debatable whether magnetic fields are actually the most interesting aspect of cool star and solar physics \citep[cp.][]{1985ARA&A..23..239M}, but it is certainly an exciting field that brings together a large variety of physical mechanisms and subtle analysis techniques. This makes it sometimes difficult to interpret observational results and compare them to theoretical expectations -- even if both are available, or even if one compares observations achieved from different techniques. A particularly interesting class of stars are cool stars of spectral type M. Covering the mass spectrum between $\sim 0.6$ and $0.1\,$M$_{\odot}$, M dwarfs are the most frequent type of stars. Within this mass range, objects can have very different physical properties. The very important transition from partly convective (sun-like) to fully convective stars happens in the M dwarf regime, probably around spectral type M3/M4. Furthermore, atmospheres of M dwarfs can be very different and both molecules and dust gain importance as the temperature drops toward late spectral types. In this article, I will concentrate on measurements of magnetic fields on M dwarfs because most of the currently available measurements of cool star magnetic fields are from M dwarfs, which is because sun-like (field) stars tend to be less active (less rapidly rotating because of shorter braking timescales) implying lower average magnetic fields that are more difficult to detect.

Our knowledge on magnetic fields in cool stars, particularly in M stars across the full convection boundary, has seen enormous progress during the last few years. Intensive observations of many M dwarfs led to the construction of Stokes~V Doppler maps, and the exploitation of the FeH molecular spectra allow a determination of the entire field from Stokes~I. We can now start to compare results from independent methods and search for the influence of stellar parameters including convective nature. The interpretation of results from different methods opens a parameter space that certainly contains deep information about the fields and their topology, but it is not yet clear what our measurements are actually telling us. Field strengths and topologies have ramifications to a broad range of astrophysics, and at spectral type late-M, we are approaching the brown dwarf regime. More field measurements, determination of molecular constants, and fundamental investigation of detectabilities are required to push the field forward, and to understand the many facets of magnetic fields in cool stars, brown dwarfs, star formation, and their links to exoplanets.

1012.1183

Magnetic fields are an important ingredient to cool star physics, and there is great interest in measuring fields and their geometry in order to understand stellar dynamos and their influence on star formation and stellar evolution. During the last few years, a large number of magnetic field measurements became available. Two main approaches are being followed to measure the Zeeman effect in cool stars; 1) the measurement of polarized light, for example to produce magnetic maps, and 2) the measurement of integrated Zeeman broadening to measure the average magnetic field strength on the stellar surface. This article briefly reviews the two methods and compares results between them that are now available for about a dozen M-type stars. It seems that we see a great variety of magnetic geometries and field strengths with typical average fields of a few kG in active M-type stars. The interpretation of geometries, however, has not yet led to a clear picture of magnetic dynamos and field configuration, and work is needed on more observational data but also on the fundamental understanding of our measurements.

12168141

2011ApJ...729...11K

1012

1012.4584_arXiv.txt

In the LCDM model, dark matter structures grow via merging and accretion. Within these dark matter halos, baryons initially agglomerate onto the centre of the potential well within the free-fall time of the halo and form stars \citep{rees77,binney77,silk77}. When the halo becomes large enough, the accreted gas can be gravitationally shock-heated and turn into diffuse hot halo gas. In the meantime, the constant attraction of gravity causes haloes and their galaxies to interact and merge, forming clusters and groups of galaxies. While orbiting within a host halo, satellite galaxies are likely to interact with other galaxies and the hot ambient gas \citep[e.g.][]{chung07,yagi10}. The evolution of the cold and hot gas contents is heavily influenced by the details of these interactions, and the exact mechanism is not well understood. It is known that galaxies in denser environments are optically redder and more quiescent than field galaxies \citep[to cite a few][]{gisler78,larson80,dressler80}. This strongly suggests that the cold gas supply (the source of star formation) is controlled by mechanisms closely associated with the environment. \citet{gunn72} demonstrated that the cold gas of a galaxy can be stripped off due to the ram pressure exerted on it by the intracluster medium as a galaxy moves within a cluster potential \citep[see also][]{abadi99,quilis00,chung07,tonnesen09,yagi10}. Tidal stripping is another important process that indirectly affects the evolution of cold gas. Satellites trapped by a larger cluster halo are tidally stripped off their dark haloes and hot gas reservoir. Throughout their orbital motions, satellites lose their interstellar/diffuse halo gas through shock heating, ram pressure and tidal stripping of their dark halo. Consequently, star formation in satellites may decreases with time. Semi-analytic models (SAMs) of galaxy formation assume that the shock heating is very efficient, and it is generally assumed to {\it instantly} remove diffuse halo gas from satellite haloes \citep[e.g.][]{kauffmann99}. Since the hot gas reservoir is removed instantly, gas cooling stops, cold gas is quickly depleted through star formation and SN feedback, and eventually star formation is quenched. Such models predict that the bulk of satellite galaxies in large haloes should be virtually passive, with only about 20\% predicted to be active \citep[e.g.][hereafter K09]{kimm09}. However, this prediction is not supported by observation. Studies based on the {\em specific star formation rates} measured from {\it GALEX} and {\it SDSS} observations \citep{salim07} have found that a much larger fraction of satellite galaxies are active \citep{kimm09}. This is sometimes referred to as the {\it satellite over-quenching problem} \citep[e.g.][]{weinmann06,baldry06,vdb08,gilbank08,fontanot09,kimm09}. The assumption of instantaneous shock heating of satellite gases is thought to be the most likely cause of this mismatch between model and data. Observationally, there is evidence for hot gas around satellite galaxies, as shown by {\it Chandra} observations of early-type galaxies in nearby clusters \citep{sun07,jeltema07}. These observations are bolstered by recent hydrodynamical simulations suggesting that a non-negligible fraction (up to $\sim$ 40\%) of diffuse hot gas remains associated with satellite galaxies for several Gyrs after merging with the host halo \citep{bekki02,kawata08,mccarthy08}. Various SAMs have included the effect of ram pressure stripping \citep{khochfar08,kang08,font08,guo10}. In a recent attempt \citet{font08} showed that inclusion of ram pressure stripping can help improve the bimodal colour distribution of satellite galaxies and the dependence of satellite colours on galaxy environment \citep[see also][]{kang08}. However, they did not take into account the decay of the satellite orbits due to dynamical friction, and hence their calculations may slightly overestimate the amount of hot gas retained in the satellite systems. Further, as we will show, the instantaneous recycling approximation for stellar mass loss used in the model may change the recent star formation history of galaxies. Alternatively, \citet{weinmann10} propose a simple {\em gradual} recipe for the stripping of gas and dark matter that reproduces the specific star formation rates as a function of clustocentric radius. Another factor that can contribute to the interstellar medium is stellar mass loss \citep[e.g.][]{bregman09}. Stellar masses vary widely and their lifetimes depend strongly on their mass. Thus, the stellar mass loss of a population is not instantaneous but changes gradually with time (though the overall change can be dramatic). Yet, most SAMs (\citealt{hatton03} is a notable exception) have used a simple approach in which roughly 30\% - 40\% of newly formed stars {\em instantaneously} evolve off and are recycled to cold gas \citep[e.g.][]{somerville08}. This may be a reasonable approximation for long-timescale star formation, but it may not be adequate for investigating short-timescale phenomena, such as the recent star formation history of satellite galaxies. In this study, we aim to investigate the effects of the two physical processes mentioned above in the context of SAMs. We first consider the orbital motion of a satellite under the influence of dynamical friction and the associated tidal stripping, both of which have been neglected in the previous studies based on EPS \citep[e.g][]{lacey93} formalism\footnote{SAMs based on N-body merger trees can follow the tidal stripping of dark matter subhaloes as long as they are robustly identified. Beyond that point, satellites are tracked based on analytic formalism.}. We also include a prescription for ram pressure stripping (\S 2.1). Next, in \S 2.2, we present further improvements on the earlier implementations of stellar mass loss. We finally describe the impact of these considerations on satellite galaxy evolution in \S 3 and discuss the implications of this in \S4.

We have investigated the recent star formation history of satellite galaxies by comparing semi-analytic models of galaxy formation with empirical data drawn from SDSS and GALEX. Based on the star formation rate measurements derived from multiband photometry \citep{salim07}, we first divided galaxies into active and passive types, and computed the fraction of passive galaxies (\fpass) as a function of galaxy stellar mass. As already shown in previous studies \citep{weinmann06,baldry06,kimm09}, satellite galaxies in theoretical models cannot reproduce the observed level of recent star formation activity. The satellite over-quenching problem is generally attributed to the strong strangulation applied in most semi-analytic models in which diffuse gas of the satellite system is instantly shock-heated from the system at the beginning of halo mergers \citep[e.g.][]{white91}. In the hope of improving the situation, we implement {\em gradual diffuse gas stripping} and {\em stellar mass loss} into our semi-analytic code. For detailed tests, we divide the sample by morphology. We have also introduced an approximation for the contribution to the ISM from stellar mass loss. Our results can be summarised as follows. \begin{itemize} \item The over-quenched satellite galaxies in SAMs are mostly late-type. \item The models with gradual diffuse gas stripping resolve much of the satellite over-quenching problem. However, they cannot account for the significant fraction of actively star-forming small late-type galaxies observed in massive haloes. \item Stellar mass loss is not a dominant source of cold gas in most satellite galaxies but enhances residual star formation. \item Our new models incorporating both effects show a significantly-improved match to the observed data. However, they still suggest that the majority of cluster late types are passive. \end{itemize} Recent SAMs including ours have shown steady progress in matching the observed data, but are missing key ingredients. For example, feedback from supernova explosions affects the recent star formation history in satellite galaxies by regulating the remaining cold disc gas but is still poorly constrained (\S 4). One might wonder whether or not simply increasing the contribution of stellar mass loss to the cold ISM can reproduce the passive late-type fraction. We have found that such an approach still shows a lack of actively star-forming late types. \citet{kaviraj07} reached a similar conclusion when trying to explain early-type galaxies with residual star formation \citep{yi05}. Recently, \citet{tonnesen09} demonstrated that a {\em weak} ram pressure does not only remove cold gas in satellites but also may enhance star formation by compressing the gas component. Provided that the cooling of cold gas onto a galaxy disc is suppressed during the interaction between the host and satellite haloes and the satellites are affected by the weak ram pressure for a long time, this may increase \fpass\ by consuming a significant fraction of cold gas. Yet, it is still unclear how gas cooling takes place in massive satellite systems during halo merging. Since cooling also relies on the density, if ram pressure compresses both cold and diffuse gas, cooling could also occur more efficiently. If this is the case, the efficient cold gas consumption might have little effect on \fpass. Detailed numerical simulations are necessary to better understand the supply of cold gas during interactions. An accurate determination of initial orbits may have an impact on \fpass. Based on numerical simulations at $z\sim0$ \citep{benson05,zentner05,khochfar06b}, we have assumed that satellite galaxies have random circularities at the time of halo mergers over the entire cosmic history. If radial orbits are more common at higher redshifts \citep[e.g.][]{wetzel10,dekel09}, gas stripping due to ram pressure would be more effective because satellites penetrate dense intra-cluster media more frequently on radial orbits. Motivated by this idea, we performed a simple experiment: we assigned an eccentricity of 1 to the haloes that experience halo merger at $z\geq 1$. In later mergers, orbital parameters were randomly chosen in terms of circularity, as done in the fiducial model. Since massive galaxies are likely to orbit several times while smaller galaxies have passed the pericentre less, the influence of ram pressure on radial orbits is more notable in massive satellite galaxies. As a result, the negative correlation of \fpass\ with \mgal\ shown by our new late-type galaxy models in massive haloes gets diminished, making the models match the data slightly better. However, the reliability of our exercise depends on the validity of the demarcation redshift and actual eccentricity distribution assumed. We also note that ignorance of the tidal disruption could have an impact on the passive galaxy fraction. \citet{taylor04} and \citet{zentner03} showed that subhaloes lose 30--40 per cent of their mass per pericentric passage. This implies subhaloes are dissolved into host haloes after several orbits unless they merge with centrals. On this basis, \citet{somerville08} implemented a prescription that satellites are disrupted when the subhaloes lose $\sim90\%$ of their initial mass. By doing so, they reproduced the luminosity and the radial distribution of Milky Way satellites \citep{maccio10}. \citet{henriques08} also present good matches to galaxy colours using a simple assumption that satellites which are not associated with subhaloes in the Millennium dark matter simulation \citep{springel05} are already disrupted by the tide in their environment. Yet, as discussed in K09, the \citet{somerville08} models still show a lack of actively star forming small satellites, implying that ignorance of tidal disruption is not the primary cause of red and dead small (late-type) satellites. Nevertheless, since it is small red and dead galaxies that are preferentially disrupted by tidal forces, the \citet{somerville08} models exhibit a slightly better agreement with empirical data than other semi-analytic models without tidal disruption (see K09 for details). In this regard, the passive galaxy fraction in the low-mass regime could be even smaller, resulting in a weaker dependence of the passive fraction on galaxy mass. Recent semi-analytic models often adopt AGN feedback to prevent hot gas cooling in massive galaxies \citep[i.e.][]{croton06}. The effect has been particularly important for ``central'' galaxies that are fuelled by ongoing cooling, in contrast to satellite galaxies that cannot retain their hot halo gas in most SAMs. However, the strong strangulation appears problematic in that even late types suffer from the lack of cold gas. For that reason, we demonstrate that external as well as internal gas supply are necessary to match the observed fraction of passive galaxies. Allowing the external supply has an interesting implication. Since the gas retained in the hot halo could funnel into black holes \citep[e.g][]{croton06}, AGN feedback is likely triggered, and possibly suppresses star formation activity in satellite galaxies with supermassive black holes. On the other hand, the effect of AGN feedback may be negligible in late-type galaxies because their black holes are not massive enough to release a large amount of energy. Instead, the evolution of late-type galaxies in clusters or groups may be more likely driven by the combination of stellar mass loss and environmental effects such as tidal and/or ram pressure stripping. We have indeed shown to first order that both mechanisms could reproduce observed levels of recent star formation, but it is still unresolved how late-type satellite galaxies show similar fractions of passive galaxies over various galaxy stellar masses for a given halo. Explaining these observations with realistic assumptions is an important challenge, especially since most of the galaxies in the universe are satellites.

1012.4584

Current semi-analytic models of galaxy formation overpredict the fraction of passive small late-type satellite galaxies in dense environments by a factor of two to three. We hypothesize that this is due to inaccurate prescriptions on cold gas evolution. In the hope of solving this problem, we apply detailed prescriptions on the evolution of diffuse hot gases in satellites and on stellar mass loss, both of which are critical in modeling cold gas evolution. We replace the conventional shock-heating motivated instant stripping with a realistic gradual prescription based on ram pressure and tidal stripping. We also carefully consider stellar mass loss in our model. When both mechanisms are included, the fraction of passive late types matches the data much more closely. However, the satellite over-quenching problem is still present in small galaxies in massive halos. In terms of the detectable residual star formation rates, gradual diffuse gas stripping appears to be much more important than stellar mass loss in our model. The implications of these results and other possibilities, such as redshift-dependent merging geometry and tidal disruption, are also discussed.

12213783

2011MNRAS.414.3617K

1012

1012.3471_arXiv.txt

Cold Dark Matter (CDM) is one of the main constituents of the current standard model in cosmology, $\Lambda$CDM. N-body simulations based on the CDM paradigm predict dark matter haloes best-described by a steep (``cuspy'', ``NFW'') power law mass density distribution \citep[e.g.][]{Navarro96a}. While the exact slope of the inner density distribution remains under debate \citep[e.g.][]{Moore99,Graham06,Stadel09,Navarro10}, simulations all agree on producing a cuspy profile with a slope of $\alpha \sim -1$ in the central part of the halo \citep[see the review of][]{deBlok10}. Observed rotation curves of galaxies typically need a massive dark matter halo with a nearly constant density core in order to describe the data well. Ideal objects for observationally studying the ``cusp-core problem'' are low surface brightness (LSB) galaxies which are late-type, gas-rich, dark matter-dominated disc galaxies. Observations using long-slit spectra in H$\alpha$ \citep[e.g.][]{deBlok02} or high-resolution two-dimensional velocity fields \citep[][hereafter K06 and K08, respectively]{Kuzio06,Kuzio08} show that a cored dark matter distribution (described e.g.~by an pseudoisothermal sphere profile) provide a better fit to the data than a cuspy NFW dark matter profile. In the few cases where NFW models do reasonably fit the rotation curve data, the predicted values for concentrations $c_{200}$ are very low and the circular velocity $V_{200}$ at the virial radius very high, inconsistent with the cosmological $(c_{200},V_{200})-relation$, \citep[see the review of][]{deBlok10}. Most observational problems such as pointing errors, centering offsets and non-circular motions are likely too small to cause cusps to be mistaken for cores in the observations \citep{deBlok10}, but halo triaxiality has been claimed to explain the presence of cores. The simulation works of \citet{Hayashi07} and \citet{Bailin07}, for example, showed that due to systematic non-circular motions introduced by a triaxal halo, cusps could be interpreted for cores if long-slit rotation curves were used. Similarly, bars inducing non-circular motions combined with projection effects can create the illusion of a constant-density core in a circular velocity analysis \citep []{Valenzuela07}. Pressure support in the rotating gas of low mass galaxies could mask the underlying dark matter distribution, but the importance of this effect remains uncertain \citep []{Dalcanton10}. \defcitealias{Kuzio06}{K06} \defcitealias{Kuzio08}{K08} The aim of this work is to investigate whether cusps in (LSB) dark matter haloes could still be mistaken for cores if high-resolution two-dimensional velocity fields data are available \citep[as e.g.~in][]{Kuzio09}. In order to do so, we use self-consistent high-resolution simulations of (LSB) galaxy formation and ``observe'' the simulated galaxies using the DensePak Integral Field Unit. We compare the mock velocity fields for different signatures from cuspy and cored haloes, spherical and triaxial potentials, and the effect of supernovae feedback. We also determine how well the underlying dark matter halo potential can be recovered. The outline of the paper is as follows. In Section 2, we present the modeling of our initial conditions, the code used for the time evolution and the halo parameters derived from the simulations. In Section 3, we discuss the process of observing the simulations. The mock velocity fields and rotation curves are presented in Section 4. In Section 5, we present the halo fits to the mock data. A discussion and summary are presented in Section 6.

We have presented mock velocity fields, rotation curves, and halo fits for simulated LSB galaxies formed in spherical and triaxial cuspy dark matter haloes and spherical cored dark matter haloes. The mock velocity fields span a range of data quality representing ideal to realistic observations. The main findings of this work are: \begin{itemize} \item The underlying halo type produces a unique signature in the velocity field. This can be used to constrain the shape of the dark matter profile (spherical or triaxial cuspy, or spherical cored) without fitting an analytic density profile to it. \item Cored and cuspy haloes can also be distinguished clearly by deriving their asymptotic inner slopes from the rotation curve data. \item Given at least one of the above information one can then successfully recover the underlying halo parameters from the rotation curve using the appropriate analytic form for the density profile (NFW or pseudoisothermal sphere). \end{itemize} \textit{This means that if an LSB galaxy were in a cuspy halo, the cusp would be observable in the data.} Given these results, we find it difficult to mistake cuspy haloes, spherical or triaxial, for cored haloes. The observed cores in dark matter-dominated galaxies are true discrepancies from the predictions of (dark matter-only) $\Lambda$CDM simulations. Systematic effects, non-circular motions, and halo triaxiality cannot explain the observed differences. Baryonic processes more effective than those we have modeled or entirely different dark matter models may be necessary to explain the observed cores. Feedback from star formation could transform the initial cusps into cores. If LSB galaxies form with a massive dark matter halo already in place and with the same low baryon fraction that we observe today, then feedback from star formation of the type we have modeled would not be enough to change a cusp into a core. In this scenario, the dark matter halos in which LSB galaxies reside cannot be cuspy CDM halos. If, however, feedback processes from star formation can affect dark matter haloes during the early stages of their formation, it is possible that cusps can be transformed into cores. This has been demonstrated by \citet{Governato10} (see also \citet{Oh10}) using a set of fully cosmological simulations of dwarf galaxy formation ($M_{vir} \sim 3 \times 10^{10} \Msun$) where strong outflows from supernovae remove low-angular-momentum gas and decrease the dark matter density to less than half in the central part \citep[see also][]{Navarro96b,Read05,Mashchenko08}. The question in this scenario then becomes one of detecting the blown-out baryons. It also remains to be seen if this kind of mechanism can act efficiently on higher mass scales, given the deeper potential wells. Alternatively, the observed cores may need to be addressed by a different dark matter model/particle that naturally produces cored haloes. Recent results, however, have found that self-interactions and warm dark matter particle properties cannot be responsible for the cores observed in dark matter-dominated galaxies \citep[e.g.][]{Kuzio10,Villaescusa10}.

1012.3471

We present mock DensePak Integral Field Unit (IFU) velocity fields, rotation curves and halo fits for disc galaxies formed in spherical and triaxial cuspy dark matter haloes and spherical cored dark matter haloes. The simulated galaxies are 'observed' under a variety of realistic conditions to determine how well the underlying dark matter halo can be recovered and to test the hypothesis that cuspy haloes can be mistaken for cored haloes. We find that the appearance of the velocity field is distinctly different depending on the underlying halo type. We also find that we can successfully recover the parameters of the underlying dark matter halo. Cuspy haloes appear cuspy in the data and cored haloes appear cored. Our results suggest that the cores observed using high-resolution velocity fields in real dark matter dominated galaxies are genuine and cannot be ascribed to systematic errors, halo triaxiality or non-circular motions.

12167490

2011ApJ...728...16H

1012

1012.1368_arXiv.txt

Cataclysmic variables (CVs) are short-period binaries in which a late-type, Roche-lobe filling secondary star transfers matter through an accretion disk onto a rotating, accretion heated primary white dwarf (WD). The standard evolutionary paradigm \citep[hereafter HNR]{HNR} postulates that CVs evolve from wide binaries of moderate orbital period and unequal masses. As the more massive component evolves off of the main sequence into a red giant, the secondary star finds itself orbiting within the atmosphere of the massive star. During this common envelope phase, the orbit of the secondary star shrinks due to interactions with the atmosphere of the more massive primary star. This shortens the binary period, until the common envelope is ejected around orbital periods of 1 day. Angular momentum is then mainly lost through an efficient magnetically constrained wind, \citep[``magnetic braking'', see][and references therein]{Cameron2002}, shrinking the orbit so much that the Roche lobe of the secondary star comes into contact with the stellar surface and mass transfer begins, signaling the birth of a long period CV. As the system evolves, the secondary star continues to lose mass, but angular momentum losses keep the secondary star in contact with its Roche lobe. During their lives as longer period CVs, orbital periods typically range from 3 to 10 hours, and mass transfer rates from $10^{-8}$ to $10^{-9}$ \msun yr$^{-1}$. This rapid mass transfer drives the secondary star out of thermal equilibrium, causing it to become bloated by $\sim$30\% compared to an isolated star at the same mass \citep{Knigge}. The secondary continues to lose mass until it reaches the fully convective boundary ($P_\mathrm{orb}\approx3$ hrs, $M_\mathrm{2}\approx0.2-0.3$ \msun), where the magnetic breaking is believed to be disrupted. The secondary star is able to regain its thermal equilibrium, and as a result shrinks and loses contact with its Roche lobe ceasing mass transfer. With the disruption of magnetic breaking from the secondary star, angular momentum in the system is lost via gravitational radiation alone, causing the orbit to decrease at a much slower rate. This continues until the orbit has shrunk sufficiently for the secondary star to overflow its Roche lobe and begin mass transfer once again ($P_\mathrm{orb}\approx2$ hrs). With the absence of mass transfer between orbital periods of 2-3 hours, the systems are dormant, and consequently, difficult to identify. This forms a gap in the observed orbital period distribution of CVs between 2-3 hours. Once mass transfer resumes, the system emerges as a ``sub-gap'' CV. It is assumed that this evolutionary sequence happens on a short enough timescale that the secondary star does not undergo any significant nuclear evolution, and therefore should retain normal abundance patterns consistent with a main sequence dwarf. Observational evidence, however, shows a growing number of systems with apparent abundance anomalies seen in the UV and/or the IR; in the UV, this is seen as unusual \ion{N}{5}/\ion{C}{4} ratios \citep{Boris}, and in IR, this is inferred through the presence of anomalously weak or absent CO features. \citet{TomLongPeriod, TomPolars, TomShorts} show that for thirteen out of the nineteen systems observed, the CO features of non-magnetic CVs above the period gap are much weaker than they should be for their spectral types. Since the water vapor features in the coolest of these stars appears normal, this result points towards a deficit of carbon. In addition, $^{\rm 13}$C appeared to be enhanced for several systems, an indication of CNO processed material in the atmosphere of the secondary star \citep{TomLongPeriod}. In contrast to these non-magnetic CVs above the gap, the majority (8/11 systems) of ``polars'', CVs with highly magnetic WD primaries, have secondary stars that appear to be completely normal \citep{TomPolars, TomPolarsSpitzer}. Pre-CV systems appear uniformly normal, with only one of nineteen systems showing any weakened CO features \citep{clausprecv, SteveRZLeo}. While the total CV sample is somewhat small and limited to the brightest objects observable with ground-based near-IR spectroscopic instrumentation, these trends are striking enough to warrant further attention and examination of all CV subtypes. Observations of short period, sub-gap CVs are extremely challenging given the very low luminosities expected for the secondary stars that must compete against that of accretion and the underlying hot white dwarf. IR spectroscopy of CVs has been possible since the early 1990's, notably with the efforts of \citet{Dhillon95} looking at systems well above the 2-3 hour period gap, where absorption lines from the secondary were clearly seen. \citet{Dhillon2000} observed five systems below the period gap, and found that the secondary stars were too faint to detect and estimated that they contribute only 10 to 30\% to the observed infrared flux. \citet{Mennickent} and \citet{MennickentRZ} focused on lower resolution studies of sub-gap systems, fitting K or M dwarf template spectra to low resolution optical and near-IR spectra. This technique was employed by \citet{pasjshortcvs} who used the Subaru Telescope to obtain $J$, $H$, and $K$-band low resolution grism spectroscopy of five CVs below the gap. Spectral component fitting allowed them to obtain rough spectral type estimates ranging from M1 to L1, but the low resolution of these data (FWHM $\sim 60$\AA) prevented examination of the CO features. To attempt to detect the secondary stars in a sample of sub-gap CVs, we have obtained moderate resolution $K$-band spectra of nine systems. This increases the sample of CVs with moderate resolution (R $\gtrsim$ 1500) near-IR spectroscopy to sixty-one systems: nineteen pre-CVs, thirty-one non-magnetic systems, and eleven magnetic or partially magnetic systems. Prior to this work, three quarters of the non-magnetic systems were {\it above} the period gap and only four below: RZ Leo, WZ Sge, GW Lib, and EI Psc. \citet{SteveRZLeo} presented the $K$-band spectrum of RZ Leo, which showed no evidence of a C deficit manifested as weakened or absent CO features. \citet{TomVYEI} obtained moderate resolution observations VY Aqr and EI Psc, finding evidence for strong deficits of C in both systems. \citet{SteveWZSge} presented observations of WZ Sge, showing this object to have both CO and H$_2$ emission from its accretion disk, the only such detection so far. In our nine observed systems, we clearly detect the presence of the secondary star in six of them. In two of the remaining cases, we are able to provide constraints on their spectral types, and in only one system we could not detect the signature of the secondary star. In contrast to the longer period CVs, the majority of these secondary stars have CO features that appear to be present at near-normal strengths. We present our observations in \S\ref{obs}, the object spectra and spectral type determinations in \S\ref{results}, a discussion of our results in \S\ref{disc}, and our conclusions in \S\ref{conc}.

\label{conc} We have performed a near-IR spectroscopic moderate resolution survey of nine sub-gap CVs, detecting a signature of the secondary star in eight systems. We demonstrate an important link between abundance anomalies seen in the UV and in the IR, as well as the reverse cases where the lack of abundance anomalies appear. Our clear detections show that future phase-resolved spectroscopy will further constrain the nature of these secondary stars by obtaining radial velocity measurements and therefore masses of the secondary stars in these systems. While our results indicated a large fraction of sub-gap CVs contain a normal CO abundance, we have insufficient objects to fully test this idea at present. Ongoing analysis and modeling of these data presented here will allow us to explore the CO bands seen in our spectra. To fully assess the $^{12}{\rm C}/^{13}{\rm C}$ ratio which gives the best indication and possibility to see if the secondary star contains any CNO processed material \citep{COCNOref}, but will require higher quality data. We thank the anonymous referee for their useful and helpful comments. This work is based on observations made with ESO Telescopes at the Cerro Paranal Observatory under program ID 081.D-0225. Some of the data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation. The authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community. We are most fortunate to have the opportunity to conduct observations from this mountain. This publication makes use of data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation. RH would like to thank the New Mexico Space Grant consortium for Graduate Research Fellowship support from 2007-09. {\it Facilities:} \facility{VLT:Antu (ISAAC)}, \facility{Keck:II (NIRSPEC)}

1012.1368

We present K-band spectroscopy of short-period, "sub-gap" cataclysmic variable (CV) systems obtained using ISAAC on the Very Large Telescope. We show the infrared (IR) spectra for nine systems below the 2-3 hr period gap: V2051 Oph, V436 Cen, EX Hya, VW Hyi, Z Cha, WX Hyi, V893 Sco, RZ Leo, and TY PsA. We are able to clearly detect the secondary star in all but WX Hyi, V893 Sco, and TY PsA. We present the first direct detection of the secondary stars of V2051 Oph, V436 Cen, and determine new spectral classifications for EX Hya, VW Hyi, Z Cha, and RZ Leo. We find that the CO band strengths of all but Z Cha appear normal for their spectral types, in contrast to their longer period cousins above the period gap. This brings the total number of CVs and pre-CVs with moderate resolution (R &gt;~ 1500) IR spectroscopy to 61 systems: 19 pre-CVs, 31 non-magnetic systems, and 11 magnetic or partially magnetic systems. We discuss the trends seen in the IR abundance patterns thus far and highlight a potential link between anomalous abundances seen in the IR with the C IV/N V anomaly seen in the ultraviolet. We present a compilation of all systems with sufficient resolution IR observations to assess the CO band strengths and, by proxy, obtain an estimate on the C abundance on the secondary star.

12102024

2010AIPC.1314..157H

1012

1012.3647_arXiv.txt

Globular clusters (GCs) host a variety of binary star systems, formed both primordially and dynamically via stellar encounters. These binaries play a crucial role in the dynamical evolution of GCs, providing an energy reservoir that can delay core collapse for many times the half-mass relaxation time \citep[\eg,][]{Fregeau07}. The dense cluster environment also dramatically alters the evolution of GC binaries \citep[\eg,][]{Ivanova06,Fregeau03,Pooley06,Fregeau08}. X-ray-emitting systems have, in particular, emerged as a promising source of information about the history of binary formation and destruction in galactic GCs. The {\it Chandra X-ray Observatory}'s high spatial resolution and resulting sensitivity to point sources makes it possible to obtain nearly complete samples of compact accreting binaries in nearby globular clusters. The ability to pinpoint sources to $< 1$\asec\ also means that the stars responsible for the X-ray emission may be recovered at other wavelengths even in the crowded fields of GCs. While the high luminosity X-ray binaries ($L_x = 10^{36-38}$ erg\,s$^{-1}$) are understood to be accreting neutron stars \citep{Brown98,Heinke03a,Heinke03b}, the low X-ray-luminosity sources are now known to comprise several distinct populations: cataclysmic variables (CVs), quiescent neutron stars (qNS, or qLMXB), millisecond pulsars (MSPs), and binaries with chromospherically active stars, \ie\ active binaries \citep[ABs;][]{Pooley02b,Heinke05,Lugger07,Haggard09}. Of these, only the quiescent NSs, with their distinctive soft X-ray spectra, can be identified uniquely on the basis of X-ray observations alone \citep{Pooley02a,Rutledge02}. For others, optical (or radio, in the case of MSPs) follow-up is essential. \wcen\ is the most massive GC in the Milky Way \citep[$4 \times 10^6$\msun;][] {Meylan02}. At 4.9 kpc, it is relatively nearby, making it possible to detect low-luminosity X-ray sources in modest exposure times with \Chandra. Its unusually complex stellar populations have prompted debate as to whether \wcen\ is a GC at all --- it is instead likely to be the remnant of a dwarf galaxy accreted by the Milky Way \citep{Bedin04, Gratton04, Piotto05, Villanova07} --- and controversy continues over the existence of an intermediate black hole in its core, \eg\ see \citet{Noyola08} vs. \citet{Anderson10}. Regardless of \wcen's origins, the binary stars that it contains play a crucial role in its dynamical evolution and can in turn shed light on the impact that a cluster has on its binary population. Here we summarize results of our search for X-ray-emitting binary stars in \wcen\ using {\it Chandra X-ray Observatory} and {\it Hubble Space Telescope} (\HST). The \Chandra\ results have been reported by \citet{Haggard09} and the complete \HST\ results will appear in \citet{Cool10}.

Using a combination of \Chandra\ and \HST\ imaging in blue, red and \ha\ filters, we have identified a total of 40 X-ray-emitting binary stars in \wcen. Five were found as X-ray counterparts of variable stars reported by \citet{Kaluzny04} and four had been found in earlier \Chandra\ and/or \HST\ studies. The remaining 31 are newly reported here. Accreting binary stars make up just over half of the identifications: one qLMXB and 21 candidate CVs. The remaining identifications are evenly split into two broad classes: 9 active binaries and 9 objects that are possible sub-subgiants. The active binaries include an assortment of different types of systems, all of whose X-ray emission is likely due to active coronae. These include two eclipsing Algol systems, three possible BY Dra stars, and a blue straggler. Of particular interest among the CV candidates are the nine very faint blue stars shown as inverted triangles in Figures 1 and 2. Given the distance modulus to \wcen, their absolute magnitudes are in the range $M_{625} = 10.9-12.7$. This is comparable to the absolute magnitudes of the short-period CVs recently uncovered in the Sloan Digital Sky Survey \citep[SDSS;][]{Gansicke09}. Thus the systems in \wcen\ could be short-period systems with very low-mass secondaries, as is expected for very old CVs. Their positions in the X-ray CMD (see Fig. \ref{xray_cmd}) generally support the CV interpretation. Alternatively, some of these objects could be helium white dwarfs with MSP companions; a few such systems are known in globular clusters \citep[\eg\,][]{Edmonds01}. Deeper \ha\ imaging and/or multiwavelength broad-band imaging is needed to distinguish between these possibilities. Given that the faintest CVs detected in this study are at the detection limit in the both the X-ray and optical images, it is likely that more CVs remain to be discovered in \wcen. How many more depends on the relative numbers of faint vs.\ bright CVs --- a ratio that depends both on the evolution of CVs and their formation history in the cluster. The present study shows that even very faint CVs can be found in the crowded environs of \wcen; deeper observations should allow a more complete census to be made. The present census of active binaries in \wcen\ is undoubtedly very incomplete. At X-ray wavelengths we are likely seeing just the tip of the iceberg. In the optical, we are hampered by the weakness of the \ha\ emission lines. Observations with a narrower \ha\ filter can help, as demonstrated by the large number of BY~Dra stars identified by \citet{Taylor01} in NGC~6397 using a \about 20 Angstrom-wide filter. Perhaps the most intriguing set of X-ray-emitting stars identified in this study are the nine that lie redward of the turnoff and giant branch. While their close proximity to the evolutionary sequences in \wcen\ suggests that they are associated with the cluster, proper motions are needed to be sure. If spectroscopic observations reveal that some or all are members of the anomalous RGB, then it will be important to understand why this subpopulation in \wcen\ is prone to producing X-ray-bright stars. If instead these stars turn out to be bonafide SSGs, then \wcen\ should provide a valuable testing ground for studying this as-yet poorly understood class of X-ray-emitting binary systems.

1012.3647

We summarize results of a search for X-ray-emitting binary stars in the massive globular cluster ω Centauri (NGC 5139) using Chandra and HST. ACIS-I imaging reveals 180 X-ray sources, of which we estimate that 45-70 are associated with the cluster. We present 40 identifications, most of which we have obtained using ACS/WFC imaging with HST that covers the central 10'×10' of the cluster. Roughly half of the optical IDs are accreting binary stars, including 9 very faint blue stars that we suggest are cataclysmic variables near the period limit. Another quarter comprise a variety of different systems all likely to contain coronally active stars. The remaining 9 X-ray-bright stars are an intriguing group that appears redward of the red giant branch, with several lying along the anomalous RGB. Future spectroscopic observations should reveal whether these stars are in fact related to the anomalous RGB, or whether they instead represent a large group of ``sub-subgiants'' such as have been seen in smaller numbers in other globular and open clusters.

12137135

2010evn..confE..96R

1012

1012.5899_arXiv.txt

The EVN has aided to obtain the deepest and highest resolution radio images ever of one of the most distant ULIRGs in the local Universe. High resolution is important, but the information provided by short baselines is also necessary to properly map the morphology of the diffuse emission in this kind of sources.

1012.5899

The European VLBI Network (EVN) provides us with the necessary sensitivity and angular resolution to study the nuclear and circumnuclear regions in Luminous and Ultraluminous Infrared Galaxies. The high Star Formation Rates (SFR) inferred for these galaxies implies both the presence of a high number of massive stars and a dense surrounding medium. Therefore, bright radio SNe are expected to occur. With the aim of estimating the SFR in ULIRGs by means of Core Collapse supernova (CCSN) detections, we started an observing campaign with the EVN on a small sample of the brightest and farthest ULIRGs in the local Universe. We present here our results from three epochs of quasi-simultaneous observations with the EVN at 6 and 18 cm towards one of the objects in our sample: IRAS 23365+3604.

3038383

2011PhRvL.106p1103O

1012

1012.1853_arXiv.txt

1012.1853

We perform 3+1 general relativistic simulations of rotating core collapse in the context of the collapsar model for long gamma-ray bursts. We employ a realistic progenitor, rotation based on results of stellar evolution calculations, and a simplified equation of state. Our simulations track self-consistently collapse, bounce, the postbounce phase, black hole formation, and the subsequent early hyperaccretion phase. We extract gravitational waves from the spacetime curvature and identify a unique gravitational wave signature associated with the early phase of collapsar formation.

12205217

2011IAUS..270..309H

1012

1012.4192_arXiv.txt

Since stars are formed within the coolest molecular material of the interstellar medium (ISM), the star-formation rate (SFR) should be determined simply by the gas reservoir and by the free-fall time $\tff$ of molecular clouds (\cite{elm02}). This, however, raises a conflict between the ISM conditions and observed SFRs in the sense that $\tff$ for a typical molecular cloud density of 100 cm$^{-3}$ amounts to 10$^{14}$ sec, i.e. $3\cdot 10^6$ yrs. For the total galactic molecular mass of $10^9 - 10^{10}\, \Msun$ the SFR should then amount to about 100 to 1000 \sfr, what is by orders of magnitudes larger than observed and the gas reservoir within the Milky Way would have been used up today. This means, that the SF timescale must be stretched with respect to collapse or dynamical timescale by introducing a SF efficiency (SFE) $\eSF$ and its definition could read: $\tSF = \eSF^{-1} \cdot \tff$. Already in 1959 Schmidt argued that the SFR per unit area is related to the gas column density $\Sg$ by a power law with exponent $n$. Kennicutt (1998) derived from the $\Ha$ luminosity of spiral galaxies $\SHa$ a vertically integrated and azimuthally averaged SFR (in \sfr $pc^{-2}$), i.e. of gas disks in rotational equilibrium, and found a correlation with $n = 1.4\pm 0.15$ holding over more than 4 orders of magnitude in $\Sg$ with a drop below a density threshold at 10 $\Mpc2$. While this relation establishes an equilibrium SFR, it is not surprising that the slope varies for dynamically triggered SF, as in starburst and merger galaxies and for high-z galaxies, when the disks form by gas infall, for the latter reaching $n = 1.7 \pm 0.05$ (\cite{bou07}). The ordinary Kennicutt-Schmidt (KS) relation can be understood by the simple analytical assumption allowing for the $\tff - \tSF$ relation and for a uniform state of the ISM. Its equilibrium on disk scales requires that heating processes counteract to the natural cooling of plasmas. Besides the possible heating processes from dissipation of dynamical effects, as there are the differential rotation of the disk, gas infall, tidal interactions, shocks, etc., to the feedback by freshly produced stars, not for all of them it is obvious, how effectively they influence the SF by the expected self-regulation or vice versa trigger it. Unfortunately, the issue of a general KS-law is confused by the similarity of slopes under various stellar feedback strengths (see e.g. sect. 2 in \cite{hen08}). \cite{koe95} demonstrated already that the SFR achieves a dependence on $\rg^2$, if the stellar heating is compensated by collisional-excited cooling radiation (e.g. \cite{bh89}). The coefficient of this relation determines $\tSF$. Obviously, the SF in galaxies, therefore, depends on both the gas content and the energy budget of the ISM. Since the most efficient stellar energy power is exerted by supernovae (SNe), and here particularly by the explosions of the shortly living massive stars as type II SNe, their feedback to the ISM is of fundamental relevance for the SF. In this paper, we overview and discuss the regulation of the ISM by the SF feedback thru SNe and more pronounced by their cumulative effect as superbubbles. Although the expression {\it feedback} of SNe also includes their release of freshly produced elements, here we only focus on the dynamical and energetic issues and refer the reader interested on the chemical evolutionary consequences to \cite{hen10}.

The dominating influence of SN explosions and superbubbles on structure, dynamics, and energy budget of the ISM are obvious and agreed. Signs and strengths of these feedback effects are, however, widely uncertain. Whether the feedback is positive (trigger) or negative (suppression) can be understood analyticly from first principles, but because of the non-linearity and the complexity level of the acting plasmaphysical processes clear results cannot be quantified reliably. In addition, the temporal behaviour varies by orders of magnitude because of the changing conditions. In summary, the energy transfer efficiencies of SNe and superbubbles to the ISM are much below unity and depend on the temporal and local conditions, but must not be overestimated. Spatially and temporarily resolved simulations of SN and superbubbles in an extended environment with varying conditions of the ISM are necessary in order to connect large-scale effects on SF clouds with the existing detailed simulations on the star-forming scales.

1012.4192

Supernovae are the most energetic stellar events and influence the interstellar medium by their gasdynamics and energetics. By this, both also affect the star formation positively and negatively. In this paper, we review the development of the complexity of investigations aiming at understanding the interchange between supernovae and their released hot gas with the star-forming molecular clouds. Commencing from analytical studies the paper advances to numerical models of supernova feedback from superbubble scales to galaxy structure. We also discuss parametrizations of star-formation and supernova-energy transfer efficiencies. Since evolutionary models from the interstellar medium to galaxies are numerous and apply multiple recipes of these parameters, only a representative selection of studies can be discussed here.

12230850

2011PrPNP..66..329S

1012

1012.4647_arXiv.txt

A great success of nuclear astrophysics was the demonstration that radioactivity powers the light curves of Type Ia supernovae (SN~Ia, most likely thermonuclear) and, at least for some events at late times, Type Ib/c and Type II (most likely core collapse) supernovae. Shortly after its explosion a supernova enters a state of homologous expansion. The temperature drops during the expansion until nuclear fusion reactions cease; radioactive decays, however, still take place, since the matter is only slightly ionized. It is now widely accepted that the energy liberated in the decay chain of radioactive \nuc{56}{Ni} is the most important nuclear source for re-heating the supernova ejecta to temperatures high enough for the spectrum to peak at optical wavelengths \cite{truran1967a,colgate1969a}. At first, the bulk of the heating is produced by the energetic $\gamma$-rays which thermalize and deposit their energy via Compton scattering and photoelectric absorption. In the homologous expansion, the column density (and therefore also approximately the Compton opacity) decreases with time as $t^{-2}$, and the ejecta become more and more transparent to these high energy photons. Once $\gamma$-rays escape, the positrons produced in the decay of \nuc{56}{Co} and \nuc{44}{Sc} were thought to be the main heating sources. Here, we draw attention to often overlooked additional leptonic heating channels: Auger and internal conversion electrons. In section \ref{sec:decay} we review the physics of nuclear decays relevant for supernova light curves. In section \ref{sec:lc} we demonstrate the impact of internal conversion electrons from the decay of \nuc{57}{Co} on different supernova light curves. We conclude in section \ref{sec:conc} with an outlook how this effect could be used to constrain supernova explosion models in the future. For the published refereed journal article that first pointed out the significance of internal conversion electrons on supernova light curves, please see \cite{seitenzahl2009d}.

\label{sec:conc} Fitting models to observed late time light curves provides a unique and independent method to directly measure the {\it isotopic} yields of prominent radioactive nuclei synthesized in the explosion (in particular \nuc{56}{Ni}, \nuc{57}{Ni}, and \nuc{44}{Ti}). We have shown that at late times, when the ejecta have become largely transparent to $\gamma$-rays, the energy carried by Auger and internal conversion electrons may constitute a significant source of heating. These additional decay channels have to be considered for reliable isotopic abundance determinations from light curves. In particular, we have shown that a re-analysis of the bolometric light curve of 1987A (taking the hitherto unconsidered effect of internal conversion electrons into account) would likely yield significantly different (smaller) \nuc{44}{Ti} and \nuc{57}{Ni} masses. The new derived values for \nuc{57}{Ni} and \nuc{44}{Ti} would allow us to gain more insight into the explosion mechanism of core collapse supernovae. In particular, observationally driven inferences about the location of the mass cut can be made, which would give us valuable information about the physical processes separating neutron star and black hole formation from massive stars. Last but not least, SNe~Ia are considered as the source of the positrons required to explain the Galactic 511 keV annihilation line observed by Integral/SPI \cite{knoedlseder2005a}. The question whether enough positrons can escape the remnant remains unanswered. An understanding of the physics that underlies the light curves of SNe is a crucial step in order to constrain the escape fraction of positrons.

1012.4647

Radioactive decays contribute significantly to the re-heating of supernova ejecta. Previous works mainly considered the energy deposited by γ-rays and positrons produced by 56<SUP>Ni</SUP>Co, 57<SUP>Ni</SUP>Co, 44<SUP>Ti</SUP>Sc. We point out that Auger and internal conversion electrons constitute an additional heat source. At late times, these electrons can contribute significantly to supernova light curves for reasonable nucleosynthetic yields. In particular, the internal conversion electrons emitted in the decay of 57<SUP>Co</SUP>Ti yields from its late time bolometric light curve.

2138141

2011MNRAS.413.1192K

1012

1012.3251_arXiv.txt

Silicates are among the most commonly-found dust species in the interstellar medium {(ISM)} of galaxies. Their presence is established through the detection of the mid-infrared resonances due to the Si-O stretching and the O-Si-O bending mode at 9.7 and 18 $\mu$m respectively. In galaxies, these bands are seen in absorption \citep{GFM_75_HII,SLT_00_galaxies}, as well as in emission \citep{2005A&A...436L...5S,2005ApJ...625L..75H}. Most of these silicates show the broad resonances characteristic of amorphous silicates, i.e.,~silicates showing a large degree of lattice defects, and it is generally assumed that silicates in the ISM are predominantly amorphous. In particular, the degree of crystallinity $x$, defined as the mass fraction of silicates that is crystalline, $x=M_X/(M_X+M_A)$, in the Galactic diffuse ISM is found to be of the order of 1\% \citep{KVT_04_GC,KVT_05_erratum}. This contrasts sharply with the much higher degree of crystallinity seen in silicates in the circumstellar environments of pre- and post-main-sequence stars \citep[see e.g.][and references herein]{MK_05_2005}. Generally speaking, the galactic cycle of dust starts with its formation in evolved stars, followed by processing in the ISM and eventually ends with incorporation in stars and planets during star formation. The silicates observed around Asymptotic Giant Branch (AGB) stars can have significant crystalline fractions, in particular for the high mass-loss rate OH/IR stars \citep[up to $\sim$20\%][]{SKB_99_ohir,KWD_01_xsilvsmdot,2010A&A...516A..86D}. For lower mass-loss rate AGB stars, such as Miras, the crystallinity is not well established, but is consistent with a value that does not vary with mass-loss rate \citep{KWD_01_xsilvsmdot}. For more massive stars, such as Red Supergiants (RSGs), the crystallinity of the silicates in the stellar ejecta is not well known, although isolated studies report high crystalline fractions \citep[e.g.~][]{MWT_99_afgl4106}, {while low crystalline fractions seem to be more common \citep{2009A&A...498..127V}. \citet{2006ApJ...638..759S} adopt a crystallinity of 15\% for RSGs and Luminous Blue Variables.} In contrast to the low crystallinity in the Galactic ISM, significant amounts of crystalline silicates have been detected in the infrared spectra of Ultraluminous Infrared Galaxies \citep[ULIRGs;][]{2006ApJ...638..759S}. In a sample of 77 ULIRGs, 12 were found to show crystallinity, with crystalline-to-amorphous silicate mass ratios ranging from 0.07 to 0.15, corresponding to crystallinities of 6.5\% to 13\%. The degree of crystallinity of a population of silicate grains provides a record of the processing history of those grains \citep[see e.g.~][]{MK_05_2005}. A high crystalline fraction points to a relatively high formation or processing temperature ($\sim 1000 - 1500$ K), while a large amorphous fraction indicates that the population of grains has undergone the damaging effects of cosmic ray hits \citep{2007ApJ...662..372B}, grain-grain collisions or atomic impacts in shocks \citep{DCL_01_He+,BSB_03_amorphisation,JFS_03_bombardment}; or that it is formed at lower temperatures. Crystalline silicates may thus form in dense circumstellar environments, where the vicinity of the central star provides the required heating. The mass-loss processes of evolved stars subsequently spread these crystalline silicates into the ISM. The fact that the silicates in the ISM of the Milky Way are almost entirely amorphous \citep{KVT_04_GC} suggests that the amorphisation processes in the ISM are more important than the injection of fresh crystalline silicates into the interstellar reservoir by evolved stars. For the Galaxy, an amorphisation time scale of 40 Myr has been derived from observations \citep{KVT_04_GC,KVT_05_erratum}, which is close to the experimental value of 70 Myr \citep{2007ApJ...662..372B}. {\citet{2006ApJ...638..759S} argue that in starburst galaxies ultraviolet (UV) photons from either the Active Galactic Nuclei (AGN) or massive stars are the only potential sources of UV photons to anneal or form the crystalline silicate mass observed. Crystallisation due to UV photons originating from the AGN is dismissed by \citet{2006ApJ...638..759S}, because of the observed lack of crystalline silicates in the inner 2 pc of NGC 1068 \citep{2004Natur.429...47J_short}, and the fact that the crystalline silicates are only seen in absorption, and can therefore not be very warm. \citet{2006ApJ...638..759S}} hypothesise that the crystalline silicates {must be} produced by massive stars originating from the starburst. In this paper, we will investigate the {viability} of the build-up of crystalline silicates {due to the} starburst {activity in ULIRGS as proposed by \citet{2006ApJ...638..759S}}, and compare the crystalline fraction to the levels observed by {these authors}. {The alternative formation of crystalline silicates due to AGN activity will be the subject of a separate future study. This future work will be based on the predicted formation of dust in the quasar wind rising from the accretion disk \citep{2002ApJ...567L.107E}, where conditions may be similar to those present in AGB star winds, a class of very efficient dust producers. Crystalline silicates may form in quasar winds if the conditions are right. Indeed, we have already observed crystalline silicates in the quasar wind of PG 2112+059 \citep{2007ApJ...668L.107M}.}

{In order to explain the high crystalline fraction of silicates in ULIRGs, as reported by} \citet{2006ApJ...638..759S} we {find} that an additional source of crystalline silicates (or crystallisation of amorphous silicates) must be present, related to the AGN itself, rather than the starburst activity. {In this case, the crystallisation will occur due to heating by (UV) photons from the AGN environment, rather than (massive) stars. A potential scenario may include the formation of (crystalline) silicates in quasar winds \citep{2002ApJ...567L.107E}.} {In addition, further observational studies are useful to better establish the crystalline fraction of silicates in ULIRGs, and validate the conclusions presented here.}

1012.3251

We present a model using the evolution of the stellar population in a starburst galaxy to predict the crystallinity of the silicates in the interstellar medium of this galaxy. We take into account dust production in stellar ejecta, and amorphization and destruction in the interstellar medium and find that a detectable amount of crystalline silicates may be formed, particularly at high star formation rates, and in case supernovae are efficient dust producers. We discuss the effect of dust destruction and amorphization by supernovae, and the effect of a low dust-production efficiency by supernovae, and find that when taking this into account, crystallinity in the interstellar medium becomes hard to detect. Levels of 6.5-13 per cent crystallinity in the interstellar medium of starburst galaxies have been observed and thus we conclude that not all these crystalline silicates can be of stellar origin, and an additional source of crystalline silicates associated with the active galactic nucleus must be present.

12201821

2011ASPC..448.1333W

1012

1012.0314_arXiv.txt

Solar and late-type stars emit X-rays through a magnetically confined plasma, or corona, at several million Kelvin. The corona is believed to be powered by the stellar dynamo, which itself is powered by differential rotation within the star \citep[e.g.][]{skum72,pall81,noye84}. The observed decrease in stellar X-ray luminosity of several orders of magnitude between the zero age main sequence \citep[e.g.][]{feig02,flac03,wrig10a} and solar age \citep[e.g.][]{pere00} has therefore been attributed to the rotational spin-down of the star, though a consistent picture has yet to emerge. X-ray observations of stellar clusters up to a billion years in age have been used to study the age -- rotation and rotation -- activity relationships \citep[e.g.][]{barn03,pizz03}. However, the dispersal of stellar clusters as they age means that very few stellar clusters older than a billion years exist, and many of these are faint and distant. Age estimates for individual main-sequence field stars are difficult to derive independently and therefore the evolution of both rotation and stellar activity for older stars is poorly understood. Attempts to overcome this problem by using large samples of field stars have often resulted in conflicting conclusions. \citet{gude97} studied a sample of nearby solar-type stars aged $1-10$~Gyr and found that the X-ray luminosities decayed as $L_X \propto t^{-1.5}$, while \citet{mice02} could find no evidence for a clear decay law over a similar age range, and \citet{feig04} estimated a decay law of $L_X \propto t^{-2}$ from a sample of faint high Galactic latitude main-sequence stars. Other studies have attempted to understand the evolution of stellar X-ray emission by comparing predictions from models \citep[e.g. XCOUNT,][]{fava92} with X-ray surveys of field stars. However, these methods have also uncovered discrepancies such as an observed excess of yellow (G-K dwarf) stars compared to model predictions \citep{fava88}, and evidence for an excess of young main-sequence stars suggesting either a recently high star formation rate \citep{mice93} or uncertainties in the binary star fraction \citep{mice07}. These discrepancies are no doubt partly influenced by the small but diverse range of stellar X-ray samples used. These include wide-field surveys from {\it Einstein} and {\it ROSAT} \citep[e.g.][]{schm04} that may be biased toward bright and nearby thin-disk stars as well as deep surveys with the {\it Chandra} and {\it XMM-Newton} observatories that include many distant or faint sources \cite[e.g.][]{feig04}. Between these two extremes deep and wide-field surveys such as the {\it Chandra} Cosmic Evolution Survey \citep[COSMOS,][]{elvi09} offer the balance necessary to produce large samples of stellar X-ray sources that are also deep enough to uncover intrinsically faint sources. In this contribution we discuss a new model for predicting the X-ray emitting stellar population in any field of view of the Galaxy. We test this model against a new and well-constrained catalog of stellar X-ray sources detected in the {\it Chandra} COSMOS survey.

To improve out understanding of the decay of stellar dynamo activity we have compiled a new model of stellar X-ray activity, XStar, which incorporates a population synthesis Galactic model, the theories of spectral-type dependent rotational spin-down and the rotation -- activity relationship. The results of the model can be compared to samples of stellar X-ray sources from a range of observatories. Here we test the model using new observations of 60 stellar X-ray sources identified in the {\it Chandra} COSMOS field. A comparison of the distribution of X-ray fluxes, spectral types, and distances for the observed and modeled populations shows a generally good agreement, but slightly under-predicts the number of observed sources. Differences between the two X-ray flux distributions suggests that the our current formulation of the rotation -- activity relationship may not be correct, and a steeper relationship may better describe the observations.

1012.0314

Existing stellar X-ray surveys suggest major problems in our understanding of the evolution of stellar magnetic activity in solar and late-type stars, reaching conflicting conclusions over the rate of decay of X-ray activity and the spectral types responsible. We are confronting these discrepancies with a new model of the expected stellar X-ray luminosity distribution, combining a Galactic population synthesis model with current theories for rotational spin- down and the rotation -- activity relation for the stellar magnetic dynamo. Here we test our model using new observations of the stellar content of the Chandra COSMOS survey, for which 60 new stellar X-ray sources are identified from the thin disk and Galactic halo populations. Our model is in approximate agreement with the observed X-ray luminosity distribution and the distribution of spectral types responsible. However, slight differences in the form of the X-ray luminosity distribution exist that may hint at problems in our understanding of stellar X-ray emission.

12167904

2011ApJ...728..123K

1012

1012.5707_arXiv.txt

Magnetic fields in molecular clouds play an important role in the early stage of star formation. Particularly, magnetic fields in molecular clouds can regulate the cloud collapse and fragmentation process. The important parameter for these processes is the mass-to-flux ratio $M/\Phi$, where $M$ is the mass and $\Phi$ is the magnetic flux of the cloud. The mass-to-flux ratio represents the relative strength of gravity and the magnetic field. There exists a critical mass-to-flux ratio $(M/\Phi)_{\rm crit}$ \citep{mes56,str66,mou76,tom88}. If $M/\Phi > (M/\Phi)_{\rm crit}$, the cloud is supercritical and is prone to collapse. On the other hand, if $M/\Phi < (M/\Phi)_{\rm crit}$, a cloud is subcritical and cannot collapse as long as magnetic flux-freezing applies. A similar condition $M/\Phi < (M/\Phi)_{\rm crit} = 1/(2\pi G^{1/2})$ is required for stability against fragmentation of an infinite uniform layer that is flattened along the direction of a background magnetic field \citep{nak78}, where $G$ is the gravitational constant. \citet{mes56} pointed out that even if clouds are magnetically supported, ambipolar diffusion will cause the support to be lost and collapse will begin. More generally, a subcritical cloud undergoes a gravitationally-driven fragmentation instability that occurs on the ambipolar diffusion timescale rather than the dynamical timescale \citep{lan78,zwe98,cio06}. The lengthscale of the instability is fundamentally the Jeans-scale in the limit of highly supercritical clouds, but can be much larger when the mass-to-flux ratio is close to the critical value \citep{cio06}. The idea that the star formation is regulated by ambipolar diffusion and the magnetic field has been considered for many years \citep[e.g.,][]{shu87,shu99,mou99}. Magnetic field strength measurements through the Zeeman effect reveal that the mass-to-flux ratios of cloud cores are close to the critical value \citep{cru04}. These observations are consistent with core formation driven by ambipolar diffusion in subcritical clouds. However, in order to assess whether the fragmentation is occurring in a subcritical or supercritical molecular cloud, magnetic field measurements of the envelope or diffuse region are important. \citet{cru09} recently attempted to do this using the Zeeman effect, and argued that the mass-to-flux ratio of the envelope (in four measured positions with beam diameters in the range of $\sim 0.5 - 1$ pc in the vicinity of four different cloud cores) is typically greater than that of the core. This contradicts the model of core formation in subcritical clouds. However, \citet{mou09} contested their statistical analysis, and argued that dropping the restrictive assumption that the magnetic field is constant across four measured envelope regions of differing morphology and density would result in the opposite conclusion, i.e., that the core mass-to-flux ratio is likely greater than that of the envelope. Future observational tests of this kind, but benefiting from increased amounts of source integration time and spatial resolution, can settle the current difficulties of interpretation \citep{cru10}. The complementary method of measurements of polarized emission from dust grains, which reveal the magnetic field morphology in the cloud, generally show that the magnetic field in cloud cores is well ordered, and application of the Chandrasekhar-Fermi method yields mass-to-flux ratios that are also near the critical value \citep[e.g.,][]{sch98,gir06}. \citet{li09} recently compared the magnetic field directions on core scales ($< 1$pc) with those on large scales ($> 200$ pc) for several molecular clouds, and found a significant correlations. \citet{alv08} distinctly shows that the magnetic field is locally perpendicular to the large filamentary structure of the Pipe Nebula. These recent observations of polarized emission indicate that the magnetic field provides a dominant force, and that core formation may have been driven by ambipolar diffusion in subcritical clouds. Most nonlinear calculations of ambipolar-diffusion-driven evolution in subcritical clouds have focused on a single axisymmetric core, but some recent models focus on a fragmentation process that results in multiple cores. Nonlinear calculations of ambipolar-diffusion-driven fragmentation in subcritical clouds were first performed by \citet{ind00}, who carried out a two-dimensional simulation of an infinitesimally thin sheet threaded by an initially perpendicular magnetic field. \citet{bas04} and \citet{bas09a} carried out two-dimensional simulations of a magnetized sheet in the thin-disk approximation, which incorporates a finite disk half-thickness consistent with hydrostatic equilibrium and thereby includes the effect of magnetic pressure. They found that the fragment spacing in the nonlinear phase agrees with the prediction of linear theory \citep{cio06}, and that the the subcritical (or critical) model had subsonic infall, while the supercritical model had supersonic infall speed. The first three-dimensional simulation of the gravitational fragmentation with magnetic fields and ambipolar diffusion was performed by \citet{kud07}, which verified some of main results of the thin-disk models: the dichotomy between subsonic infall speeds in subcritical clouds and somewhat supersonic speeds in supercritical clouds, for example. Besides the magnetic field, the supersonic turbulence in molecular clouds is also an important component in the early stage of star formation \citep{mac04,mck07}. The inclusion of supersonic turbulent initial conditions to a fragmentation model in subcritical clouds was studied by \citet{li04} and \citet{nak05}, adopting the thin-disk approximation. They found that a mildly subcritical cloud can undergo locally rapid ambipolar diffusion and form multiple fragments because of the initial supersonic motion in which the large scale wave mode dominates the power spectrum. The core formation occurs on the order of turbulence crossing time over the simulation box, which is comparable to the dynamical time scale. \citet{bas09b} carried out a parameter study of the fragmentation regulated by gravity, magnetic fields, ambipolar diffusion, and nonlinear turbulent flows. They confirmed the onset of runaway collapse in subcritical cloud is significantly accelerated by the initial nonlinear flows. \citet{kud08} also verified that the mode of turbulence-accelerated magnetically-regulated star formation occurs in a fully three-dimensional simulation. \citet{nak08} also carried out three-dimensional MHD simulations of subcritical clouds including star star formation and bipolar outflows, and applied their model to the Taurus molecular cloud. Since three-dimensional simulations of subcritical clouds are resource-limited, large parameter studies have not yet been performed. In this paper, we study the fragmentation process in subcritical clouds, including ambipolar diffusion, by fully three-dimensional simulations. Our study follows the previous ones of \citet{kud07} and \citet{kud08}. We focus on the early stage of core formation and evolution, and do not follow evolution past the runaway collapse of a core. We carry out a parameter study by running a large number of models, and with higher spatial resolution than in our previous studies. We are especially interested in the effect of the large-amplitude nonlinear initial perturbations on the time evolution of the cloud fragmentation, and discuss the mechanism of turbulence-accelerated star formation in subcritical clouds. Our paper is organized in the following manner. The numerical model is described in Section 2, results are given in Section 3, and a discussion of the results is given in Section 4. We summarize our results in Section 5.

\subsection{Timescale of core formation} One of the problems for core formation in subcritical clouds is that the timescale of the core formation is slower than that is expected from observations \citep[e.g.,][]{jij99}. In order to solve this problem, \cite{zwe02} and \cite{fat02} pointed out that the ambipolar diffusion rate can be enhanced in a turbulent medium with a fluctuating magnetic field. In a different way, \citet{li04} and \citet{nak05} found that the compression of the cloud by the large-scale turbulent flow shortened the timescale of core formation even in subcritical clouds. \cite{kud08} and this paper followed the latter idea using three-dimensional simulations without a thin-disk approximation. Here, we discuss how the compression shortens the timescale of core formation in subcritical clouds. Figure 16 indicates that the core formation time is nearly proportional to $1/\sqrt{\rho_{peak}}$, where $\rho_{peak}$ is the value of the density peak during the first compression in the time evolution of the maximum density at $z=0$. The ambipolar diffusion time ($\tau_{\rm ad}$) is estimated from equation (\ref{eq:induction2}) as $\tau_{\rm ad} \sim 4\pi (\sqrt{2\pi G}/\alpha) \rho^{3/2}L^2/B^2$, where $L$ is the gradient length scale introduced by the turbulent compression. Because the compression by the nonlinear flow is nearly one-dimensional, the magnetic field scales roughly as $B \propto L^{-1}$ within the flux-freezing approximation. The surface density $\Sigma$ also scales as $\Sigma \propto L^{-1}$. If the compression is rapid enough that vertical hydrostatic equilibrium cannot established, then $\rho \propto \Sigma \propto L^{-1}$, and $\tau_{ad} \propto L^{5/2} \propto \rho^{-5/2}$ \citep{elm07,kud08}. This means that diffusion can occur quickly if the turbulent compression creates small values of $L$ or large values of $\rho$. This would lead to a rapidly rising value of $\beta$ in Figure 4 at $t/t_0 \sim 1$. If diffusion is so effective during the first turbulent compression that a dense region becomes supercritical, then it will evolve directly into collapse. Alternatively, the stored magnetic energy of the compressed (and still subcritical) region may lead to reexpansion of the dense region. If reexpansion of the initial compression does occur, then there is enough time for the vertical structure settle back to near-hydrostatic equilibrium. In this case, the density scales roughly like $\rho \propto \Sigma^2 \propto L^{-2}$, and $\tau_{\rm ad} \propto L \propto \rho^{-1/2}$. The ambipolar diffusion time is also estimated by \citet{mou99} as $\tau_{\rm ad} \propto \rho^{-1/2}$ in quasistatically contracting magnetically supported cores, assuming a radial force balance between gravity and the magnetic Lorentz force. In our simulation, force balance is not exactly achieved, but the time average (of the oscillations) in the cores can be approximately in force balance. Since the first strong compression leads to rapid magnetic flux loss, a higher density region (than the initial value) is rapidly attained and then settles into an approximate force balance. Since the phase of continuing ambipolar diffusion in the oscillating high density region takes longer than the initial compression, the overall time scale of the core formation is approximately proportional to $\rho_{\rm peak}^{-1/2}$. Here, instead of the time-average density, $\rho_{\rm peak}$ is used as a representitive density of the compressed gas, since it is complicated to determine the time-average density of the moving, oscillating and collapsing high density region. It is interesting that the relation attains even when the $\rho_{\rm peak}$ is used. We expect that the average density of the compressed overdense region would nearly proportional to the peak density. The relation means that the density dependence of $\tau_{\rm ad} \propto \rho^{-1/2}$ is the same as that of the free fall time. Figure 16 shows that the actual time of the core formation is about 30 times longer than the free fall time of gas with $\rho_{peak}$ when $\alpha=0.11$ and $\beta_0=0.5$. The vale of the dimensionless coefficient ($\sim 30$) is expected to be inversely proportinal to $\alpha$ (section 3.3), and not to be strongly dependent on the mass-to-flux ratio except when it is nearly critical (section 3.4). Even when the time-average density is used instead of the peak density, the value of the dimensionless coefficient would not be so different because of the weak dependence of $\rho$ ($\propto \rho^{-1/2}$). In some cases, collapse may occur during the first compression itself or soon thereafter, and the approximate scaling of the core formation time with $\rho_{\rm peak}^{-1/2}$ will not hold. Whether or not this occurs depends not only on the strength of the turbulent compression, but also on the initial mass-to-flux ratio of the cloud (related to $\beta_{0}$) and the ambipolar diffusion coefficient ($\alpha$). In model B2 (a subcritical model that is closest to critical, with normalized initial mass-to-flux ratio $\simeq 0.6$), a reexpansion does occur but the collapse starts relatively quickly, in the third oscillation (Fig. 20). In model A2 ($\alpha=0.2$), which has the poorest neutral-ion coupling of all models, the collapse starts almost in the first compression (Fig. 17). An overall conclusion is that the core formation occurs more rapidly than it would in the initial state due to the elevated value of $\rho$ in the compressed but oscillating region. Since the core formation time is approximately proprtional to $\rho^{-1/2}$, $10-100$ times desity enhancement is needed to get $3-10$ times shorter core formation time than that of the standard core formation model in subcritical clouds \citep[e.g.][]{jij99}. The observed non-thermal velocity ($3-10$ times sound speed) is eligible to make such enhancemenr by the compression in the isothermal clouds, if its scale is larger than the Jeans length. \subsection{Structure of cores} Even though the core formation time is accelerated by the nonlinear flows, the density, velocity, and magnetic field structure of a core does not strongly depend on the initial strength of the velocity fluctuation (e.g., Fig. 3, Fig. 8, and Fig. 12). The infall velocities of the cores are subsonic and the magnetic field lines show weak hourglass shapes. This result may be consistent with the result that the core formation time is proportional to $\rho_{\rm peak}^{-1/2}$. Even when core formation is initiated by the initial turbulence, the core properties can be similar to the quasistatically contracting magnetically supported cores discussed by \citet{mou99}. The initial turbulence accelerates core formation, but it eventually dissipates in the dense region when the collapsing core is formed. Subsonic infall motions were found by \cite{lee01} in an observational survey of starless cores. They found that the typical infall radii are $0.06 - 0.14$ pc and that the infall speed lies in the range of $0.05 - 0.09$ km s$^{-1}$. These values are consistent with our results. The subsonic infall is the common feature of core formation in subcritical clouds, which has been pointed out by 1D axisymmetric \citep{cio00} and 2D thin-disk \citep[]{bas04,bas09a,bas09b} models. Because of the subcritical infall, the hourglass shapes are a little weaker than that of an initially supercritical cloud that shows trans-sonic infall \citep{kud07}. More quantitative analysis of the field morphological difference and the infall speeds may distinguish the mass-to-flux ratio in a molecular cloud. \subsection{Energy dissipation} In Figure 19, we showed the time evolution of the total turbulent kinetic energy in isothermal subcritical clouds. It shows that the kinetic energy dissipates efficiently even in the flux-freezing limit ($\alpha=0$). This is different from the result of \citet{bas10}. They found that the turbulent kinetic energy persists in an isothermal subcritical cloud in the flux-freezing limit. They showed that the energy dissipation occurs effectively only in the stage of the first nonlinear compression. After that, the turbulent energy persists, and oscillates about an average value. The average value is estimated to be about half of the initial turbulent kinetic energy for the same parameters as those of our fiducial model. On the other hand, our simulation shows that energy dissipation occurs effectively even in the flux-freezing limit; the energy dissipation occurs almost exponentially during several oscillations, although about 12 percent of the initial kinetic energy remains at the final stage. The difference is likely due to the different vertical boundary conditions on the molecular cloud. \citet{bas10} use a thin-disk approximation of the cloud, and adopt the current-free condition of the magnetic field outside of the disk, assuming the outside density is negligibly small. We study the problem by three-dimensional MHD and without a thin-disk approximation, but assume that the molecular cloud is surrounded by warm gas whose density is 10 times smaller than the molecular cloud. The computational region in our model is $0<z<4H_0$, so does not cover a very large dynamic range of densities. Therefore, future three-dimensional simulations with a larger computational region along the $z$-direction and including low density gas outside of the disk may make for more revealing comparisons with the result of \citet{bas10}. \subsection{Additional parameter surveys for numerical accuracy} In addition to the parameter study in Table 1, we performed simulations of models with the same parameters as models V4 and V1, but different random realizations of the initial velocity fluctuations. The core formation time for different realizations varies by about 25\% in these samples. The overall evolution and the result of the accelerated core formation by nonlinear flows were not altered by the different initial random realizations. We performed simulations with different spatial resolutions for the model V4. In the case of the highest resolution of $(N_x, N_y, N_z)=(512,512,40)$, the core formation time is $t_{\rm core} = 14.2t_0$. There is a slight tendency that the core formation time becomes a little shorter for the high resolution cases \citep{kud09}. However, we should note that a realization of the random perturbations for the initial velocity fluctuation is also not the same for the different resolutions, because the random perturbations are input on all scales down to the grid scale. At the least, we can say that the result of the accelerated core formation by nonlinear flows is not altered significantly by spatial resolution. We also performed a simulation with a different boundary condition of magnetic fields for the model V4. Instead of the symmetrical boundaty at $z=z_{\rm out}$, we used the vertical field condition for the magnetic field, i.e., $B_x=B_y=0$ for $z \ge z_{\rm out}$. The overall evolution and the result of the accelerated core formation were not also altered by this. The core formation time for the different baounday varies by about 0.2\% in this sample.

1012.5707

We employ three-dimensional magnetohydrodynamic simulations including ambipolar diffusion to study the gravitationally driven fragmentation of subcritical molecular clouds, in which the gravitational fragmentation is stabilized as long as magnetic flux-freezing applies. The simulations show that the cores in an initially subcritical cloud generally develop gradually over an ambipolar diffusion time, which is about a few ×10<SUP>7</SUP>yr in a typical molecular cloud. On the other hand, the formation of collapsing cores in subcritical clouds is accelerated by supersonic nonlinear flows. Our parameter study demonstrates that core formation occurs faster as the strength of the initial flow speed in the cloud increases. We found that the core formation time is roughly proportional to the inverse of the square root of the enhanced density created by the supersonic nonlinear flows. The density dependence is similar to that derived in quasistatically contracting magnetically supported clouds, although the core formation conditions are created by the nonlinear flows in our simulations. We have also found that the accelerated formation time is not strongly dependent on the initial strength of the magnetic field if the cloud is highly subcritical. Our simulation shows that the core formation time in our model subcritical clouds is several ×10<SUP>6</SUP> yr due to the presence of large-scale supersonic flows (~3 times sound speed). Once a collapsing core forms, the density, velocity, and magnetic field structure of the core do not strongly depend on the initial strength of the velocity fluctuation. The infall velocities of the cores are subsonic and the magnetic field lines show weak hourglass shapes.

1408504

2011ApJ...728..102W

1012

1012.2127_arXiv.txt

\label{sec:Intro} Under the current paradigm of galaxy formation galaxies build via a hierarchical process and our Galaxy is deemed no exception. Relics of formation are observed as spatial and kinematic substructures in the Galaxy's stellar halo. Recent observations such as those from the Sloan Digital Sky Survey (SDSS) have brought a large increase in the detections of substructures within the outer reaches of the halo (out to $d<80$ kpc). These streams have usually been detected as spatial overdensities from photometry e.g., \citet{Yanny2000, Majewski2003, Belokurov2006, Newberg2009}. Many of these structures have been identified as belonging to the debris of the Sagittarius dwarf spheroidal galaxy (Sgr dSph), which traces the polar orbit of this galaxy as it merges with the Milky Way. Furthermore, after subtracting such prominent substructures \citet{Bell2008} observed a dominant fraction of the halo to deviate from a smooth distribution, consistent with being primarily accretion debris. Closer to the Sun the spatial coherence of streams and substructures is not so easily discernible and most streams of stars are visible only as velocity structures, such as the \cite{Helmi1999} stream. Indeed, \citet{Helmi2009} has shown that only at distances greater than $\sim10\,\kpc$ do we expect that the structures associated with tidal debris to be observable as spatial overdensities. Therefore, if we wish to identify and study structures within the inner reaches of the halo - where they are most accessible for high resolution follow-up observations - we must search utilizing kinematic data. Kinematic surveys of the solar neighbourhood are therefore ideal to detect substructures in the nearby regions of the Galaxy's halo. RAVE (RAdial Velocity Experiment) is an ambitious program to conduct a 17,000 square degree survey measuring line-of-sight velocities, stellar parameters, metallicities and abundance ratios of up to 1 million stars \citep{Steinmetz2006}. RAVE utilizes the wide field ($30$ deg$^2$) multi-object spectrograph 6dF instrument on the 1.2-m UK Schmidt Telescope of the Anglo-Australian Observatory (AAO). RAVE's input catalogue for the most part\footnote{Red giants in the direction of rotation were also targeted between $225^\circ < l < 315^\circ$, $5^\circ < |b| < 25^\circ $ with $J-K > 0.5$. This region is not discussed in this paper however.} has only a magnitude selection criterion of $9<I<13$, thus creating a sample with no kinematic biases. The observations are in the Ca-triplet spectral region at 840 nm to 875 nm with an effective resolution of $R=7500$. Starting in April 2003, at the end of 2009 RAVE had collected more than 400,000 spectra. RAVE's radial velocities are accurate to $1.3\,\kms$ when compared to external measurements, while the repeat observations exhibit an accuracy of $2\,\kms$ \citep{Zwitter2008}. These highly accurate radial velocities make RAVE ideal to search for kinematic substructures in an extended region around the sun. Indeed, with RAVE we now move away from studying the \textit{solar neighbourhood} (e.g. \citet{Nordstrom2004}: $d < 0.2\,\textrm{kpc}$) to examining the \textit{solar suburb} ($d < 4\,\kpc$). Using RAVE's highly accurate radial velocities, we have discovered a stream that lies mostly within the constellation of Aquarius at a distance of $0.5\lesssim d\lesssim 10\,\kpc$, in the direction $(l,\ b)\sim(55^\circ,\ -60^\circ)$ and at $\RV=-200\,\kms$. The velocity places the stream as part of the Galaxy's halo. As it lies in the direction of the constellation of Aquarius we have named it the Aquarius stream. The detection of this stream is described in Section \ref{sec:Stream}. In Section \ref{sec:Bes} we compare the RAVE data to mock data from the Besan\c con Galaxy mode and the newly-developed galaxy modelling code Galaxia, which offers a number of significant advantages. Using these models we determine the significance of the detection and constrain its localization. In Section \ref{sec:Pop} we use RAVE's stellar parameters combined with 2MASS ($JHK$) photometry to infer basic properties of the stream population and derive distance estimates. We also use Reduced Proper Motions to obtain another estimate of the distances. The stream appears to be highly localized on the sky which is interesting considering the apparent proximity of the stream. In Section \ref{sec:Pos} we explore possible connections of the Aquarius stream to other known spatial and kinematic streams, finding that it is not linked to any previously reported structure. In Section \ref{sec:Nat} we investigate possible connections to other (marginal) over-densities in the RAVE dataset, and conclude that the stream is unlikely to be associated with any of them. A simple model of the recent disruption of a satellite in the Galaxy's potential is able to account for the observed localization. The Aquarius stream thus is a new and nearby enigma in the Milky Way's halo. \begin{figure*}[!tb] \epsscale{1.0}\plotone{f1.jpg} \caption{(a) $\vrad$ as a function of galactic latitude for RAVE data with $-70<b<-50$, $J>10.3$. The Aquarius Stream is identified as an overdensity of stars with $-250<V_\mathrm{los}<-150\,\kms$, $30^\circ<l<75^\circ$, as delimited by the red box. (b) The histogram of $\vrad$ with the additional constraint $30^\circ<l<75^\circ$ clearly shows the stream as an anomalous feature in the wings of the velocity distribution. The grey shading displays the $\pm 1\sigma$ limits. \label{f1}} \end{figure*}

In this paper we report the detection of a new halo stream found as an overdensity of stars with large heliocentric radial velocities in the RAVE data-set. The detection is enabled by RAVE's selection criteria creating no kinematic biases. The fifteen member stars detected have $\RV=-199\pm27\,\kms$ and lie between $-70^\circ< b <-50^\circ$, $30^\circ<l<75^\circ, J>10.3$ in the constellation of Aquarius. We established the statistical significance of the stream by comparing the RAVE data, in the Galactic latitude range $-70^\circ<b<-50^\circ$, to equivalent mock samples of stars created using the Besan\c con Galaxy model and the code Galaxia. For different cell sizes, $\Delta l \times \Delta \vrad$, we compared the number of stars in the data and models, finding that for the majority of cell sizes the region around the Aquarius stream exhibited a $4\sigma$ overdensity in the data, irrespective of the dust modelling. Searching for additional overdensities in neighboring latitude regions yields no structures of the same level of significance (other than the LMC), though two regions are identified as being marginally overdense. For most of the Aquarius stars RAVE stellar parameter estimates are also available. The member stars are metal-poor with $\mh=-1\pm 0.4$ and we derive a preliminary isochrone fit in the $\teff$-$\log g$ plane with an population age of $10\,\gyr$. Both the $\vrad$ and metallicity are consistent with the group being within the stellar halo. We further use a Reduced Proper Motion Diagram to derive the transverse velocity for the stream, finding $v_T=250\pm100\,\kms$ for the group. This again places it within the Galaxy's halo. We use the isochrone fits and the RPMD to provide distance estimates to the stars, where we prefer the latter as they give more consistent kinematics. We investigated the relation of the stream to known substructures. We first discussed the probability of the stream being with debris from the Sagittarius dwarf. This is a priori plausible because the stream does not fall far from the orbital plane of the Sgr dwarf and the stream's metallicity is consistent with that of the dwarf. A comparison to the models of Sagittarius dwarf debris from \citet{Law2005} and \citet{Law2009}, shows that although the majority of models do not yield a good fit, a certain selection of nearby stars in the oblate model provides a reasonable fit in the $\Lambda_\sun$-$\vgal$ plane. This is most likely just coincidental however: the distributions in both distance and $\VPHI$ are clearly inconsistent with those of the Sgr stream. Also, the oblate model is the least favoured of all the models when compared to the most recent data for the Sagittarius stream. We thus conclude that the Aquarius stream is most likely not associated with the Sagittarius dwarf. A search of other known substructures both in the solar neighbourhood (e.g. Kapteyn group) and in the solar suburb (e.g. Canis Major and Virgo overdensities, Monoceros stream, Hercules-Aquila cloud) yielded no positive identifications. Finally, to understand better how the stream is both local and localized on the sky, we performed simple dynamical simulations of a model satellite galaxy dissolving in the Galactic potential. We presented simulations for two time-scales, one where the satellite is dissolving and the other when it is completely phase mixed. We compared the distribution in $l,\ b, \vrad$ space of nearby tracer particles at the present day to that of the Aquarius stream stars plus the two other marginally overdense regions found in the RAVE data. The model in which the progenitor has had time to become phase mixed predicts over-densities in places were the data show none. By contrast, the dissolving, not-yet-phase-mixed scenario was able to account for the localization as well as reproducing the observed structure of the Aquarius stream. We therefore suggest that the stream is dynamically young: the localization could be explained as a recent disruption event of a progenitor whereby the stream has yet to occupy the available phase-space. The progenitor could either be a globular cluster or a dwarf galaxy, which may or may not have survived to the present day. We make no positive identification of with any globular clusters, though there could be a possible link with likely dwarf galaxy remnant, $\omega$ Cen, and the associated Kapteyn group. Follow-up high-resolution abundances would elucidate this possible connection. Further, more sophisticated simulations of Aquarius are required. This will enable a better understanding of this interesting, new halo stream which places hierarchical formation right on our proverbial doorstep.

1012.2127

We identify a new, nearby (0.5kpc &lt;~ d &lt;~ 10 kpc) stream in data from the RAdial Velocity Experiment (RAVE). As the majority of stars in the stream lie in the constellation of Aquarius, we name it the Aquarius Stream. We identify 15 members of the stream lying between 30° &lt; l &lt; 75° and -70° &lt; b &lt; -50°, with heliocentric line-of-sight velocities V <SUB>los</SUB> ~ -200 km s<SUP>-1</SUP>. The members are outliers in the radial velocity distribution, and the overdensity is statistically significant when compared to mock samples created with both the Besançon Galaxy model and newly developed code Galaxia. The metallicity distribution function and isochrone fit in the log g-T <SUB>eff</SUB> plane suggest that the stream consists of a 10 Gyr old population with [M/H] ~ -1.0. We explore relations to other streams and substructures, finding that the stream cannot be identified with known structures: it is a new, nearby substructure in the Galaxy's halo. Using a simple dynamical model of a dissolving satellite galaxy, we account for the localization of the stream. We find that the stream is dynamically young and therefore likely the debris of a recently disrupted dwarf galaxy or globular cluster. The Aquarius stream is thus a specimen of ongoing hierarchical Galaxy formation, rare for being right in the solar suburb.

4886212

2011ApJ...727...59B

1012

1012.0855_arXiv.txt

Supergiant Fast X-ray Transients \citep[SFXTs:][]{Smi04,Sgu05,Sgu06,Neg06a} form a new subclass within the group of high-mass X-ray binaries (HMXBs). Previously, HMXBs were catalogued \citep[e.g.][and references therein]{Liu00} into one of either two groups: the transient sources with B-emission line stellar companions (BeXBs); or the persistent, but variable, HMXBs with supergiant O or B-type companions (SGXBs). The SFXTs share properties from both groups: like BeXBs, they are characterized by outbursts of short (hours) to long (days) duration whereby the peak intensity is several orders of magnitude greater than during quiescence; and like SGXBs, they are systems in which a compact object (usually a neutron star) is paired with a supergiant OB star. Their study is therefore important because SFXTs could represent an evolutionary link between BeXBs and SGXBs \citep[e.g.][]{Neg08,Cha10,Bod10,Liu10}. \object{XTE\,J1739$-$302} \citep{Smi98} was the first hard X-ray transient to display behavior that was identified as characteristic of the class of SFXTs \citep{Neg06a}. It was listed as \object{AX\,J1739.1$-$3020} in the \asca\ catalog of X-ray sources in the Galactic Center region \citep{Sak02}. Five years after its discovery, \integ\ detected a transient source named \igr\ which, because of the positional coincidence and similar emission behavior, was shown to be the hard $X$-ray (18--50\,keV) counterpart to the \xte\ source \citep{Sun03,Smi03}. A refined X-ray position was obtained with \chan\ \citep{Smi06}, which led to the identification of \object{USNO-B1.0\,0596-0585865} ($=$ \object{2MASS\,J17391155$-$3020380}), a supergiant star of spectral type O8\,Iab located $\sim$2.7\,kpc away, as the optical/IR counterpart \citep{Neg06b,Cha08,Rah08}. This source has been detected numerous times in outburst and in quiescence thanks to monitoring campaigns with \integ\ \citep{Lut05,Sgu05,Tur07,Bla08,Che08,Rom09a,Duc10} and with \swift\ \citep{Rom08a,Rom08b,Sid08,Sid09a,Sid09b,Rom09b,Rom10}. Most of the outbursts are short ($\sim$0.5\,h) and weak ($\sim$10\,cps in 20--40\,keV as seen with \integ), while a few flares last longer than an hour (sometimes days as seen in the 0.5--10\,keV band with \swift) and are rather intense ($\sim$60\,cps or $\sim$400\,mCrab with \integ) \citep{Bla08}. The photoelectric absorption is variable from outburst to outburst \citep{Smi06}, reaching up to $(13^{+4}_{-3})\times10^{22}$\,\cmsq\ during the rising portion of certain flaring episodes \citep{Sid09a,Sid09b}. The dynamic range between quiescence and peak outburst is $\sim10^{4}$ \citep{Int05}. Based on multiple observations of \igr\ gathered over the course of a year with \swift, \citet{Rom09b} showed that bright outbursts account for only 3--5\% of the entire observation set, while the source spends 39$\pm$5\% of the time in an inactive state near quiescence. Quiescence, i.e., a flux in the 0.5--10\,keV range below the \swift\ detection limit of $1.6\times10^{-12}$\,\ergcms\ for a $\sim$1\,ks exposure, is a very rare state \citep{Rom09b}. What triggers the outbursts is still an open question. Several mechanisms have been proposed to explain the peculiar X-ray emission behavior of SFXTs. The outbursts could be due to one or a combination of the following effects: the intermittent accretion of clumpy wind material by a compact object in a larger orbital radius than in classical SGXBs; passage of the compact object in an eccentric orbit through anisotropic winds emanating from some supergiant stars; and the dynamics of wind accretion onto highly-magnetized neutron stars including gating effects, photo-ionization, transient disks, and Rayleigh-Taylor instabilities \citep[e.g.][and references therein]{Int05,Wal07,Sid07,Neg08,Boz08,Duc09}. While the identity of the stellar companion is known, the nature of the compact object has yet to be conclusively determined. The spectral properties of \igr, i.e. its hard power-law shape and possible spectral cutoff around 13\,keV \citep{Sid09a}, suggest a neutron star primary. However, coherent pulsations or cyclotron absorption lines that would firmly establish the presence of a neutron star have not been detected thus far \citep{Smi98,Sgu05,Bla08}. Recently, the orbital period of the system was shown to be 51.47(2)\,d with hard X-ray outbursts ($\gtrsim$20\,keV) and quiescent epochs coexisting at periastron suggesting clumpy, anisotropic winds \citep{Dra10}. \begin{figure*}[!t] \centering \includegraphics[width=\textwidth,angle=0]{fig1.eps} \caption{Image of IGR\,J17391$-$3021 as captured by the XIS1 detector on {\it Suzaku}. Pixels have been grouped in blocks of 4 (bin: group4 in \texttt{ds9}). The circular source and square background selection regions are also shown. Coordinates are given as equatorial R.A. and Dec. (J2000) where North is up and East is left. } \label{img} \end{figure*} In February, 2008, \suz\ observed \igr\ for 37\,ks of effective exposure time. The data from this observation are described in Section \ref{sec_obs}. Results from timing and spectral analyses are reported in Sections \ref{sec_time} and \ref{sec_spec}, respectively. We discuss our findings and present our conclusions in Sections \ref{sec_disc}--\ref{sec_conc}.

\label{sec_conc} The main results from our \suz\ observation of the SFXT \igr\ can be summarized as follows: \begin{itemize} \item[-]For most of the observation, \igr\ is in what we believe represents the quiescent state. The unabsorbed luminosity in the 0.5--10\,keV range is $1.3\times10^{33}$\,\ergs, and as low as $8.7\times10^{32}$\,\ergs\ in the flux-limited spectrum. The spectrum from this epoch is well described by an absorbed, hard power law with \nh\ $=(1.0_{-0.5}^{+0.6})\times10^{22}$\,cm$^{-2}$ and $\Gamma = 1.0\pm0.3$, or by an absorbed thermal blackbody with a temperature of $1.57_{-0.15}^{+0.16}$\,keV. \item[-]Around 33\,ks into the observation (MJD\,54518.92) the source became slightly more active (``low state'') culminating in a series of weak flares each lasting $\sim$3\,ks whose peak luminosity is only a factor of $\sim$5 greater than the quiescent emission that preceded it. During this low state, the \nh\ reached $(4.1_{-0.4}^{+0.5})\times10^{22}$\,cm$^{-2}$ which is a factor of $\sim$2--4 times that of the expected interstellar value, and which is $\sim$2--9 times the \nh\ measured during quiescence. The accretion of small clumps of stellar wind material can be responsible for the weak flares, while the increase in \nh\ can be explained by the passage of the clump in the line-of-sight to the X-ray source prior to accretion. \item[-]The average luminosity of the entire observation is $4.8\times10^{33}$\,\ergs. When placed in the full light curve from the \swift-XRT monitoring campaign by \citet{Rom09b}, we find that this luminosity is on par with the low-intensity XRT detections (as well as a few of the 3$\sigma$ upper limits from non-detections). Furthermore, we suspect that many of the faintest detections from the XRT contain similar low states (potentially including weak flares) that are just above the quiescent flux. We estimate that the duty cycle of the low state is 60$\pm$5\%, which makes it the most common state for this source. \end{itemize} As a prototypical member of the class of SFXTs, \igr\ holds valuable clues to the accretion processes of these intriguing transients. Continued monitoring of this source by sensitive instruments such as \suz\ should fuel the ongoing debate concerning the mechanisms that are responsible for their peculiar emission behavior.

1012.0855

We present an analysis of a 37 ks observation of the supergiant fast X-ray transient IGR J17391-3021 (= XTE J1739-302) gathered with Suzaku. The source evolved from quiescence to a low-activity level culminating in three weak flares lasting ~3 ks each in which the peak luminosity is only a factor of five times that of the pre-flare luminosity. The minimum observed luminosity was 1.3 × 10<SUP>33</SUP> erg s<SUP>-1</SUP>(d/2.7 kpc)<SUP>2</SUP> in the 0.5-10 keV range. The weak flares are accompanied by significant changes in the spectral parameters including a column density (N <SUB>H</SUB> =(4.1<SUP>+0.4</SUP> <SUB>-0.5</SUB>) × 10<SUP>22</SUP> cm<SUP>-2</SUP>) that is ~2-9 times the absorption measured during quiescence. Accretion of obscuring clumps of stellar wind material can explain both the small flares and the increase in N <SUB>H</SUB>. Placing this observation in the context of the recent Swift monitoring campaign, we find that weak-flaring episodes, or at least epochs of enhanced activity just above the quiescent level but well below the moderately bright or high-luminosity outbursts, represent more than 60% ± 5% of all observations in the 0.5-10 keV energy range making this the most common state in the emission behavior of IGR J17391-3021.

12130635

2010arXiv1012.0587G

1012

1012.0587_arXiv.txt

In 2008-2009, the cosmic ray (CR) observation experiments PAMELA \cite{Adriani:2008zr} and FERMI \cite{Abdo:2009zk} both recorded an anomalous positron flux in excess of the expected astrophysical background. Similar excesses were also reported by the ATIC \cite{:2008zzr} and HESS \cite{Aharonian:2009ah} experiments at different energies. These observations initiated a flurry of activity aimed at explaining the discrepancy. Proposed solutions include both astrophysical sources, such as local pulsars \cite{Hooper:2008kg, Fujita:2009wk, Profumo:2008ms}, as well as more exotic sources such as annihilating or decaying dark matter particles \cite{Feng:2010gw, Ishiwata:2010am, Ibarra:2008jk}. The amount to which the latter contributes is extremely important information, as many models of new physics that extend the Standard Model predict the existence of particles that can make excellent candidates for dark matter. Inherent to any of the possible solutions is an assumption for the underlying physics describing how the CR particles are transported from their sources to Earth. At present, our understanding of how CR propagate throughout the galaxy is mature but still incomplete. The most realistic description of CR propagation is currently obtained from models that assume CR transport occurs via a combination of spatial diffusion (resulting from interaction with the random and complicated galactic magnetic field) and convection (resulting from interaction with large-scale streaming of CR away from the galactic disk). Explicit implementations have been presented in the literature. These include both semi-analytic \cite{Putze:2010zn, Maurin:2001sj} as well as fully numerical treatments \cite{Strong:1998pw,Strong:2009xj, Evoli:2008dv, DiBernardo:2009ku}. In the canonical propagation scenario, the diffusion behavior is assumed to be identical everywhere in the galaxy: both in the gaseous disk as well as the surrounding region known as the halo. This isotropic diffusion model is capable of reproducing numerous astrophysical observations well, and this has led to its adoption as the model from which astrophysical backgrounds are normally derived. However, in isotropic diffusion models the convective wind velocity is limited to no more than about 10-20 km/s. Assuming a higher wind velocity results in the incorrect prediction of several existing observations, most notably measurements of secondary-to-primary ratios such as the boron-to-carbon (B/C) ratio, as well as the ratio of radioactive beryllium isotopes ($^{10}$Be/$^9$Be) \cite{Strong:2007nh, Maurin:2002hw, jones:1}. These ratios, in particular, are highly sensitive to the parameters that control the various CR transport mechanisms, and are often used to benchmark predictions of a given model. While very high velocity galactic winds on the order of 1000 km/s have been observed in other galaxies \cite{Veilleux:2005ia}, until recently it was thought that the Milky Way did not exhibit a wind, or that the wind velocity was limited to very low values. However, in 2007 the ROSAT satellite observed x-ray emission from within the Milky Way that is consistent with the presence of a galactic wind having a velocity of a few hundred km/s \cite{Breitschwerdt:2008na, Everett:2007dw}. The measurements are also consistent with a wind velocity profile that exhibits spatial dependence, and follows the radial distribution of supernova remnants (SNR) in the galaxy. The presence of high-velocity, radially dependent galactic winds appear to be natural and may also explain the large bulge-to-disk ratio observed by INTEGRAL \cite{2005A&A...441..513K, 2007ESASP.622...25W}. Recently, a propagation model has been introduced that employs both a radially dependent convective wind profile and \textit{anisotropic} diffusion, where the diffusion behavior is not uniform throughout the galaxy, but instead varies as a function of position \cite{Gebauer:2009hk}. Specifically, the model assumes that the intensity of CR diffusion increases with distance from the galactic plane. This model is capable of supporting a convective wind velocities of several hundred km/s, which is consistent with the ROSAT measurement, while simultaneously reproducing benchmark measurements such as the B/C and $^{10}$Be/$^9$Be ratios with similar accuracy to the traditional isotropic diffusion model. A crucial prediction of this scenario is that the astrophysical anti-proton background flux appears lower than what is predicted by isotropic diffusion models \cite{deBoer:2009tz}. Such an outcome has significant implications for the indirect detection of particle dark matter, as it would reduce constraints on models for new physics that tend to over-predict the antiproton flux, assuming no excess is currently observed in the data \cite{Donato:2008jk}. In this paper we implement the anisotropic propagation model of \cite{Gebauer:2009hk} in the public CR propagation code GALPROP v50.1 \cite{Strong:2009xj}, and explore varying levels of diffusion anisotropy by adjusting the rate at which the diffusion intensity increases with distance from the galactic plane. We perform a $\chi^2$ analysis for B/C, $^{10}$Be/$^{9}$Be and the recent PAMELA $\bar{p}/p$ data and find the models that fit the data best all exhibit spatially-dependent, anisotropic diffusion. Further, we find these models predict a $\bar{p}/p$ flux that is significantly lower than the observed flux between 2 GeV and 20 GeV. When we attempt a combined fit including the $\bar{p}/p$ data, we find an increase in $\chi^2_{\mathrm{min}}$, by $\Delta \chi^2=35.4$ for the 23 PAMELA data points considered.

We consider the anisotropic propagation model of \cite{Gebauer:2009hk}, which assumes that CR diffusion increases linearly with distance from the galactic disk, and radially dependent convective wind flows with maximum velocity of $\mathcal{O}(100)$ km/s suggested by the ROSAT measurement. This model is capable of reproducing key observations such as the B/C and Be/Be flux ratios, whereas traditional models based on isotropic diffusion cannot support convective wind velocities of more than approximately 20 km/s. We implement this model in GALPROP, and introduce a new parameter, $\alpha$, which controls the dependence of the diffusion coefficient $D_{xx}$ on the distance $|z|$ above the gaseous disk; $\alpha=0$ corresponding to isotropic diffusion, while $\alpha=1$ a linear increase as in \cite{Gebauer:2009hk}. We then perform a $\chi^2$ analysis for the B/C, $^{10}$Be/$^9$Be, and the recent PAMELA $\bar{p}/p$ data sets in order to constrain the transport parameters. For a combined fit to the B/C and $^{10}$Be/$^9$Be flux only, we observe $\chi^2_{\mathrm{min}}/n_{\mathrm{dof}}=38.3/31$, indicating a good fit to both data sets. We find $\alpha > 0.27$, $4.3\times10^{28}\;\mathrm{cm}^2/\mathrm{s}\le D_0 \le 5.1\times10^{28}\;\mathrm{cm}^2/\mathrm{s}$, $0.35\le \delta \le 0.43$, and $47\;\mathrm{km/s} \le v_A \le 58\; \mathrm{km/s}$ at $3\sigma$ for the range of parameter space we explored. The Alfv\'en velocity is higher than expected, and due to a large wind velocity gradient, $dV/dz=35$ km/s/kpc, required to model the wind profile suggested by the ROSAT measurements. The isotropic diffusion case $\alpha=0$ is excluded at $4\sigma$. Models within the 68\% CL region predict a lower $\bar{p}/p$ flux than the PAMELA observation between 2 GeV to 20 GeV. For a combined fit to B/C, $^{10}$Be/$^9$Be and also the PAMELA $\bar{p}/p$ data we observe an increase in $\chi^2_{\mathrm{min}}$, by $\Delta \chi^2=35.4$ for the additional 23 PAMELA data points considered, indicating the overall fit is marginal. In this case we find $\alpha > 0.51$, $4.1\times10^{28}\;\mathrm{cm}^2/\mathrm{s}\le D_0 \le 4.8\times10^{28}\;\mathrm{cm}^2/\mathrm{s}$, $0.33\le \delta \le 0.38$, and $43\;\mathrm{km/s} \le v_A \le 49\; \mathrm{km/s}$ at $3\sigma$, which are consistent with expectations. Models within the 68\% CL region give an improved fit to the $\bar{p}/p$ flux, however the predicted distributions are still below the PAMELA observation between approximately 2 GeV to 15 GeV. These results suggest that spatially dependent diffusion represents a viable solution to the problem of accommodating $\mathcal{O}(100)$ km/s convective wind flows in models of CR propagation, and that in this case the observed antiproton flux may exhibit an excess above the astrophysical background. An anomalous excess in the antiproton flux would have many implications for the indirect detection of particle dark matter (DM). For example, pair annihilation of light neutralino DM can explain the positron excess observed by e.g. PAMELA \cite{Grajek:2008pg}. These annihilations necessarily produce hadrons in addition to leptons, and so a contribution to the antiproton flux is also expected from this scenario.

1012.0587

Recently a cosmic ray propagation model has been introduced, where anisotropic diffusion is used as a mechanism to allow for $\mathcal{O}(100)$ km/s galactic winds. This model predicts a reduced antiproton background flux, suggesting an excess is being observed. We implement this model in GALPROP v50.1 and perform a $\chi^2$ analysis for B/C, $^{10}$Be/$^{9}$Be, and the recent PAMELA $\bar{p}/p$ datasets. By introducing a power-index parameter $\alpha$ that dictates the dependence of the diffusion coefficient $D_{xx}$ on height $|z|$ away from the galactic plane, we confirm that isotropic diffusion models with $\alpha=0$ cannot accommodate high velocity convective winds suggested by ROSAT, while models with $\alpha=1$ ($D_{xx}\propto |z|$) can give a very good fit. A fit to B/C and $^{10}$Be/$^{9}$Be data predicts a lower $\bar{p}/p$ flux ratio than the PAMELA measurement at energies between approximately 2 GeV to 20 GeV. A combined fit including in addition the $\bar{p}/p$ data is marginal, suggesting only a partial contribution to the measured antiproton flux.

1372174

2011A&A...528A.102M

1012

1012.4793_arXiv.txt

Galaxy clusters form by the hierarchical accretion of cosmic matter. They reach the virial equilibrium over a volume that defines the regions where the pristine gas accretes on the dark matter (DM) halo through gravitational collapse and is heated up to millions degrees through adiabatic compression and shocks. The end products of this accretion process exhibit in the X-ray band similar radial profiles of surface brightness \citep{Vikhlinin99, Neumann05, Ettori09}, plasma temperature \citep[e.g.]{Allen01, Vikhlinin05, Leccardi09} and gravitational mass distribution \citep[e.g.]{Pointecouteau05}. \begin{figure*} \begin{tabular}{ccc} \includegraphics[width=5.75cm]{ima_coma.ps} & \includegraphics[width=5.75cm] {ima_a1795.ps} & \includegraphics[width=5.75cm]{ima_abell2029.ps} \\ \includegraphics[width=5.75cm]{ima_pks0745.ps} & \includegraphics[width=5.75cm]{ima_1557+35.ps} & \includegraphics[width=5.75cm]{ima_0847+13.ps} \\ \end{tabular} \caption{Images for the six objects in our sample, with detected sources excised. Data are smoothed with a 30\arcsec Gaussian filter. Superimposed are the circular extraction regions used for the spectral analysis. Circular regions around the detected sources are excised.} \label{fig:imaprof} \end{figure*} \noindent The measurement of the properties of the ICM have been enormously improved thanks to the arcsec resolution and large collecting area of Chandra and XMM-Newton, but still remain possible only where the X-ray emission can be well resolved against the background (both instrumental and cosmic). While the X-ray surface brightness and gas density can be estimated in few cases above $0.7 r_{200}$ \citep[e.g.]{Vikhlinin99, Neumann05, Ettori09}, the ICM temperature, requiring more than an order of magnitude in net counts than the surface brightness to be firmly measured, can be reasonably well constrained up to a fraction ($\sim 0.5-0.6$) of the virial radius \citep[e.g.]{Vikhlinin06, Leccardi09, Ettori10, Arnaud10}. The regions not-yet observed are expected to retain most of the information on the processes that characterize the accretion and evolution within the cluster of the main baryonic component \citep{Roncarelli06, Rasheed10}. \begin{table*} \begin{center} \caption[]{Cluster sample. Mean spectroscopic temperatures and redshift of the first four clusters are from the references reported in the fifth column, through a query to BAX archive (\texttt{http://bax.ast.obs-mip.fr/}); the remaining 2 are calculated from our data (see text). r$_{200}$ are estimated by adopting the scaling relations \citep{Arnaud05} and here are expressed in units of arcmin and Mpc. r$_{max}$ is the center of the most external annulus for which we measured the temperature, here expressed in units of r$_{200}$. } \begin{tabular}{|l|cccccccc|} \hline Name & RA, Dec & z & kT & references & r$_{200}$ &r$_{max}$ & Exp.& N$_{\rm H}$ \\ & [deg] & & [keV] & &[$\arcmin$, Mpc]&[r$_{200}]$& [ks] &10$^{20}$cm$^{-2}$\\ \hline Coma &194.9392 +27.9429&0.023 & 8.25$\pm$0.01 &\cite{Arnaud01} & 78.2 [2.18] &0.17 & 43.7 & 0.80\\ Abell 1795 &207.2183 +26.5903&0.062 & 6.12$\pm$0.05 &\cite{Vikhlinin06}& 25.7 [1.84] &0.41 & 20.3 & 1.32\\ Abell 2029 &227.7336 +5.74440&0.077 & 8.47$\pm$0.09 &\cite{Vikhlinin06}& 26.7 [2.16] &0.43 & 42.3 & 3.25 \\ PKS0745-19 &116.8799 -19.2948&0.102 & 7.97$\pm$0.28 &\cite{Arnaud05} & 18.1 [2.06] &0.58 & 63.1 & 41.8\\ SWJ1557+3530 &239.4287 +35.5073&0.153$\pm0.006$& 6.79$\pm$0.25& this work& 11.5 [1.89] &0.62 & 180.1 & 0.20\\ SWJ0847+1331 &131.9550 +13.5278&0.358$\pm0.005$& 6.02$\pm$0.34& this work& 5.3 [1.56] &0.50 & 203.1 & 3.23\\ \hline \end{tabular} \label{tab:obse} \end{center} \end{table*} It is therefore crucial to obtain direct measurements of the cluster properties at these large radii where very important processes for the evolution of the clusters take place. Very recently, Suzaku, thanks to its low background and high sensitivity, has been able to map roughly (i.e. with a spatial resolution limited to $>$4$\arcmin$) the regions close to the virial radius, providing the first estimate of the gas temperature in 6 objects \citep{Fujita08, George09, Reiprich09, Bautz09, Kawaharada10, Hoshino10}. The aim of this paper is to show that the X-ray telescope (XRT) \citep{Burrows05} on board the Swift satellite \citep{Gehrels04} can improve the accuracy of these measurements with fairly deep observations. To this end, we present the first spectral analysis of the archived Swift observations of nearby clusters (Section~\ref{sect:data}). Since this is a non-standard analysis for the Swift - XRT data and is used and presented here for the first time, we discuss in detail the techniques adopted. In particular in Section~\ref{sect:backg} we describe the procedure we developed to estimate the background and its systematic uncertainty. To calculate how much this uncertainty affects the temperature measurement we performed a series of Monte Carlo simulations of thermal spectrum at different level of surface brightness (described in Section~\ref{sect:ssist} ). Finally in Section~\ref{sect:tsimu} we used these results to properly and robustly evaluate the expected uncertainties on the gas temperature measurements at r$_{\rm200}$ \footnote{The radius that defines the sphere enclosing a mean cluster density that is 200 times the critical value at the cluster's redshift. Here we use $r_{\rm200}$ and virial radius indifferently and we calculate r$_{\rm200}$ using the scaling relations given by \citet{Arnaud05} $r_{200}=1714\frac{(kT/5)^{0.5}}{hz}$ [kpc] where hz=$\sqrt{((1+z)^3 \Omega_m+\Omega_\Lambda)} $} on a simulated 300 ks observation of the well studied cluster Abell 1795. \noindent Throughout this paper we assume H$_{\rm0}$=70 km s$^{\rm-1}$ and $\Omega_\lambda$=0.73 and $\Omega_{\rm m}$=0.27, which are the default values in the \texttt{XSPEC}(v12.5) software. All errors are quoted at 68\% confidence level for one parameter of interest, unless otherwise specified.

\noindent In the outer regions of nearby clusters, the ICM emission is only a small fraction of the whole signal collected by the detector. The regime, where the background systematics, affect the spectroscopic measurements much more than the statistical error, is easily reached. While XMM-Newton and Chandra are not suitable for this kind of observation due to the high level of particle background, in all the published works presenting Suzaku observations of cluster outskirts the evaluation of the XRB and its variance is the main issue. Different approaches have been pursued. \cite{Bautz09} and \cite{Hoshino10}, in their studies of Abell 1795 and Abell 1413 respectively, used the signal from external regions at $\sim$ 1.2 $\times$r$_{200}$ as XRB, in the assumption that the cluster emission is negligible at that distance from the center. To study the temperature profile of PKS0745-19, \cite{George09} used the Lockman Hole observation which was performed just few days before the cluster observation. Interestingly they found significant emission from the cluster ICM at distance $>$ 1.5 $\times$r$_{200}$. \cite{Kawaharada10} used the two closest observations (at $\sim$ 8$^\circ$ of distance) among the ones suitable as blank fields in the Suzaku archive. \cite{Reiprich09} and \cite{Fujita08} used the classical models from literature to fit the cluster spectrum together with the XRB (using a different choice of free parameters). \noindent In this paper we presented for the first time the analysis of the Swift-XRT observations of a sample of 6 galaxy clusters. We measured the temperature profiles as far as $\sim$ 0.5 R$_{200}$ for all (but Coma) sample clusters. To estimate the cosmic background we used a statistical approach, modeling a representative sample of blank fields, with a sky-coverage of $\sim$ 15 deg$^2$. We used the blank field median spectrum as XRB model to fit our cluster spectra. We calculated the systematics of this approach by simulating realistic clusters with different temperatures and surface brightnesses summed to real XRB data extracted from blank fields with different sizes and exposure times. With this approach in the systematics calculation, exploiting a statistically fair sample of BF, we directly accounted for the XRB variance (for both CXRB and GXRB ). Moreover we presented a new way to calculate the uncertainties in the temperature measurements significantly refining the current approach in literature. Given a measure T$_m$, we accounted for the possibility that T$_m$ is produced by a T$_i \neq$ T$_m$. This allowed us to realistically simulate the temperature measurement in the outer regions of Abell 1795 which would be provided by a deep XRT observation. We showed that, thanks to an unprecedented combination of low background, good PSF the Swift XRT would be able to significantly improve the current accuracy of the temperature measurements in the outer regions of nearby clusters. \noindent The ideal telescope for cluster outskirts observation would be a large grasp (wide field and large collecting area) and low background telescope such the proposed WFXT \citep{Murray10}. In the next decade eRosita will be the only mission operative with these characteristics \citep{Predehl10}, a grasp 10 times larger than XMM (100 times larger than Swift-XRT). Interestingly, eRosita, with an effective area of $\sim$ 1500 cm$^2$ at 1.5 keV ($\sim$10 times larger than XRT) and an expected background of $\sim$ 9 counts s$^{-1}$ deg$^{-2}$ ($\sim$10 times larger than XRT) will have the same source / background ratio of the Swift-XRT when observing extended sources. If these numbers will be confirmed in flight, and the NXB will be reproduced with the same accuracy ($\lesssim$ 3\%), at a given value of surface brightness, eRosita will have systematic errors on temperature measurements which will be very close to the ones we found for Swift-XRT. In this case, our proposed XRT observations would represent a pilot for the eRosita mission; on the other hand, if the NXB of eRosita will be higher than expected (eRosita will be the first X-ray telescope positioned in L2) or it will be impossible to reproduce it with the same accuracy, the XRT observation would remain the only way to improve our knowledge of the cluster outskirts physics at least for the next decade. \noindent %

1012.4793

<BR /> Aims: We investigate the possibility of using the X-ray telescope (XRT) on board the Swift satellite to improve the current accuracy of the intra-cluster medium (ICM) temperature measurements in the region close to the virial radius of nearby clusters. <BR /> Methods: We present the spectral analysis of the Swift XRT observations of 6 galaxy clusters and their temperature profiles in the regions within 0.2-0.6 r<SUB>200</SUB>. Four of them are nearby famous and very well studied objects (Coma, Abell 1795, Abell 2029 and PKS0745-19). The remaining two, SWJ1557+35 (Abell 2141) and SWJ0847+13, at redshift z = 0.16 and z = 0.36, were serendipitously observed by Swift-XRT. We accurately quantify the temperature uncertainties, with particular focus on the impact of the background scatter (both instrumental and cosmic). We extrapolate these results and simulate a deep observation of the external region of Abell 1795 which is assumed here as a case study. In particular we calculate the expected uncertainties in the temperature measurement as far as r<SUB>200</SUB>. <BR /> Results: We find that, with a fairly deep observation (300 ks), the Swift XRT would be able to measure the ICM temperature profiles in the external regions as far as the virial radius, significantly improving the best accuracy among the previous measurements. This can be achieved thanks to the unprecedented combination of good PSF over the full field of view and very accurate control of the instrumental background. <BR /> Conclusions: Somehow unexpectedly we conclude that, among currently operating telescope, the Swift-XRT is the only potentially able to improve the current accuracy in plasma temperature measurement at the edges of the cluster potential. This will be true until a new generation of low-background and large field of view telescopes, aimed to the study of galaxy clusters, will operate. These observations would be of great importance in developing the observing strategy for such missions.

12168270

2011ApJ...732L..21G

1012

1012.1349_arXiv.txt

Charge-exchange (CX) collisions between highly charged Solar Wind (SW) ions and neutral gas were recently identified as an efficient mechanism for production of EUV and soft X-ray emissions \citep{2001SSRv...97..401R,2006A&A...460..289K,2009SSRv..143..217K,2007P&SS...55.1135B,2009Snowden}. There are indications that the CX collisions in the Heliosphere and Geocorona yield between 50\% and 80\% of the observed soft (below 1 keV) X-ray photons, making them a significant contributor to the soft X-ray background. Heliospheric CX X-ray emissions are sensitive to the parameters of the SW plasma and strong correlations between variations in the SW intensity and composition, and intensity of the soft X-ray background were observed and analyzed \citep{2001SSRv...97..401R,2008AGUFMSH21B1596R,2009SSRv..143..217K}. These findings indicate that the heliospheric CX X-rays could be used for diagnostics of the solar wind composition and velocities, as well as an independent probe of spatial distribution of the heliospheric neutral gas. The X-ray emission, if measured, could provide additional insight into interaction of the SW plasma with the neutral heliospheric gas. Parameters of this interaction are critical for modeling of the interstellar gas distribution and for interpretation of observational data on the diffuse soft X-ray emission from the Local Bubble \citep{2005Sci...307.1447L,2008ApJ...676..335H,2009ApSS.323....1W}. While polarization of optical emissions is readily used in investigation of various astrophysical objects, detection of X-ray polarization remains a technically challenging feat. Currently available observational data do not include polarization. However, the angular distribution of polarized emissions is known to be anisotropic \citep{1973RvMP...45..553F,2000JPhB...33.5091T}. If the CX X-ray emissions are strongly polarized, the total intensities measured in current satellite observations should be redefined to adjust for the anisotropy. It is realistic to expect that X-ray polarization will be investigated in future space missions, making an accurate theoretical consideration of these phenomena necessary. Currently, polarization data are available from laboratory experiments for some CX collisions of atoms and ions of astrophysical interest. For example, polarization spectroscopy of ${\rm O}^{5+} (1s^2 3p)$ produced in collisions of ${\rm O}^{6+}$ with He and ${\rm H}_2$, for projectile velocities from 740 to 1200 km s$^{-1}$, showed that the CX X-rays are polarized and strongly dependent on the projectile velocity \citep{2000JPhB...33.5091T}. We expect CX X-ray emissions produced in astrophysical environments, such as Jupiter heavy ion auroras \citep{2006GeoRL..3311105K,2008JGRA..11308229K} and coronal mass ejections, to exhibit similar properties. In this Letter, we report the results of our theoretical study of polarization of the X-rays emitted in charge-exchange collisions between fully stripped SW ions and neutral heliospheric hydrogen. In particular, we investigate the dependence of the X-ray polarization on distribution of the SW plasma, heliospheric neutral gas and SW ion velocity.

In this work, we present the first calculations of polarization of the heliospheric X-rays induced by the CX collisions with fully stripped oxygen and carbon ions, O$^{8+}$ and C$^{6+}$. To analyze its velocity and directional dependence, we calculated the polarization in the ecliptic plane as a function of the LOS for ion velocities from 200-2200 km s$^{-1}$, where the heliospheric plasma was described using two models of different complexity \citep{2004A&A...418..143L,2000JGR...10527419M}. Our calculations indicate that the CX heliospheric X-rays are mildly polarized and that the polarization depends on the SW ion velocity. While this study was restricted to C$^{6+}$ and O$^{8+}$ ions in a simplified geometry, it nevertheless, illustrates a rather general property of the CX-induced radiation. In the considered interval of ion velocities colliding ion and atom form an intermediate quasi-molecule whose projection of electronic angular momentum is quantized along the direction of the quasi-molecular axis. In collisions, quasi-molecular axes of different ion-atom pairs rotate until, at the end of encounter, they are oriented along the ion velocity vector. Thus, CX collisions yield an ensemble of aligned excited ions that emit polarized photons. The polarization reflects a directional orientation of local collisional velocities in the regions of the SW plasma that are mostly responsible for production of the X-ray flux. Analogously, we may expect a relatively high level of polarization of the CX X-ray emissions from cometary atmospheres or in Jovian polar X-ray auroras, where the ion velocity co-orientation in regions of the X-ray-producing CX collisions may be very high. Fast CMEs, propagating in the interplanetary gas, could also produce highly polarized CX X-ray emissions, particularly if the X-ray detector is located close to the CME's plasma stream. Polarization measurements could be used as a powerful tool to supplement the CX spectra and provide additional insight into underlying astrophysical processes. An example of such an environment, directly related to this work, is the Local Bubble. Several recent studies backed by new observational data questioned the validity of the current Local Bubble picture, arguing that the contribution of heliospheric X-rays to the soft X-ray background may be higher than previously thought \citep{2001SSRv...97..401R,2007A&A...475..901K,2008AGUFMSH21B1596R,2008ApJ...676..335H,2009ApSS.323....1W}. The polarization measurements could help identify the contribution from the heliospheric CX radiation in the diffuse X-ray background, especially if performed for different levels of solar activity. As a continuation of this work, a detailed map of the polarization of the SW induced X-rays in the Solar System should be constructed using results of an elaborate 3D MHD model of the heliospheric plasma and neutral gas. If computed for different solar conditions, such maps could be used as an independent method for determining velocity and spatial distributions of the solar wind plasma. Finally, the investigation of CX-induced X-ray polarization may be extended to other interesting topics, including the He focusing cone, dependence of polarization on the solar activity, or regions of the SW plasma turbulence \citep{2008ApJ...682.1404B}. VK acknowledges support by NASA grant NNX08AH51G and HRM acknowledges support by NSF grant AST-0607641.

1012.1349

We report results of a theoretical investigation of polarization of the X-ray emissions induced in charge-exchange collisions of fully stripped solar wind (SW) ions C<SUP>6 +</SUP> and O<SUP>8 +</SUP> with the heliospheric hydrogen atoms. The polarization of X-ray emissions has been computed for line-of-sight observations within the ecliptic plane as a function of SW ion velocities, including a range of velocities corresponding to the slow and fast SW, and coronal mass ejections. To determine the variability of polarization of heliospheric X-ray emissions, the polarization has been computed for solar minimum conditions with self-consistent parameters of the SW plasma and heliospheric gas and compared with the polarization calculated for an averaged solar activity. We predict the polarization of charge-exchange X-rays to be between 3% and 8%, depending on the line-of-sight geometry, SW ion velocity, and the selected emission lines.

12168210

2011ApJ...730L..16M

1012

1012.4931_arXiv.txt

\label{introduction} It is widely accepted that metallicity is the main parameter governing the horizontal branch (HB) morphology of globular clusters (GC), i.e.\ metal-rich GC stars preferentially populate the red side of the RR-Lyrae instability strip while metal-poor GCs tend to have blue HBs (e.g.\ Arp et al.\ 1952). However, since the 1960's several observations revealed that some GCs with similar metallicity exhibit different HB morphologies suggesting that a second parameter is needed to fully account for the HB extent (Sandage \& Wallerstein 1960; Sandage \& Wildey 1967). Despite many efforts to understand the nature of the second parameter, it still remains one of the open issues of modern astrophysics. Several solutions have been proposed, including age, stellar mass, cluster central density, cluster mass (see Dotter et al.\ 2010 and references therein). Interestingly, a difference in helium could explain the observed HB morphology in GCs (van den Bergh 1967). In this context, the recent discovery of multiple stellar populations in GCs (Piotto et al.\ 2007; Milone et al.\ 2008, 2010; Marino et al.\ 2009) allows us to look at the second parameter problem from a new point of view, as the multiple HB components of several GCs have been tentatively assigned to different stellar generations by many authors (e.g.\ D'Antona et al.\ 2005, Milone et al.\ 2008). The aim of this paper is to test whether stellar populations may be linked to the HB morphology in the nearby GC M4 (=NGC~6121). This cluster is an ideal target, as both spectroscopic and photometric studies have provided evidence that it experienced multiple star-formation events. Chemical composition analysis of bright red giant branch (RGB) stars showed star-to-star variations in the C, N, O, Na, and Al abundances (Gratton et al.\ 1986, Brown et al.\ 1990, Brown \& Wallerstein 1992, Drake et al.\ 1992), that define a clear CN distribution bimodality (Norris 1981), a Na-O anticorrelation, and a Na-Al correlation (Ivans et al.\ 1999; Marino et al.\ 2008, hereafter M08). A bimodality in the distribution of stars on the Na-O anticorrelation in M4 has been found the recent analysis of high resolution spectra of 105 RGB stars by M08, who identified two stellar populations with a strong dichotomy in Na and O abundance. Na-rich stars are also CN-strong and are more Al-enriched than Na-poor ones. Neither any Mg-Al anticorrelation nor significant star-to-star variations in iron, iron-peak elements, {\it s}-process, or $\alpha$ elements were detected. The presence of two stellar populations in M4 is confirmed by photometry. The two groups of Na-rich/O-poor and Na-poor/O-rich stars define two distinct sequences in the $U$-$(U-B)$ color-magnitude diagram (CMD). Na-rich stars populate a sequence on the red side of the RGB, while Na-poor ones define a bluer, broader sequence. Interestingly, M4 has a bimodal HB, well populated both on the blue and the red side of the RR-Lyrae gap. The existence of a physical connection between the CN-dichotomy observed on the RGB of this cluster and the HB morphology (the so-called second parameter problem) was already suggested by Norris (1981). The first evidence of such a connection came from Smith \& Norris (1993), who found that the majority of the red-HB stars in M4 have a similar strong CN content. Nevertheless, from the Smith \& Norris (1993) study one could tentatively associate the red-HB with the N-rich/C-poor/Na-rich stars, the analysis of one blue-HB star done by Lambert et al.\ (1992), reveals that it could be the counterpart of a C-poor RGB star. We note, however, that Smith \& Norris (1993) expressed caution concerning their absolute measures of the CN abundances, but stated that the CN abundance was similar in the analyzed red-HB stars. In this paper we investigate the possibility that the HB bimodality in M4 could be related with the two stellar populations identified by M08 by investigating chemical abundances directly on HB stars.

\label{discussion} In this paper we studied a sample of HB stars in the GC M4, with the main purpose of measuring their abundance of sodium and oxygen. Our targets exhibit a bimodal Na distribution very similar to that found for RGB stars, with red-HB stars having roughly solar-scaled [Na/Fe] and blue-HB stars being all sodium enhanced. Blue-HB stars are also oxygen depleted, while the two variables are oxygen rich, and sodium poor as well as the red-HB stars. The Na-O anticorrelation of HB stars is almost the same as we found for RGB stars, where we detected two distinct stellar populations with different mean Na and O. {\it On the basis of their position in the Na-O plane, the blue-HB stars can be clearly associated with the second (O-poor/Na-rich) stellar generation,while the red-HB and the RR-Lyrae must be the progeny of the O-rich/Na-poor first generation.} As already mentioned in Sect.~\ref{introduction}, the Na-O anticorrelation is a powerful tracer of the star formation history in a GC. It has been observed among unevolved stars in GCs (Gratton et al.\ 2001), demonstrating that this pattern is not the result of some mixing process but is rather produced by the ejecta of a first stellar generation. It is not a property of a few peculiar objects but has been observed in all the GCs studied so far. While it is now widely accepted that it indicates the presence of material gone through high temperature H-burning processes, which include the CNO, NeNa, and MgAl cycles, the nature of the polluters is still unclear. These processes might have occurred in intermediate-mass asymptotic giant branch stars (D'Antona et al.\ 2002) or in fast-rotating massive stars (Decressin et al.\ 2007). In any case, helium is the main product of H-burning and is expected to be related to Na and O abundance, with stars that are Na-enhanced and O-depleted being also He-enriched. According to this scenario, the present day He-rich (Na-rich, O-poor) stars should be less massive than He-poor (Na-poor, O rich) ones because the latter evolve more slowly. As a consequence of this, the blue-HB should consist of He/Na-rich, O-poor stars, while the He/Na-poor, O-rich stellar population should end up on the red-HB (D'Antona et al.\ 2002). The abundances of sodium and oxygen in HB stars presented in this paper strongly support this picture. Our results show that in M4 the spread of light elements is strongly correlated with the HB morphology. It is tempting to affirm that, at least in this cluster, this spread is the {\it second parameter}, with He possibly being the main driver of its HB morphology after Fe. {\it Acknowledgements} We thank the referee whose suggestions have improved significantly the paper, M.\ Bergemann, L.\ Mashonkina, F.\ D'Antona for useful discussions. AFM, APM, and GP acknowledge the support by MIUR under the PRIN2007 (prot.20075TP5K9).

1012.4931

The horizontal branch (HB) morphology of globular clusters (GCs) is mainly governed by metallicity. The second parameter problem, well known since the 1960s, states that metallicity alone is not enough to describe the observed HB morphology of many GCs. Despite many efforts to resolve this issue, the second parameter phenomenon still remains without a satisfactory explanation. We have analyzed blue-HB, red-HB, and RR-Lyrae stars in the GC M4 and studied their Fe, Na, and O abundances. Our goal is to investigate possible connections between the bimodal HB of M4 and the chemical signatures of the two stellar populations recently discovered among red giants of this cluster. We obtained FLAMES-UVES/GIRAFFE spectra of a sample of 22 stars covering the HB from the red to the blue region. While iron has the same abundance in both the red-HB and blue-HB segments, the red-HB is composed of stars with scaled-solar sodium abundances, while the blue-HB stars are all sodium enhanced and oxygen depleted. The RR-Lyrae are Na-poor, as the red-HB stars, and O-rich. This is what we expect if the blue-HB consists of a second generation of stars formed from the ejecta produced by an earlier stellar population through high-temperature hydrogen-burning processes that include the CNO, NeNa, and MgAl cycles and are therefore expected to be He-rich. According to this scenario, the sodium and oxygen pattern detected in the blue-HB and red-HB segments suggests helium as the second parameter that rules the HB morphology in M4. <P />Based on data collected at the European Southern Observatory with the VLT-UT2, Paranal, Chile.

11492452

2011PhRvD..83f4010B

1012

1012.2456_arXiv.txt

In the last few years numerical and analytical relativity have demonstrated how to use information from the strong-field--fast-motion regime of coalescing black-hole binaries to build accurate analytical models of their dynamics and of the gravitational radiation emitted~\cite{Buonanno:2007pf,Damour:2007yf,Damour:2007vq,Damour:2008te,Boyle:2008ge,Damour:2009kr,Damour:2009ic,Buonanno:2009qa,Damour:2008gu,Damour:2009sm,Pan:2009wj}. Although numerical-relativity (NR) simulations of binary black holes have reached a high degree of accuracy and flexibility~\cite{Pretorius:2005gq,Gonzalez:2008bi,Pollney:2009yz,Lousto:2010tb,Lousto:2010ut}, a comprehensive spanning of the multidimensional parameter space remains prohibitive. Analytical models are thus of fundamental importance to set up the bank of gravitational wave (GW) templates for detection. The limiting case is given by extreme-mass-ratio inspirals (EMRIs) and mergers; NR simulations simply can not access such regime and post-Newtonian (PN) techniques are inaccurate at such velocities. We need analytical models for the GW emission from EMRI systems because they are primary target sources for the Laser Interferometer Space Antenna (LISA) and because their parameter space is very large~\cite{Gair:2004iv}. The only analytical approach currently capable of accurately following the complete dynamics and providing waveforms (inspiral-plunge-merger-ringdown) of coalescing black-hole binaries is the effective-one-body (EOB) approach to the general relativistic two-body dynamics~\cite{Buonanno:1998gg,Buonanno:2000ef,Damour:2000we,Damour:2001tu,Buonanno:2005xu,Damour:2009ic}. The EOB formalism employs \emph{resummed} PN results (for dynamics and waveforms) in order to extend their validity in the strong-field--fast-motion regime, \ie in a region where they are inaccurate in their standard Taylor-expanded form. In brief the analytical construction is based on (i) a dynamics governed by a resummed Hamiltonian and an expression for the mechanical angular momentum loss (the \emph{radiation reaction}) and (ii) a waveform-generating algorithm which combines a prescription to resum the Taylor-expanded PN multipolar waveform up to the merger and a matching procedure to the quasinormal-mode (QNM) waveform to describe the postmerger phase (an oscillating black hole). One key aspect of the EOB approach is its {\it flexibility}~\cite{Damour:2002vi}. Although the formalism is based on analytical results known only at a given PN order, it is possible to take into account (yet uncalculated) higher-order effects by means of suitable {\it flexibility parameters}. These parameters may be determined (or just constrained) by comparison with results from numerical-relativity simulations valid in the strong-field--fast-motion regime. Several recent works~\cite{Damour:2007yf,Damour:2007vq,Damour:2008gu, Damour:2009kr,Buonanno:2009qa,Pan:2009wj} have shown how this tuning can be implemented to obtain analytical waveforms that match the numerical ones within numerical errors. The tuned EOB formalism can then be used for parametric studies. The Regge-Wheeler-Zerilli (RWZ) metric perturbation theory~\cite{Regge:1957td,Zerilli:1970se,Martel:2005ir,Nagar:2005ea}. is the natural tool to compute the GW emission from a system of two nonspinning black holes, of masses $\mu$ and $M$, in the extreme-mass-ratio limit (EMRL) $\nu\equiv \mu/M\ll 1$. In this regime several numerical results can be used to calibrate the EOB dynamics and waveforms~\cite{Damour:1997ub,Damour:2007xr,Damour:2008gu,Yunes:2009ef,Yunes:2010zj}. In particular, recent gravitational--self-force calculations~\cite{Barack:2009ey,Barack:2010tm} helped in putting constraints on the functions entering the EOB conservative dynamics~\cite{Damour:2009sm,Barack:2010ny}. The Regge-Wheeler-Zerilli perturbation theory has been used for many years in the {\it Fourier domain} (see, for example, \cite{Berti:2010ce} and references therein) and {\it neglecting radiation-reaction effects}, since Davis, Ruffini, and Tiomno computed the waveform emitted by a particle radially plunging into the black hole~\cite{Davis:1972ud}. Only recently the RWZ approach has been extensively developed in the {\it time domain}~\cite{Martel:2001yf,Martel:2003jj,Nagar:2004ns,Nagar:2005cj, Sopuerta:2005gz,Canizares:2010ah} with the inclusion of the radiation-reaction force~\cite{Nagar:2006xv,Barack:2009ey,Barack:2010tm,Sundararajan:2010sr,Bernuzzi:2010ty}. Time-domain simulations using the perturbation theory are efficient and accurate and complement NR simulations in EMRL. The first calculation of the complete gravitational waveform emitted during the transition from inspiral to plunge, merger and ringdown in the EMRL was performed in Ref.~\cite{Nagar:2006xv}, thanks to the combination of the RWZ perturbation theory and the 2.5PN accurate (analytical) Pad\'e-resummed radiation-reaction force~\cite{Damour:1997ub}. Reference~\cite{Damour:2007xr} used that result as a target waveform to assess the performances of the corresponding EOB (resummed) analytical waveform. The comparison was restricted to the quadrupole case, $m=\ell=2$. The knowledge from that study was useful in subsequent EOB/NR waveform comparisons. The treatment of the analytical radiation reaction in the strong-field--fast-motion regime has been improved since then, thanks to a resummed and factorized form of the PN multipolar waveform~\cite{Damour:2007xr,Damour:2008gu,Fujita:2010xj}. In~\cite{Bernuzzi:2010ty} (hereafter paper I) two of us presented an accurate computation of the gravitational radiation generated by the coalescence of two circularized nonspinning black holes in the EMRL. The results were obtained with an improved version of the finite-difference code of~\cite{Nagar:2006xv,Damour:2007xr}, which implements the expression of the radiation-reaction force based on the (5PN-accurate) analytical waveform resummation of~\cite{Damour:2008gu}. The knowledge of the ``exact'' RWZ multipolar waveform opened the way to two main conclusions, extensively discussed in paper I: first, the computation of the final kick velocity imparted to the system by GW emission, $v^{\rm kick}/(c\nu^2)=0.0446$. This value proved consistent with the corresponding one extrapolated from a sample of numerical-relativity simulations~\cite{Gonzalez:2006md} (see Fig.~7 and Tables~IV and V in paper I), as well as with the outcome of an independent calculation that relies on a different treatment of the radiation reaction~\cite{Sundararajan:2010sr}. Second, it was possible to show a very good agreement (at the $10^{-3}$ level) between the mechanical angular momentum loss provided by the analytical expression of the radiation reaction and the GW angular momentum flux computed from the RWZ waveforms. This second result supports the consistency of our approach. Notably, the agreement between the two functions was excellent also {\it below} the last stable orbit (LSO) and almost along the entire plunge phase up to merger (see Figs.~8 and 9 of paper I). The results of~\cite{Bernuzzi:2010ty} also turned out to be compatible with the first NR computation of binary black hole coalescence in the large-mass-ratio regime (1:100)~\cite{Lousto:2010ut}. Recently~\cite{BNZ:2010} we further improved the RWZ approach of paper I by combining it with the hyperboloidal layer method~\cite{Zenginoglu:2010cq}. This approach brings two main benefits. First, it allows us to extract GWs at null infinity ($\I^+$), thereby eliminating the gauge effects related to the GW extraction at a finite radius. In addition, because we evolve the RWZ equations on a smaller coordinate domain, we substantially improve the efficiency of our code. The aim of this paper is to perform a detailed comparison, multipole by multipole, between the RWZ and the corresponding analytical waveforms computed within the EOB approach. For the particle dynamics both codes (RWZ and EOB) implement the resummed radiation-reaction force ${\cal F}_\varphi$ of~\cite{Damour:2008gu} updated to include 5PN-accurate terms also for subdominant multipoles. The latter come from the 5.5PN-accurate (Taylor-expanded) circularized multipolar waveform computed by Fujita and Iyer~\cite{Fujita:2010xj}. The particle dynamics is computed within this 5PN-accurate (resummed) approximation and is the same in both codes. For simplicity, we decided not to improve it further by tuning the resummed flux entering the radiation reaction~\cite{Yunes:2009ef,Yunes:2010zj}. For the waveform we compare the full multipolar structure up to $m=\ell=4$, going beyond the simple quadrupole contribution. The waveform comparison brings new knowledge with respect to the flux comparison of paper I (see also~\cite{Damour:2007xr}) for two main reasons. First, we assess the performance of the resummed EOB waveform in describing the {\it phase} of each multipole. Second, we perform detailed analyses of the next-to-quasicircular (NQC) corrections that are needed in the late-plunge phase. NQC effects are actually responsible for the differences in the EOB and RWZ fluxes in the strong-field--fast-motion regime, as it was pointed out in paper I (see Fig.~8 there and also the related discussion in~\cite{Damour:2007xr}). At the waveform level, several studies~\cite{Damour:2007xr,Damour:2009kr,Buonanno:2009qa,Pan:2009wj} have demonstrated that NQC corrections to the EOB waveform (and radiation reaction) are needed to improve its agreement with the numerical one during the late-plunge and merger phase. Previous works were restricted to the quadrupole case. Two central new benefits of this paper are (i) the assessment of the complete multipolar EOB waveform in the EMRL during the transition from inspiral to plunge and merger; and (ii) the development of a robust procedure to tune NQC corrections to the gravitational wave amplitude and {\it phase}. The paper is organized as follows. In Sec.~\ref{sec:RWZ} we summarize the main features of our RWZ numerical target waveform described in detail elsewhere~\cite{Bernuzzi:2010ty,BNZ:2010}. In Sec.~\ref{sec:EOBwaves} we describe the structure of the multipolar EOB waveform of~\cite{Damour:2008gu,Fujita:2010xj}, giving all the details of the implementation used here. In Sec.~\ref{sec:compare} we first present our results for the inspiral phase and then describe the procedure to tune NQC parameters necessary to improve the EOB waveform at merger. The discussion is based mainly on the $\ell=2$ multipoles. In Sec.~\ref{sec:htotal} we assess the quality of the complete EOB multipolar waveform, discussing explicitly multipoles up to $\ell=4$. We finally put together some concluding remarks in Sec.~\ref{sec:conclusions}. Two appendixes are included to complement the information given in the main text. Throughout this paper we use geometrized units with $c=G=1$.

\label{sec:conclusions} In this paper we have discussed the properties of the gravitational radiation emitted during the transition from quasicircular inspiral to plunge of two nonspinning black holes in the EMRL within the EOB framework. We considered for the first time the whole multipolar structure of EOB-resummed waveforms and we compared and calibrated them against recently calculated Regge-Wheeler-Zerilli numerical waveforms~\cite{Bernuzzi:2010ty,BNZ:2010}. The target numerical waveforms are extracted at null infinity via the hyperboloidal layer method~\cite{BNZ:2010,Zenginoglu:2007jw,Zenginoglu:2009ey,Zenginoglu:2010cq}. The binary dynamics is modeled in both cases for a point particle moving on a Schwarzschild background under the action of ${\cal O}(\nu)$ dissipative radiation-reaction force computed using analytically resummed 5PN results~\cite{Damour:2008gu,Fujita:2010xj}. As a paradigmatic example, we have considered a binary with mass ratio $\nu=10^{-3}$, initially at a relative separation of $7M$, that inspirals for about 37 orbits before plunging into the black hole. The setup of our point-particle ``laboratory'' is sufficiently general to allow us to gather information that can be useful both in the analytical models of waveforms emitted by EMRI (target sources for LISA) as well as for the coalescence of comparable mass black holes (target sources for ground-based interferometers). Our results can be summarized as follows. \emph{Quasiadiabatic inspiral.---}At the beginning of the inspiral, the phase difference between the complete EOB and RWZ waveforms (computed without allowing for any relative time and phase shift) is very small: $\Delta\phi^{\rm EOBRWZ}\sim 5\times 10^{-4}$ rad. This value is consistent with the estimated uncertainty related to the residual phases $\delta_{\ell m}$ entering the EOB waveform known only up to 4.5PN level. During the $\sim 33$ orbits of the inspiral after the junk radiation ($t<500M$) up to the LSO crossing (corresponding to $\sim 420$ rad of total GW phase) the system accumulates only $-2.5\times 10^{-3}$ rad, \ie $\Delta\phi^{\rm EOBRWZ}/\phi^{\rm RWZ}=-5.95\times 10^{-6}$. Such remarkable phase coherence that is obtained with the EOB insplunge waveform ---without any tunable parameter--- strongly indicates the aptitude of EOB waveforms to model EMRIs for LISA. Our conclusions are compatible with those of Refs.~\cite{Yunes:2009ef,Yunes:2010zj} although there are two important differences. First, our two waveforms are computed from the same dynamics in order to focus {\it only} on waveform comparison so to test the efficiency of the resummation of the EOB waveform. Second, we do not further calibrate the $\nu=0$ EOB-resummed flux (and thus the radiation reaction ${\cal F}_\varphi$) to circularized exact data~\cite{Fujita:2009us}. We believe that an additional tuning of higher PN contributions to the flux, though necessary for dynamical accuracy, would have only a marginal influence on our results. We finally remark that our setup and the accuracy of our data permitted us to assess the quality of the approximation to the $\delta_{\ell m}$'s residual phases entering the EOB waveform. Note that we have used the $\delta_{\ell m}$'s in their standard Taylor-expanded form~\cite{Damour:2008gu,Fujita:2010xj} and we have not attempted to further resum them using nonpolynomial expressions. This might certainly be interesting to explore to further reduce the (small) phase gap we have at the beginning of the evolution. \emph{Transition from inspiral to plunge, merger and ringdown.---}For the first time we have explored the impact of NQC corrections to the complete multipolar waveform, including higher multipoles with $2\leq \ell\leq 4$. The addition of NQC corrections is important to improve the EOB and RWZ modulus and phase agreement towards merger. We have proposed a simple procedure to determine NQC corrections on both the phase and amplitude for each multipole; four parameters are required, two for the amplitude and two for the phase. They are determined by imposing compatibility between EOB and RWZ waveform amplitude, frequency, and their first derivatives at the light ring, \ie the maximum of the orbital frequency. The procedure is robust and applies directly to all multipoles (including those with $\ell>4$, that we have not explicitly discussed in the text). The complete EOB gravitational waveform (summed up to $\ell=4$) shows a remarkably good phasing and amplitude agreement with the numerical one up to merger ($\pm 0.015$ rad). After the light ring crossing we have a total phase difference of about $0.25$ due to the approximate treatment of the ringdown (via matching to QNMs), although it mainly oscillates around zero. The maximum relative amplitude difference is about $2.5\%$ just before the light-ring. We emphasize that the exquisite phase agreement that we find at merger crucially relies on the calibration of NQC corrections to the phase. \begin{figure*}[t] \begin{center} \includegraphics[width=0.45\textwidth]{fig08a.eps} \hspace{5 mm} \includegraphics[width=0.45\textwidth]{fig08b.eps}\\ \vspace{5 mm} \includegraphics[width=0.45\textwidth]{fig08c.eps} \hspace{5 mm} \includegraphics[width=0.45\textwidth]{fig08d.eps} \caption{\label{fig:l2shifted} Improvement of the EOB $\ell=2$ ringdown waveform when matching the QNMs at $t_{\rm match}=t_{\rm LR}+3M$. As in Fig.~\ref{fig:l2}, the (light) dashed lines refer to the bare insplunge waveform, without the addition of NQC corrections (dash-dotted line, blue online) nor of QNM ringdown (dashed line, red online). The vertical dashed line locates the maximum of $M\Omega$.} \end{center} \end{figure*} The procedure discussed here to determine the NQC corrections can be directly applied to EOB and NR comparisons for comparable mass ratios generalizing current techniques. However, when $\nu\neq 0$, the procedure is more complicated due to the dependence of the EOB dynamics, notably of the Hamiltonian, on other flexibility parameters that are also required to be determined (or constrained) by NR data. In particular, one of the most evident physical effects entailed by these corrections on the dynamics is to displace the location of the ``EOB-light ring'' (\ie the maximum of $\Omega$) and thus the location of the matching time $t_m$. In our $\nu=0$ setting we can identify on the waveform unambiguously the time $t_{\rm LR}$ that corresponds to the crossing of the light ring, because of the very good ($\sim 10^{-4}$ rad) phase alignment of the waveforms at early times and because the underlying dynamics is the same. This allows us to measure the useful RWZ information at the correct location. As we emphasized in the text $t_m$ {\it does not} coincide with the time when $A_{22}$ peaks, but it occurs $2.56M$ earlier. On the other hand, when dealing with NR data, the exact dynamics is not evidently available and one can rely only on waveform information. The peak of the exact orbital frequency, if it existed, should occur slightly after the peak of the $A_{22}$ metric waveform amplitude. As a consequence, to apply the same discussed here to fix the NQC parameters and to keep the maximum of $\Omega$ as the matching point, one should measure the four numbers per multipoles slightly (say by $\sim 1M$) after the peak of $A_{22}$. This method is different from current methods in EOB and NR comparisons, \ie fix the NQC amplitude corrections imposing that the EOB and NR $A_{22}$ peaks {\it coincide} at the maximum of the EOB orbital frequency. Even if this procedure is not \emph{a priori} incorrect, we stress that if we were following this prescription in our setup we would have obtained a significantly larger phase difference ( $\sim~-0.2$ rad), accumulated (starting from the $10^{-2}$ level) in the last $50M$ before merger. This suggests that a more detailed analysis of the impact of NQC corrections to EOB waveforms in the comparable mass case might be needed in the future. As a last remark, we emphasize that extraction of numerical waveforms at null infinity convinced us to avoid any further (arbitrary) phase and time shift, providing clean information about the accuracy of the analytical modeling of the GW phase in the EOB waveform. On the contrary, we have shown that waveforms extracted at the finite radius $r_*=1000M$ yield initial phase differences (with the EOB ones) during the inspiral that are $\sim 3\times 10^{-2}$ rad for the (2,2) multipole and about twice as much, $\sim 5.5\times 10^{-2}$ rad, for the (2,1) multipole (and even larger for the subdominant multipoles~\cite{BNZ:2010}). This fact strongly indicates, once more, that in any EOB and NR comparison it is {\it necessary} to work either with NR waveforms extrapolated to infinite extraction radius or evolved up to null infinity~\cite{Reisswig:2009us,Reisswig:2009rx}.

1012.2456

We discuss the properties of the effective-one-body (EOB) multipolar gravitational waveform emitted by nonspinning black-hole binaries of masses μ and M in the extreme-mass-ratio limit μ/M=ν≪1. We focus on the transition from quasicircular inspiral to plunge, merger, and ringdown. We compare the EOB waveform to a Regge-Wheeler-Zerilli waveform computed using the hyperboloidal layer method and extracted at null infinity. Because the EOB waveform keeps track analytically of most phase differences in the early inspiral, we do not allow for any arbitrary time or phase shift between the waveforms. The dynamics of the particle, common to both wave-generation formalisms, is driven by a leading-order O(ν) analytically resummed radiation reaction. The EOB and the Regge-Wheeler-Zerilli waveforms have an initial dephasing of about 5×10<SUP>-4</SUP>rad and maintain then a remarkably accurate phase coherence during the long inspiral (∼33 orbits), accumulating only about -2×10<SUP>-3</SUP>rad until the last stable orbit, i.e. Δϕ/ϕ∼-5.95×10<SUP>-6</SUP>. We obtain such accuracy without calibrating the analytically resummed EOB waveform to numerical data, which indicates the aptitude of the EOB waveform for studies concerning the Laser Interferometer Space Antenna. We then improve the behavior of the EOB waveform around merger by introducing and tuning next-to-quasicircular corrections in both the gravitational wave amplitude and phase. For each multipole we tune only four next-to-quasicircular parameters by requiring compatibility between EOB and Regge-Wheeler-Zerilli waveforms at the light ring. The resulting phase difference around the merger time is as small as ±0.015rad, with a fractional amplitude agreement of 2.5%. This suggest that next-to-quasicircular corrections to the phase can be a useful ingredient in comparisons between EOB and numerical-relativity waveforms.

12163023

2011A&A...526L...5D

1012

1012.0179_arXiv.txt

Mid-infrared stellar interferometry of young massive stars has the potential to bring significant new insights into the process of massive star formation. The technique overcomes the two age-old problems of low angular resolution and high extinction. In de Wit et al. (2007\nocite{2007ApJ...671L.169D}, 2010\nocite{2010A&A...515A..45D}), we presented spectrally dispersed mid-infrared (mid-IR) interferometric observations of the massive young stellar object (MYSO) W33A using the Very Large Telescope Interferometer (VLTI). We found that the $N$-band emission is dominated by warm dust on 100\,AU scales, located in the walls of the outflow cavity. Recently, VLTI observations in the near-infrared (near-IR) using the AMBER instrument show unambiguously the presence of a disk-like structure on size scales of $\sim 15$\,AU (Kraus et al. 2010\nocite{2010Natur.466..339K}). In this letter, we focus on the MYSO CRL\,2136. The dominant infrared source IRS\,1 ($L\sim7\,10^{4}L_{\odot}$) is the driving force of an arcminute scale bipolar CO outflow with a P.A. of $\sim 135\degr$ (Kastner et al. 1994\nocite{1994ApJ...425..695K}). Weak, optically thick radio emission originating from IRS\,1 was detected by Menten \& Van der Tak (2004\nocite{2004A&A...414..289M}). Near-IR polarimetric observations demonstrate the presence of a polarization disk (Minchin et al. 1991\nocite{1991MNRAS.251..522M}; Murakawa et al. 2008\nocite{2008A&A...490..673M}). These credentials render CRL\,2136 a promising target to probe for the geometry of the material on milli-arcsecond angular scales in the harsh environment close to the stellar surface. We perform VLTI observations in the $N$-band, where MYSOs are bright and warm dust in the envelope and possibly in a circumstellar disk should still contribute significantly to the total flux.

\label{conclusions} We have analysed mid-IR interferometric observations of the massive young stellar object CRL 2136 obtained with the VLTI and MIDI. The dispersed visibilities show that there is a strong change in character of the emitting region at $\lambda=8.5\,\mu$m. We have found that either a cool star or an accretion disk interior to a dusty envelope can explain the presented spatial and spectral observations. At present, neither scenario can be excluded based on the single MIDI baseline discussed in this paper, but clearly multi baselines will help to distinguish between them. We confidently conclude that the rim of the dust envelope is found at about 7 times the formal dust sublimation radius. The central dust-free zone may have been evacuated by the ionized stellar wind seen by Menten \& van der Tak (2004\nocite{}), which occupies a similar volume. Our present work has illustrated the potential of MIDI to detect compact emission in massive YSOs, which is of particular interest given the large numbers of massive YSOs that can be studied in this way. Mid-IR interferometry can therefore provide a significant contribution to the characterisation of the accretion environment of massive young stellar objects.

1012.0179

Context. Establishing the importance of circumstellar disks and their properties is crucial to fully understand massive star formation. <BR /> Aims: We aim to spatially resolve the various components that make-up the accretion environment of a massive young stellar object (⪉100 AU), and reproduce the emission from near-infrared to millimeter wavelengths using radiative transfer codes. <BR /> Methods: We apply mid-infrared spectro-interferometry to the massive young stellar object CRL 2136. The observations were performed with the Very Large Telescope Interferometer and the MIDI instrument at a 42 m baseline probing angular scales of 50 milli-arcseconds. We model the observed visibilities in parallel with diffraction-limited images at both 24.5 μm and in the N-band (with resolutions of 0.6´´and 0.3´´, respectively), as well as the spectral energy distribution. <BR /> Results: The arcsec-scale spatial information reveals the well-resolved emission from the dusty envelope. By simultaneously modelling the spatial and spectral data, we find that the bulk of the dust emission occurs at several dust sublimation radii (approximately 170 AU). This reproduces the high mid-infrared fluxes and at the same time the low visibilities observed in the MIDI data for wavelengths longward of 8.5 μm. However, shortward of this wavelength the visibility data show a sharp up-turn indicative of compact emission. We discuss various potential sources of this emission. We exclude a dust disk being responsible for the observed spectral imprint on the visibilities. A cool supergiant star and an accretion disk are considered and both shown to be viable origins of the compact mid-infrared emission. <BR /> Conclusions: We propose that CRL 2136 is embedded in a dusty envelope, which truncates at several times the dust sublimation radius. A dust torus is manifest in the equatorial region. We find that the spectro-interferometric N-band signal can be reproduced by either a gaseous disk or a bloated central star. If the disk extends to the stellar surface, it accretes at a rate of 3.0 × 10<SUP>-3</SUP> M<SUB>⊙</SUB> yr<SUP>-1</SUP>. <P />Based on observations with the VLTI, proposal 381.C-0607.

4862499

2011A&A...527A..50E

1012

1012.2383_arXiv.txt

Plans are well advanced for the next generation of optical and infrared (IR) ground-based telescopes, the extremely large telescopes (ELTs). Their science cases are broad, ranging from direct imaging of exoplanets, to studies of resolved stellar populations in external galaxies, and spectroscopy of distant `first light' galaxies at the highest redshifts \citep[e.g.][]{hdg06}. The ELTs are an increasingly global effort, with three projects under detailed study -- the Giant Magellan Telescope (GMT), the Thirty Meter Telescope (TMT), and the European Extremely Large Telescope (E-ELT). With primary apertures in excess of 20\,m, they will deliver a huge gain in our capabilities via a combination of unprecedented sensitivity and exquisite angular resolution. In parallel to the design of the observatories, significant effort has also been invested in studies of ELT instrumentation \citep[see][]{j10,s10,skr10}. Over the past few years, deep imaging from ground-based telescopes and the {\em Hubble Space Telescope (HST)} has provided us with new and unique views of the outer regions of large galaxies beyond the Milky Way for the first time, such as M31 \citep[e.g.][]{f05} and M33 \citep[e.g.][]{bark07}. From analysis of the resulting photometry and colour-magnitude diagrams we can determine star-formation and assembly histories for external galaxies, enabling tests of theoretical models of galaxy evolution \citep[e.g.][]{bj05}, i.e. using resolved stellar populations as a tracer of the processes which have shaped the evolution of their host systems. The imaging from the Advanced Camera for Surveys (ACS) Nearby Galaxy Survey Treasury \citep[ANGST;][]{angst} provides an excellent illustration of the diversity of galaxies beyond the Local Group. Crucially, the Local Volume includes a wide variety of morphological types -- massive ellipticals, large metal-poor irregulars, lower-mass late-type spirals, interacting systems, dwarf starbursts -- providing an excellent opportunity to quantify the effects of environment on galaxy evolution. Photometric methods are powerful when applied to extragalactic stellar populations but follow-up spectroscopy can aid our understanding significantly via precise chemical abundances and stellar kinematics. For example, results from the spectroscopy of luminous blue supergiants in external galaxies obtained by the Araucaria project \citep{araucaria}. However, in pursuit of spectroscopy of evolved stellar populations, the 8-10m class telescopes are already near their limits beyond a few hundred kpc. For instance, observations with Keck-DEIMOS struggled to yield useful spectra below the tip of the red giant branch (TRGB) in M31 at $I$~$>$~21.5 \citep{c06}. If we aspire to spectroscopy of individual evolved stars in galaxies beyond the Local Group, we require the increased sensitivity of the ELTs, likely combined with some degree of correction from adaptive optics (AO) to mitigate the effects of crowding. Indeed, for stellar spectroscopy with the ELTs there will be a fine balance in sensitivity between the improved image quality from AO as one goes to longer wavelengths (where the wavefront errors become less significant compared to the observed wavelengths) versus the increased sky background. Moreover, work over the past decades has provided an excellent understanding of many of the optical spectral lines available to us, enabling robust estimates of physical parameters and chemical abundances. To exploit the best performance (in terms of angular resolution) from the ELTs, we need to improve our knowledge of near-IR diagnostics, and also to pin-down the `sweet spot' in terms of the gain in sensitivity from AO versus the background contribution. A recent study by Davies, Kudritzki \& Figer (\citeyear{dkf10}; hereafter DKF10) suggested new $J$-band diagnostics, spanning 1.15-1.22\,$\mu$m (including lines from Mg, Si, Ti, and Fe), as a means to obtain stellar metallicities in extra-galactic red supergiants (RSGs). This wavelength region is relatively unexplored in this regard, only previously considered by \citet{o04}. DKF10 noted that this method could be very powerful with new IR instruments under construction for 8-10\,m class telescopes (such as VLT-KMOS and Keck-MOSFIRE), providing precise stellar abundances in galaxies out to a potential distance of $\sim$10\,Mpc, complementing those from luminous blue supergiants \citep[e.g.][]{b01,kud08}. An even more compelling future prospect is in the context of ELT observations, with DKF10 noting the potential of direct abundance estimates for individual RSGs in galaxies even beyond the Local Volume (subject to crowding). A second possible application of this spectral region is in observations of evolved red giant branch (RGB) stars -- the long-lived descendants of much lower-mass stars. The {\sc marcs} model atmospheres \citep{g08} used by DKF10 are actually of red giants (see their Section~2.2 for a discussion of the applicability of these models over a large range of stellar luminosities). The potential of a direct abundance diagnostic for $J$-band ELT observations of both young (RSG) and old (RGB) stellar populations in external galaxies -- with the benefit of better AO correction compared to shorter wavelengths -- warrants further exploration. EAGLE is a conceptual design of an AO-corrected, multi-IFU, near-IR spectrograph, undertaken as one of the Phase~A instrument studies for the E-ELT \citep{cuby10}. In this article we employ tools developed in the course of the EAGLE study as a proof-of-concept for quantitative $J$-band spectroscopy with the ELTs. We simulate EAGLE $J$-band observations with two objectives. First, to validate the technique of DKF10 for metal-poor spectral templates, and secondly, to explore the distances to which we could obtain robust abundance estimates (for both RGB stars and RSGs) with ELT observations. To compare these results with diagnostics already in common use, we also consider simulations in the $I$ band. Section~\ref{models} describes the tools and assumptions used in the simulations, with the analysis presented in Section~\ref{analysis}. In Sections~\ref{discussion} and \ref{rsgs} we discuss the results in the context of the scientific potential of the E-ELT and other upcoming facilities.

\label{discussion} We now consider the simulation results in the context of the (unreddened) distances to which we can recover accurate metallicities from $J$- and $I$-band spectroscopy with the E-ELT. For reference, absolute $I$-band magnitudes ($M_I$) for RGB stars span $-$1\,$<$\,$M_I$\,$<$\,$-$3, with typical colours of ($I - J$)$_{\rm 0}$\,$\sim$\,0.7-0.8 \citep[e.g.][]{g02,m08}. At the tip of the RGB absolute magnitudes extend up to $M_I$\,$=$\,$-$4. \subsection{Potential for E-ELT \& EAGLE} We require S/N\,$\gtrsim$\,50 for accurate metallicities from the $J$-band. At $R$\,$=$\,4000, the results in Table~\ref{Jsims} provide sufficient S/N down to $J$\,$=$\,22.75 to 23.25, depending on the NGS configuration and seeing. Taking the bright end of the RGB (adopting $M_J$\,$=$\,$-$3.75 from the colours above), this corresponds to an unreddened distance modulus of 27, or a distance of $\sim$2.5\,Mpc. Thus, the DKF10 methods could provide a direct metallicity indicator in RGB stars out to a volume which encompasses a diverse range of galaxies. For instance, as we move beyond the Local Group, there are 32 known galaxies with distances in the range 1.0\,$<$\,d\,$\le$\,2.65\,Mpc \citep{k04}. Most notably these include NGC\,55 and NGC\,300, the two spirals in the closer part of the Sculptor `Group' at 1.9\,Mpc \citep{gieren05,gieren08,p06}. NGC\,55 is seen to have `LMC-like' abundances from spectroscopy of luminous supergiants \citep[][\& in prep]{castro08} and NGC\,300 has a radial oxygen gradient that ranges from 0.1\,dex below Solar to near LMC levels at a radius of $\sim$5\,kpc \citep[][and references therein]{b09}. This range of distances also includes lower-mass galaxies such as Sextans~A, NGC\,3109, and GR\,8, which are all depleted in metals (as traced by oxygen) by approximately 1\,dex \citep[][respectively]{kaufer04,e07,vsh06}. Of course, some galaxies within 2.5\,Mpc, in addition to the faint dwarfs in the Local Group, will have even lower metallicities, providing motivation to further explore these diagnostics in the very metal-poor regime. Scaling these results for TRGB stars ($M_J$\,$\sim$\,$-$5), good metallicities should be recoverable out to distances of 4 to 5\,Mpc, depending on the exact AO correction. This opens-up an even wider range of targets, encompassing over 200 galaxies \citep[][]{k04}, including those in the Centaurus~A Group \citep[e.g.][]{k02}. Direct metallicities for TRGB stars will play an important role in calibration of the distance scale in the local universe, reducing the uncertainties associated with photometric distance determinations from the TRGB method \citep[see discussion by][]{kud10}, which can lead to systematic uncertatinies of $\sim$5\% \citep[e.g.][]{mmf08}. \subsection{Comparison with $I$-band performances} To compare the $J$-band results with the commonly used CaT diagnostic, we now consider the $I$-band performances given in Table~\ref{Isims}. To determine the limiting magnitude for metallicities from our $I$-band simulations, we adopt a requirement of S/N\,$\ge$\,20 (per two-pixel resolution element). \citet{bit08} consider S/N\,$\ge$\,20 (per \AA) as the threshold to obtain metallicities to within 0.1\,dex from CaT spectroscopy at $R$\,$=$\,6500 (from VLT-FLAMES). At the centre of our $I$-band template 1\,\AA\/ is equivalent to 2.33 pixels, so the S/N results in Table~\ref{Isims} are slightly conservative in this regard, but they are also at greater spectral resolving power. However, the dominant factor here is the lower AO performance in the $I$-band than at longer wavelengths. The EE in the simulated PSFs is already considerably lower than in the $J$-band (Table~\ref{psfs}), but the additional error terms are also larger (i.e. a factor of four compared to a factor of two in the $J$-band). Depending on the seeing conditions and NGS configuration, the $I$-band simulations yield S/N\,$\ge$\,20 for 22.1\,$<$\,$I$\,$<$\,23.1, comparable to the $J$-band limit discussed in the previous section (S/N\,$\gtrsim$\,50 at $R$\,$=$\,4000). Given that RGB/TRGB stars are moderately red ($I - J$\,$=$\,0.6-1.0), the DKF10 technique is probing deeper in terms of absolute luminosity of a given star. The $J$-band technique exploits the improved AO correction at longer wavelengths. It requires greater S/N to recover a good estimate of stellar metallicity but, given the final EAGLE instrument throughputs and current expectations of AO performance shortwards of 1\,$\mu$m, outperforms the CaT observations for a fixed exposure time of a given RGB star. In addition to a {\em direct} metallicity measurement (rather than the CaT calibration, albeit thought to be well understood), $J$-band observations will have the advantage of reduced extinction compared to observations at shorter wavelengths. Although we have limited ourselves to consideration of extra-galactic sources, working at longer wavelengths would be attractive for observations in, for example, obscured Galactic clusters. \subsection{Potential for TMT-IRMS} Our simulations can also be considered in the context of the Infra-Red Multi-Slit Spectrometer (IRMS) being developed for the TMT \citep{s10}, which is a clone of the MOSFIRE instrument for the Keck\,I telescope. IRMS will be located behind the TMT multi-conjugate AO system (NFIRAOS), taking advantage of excellent image quality across a 2$'$ field. Its planned wavelength coverage is 0.9-2.5\,$\mu$m and, depending on the slit-widths, the delivered resolving powers will be $R$\,$=$\,3000 to 5000. The minimum (Nyquist sampled) IRMS slit-width will be 160\,mas, i.e. coarser spatial sampling and with a greater sky contribution for a point-source target than from EAGLE. The expected $J$-band {\em encircled} energies for a slit-width of 160\,mas range from approximately 30 to 45\,\%, depending on the position within the 2$'$ field (Dr.~Brent Ellerbroek, private communication). These simulations include some of the additional error terms factored into the discussion of our MOAO PSFs, i.e., when the larger IRMS slit-widths are taken into account, they provide a comparable range of image quality. Assuming that the sensitivity varies as a factor of primary aperture (combined with a greater background flux from a larger effective slit), this corresponds to observations of RGB stars out to distances of $\sim$1.5\,Mpc and stars at the TRGB out to $\sim$3\,Mpc (depending on limitations of the NFIRAOS performance estimates cf. those discussed in Section~\ref{moaopsfs}). Tailored performance simulations are warranted for stellar spectroscopy specific to the TMT design, instrumentation and site (selected to be Mauna Kea in Hawaii). Nevertheless, we highlight the potential of the DKF10 methods for TMT spectroscopy of stellar populations beyond the Local Group. For example, in NGC\,1569, a northern starburst galaxy at 1.95\,Mpc \citep{k04} with a metal abundance that appears to be intermediate to those of the LMC and SMC \citep[\citeauthor{drd97} 1997, \citeauthor{ks97} 1997, cf., for example,][]{t07}. The TMT will also be well placed to explore the wealth of faint targets in northern hemisphere galaxies that are closer to home, e.g., M31 and M33. We have simulated $J$-band EAGLE/E-ELT observations of an M0 giant (and, by extrapolation of the model atmospheres, an M0 supergiant). We extended the work of DKF10 for the 1.15-1.22\,$\mu$m diagnostic region, finding: \begin{itemize} \item{Robust stellar metallicities can be obtained from analysis of $J$-band spectra of metal-poor stars (\logz\,$=$\,$-$1.0).} \item{Accurate metallicities ($\pm$0.1\,dex, at \logz\,$=$\,$-$1.0 and 0.0) require: S/N\,$\ge$\,55 at $R$\,$=$\,4000, and S/N\,$\ge$\,45 at $R$\,$=$\,10000.} \item{The Phase~A design of EAGLE on a 42\,m E-ELT has the potential of {\em direct} metallicity estimates for RGB stars out to distances of $\sim$2.5\,Mpc, and for RSGs well beyond the Virgo and Fornax Clusters (subject to crowding).} \item{For our instrument assumptions and simulations of relative AO performance, the $J$-band diagnostics are slightly more sensitive than the CaT ($\Delta$ mag $\sim$0.5-1.0) for recovering metallicities of a given M0-type star. The $J$-band method also has the advantage of reduced extinction compared to the CaT region.} \end{itemize} The first two of these points are not specific to EAGLE (nor dependent on the simulated PSFs used here), serving as a proof-of-concept for broader application of the methods advanced by DKF10. This part of the $J$-band is relatively free of OH emission lines compared to other parts of the near-IR, enhancing its potential for quantitative spectroscopy of extragalactic stellar populations with ELTs. This diagnostic range also has great potential for other upcoming facilities such as VLT-KMOS, {\em JWST}-NIRSpec, and TMT-IRMS. One of the prime drivers for inclusion of the $I$-band in the design of near-IR spectrographs for ELTs is to study extragalactic stellar abundances via the CaT. If comparable or better results can be obtained from the $J$-band, inclusion of the $I$-band would become less critical, leading to simpler and potentially cheaper instrument designs (e.g. fewer gratings), and avoiding compromises which may detract from the potential performance at longer wavelengths. As noted by DKF10, an important next step is quantitative analysis of RSGs in the Magellanic Clouds to refine these techniques over a range of metallicities, and to also extend this work to RGB stars in the Clouds. These methods should also be explored at yet lower metallicities (relevant to the study of metal-poor galaxy halos), both via analysis of metal-poor observations, and by use of a larger grid of theoretical calculations. Diagnostics in the $H$- and $K$-bands (in which the AO correction will be even better) also warrant consideration in this context. For instance, use of $K$-band equivalent-width indices to estimate stellar metallicities \citep{f01}, and detailed abundance analysis of red giants \citep{ors02,or04}, and supergiants \citep{d09a,d09b} at greater resolving powers in the $H$-band. \vspace{0.2in} \noindent{\sc acknowledgments:} We thank Fran\c{c}ois Ass\'{e}mat for his work on the EAGLE AO study, Brent Ellerbroek for updated information on the TMT simulations, and the referee for their constructive comments on the manuscript. BD is funded by a fellowship from the Royal Astronomical Society. RPK acknowledges support from the Alexander-von-Humboldt Foundation. MP, JGC and GR acknowledge support from the Agence Nationale de la Recherche under contract ANR-06-BLAN-0191. DF is supported by NASA under award NNG 05-GC37G, through the Long Term Space Astrophysics program, and by a NYSTAR Faculty Development Program grant.

1012.2383

We present simulated J-band spectroscopy of red giants and supergiants with a 42 m European Extremely Large Telescope (E-ELT), using tools developed toward the EAGLE Phase A instrument study. The simulated spectra are used to demonstrate the validity of the 1.15-1.22 μm region to recover accurate stellar metallicities from Solar and metal-poor (one tenth Solar) spectral templates. From tests at spectral resolving powers of four and ten thousand, we require continuum signal-to-noise ratios in excess of 50 (per two-pixel resolution element) to recover the input metallicity to within 0.1 dex. We highlight the potential of direct estimates of stellar metallicites (over the range - 1 &lt; [Fe/H] &lt; 0) of red giants with the E-ELT, reaching out to distances of ~ 5 Mpc for stars near the tip of the red giant branch. The same simulations are also used to illustrate the potential for quantitative spectroscopy of red supergiants beyond the Local Volume to tens of Mpc. Calcium triplet observations in the I-band are also simulated to provide a comparison with contemporary techniques. Assuming the EAGLE instrument parameters and simulated performances from adaptive optics, the J-band method is more sensitive in terms of recovering metallicity estimates for a given target. This appears very promising for ELT studies of red giants and supergiants, offering a direct metallicity tracer at a wavelength which is less affected by extinction than shortward diagnostics and, via adaptive optics, with better image quality.

12163436

2011A&A...528A..26T

1012

1012.1969.txt

\label{sect_int} %Silicon-containing molecules %formation chemistry depletion in grains The Orion BN/KL \citep{Bec67, Kle67} nebula is one of the most studied star formation regions in the Milky Way. At a distance of 414 pc \citep{Men07} the nebula is embedded in a giant molecular cloud harboring practically all phases of the interstellar medium, from hot and diluted plasma, to PDRs, protostellar cores, molecular outflows, SiO and H$_2$O masering regions, high density cores, intermediate and high mass star formation, protoplanetary disks, and proplydes (see, e.g., \citealp{Gen80, Gen89, Wri95, Cer90, Cer94, Pla09}). Together with Sgr B2, Orion BN/KL nebula exhibits a rich spectrum (see, e.g., Tercero et al., 2010, hereafter Paper I, and references therein) produced by complex organic molecules which are formed through reactions on the grain surfaces during the collapse phase followed by evaporation when radiation from a newly formed star becomes available. Due to the high temperature of the gas, molecular lines are particularly strong in Orion, allowing several line surveys of this source over the last 20 years. Recently, we have performed a line survey towards Orion IRc2 source between 80 and 280 GHz (Paper I), not limited by sensitivity but only by line confusion. The data provide a significant number of transitions for all molecules detected so far towards this source. Although physical structure of the Orion is rather complex, the large number of transitions observed for each species allows to model the different cloud components and to derive reliable physical parameters. In addition, the line survey provides a deep insight on the chemistry of the Orion KL region and allows to refine our knowledge of its chemical structure by searching for new molecular species and new isotopologues and vibrationally excited states of molecules already known to exist in this source. In a first paper we have presented the line survey and analyzed the CS-bearing species deriving the abundance of CS, OCS, CCS, CCCS, H$_2$CS, and HCS$^+$ (Paper I). In this paper we will analyze the Silicon-bearing species SiO and SiS. SiO has been observed with single dishes and interferometers (see, e.g., \citealp{Pla09} and references therein), while only a few observations are available for SiS \citep{Dic81, Ziu88, Ziu91}. SiO lines in Orion show a complex pattern of thermal and maser emission. The masers seem to arise from a small region around the radio continuum \object{source $I$} \citep{Chu87}, a young star with a very high luminosity without infrared counterpart, $\simeq$10$^5$ L$_\odot$, \citep{Gez98, Gre04}. This source is also driving the low % and %high velocity outflow observed in SiO \citep{Beu05, Pla09} and in other molecular species (Paper I). Recent studies have discussed the driving source of the high velocity outflow: whereas \citet{Beu08} claimed that \object{SMA1} (a sub-millimeter source not detected at IR and centimeter wavelengths, predicted by \citealt{deV02} and detected by \citealt{Beu04}) is the host of the high velocity outflow (based on combined observations of $J$=2-1 C$^{18}$O from the SMA and the IRAM 30m telescope), \citet{Pla09} defended that this outflow is a continuation of the low velocity outflow (based on CARMA observations of SiO {\it{v}}=0 $J$=2-1). The SiO $v$=1 maser emission was modeled in early interferometric observations as arising from a rotating and expanding disk \citep{Pla90}. However, maser emission was also found in the $v$=0 $J$=2-1 line by \citet{Wri95} adding more complexity to the modeling of the structure of the emitting source. Recently, \citet{God09} have found maser emission in the $v$=0 $J$=1-0 transition of $^{29}$SiO and $^{30}$SiO associated with source $I$. In addition, observations of \citet{Pla09} clearly show that the emission of this line arise essentially from an outflow driven by source $I$. Although our single dish data cannot provide a view of the spatial structure of the thermal and maser emission around that source, the observed maser and thermal lines can provide useful constraint on the physical conditions of the gas. In addition to the study of SiO and SiS we derive upper limits to the abundance of SiC, SiC$_2$, c-SiC$_3$, SiC$_4$, SiN, SiCN, SiNC, ob-SiC$_3$, l-SiC$_3$, Si$_3$, SiCCO, SiCCS, SiH$_2$, H$_2$CSi, and the different isomers of Si$_2$H$_2$. The observations are described in Sect. \ref{sect_obs}. The results for SiO and SiS are analyzed in Sect. \ref{sect_res}. Section \ref{sect_phy} is devoted to the modeling of the observed lines. All these results are discussed in Sect. \ref{sect_dis} in terms of comparisons with chemical models predictions for Silicon-bearing species. The effect of the new collisional rates for SiO and SiS is analyzed in the Appendix.

\label{sect_dis} %We have observed several rotational lines of Silicon-bearing molecules toward %Orion KL and our models well reproduce observations. SiO is a key tracer of shocked emission. Many interferometric observations show that thermal and maser SiO emission depicted the low velocity outflow centered in source $I$ in Orion KL \citep{Bla96, Wri96, Beu05, Pla09, God09, Zap09}. In addition, SiO traces many other molecular outflows in different sources \citep{Jim04, Gib07, Deb09, Zap09}. \citet{Moo07} have not found emission from SiO at the position of the hot molecular core in G34.26+0.15. This hot core does not have a central source but rather it is externally heated, similar to the Orion compact ridge, by shocks, ionization fronts and stellar winds. \citet{Moo07} pointed out that the lack of SiO in this hot core rules out any significant role played by shocks in determining the hot core chemistry. Observations of SiO in the L1448-mm outflow permit to distinguish between the shock precursor and the postshock components \citep{Jim05}; they observed an enhancement in the abundances of SiO (and another shock tracers) by one order of magnitude in the shock precursor component and three orders of magnitude in the postshock gas (leading to the broadening of the line profiles), evidence of recent ejection of SiO from grains \citep{Flo96}. %\subsection{SiO/SiS Column Density Ratio} %In order to compare the ratio of a column densities, we have to %assumes the region of the line formation is the same for each molecules %and the excitation tempareture is similar for both species. %Therefore we can wonder which %particular conditions could lead to such a high abundance of the SiS %molecule and how it has been formed here. %Chemical models cannot explain this ratio. Silicon is much locked up in dust %grains, and is believed to be easily trapped in SiO molecules and to form few %SiS. %Silicon monosulfide radiative association have been studed by %\cite{And07} and %they found a more rapid formation than SiO radiative association. %XXXXXXXXabundancias con sic, sic3XXXXXXXXXX %\subsection{Silicon chemistry and high temperature} %Most of the available models of the formation of Silicon-bearing molecules in %astrophysical objects focus on SiO and don't predict anything for SiS. It is %the case for shocks models (Schilke 97, Gusdorf 2008). In the last one, %Gusdorf et al. found a high column density value of SiO, as we do, and they %point out that the lines are optically thick in the emission regions of the %shock wave. It seems to be consistent with our model of Orion-KL in which the %emission arrises mostly from the high velocity outflow. \subsection{Molecular abundances} \label{sect_dis_abu} Molecular abundances were derived using the H$_2$ column density calculated by means of the C$^{18}$O column density (1.5$\times$10$^{16}$, 1.5$\times$10$^{16}$, 1$\times$10$^{17}$, 5$\times$10$^{16}$ cm$^{-2}$, and 2$\times$10$^{17}$ cm$^{-2}$ for the extended ridge, compact ridge, plateau, high velocity plateau, and hot core, respectively) and the isotopic abundance $^{16}$O/$^{18}$O=250, both provided in Paper I, assuming that CO is a robust tracer of H$_2$ and therefore their abundance ratio is roughly constant, ranging from CO/H$_2$ $\simeq$ 5$\times$10$^{-5}$ (for the ridge components) to 2$\times$10$^{-4}$ (for the hot core and the plateau). In spite of the large uncertainty in this calculation, we include it as a more intuitive result for the molecules described in the paper. We obtained N(H$_2$) = 7.5$\times$10$^{22}$, 7.5$\times$10$^{22}$, 2.1$\times$10$^{23}$, 6.2$\times$10$^{22}$, and 4.2$\times$10$^{23}$ cm$^{-2}$ for the extended ridge, compact ridge, plateau, high velocity plateau, and hot core, respectively; for the 15.5 km s$^{-1}$ component we assume $N$(H$_2$)=1.0$\times$10$^{23}$ cm$^{-2}$ as an average value in Orion KL. In addition, we assume that the H$_2$ column density spatially coincides with the emission from the species considered. Our estimated source average abundances for each Orion KL component are summarized in Table \ref{tab_abun} (only available online), together with comparison values from other authors (\citealt{Sut95}, \citealt{Per07}, and \citealt{Ziu88}). The differences between the abundances shown in Table \ref{tab_abun} are mostly due to the different H$_2$ column density considered, to the assumed cloud component of the molecular emission and discrepancies in the sizes of these components. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%SERENA VITI \subsection{On the origin of the SiO and SiS emission} \begin{figure*}%f8 \includegraphics[angle=270,scale=.60]{plot_siosis.eps} \caption{Column densities of SiO (solid line) and SiS (dotted line) as a function of time from Phase II of the hot core model with an initial elemental abundance of 1x10$^{-6}$ for Sulphur and 8x10$^{-8}$ for Silicon. The different plots show different efficiencies for the formation of SiS on the grain, namely: 100\% (top left); 0\% (top right); 30\% (bottom left); 10\% (bottom right).} \label{fig_serena} \end{figure*} In order to qualitatively investigate the origin of the SiS and SiO emissions in Orion KL we ran a grid of models using the chemical model UCL\_CHEM \citep{Vit04a, Vit04b, Ler10}, a time and depth dependent gas-grain model. We modeled the hot core and the plateau separately. Both models are two phase calculation. In Phase I we follow the chemical and dynamical evolution of a collapsing core up to a final density of 5$\times$10$^7$ cm$^{-3}$ for the hot core component and 10$^6$ cm$^{-3}$ for the plateau component as derived in Sect. \ref{sect_phy}. The initial gas is at typical densities of $\sim$ 200 cm$^{-3}$ and in atomic form (apart from a fraction of hydrogen which is already in molecular form) and it undergoes a free-fall collapse \citep{Raw92} until the final densities are reached. During this time, atoms and molecules from the gas freeze on to the dust grains and they hydrogenate where possible. Note that the advantage of this approach is that the ice composition is not assumed but it is derived by a time dependent computation of the chemical evolution of the gas/dust interaction process. However, the ice composition does depend on the percentage of gas depleted on to the grains during the collapse, and this in turns depend on the density as well as on the sticking coefficient and other properties of the species and of the grains (see \citealt{Raw92}). In our model we can vary such percentage (reflecting the uncertainty on the grain properties and sticking probabilities) and the degree (or efficiency) of depletion (as well as the viability of different surface reactions) is explored in this study. Our initial elemental abundances for Phase I are as in \citet{Bel06} (see Table 1). We also ran some models where the initial abundances of both S and Si were either depleted or enhanced by a factor of 10 with respect to these values, reflecting the uncertainty of their degree of depletion onto dust. In Phase II, we follow the chemical evolution of the remnant core. For the hot core models, we simulate the effect of the presence of an infrared source in the center of the core or in its vicinity by subjecting the core to an increase in gas and dust temperature, up to T = 300 K. This increasing of temperature is based on the luminosity of the protostar by using the observational luminosity function of \citet{Mol00}. The models we use have been published before: for the models describing the hot core we refer the reader to \citet{Vit04a} and \citet{Ler10} while for the models describing the plateau we refer the reader to \citet{Ler08}, \citet{Ler10} and \citet{Vit04b}. In both models the presence of a non dissociative C-shock (modeled as in \citealt{Ber97}) can be simulated. If a shock is included in the model then sputtering also occurs and is faster than thermal evaporation. We have ran a total of 6 models for the hot core component and 2 models for the plateau component. For the hot core model we investigated the sensitivity of the chemical abundances to the degree of gas depleted on to the grains during the formation of the core; the branching ratios of surface reactions relevant to the formation of SiO and SiS; the initial abundances of Sulphur and Silicon; whether the gas is subjected to a non-dissociative shock during the hot core lifetime. For the plateau models, we only varied the initial Sulphur and Silicon abundances. We were able to reproduce the observed column density of SiO with most models. The only constrain we found was that the temperature of the gas must be at least $\sim$ 100 K or, alternatively, must have undergone a shock. SiS, on the other hand, is difficult to produce: surface reactions (and subsequent evaporation or sputtering of the mantles) seem to be necessary. We find that the only models that succeed in reproducing the data are those where a percentage (even as small as 5\%) of Sulphur on the grains react with Si to form SiS: Figure \ref{fig_serena} shows the column density of SiO and SiS as a function of time during Phase II of the hot core for models differing only in the mantle formation efficiency of SiS (i.e on how efficient Si bonds with Sulfur). A qualitative match with the observation can be achieved, at early times, by those models where only 5\%-10\% of Sulphur on the grains react with Si to form SiS, or at late times if a higher percentage of Sulphur is involved in the formation of SiS. Note, however, that exact ages can not be derived from these considerations as the relationship of the time dependencies with the efficiency of SiS formation on grains will depend on the desorption times of SiO and SiS. In conclusions, while it is possible to reproduce SiO in the gas phase (as well as on the grains), our models indicate that SiS is a product of surface reactions, most likely involving direct reactions of Sulphur with Silicon.

1012.1969

<BR /> Aims: We present a study of the silicon-bearing species detected in a line-confusion limited survey towards Orion KL performed with the IRAM 30-m telescope. The analysis of the line survey is organized by families of molecules. Our aim is to derive physical and chemical conditions for each family taking all observed lines into account from all isotopologs of each species. The large number of transitions in different vibrationally excited states covered by our data, which range from 80 to 280 GHz, let us provide reliable source-average column densities (hence, isotopolog abundances and vibrational temperatures) for the detected molecules. In addition, we provide a wide study of the physical properties of the source based on the different spectral components found in the emission lines. <BR /> Methods: We modeled the lines of the detected molecules using a radiative transfer code, which permit us to choose between large velocity gradient (LVG) and local thermodynamic equilibrium (LTE) approximations depending on the physical conditions of the gas. We used appropriate collisional rates for the LVG calculations. To qualitatively investigate the origin of the SiS and SiO emissions in Orion KL we ran a grid of chemical models. <BR /> Results: For the v = 1 state of SiO, we detected the J = 2-1 line and, for the first time in this source, emission in the J = 4-3 transition, both of them showing a strong masering effect. For SiO v = 0, we detected <SUP>28</SUP>SiO, <SUP>29</SUP>SiO, and <SUP>30</SUP>SiO; in addition, we have mapped the J = 5-4 SiO line. For SiS, we have detected the main species, <SUP>29</SUP>SiS, and SiS v = 1. Unlikely other species detected in Orion KL (<ASTROBJ>IRc2</ASTROBJ>), the emission peak of SiS appears at a velocity of ≃ 15.5 km s<SUP>-1</SUP>. A study of the 5-4 SiO line around IRc2 shows this feature as an extended component that probably arises from the interaction of the outflow with the ambient cloud. We derive an SiO/SiS column density ratio of ≃ 13 in the plateau component, four times lower than the cosmic O/S ratio ≃ 48. In addition, we provide upper limits to the column density of several non-detected silicon-bearing species. The results of our chemical models show that while it is possible to reproduce SiO in the gas phase (as well as on the grains), SiS is a product of surface reactions, most likely involving direct reactions of sulfur with silicon. <P />Appendices are only available in electronic form at <A href="http://www.aanda.org">http://www.aanda.org</A>This work was based on observations carried out with the IRAM 30-m telescope. IRAM is supported by INSU/CNRS (France), MPG (Germany), and IGN (Spain).

12133656

2010arXiv1012.6039F

1012

1012.6039_arXiv.txt

Measurements of the cosmic microwave background (CMB) and large scale structure, such as those provided by the WMAP satellite or the Sloan Digital Sky Survey (SDSS), agree well with the predictions of standard single field slow-roll inflation. In particular the power spectrum verifies the prediction of a nearly scale invariant spectrum of adiabatic perturbations with a Gaussian distribution. However, there remains the prospect that significant non-Gaussianities may be produced by well-motivatived cosmological models which are consistent with the measured power spectrum. In order to quantify this non-Gaussianity we need to measure higher order correlators, beyond the two-point function or power spectrum. In \cite{Regan:cn2010}, we presented an optimal estimation methodology to obtain the four-point correlator or trispectrum. This was developed from the general bispectrum estimator of \cite{FLS09,FLS10} which used separable mode expansions to investigate a much wider class of models than previously had been investigated, as well as to reconstruct the full CMB bispectrum of the Universe. These general bispectrum results were consistent with a Gaussian distribution. However, it is possible for the three-point correlator to remain small but for there to be a large four-point correlator (see, for example, some inflationary models \cite{aChen} and cosmic strings \cite{Regan}). It is our purpose here to continue to test the standard inflationary paradigm through an optimised and expanded search for a trispectrum signal in the WMAP data. \par In this paper we consider the class of four-dimensional trispectra which are independent of the diagonal. Such models include the cubic term of the local model, the equilateral model and the so-called `constant' model. However, as we shall discuss, a general search for non-diagonal trispectrum shapes encompasses a much wider class of models, including most trispectra currently discussed in the literature. We demonstrate this for the non-diagonal equilateral shape which exhibits a high degree of correlation with the (apparently) higher dimensional shapes predicted by DBI and K inflation. Another example is the cosmic string trispectrum which can be reduced to a closely correlated non-diagonal shape. Such dimensional reduction is an important first step in testing even truly diagonal shapes because trispectrum estimation in the non-diagonal case is much more straightforward. Exploiting the use of a separable expansion, as we do here, ensures the computation is tractable, reducing the complexity from $\mathcal{O}(l_{\rm{max}}^7)$ to $\mathcal{O}(l_{\rm{max}}^4)$, and it also ensures the stability of the algorithms used in the analysis without the need to correct for pathological terms commonly present in other approaches. \par The constraints obtained here on the cubic term of the local model, the equilateral model and the `constant' model result from comparison to year 5 WMAP data out to $l=500$ together with a pseudo-optimal analysis of inhomogeneous noise and masking contributions. The estimators and other algorithms employed here were outlined in detail in \cite{Regan:cn2010}, but are expressed in this paper with the simplifying assumption that the trispectrum is independent of the diagonal term. We validate our results by using known analytic results in the large angle limit \cite{RSF10} where the signal-to-noise can be calculated explicitly. This is important because we are able to show that previous trispectrum forecasts using this Sachs-Wolfe approximation were over-optimistic (see, for example, the interesting analysis of trispectrum forecasts in refs~\cite{Kogo,DS}). We note that we will describe the implementation of these methods and the reconstruction of the CMB trispectrum in much more detail in a longer accompanying paper \cite{FRS11}. \par Some alternative approaches to extracting information about the primordial trispectrum from the WMAP data have been explored in the literature. In \cite{VIELVA} the constraint $-5.6\times 10^5<g_{NL}<6.4\times 10^5$ was obtained by analysing the $N$-point probability distribution of CMB anisotropies (a non-optimally-weighted method, see \cite{Regan:cn2010}). This work assumed a local perturbative model, $\Phi=\Phi_L+f_{NL}(\Phi_L^2+\langle \Phi_L^2\rangle)+g_{NL}\Phi_L^3$, while using a Sachs-Wolfe approximation out to $\theta < 1^\circ$. However, we note that our calculation of the {\it optimal} WMAP variance at $\lmax \approx 500$ has error bars of $\pm 10.7 \times 10^5$ (at 95\% confidence). In ref.~\cite{Cooray,Cooray2}, a local trispectrum constraint $-7.4 \times 10^5 < g_{\rm NL} < 8.2 \times 10^5$ is estimated from WMAP using a kurtosis power spectrum analysis (along with a constraint on $\tau_{\rm NL}$). However, this approach does not directly subtract the effect of anisotropic noise and other systematic effects using the quadratic terms in the optimal trispectrum estimator \cite{Regan:cn2010}; we know from the present work that these are important in obtaining an accurate and optimized result. We will discuss and compare these alternative approaches in more detail in our longer paper. It is also appropriate at this stage to mention other earlier proposals for measuring the four-point correlator using a harmonic analysis on COBE\cite{0111250} and a wavelet approach \cite{0301220}. \par n section II we review results relating primordial and CMB trispectra and their optimal estimation, here, focusing on primordial (and hence CMB) trispectra which are independent of the diagonal. The eigenmode decomposition of the trispectrum constitutes the basis of our method and is reviewed in section III. In section IV we obtain constraints on the cubic local model term, the equilateral model and the constant model. The accuracy of this approach is also verified, using large-angle analytic results for the local $\gnl$ model. We emphasise the close correlation of our non-diagonal equilateral model with a broader set of equilateral-type models. We demonstrate the application of this method also at late times, when we constrain the cosmic string trispectrum and discuss forecasts for the Planck data. Finally we discuss our results and present our conclusions in section V.

We have presented results from an implementation of an optimal CMB trispectrum estimator which employs separable mode expansion applicable to a wide class of isotropic models that are diagonal independent, i.e.\ that only depend on the wavenumbers $k_1,k_2,k_3,k_4$. We have obtained constraints on the cubic term for the local model, the constant model and a notable new constraint on equilateral-type models. We found no evidence for primordial non-Gaussianity for these trispectrum shapes (or models closely correlated with them) at the 95\% confidence level. The constraints on the parameter $g_{NL}^{local}$ represent a constraint on the self-interaction term of the local model and, in addition, allows for a qualitative classification of models of the local type \cite{1009.1979}. The constraints presented here on the equilateral model are entirely new results. The importance of finding bounds on such models is, not least, as a consequence of the high correlation between the equilateral model, $K$-inflation and DBI inflation. \par In addition we have demonstrated that this estimator methodology can be applied to late-time models by searching for the characteristic trispectrum shape induced by cosmic strings. We have also obtained a conservative bound on cosmic strings $G\mu\lesssim 1.1\times 10^{-6} $(at $95\%$ confidence). Using forecasts for the signal to noise at Planck resolution we have established that the trispectrum of cosmic strings is expected to give comparable constraints to the power spectrum in the near future. This constraint has the advantage of not exhibiting degeneracies with background cosmological parameters, while also being largely independent of the underlying properties of the cosmic strings (unlike gravitational wave constraints). \par This paper marks a significant first step in a general analysis of trispectrum models. Among the advantages of the modal approach pursued here are the exploitation of efficiencies arising from separability and the transfer function transformation, the absence of pathologies when representing shapes, and the opportunity to investigate non-separable models which were previously deemed intractable. We also note that this general mode expansion allows for an accurate characterisation of the noise and foregrounds which must be subtracted to achieve an optimal estimation measure. While the constraints published in this paper are consistent with Gaussianity, the extension of the analysis to general trispectra represents an important development \cite{FRS11}. Such an analysis will, for instance, allow a measurement of the local trispectrum parameter $\tau_{NL}$ providing a test of local inflation which requires $\tau_{NL}\geq(6f_{NL}/5)^2$. An implementation of the late-time modal estimator outlined in \cite{RSF10} to include trispectra dependent on the diagonal term will allow identification of any trispectrum whether generated at primordial times like inflation or late-times like gravitational lensing or second-order gravitational effects. This, in conjunction with recent results classifying CMB bispectrum constraints \cite{FLS10}, offers the hope of a comprehensive test for non-Gaussianity. \par We conclude by noting that we have obtained near-optimal trispectrum constraints on primordial local $\gnl$ and equilateral (and correlated) models, together with late-time constraints on cosmic strings. We find no evidence for a significant trispectrum in these non-diagonal cases, about which we will provide greater detail elsewhere \cite{FRS11}, together with a reconstruction of the CMB trispectrum obtained from the WMAP data. The new results here represent a significant test of the simplest standard model of inflation, further affirming the Gaussian hypothesis for primordial fluctuations.

1012.6039

We present an implementation of an optimal CMB trispectrum estimator which accounts for anisotropic noise and incomplete sky coverage. We use a general separable mode expansion which can and has been applied to constrain both primordial and late-time models. We validate our methods on large angular scales using known analytic results in the Sachs-Wolfe limit. We present the first near-optimal trispectrum constraints from WMAP data on the cubic term of local model inflation $ g_{\rm NL} = (1.6 \pm 7.0)\times 10^5$, for the equilateral model $t_{\rm NL}^{\rm{equil}}=(-3.11\pm 7.5)\times 10^6 $ and for the constant model $t_{\rm NL}^{\rm{const}}=(-1.33\pm 3.62)$. These results, particularly the equilateral constraint, are relevant to a number of well-motivated models (such as DBI and K-inflation) with closely correlated trispectrum shapes. We also use the trispectrum signal predicted for cosmic strings to provide a conservative upper limit on the string tension $G\mu \le 1.1\times 10^{-6}$ (at 95% confidence), which is largely background and model independent. All these new trispectrum results are consistent with a Gaussian Universe. We discuss the importance of constraining general classes of trispectra using these methods and the prospects for higher precision with the Planck satellite.

12213314

2011MNRAS.412L.108S

1012

1012.0045_arXiv.txt

\label{sec:intro} It is well known that NS cores contain superdense matter whose properties are still uncertain (see, e.g., \citealt{hpy07,lattimerprakash07}). One can explore these properties by studying the cooling of isolated NSs \citep[see, e.g.,][for review]{pethick92,yakovlevpethick04,pageetal06,pageetal09}. We analyse observations of the NS in the supernova remnant Cassiopeia~A (Cas~A). The distance to the remnant is $d=3.4^{+0.3}_{-0.1}$~kpc \citep{reedetal95}. The Cas~A age is reliably estimated as $t \approx 330\pm 20$~yr from observations of the remnant expansion \citep{fesenetal06}. The compact central source was identified in first-light \Chandra\ X-ray observations \citep{tananbaum99} and studied by \citet{pavlovetal00,chakrabartyetal01,pavlovluna09} but its nature has been uncertain. The fits of the observed X-ray spectrum with magnetized or non-magnetized hydrogen atmosphere models or with black-body spectrum revealed too small size of the emission region (could be hot spots on NS surface although no pulsations have been observed, e.g., \citealt{pavlovluna09}). Recently \citet{hoheinke09} have shown that the observed spectrum is successfully fitted taking a carbon atmosphere model with a low magnetic field ($B \lesssim 10^{11}$ G). The gravitational mass of the object, as inferred from the fits, is $M \approx 1.3-2\,M_\odot$, circumferential radius $R \approx 8-15$ km, and the non-redshifted effective surface temperature $\Ts \sim 2 \times 10^6$~K \citep{yakovlevetal10}. These parameters indicate that the compact source is an NS with the carbon atmosphere. It emits thermal radiation from the entire surface and has the surface temperature typical for an isolated NS. It is the youngest in the family of observed cooling NSs. \citet{yakovlevetal10} compared these observations with the NS cooling theory. The authors concluded that the \casa\ has already reached the stage of internal thermal relaxation. It cools via neutrino emission from the stellar core; its neutrino luminosity is not very different from that provided by the modified Urca process. Following \citet{hoheinke09}, \citet{heinkeho10} analysed \Chandra\ observations of the \casa\ during 10 years and found a steady decline of $\Ts$ by about 4\%. They interpret it as direct observation of \casa\ cooling, the phenomenon which has never been observed before for any isolated NS. These results are confirmed by new observations we report below. We interpret them as a manifestation of neutron superfluidity in the \casa. When this paper was nearly completed we became aware of the paper by \citet{pageNew10} who proposed similar explanation of the \casa\ observations. However the two papers are different in details, and can be regarded as complementary. In particular, we discuss the dependence of cooling curves on the poorly known efficiency of neutrino emission due to CPF process and the possibility to interpret observations of all cooling stars by one model of superdense matter. It is important that we report the new observation.

\label{sec:concl} We report a new (November, 2010) \Chandra\ observation of the young \casa\ that confirms the observed \citep{heinkeho10} steady decline of the surface temperature $\Ts$ (by 4\% over 10 years). We propose a natural explanation of the observed decline. We assume that the \casa\ underwent the traditional internal crust-core relaxation some time ago and now demonstrates the second temperature drop due to the onset of triplet-state neutron superfluidity in its core and associated neutrino emission. We can explain the \casa\ observations under the following conditions: \begin{itemize} \item The maximum critical temperature for triplet-state pairing of neutrons should be $T_{\rm cn\; max} \approx (7-9) \times 10^8$~K. Otherwise the second temperature drop occurs earlier or later than required by observations. \item The $T_\mathrm{cn}(\rho)$ profile over the NS core should be rather wide for the CPF neutrino emission to gain enough strength. \item For the same reason the suppression of the CPF process by collective effects cannot be too strong ($q > 0.4$). \item The neutrino emission of the star before the second temperature drop should be 30--100 times lower than due to the modified Urca process (e.g., the modified Urca can be suppressed by strong proton superfluidity). Otherwise the second temperature drop would not be pronounced. \end{itemize} When these criteria are met, we can still locate $T_\mathrm{cn}(\rho)$ profiles in different parts of the NS core. If, however, we wish to explain all current observations of isolated NSs with one and the same $T_\mathrm{cn}(\rho)$-profile, we will be forced to push this profile deeper in the core (Fig.\ \ref{fig:Cool}). Alternatively, we could employ broader profiles but with density-dependent factor $q$ (which can increase within the core as for singlet-state pairing, e.g., \citealt{kv08}). This would shift the efficiency of the CPF process to higher $\rho$ in superfluid matter. We have taken one EOS and focused on $1.65 \, M_\odot$ neutron star model but our basic conclusions will not change for a large variety of EOSs and masses $M$. % For instance, taking the same EOS we have considered the \casa\ models with $M$ from $1.4\,M_\odot$ to $1.9\,M_\odot$. By slightly changing $T_\mathrm{cn}(\rho)$ profiles we are able to explain the data for any $M$ from this range. Our calculations indicate that the second temperature drop lasts for a few tens of years and \casa\ is at its active CPF neutrino emission stage. These models would be inconsistent with a sharp stop of the temperature decline in a few years, which can be verified with future observations. After the second temperature drop the \casa\ is expected to become a rather cold slowly cooling NS.

1012.0045

According to recent results of Ho &amp; Heinke, the Cassiopeia A supernova remnant contains a young (≈330-yr-old) neutron star (NS) which has carbon atmosphere and shows notable decline of the effective surface temperature. We report a new (2010 November) Chandra observation which confirms the previously reported decline rate. The decline is naturally explained if neutrons have recently become superfluid (in triplet state) in the NS core, producing a splash of neutrino emission due to Cooper pair formation (CPF) process that currently accelerates the cooling. This scenario puts stringent constraints on poorly known properties of NS cores: on density dependence of the temperature T<SUB>cn</SUB>(ρ) for the onset of neutron superfluidity [T<SUB>cn</SUB>(ρ) should have a wide peak with maximum ≈ (7-9) × 10<SUP>8</SUP> K]; on the reduction factor q of CPF process by collective effects in superfluid matter (q &gt; 0.4) and on the intensity of neutrino emission before the onset of neutron superfluidity (30-100 times weaker than the standard modified Urca process). This is serious evidence for nucleon superfluidity in NS cores that comes from observations of cooling NSs.

12163116

2011A&A...528A.100S

1012

1012.2871_arXiv.txt

We know that quasars and active galactic nucleii (AGN) are powered by matter accreted onto a supermassive black hole. The accretion of material in the direct vicinity of the central black hole releases most of the quasar energy in the form of power-law continuum emission. Ionised gas surrounds the central accretion disc and gives rise to broad emission lines, which are used as footprints that allow the identification and classification of quasars. Our knowledge of the kinematics and physical conditions prevailing in the BLR gas remain elusive, especially because the nuclear region of quasars is still spatially unresolved with existing instrumentation. Current insights into the BLR come from various kind of studies: empirical modelling of the line shape with kinematical models, use of photo-ionisation codes to reproduce the observed flux ratios between spectral lines, spectropolarimetric observations, statistical study of the width and asymmetry of the lines, use of the principal component analysis technique, and velocity resolved reverberation mapping (e.g. Boroson \& Green~\cite{BOR92}, Sulentic et al.~\cite{SUL00}, Smith et al.~\cite{SMI05}, Marziani et al.~\cite{MAR06}, Zamfir et al.~\cite{ZAM08}, Gaskell~\cite{GAS09}, ~\cite{GAS10b}, Bentz et al.~\cite{BEN10}). Despite the development and many successes of these methods, as briefly summarised below, we still do not completely understand the structure and kinematics of the BLR. The microlensing of the broad emission lines in multiply imaged lensed AGNs provides us with a powerful alternative technique for looking at the BLR, measure its size even in high luminosity distant quasars, and get hints of the structure and geometry of both emission and intrinsic absorption within the BLR (e.g. Schneider \& Wambsganss~\cite{SCH90}, Hutsem\'ekers et al.~\cite{HUT94}, Lewis \& Belle~\cite{LEW98}, Abajas et al.~\cite{ABA02}, Popovi{\'c} et al.~\cite{POP03}, Lewis \& Ibata~\cite{LEW04}, \cite{LEW06}, Richards et al.~\cite{RIC04}, Abajas et al.~\cite{ABA07}, Sluse et al.~\cite{SLU07},~\cite{SLU08}, Hutsem\'ekers et al.~\cite{HUT10}). \subsection{Phenomenology of the line profiles} Our primary clue on the BLR comes from the shape of the broad emission lines. The most detailed studies of broad emission lines have focused on two lines: \ion{H}{$\beta$}\,$\lambda$4863 and \ion{C}{IV}$\,\lambda$1549 (Sulentic et al.~\cite{SUL00} for a review). Of direct interest for the present work is \ion{C}{IV}$\lambda$1549. The \ion{C}{IV} profile shows a broad variety of shapes from strongly asymmetric to symmetric (e.g. Wills et al.~\cite{WIL93}, Baskin \& Laor \cite{BAS05}), has an equivalent width anti-correlated with its intensity (Francis et al.~\cite{FRA92}), and shows greater variability in the wings than in the core (Wilhite et al.~\cite{WIL06}). In addition, the \ion{C}{IV} line is also found to be systematically blueshifted by several hundred to a few thousand kilometres per second compared to the low ionisation lines (Gaskell~\cite{GAS82}, Corbin~\cite{COR90}, Vanden Berk et al.~\cite{VAN01}). The analysis of about 4000 SDSS quasars by Richards et al. (\cite{RIC02}) suggests that this shift is caused by a lack of flux in the red wing of the line profile and correlates with the quasar orientation. Variability studies and emission-line decomposition techniques (principal component analysis and analytical fitting) indicate that the region emitting \ion{C}{IV} could be built up with two components, a ``narrow'' emission core of FWHM $\sim$ 2000 km\,s$^{-1}$ emitted in an intermediate line region (ILR) possibly corresponding to the inner part of the narrow line region and a very broad component (VBC) with FWHM $\sim$7000 km\,s$^{-1}$ producing the line wings (Wills et al.~\cite{WIL93}, Brotherton et al.~\cite{BRO94a}, Sulentic~\cite{SUL00}, Wilhite et al.~\cite{WIL06}). In radio-quiet objects, the VBC is observed to be systematically blueshifted by thousands of km\,s$^{-1}$ with respect to the narrow core (Brotherton et al.~\cite{BRO94a}, Corbin~\cite{COR95}) suggesting it is associated with outflowing material. The ILR component disclosed in \ion{C}{IV} is probably different from the one recently uncovered in \ion{H}{$\beta$} (Hu et al.~\cite{HU08}), but its exact nature and the physical conditions in this region are still being debated (Brotherton et al.~\cite{BRO94b}, Sulentic \& Marziani~\cite{SUL99}, Marziani et al.~\cite{MAR06}). The properties of the \ion{C}{III]} emission line have received less attention in the literature, probably mostly because the line is blended with \ion{Al}{III}\,$\lambda$1857, \ion{Si}{III]}\,$\lambda$1892 and, an \ion{Fe}{II+III}\,$\lambda$1914 complex. In their study of the \ion{C}{IV} and \ion{C}{III]} emission, Brotherton et al. (\cite{BRO94a}) find that the 2 lines often have different profiles and that, for some objects, a two-component decomposition (i.e. narrow core+very broad wings) does not provide a good model of \ion{C}{III]}, so a third component is required. They also find that the VBC needed to reproduce \ion{C}{III]} has to be larger than the corresponding component of \ion{C}{IV}. This supports the idea that the region emitting \ion{C}{III]} is different from the one emitting \ion{C}{IV} (Snedden \& Gaskell~\cite{SNE99}). Recently, Marziani et al. (\cite{MAR10}), based on the line profile decomposition of a small sample of AGNs selected in the 4D Eigenvector 1 context (e.g. Boroson \& Green~\cite{BOR92}, Zamfir et al.~\cite{ZAM08}), suggest that all the broad emission line profiles are composed of three components of variable relative intensity and centroid shift (from line to line in a given object and between objects). They suggest a classical unshifted broad component (FWHM=600-5000 km\,s$^{-1}$), a redshifted very broad component and a blueshifted component mostly visible in the so-called Population A objects (i.e. objects with FWHM $<$ 4000 km\,s$^{-1}$). Based on line ratios, they also tentatively infer that these components arise from different emitting regions. The smoothness of the line profiles and the physical conditions derived from photo-ionisation models allowed several authors to put constraints on the ``structure'' of the BLR gas. Two popular models remain. The first one considers that the BLR is a clumpy flow composed of small gas clouds, and the second one assumes a smooth gas outflow originating in an accretion disc (Elvis ~\cite{ELV00}, Laor ~\cite{LAO07} and ref. therein). Several geometries have been considered for the BLR gas, the most popular ones being disc-like, spherical and biconical models (e.g. Chen \& Halpern~\cite{CHE89}, Robinson \cite{ROB95}). Murray and Chiang (\cite{MUR97}) demonstrate that a continuous wind of gas originating in an accretion disc can successfully reproduce the profile and systematic blueshift of the \ion{C}{IV} emission line. Other indications that a fraction of the BLR material has a disc-like geometry comes from statistical studies of the broad line profiles in samples of (mostly radio-loud) AGNs (Vestergard et al. ~\cite{VES00}, McLure \& Dunlop~\cite{MCL02}, Jarvis \& McLure~\cite{JAR06}, Decarli et al.~\cite{DEC08}, Risaliti et al.~\cite{RIS10}), and from spectro-polarisation observations of Balmer lines in AGNs (Smith et al.~\cite{SMI05}). Some of these studies also suggest there is a second, spherically symmetric component with Keplerian motion. The spectropolarimetric observations of PG 1700+518 by Young et al. (\cite{YOU07}) support a disc+wind model for the \ion{H}{$\alpha$} emission. On the other hand, biconical models of the BLR seem needed to explain the variability of the rare double-peaked AGNs (i.e. AGNs showing emission lines with two peaks) and their polarisation properties (Sulentic et al.~\cite{SUL95}, Corbett et al.~\cite{COR98}). These studies show that there is no consensus on the geometry of the BLR. From an observational and theoretical perspective, it seems nevertheless to depend on the ionisation degree of the line and on the radio properties of the object. \subsection{The radius-luminosity relationship} \label{subsec:RL} The ``size'' of the the broad line region $R_{BLR}$ has been measured in about 40 AGNs by use of the reverberation mapping technique (e.g. Krolik et al.~\cite{KRO91}, Horne et al.~\cite{HOR04}). The empirical relation $R_{BLR} \propto L^\alpha$ ($\alpha\sim$0.5-0.6, Kaspi et al.~\cite{KAS05}, Bentz et al.~\cite{BEN06}), combined with the virial theorem, allows one to derive a relation linking the black hole mass $M_{BH}$, the AGN luminosity $L$ and the FWHM of the emission line. This relation is one of the most popular methods used to measure black hole masses based on single epoch spectroscopic data and study their growth, evolution and correlation to other AGN properties. The $R_{BLR} \propto L^\alpha$ relation has been derived quite accurately for the broad component of the \ion{H}{$\beta$} line, but only a few objects have $R_{BLR}$ measurements for high ionisation lines like \ion{C}{IV} (e.g. Kaspi et al. ~\cite{KAS07}). This is a severe problem for black hole mass measurements of high redshift objects. The use of the \ion{C}{IV} line to derive black hole masses is desirable but faces several problems related to our understanding of the structure and geometry of the region emitting this line (see e.g. Marziani et al. ~\cite{MAR06} for a review). The existence and contamination of a narrow emission component in \ion{C}{IV} which could bias FWHM$_{\ion{C}{IV}}$ (Bachev et al.~\cite{BAC04}) and the possible absence of virial equilibrium, especially in sources showing large blueshifts (Richards et al.~\cite{RIC02}), are major concerns. Velocity resolved reverberation mapping (Horne et al.~\cite{HOR04}) should allow one to get insights on these problems, but the technique is still under development and has not yet been applied to the \ion{C}{IV} line. Most of the recent advances in velocity resolved echo-mapping are based on the analysis of the \ion{H}{$\beta$} emission line (Denney et al.~\cite{DEN09}, Bentz et al.~\cite{BEN09b},~\cite{BEN10}). Although the kinematics of the Balmer gas was thought to be relatively simple, the technique provided puzzling results as both keplerian rotation, inflow and outflow signatures appear in various objects. It is still unclear whether this reveals a wide variety of kinematics, a biased interpretation of the observed signal, or a spurious signal introduced by observational artifacts. The past evidence of inflow and outflow signatures in NGC 5548 (Clavel et al.~\cite{CLA91}, Peterson et al.~\cite{PET91}, Kollatschny \& Dietrich~\cite{KOL96}) might indicate that inflow/outflow signatures are not unambiguous. A possible explanation of the variety of observed signals might be off-axis illumination of the BLR (Gaskell~\cite{GAS10a}, \cite{GAS10b} and references therein). \subsection{Microlensing in \obj} The many open questions concerning the BLR we outlined above motivate the interest in developing new techniques of probing this region. In this paper we investigate the constraints on the broad lines provided by means of the microlensing study of the lensed quasar \obj, for which long-term spectro-photometric monitoring has been carried out (Eigenbrod et al.~\cite{EIG07}, hereafter {\it{Paper I}}). The gravitational lens \obj, also known as ``Huchra's lens'' or the ``Einstein Cross'', was discovered by Huchra et al. (\cite{HUC85}) during the Center for Astrophysics Redshift Survey. It consists of a $z_s=1.695$ quasar gravitationally lensed into four images arranged in a cross-like pattern around the nucleus of a $z_l=0.0394$ barred Sab galaxy. The average projected distance of the images from the lens centre is only 700\,pc $\sim$ 0.9\arcsec, such that the matter along the line of sight to the lensed images is mostly composed of stars. The symmetric configuration of the lensed images around the lens-galaxy bulge ensures a time delay of less than a day between the lensed images, such that intrinsic flux variations should be seen quasi- simultaneously in the four images (Rix et al.~\cite{RIX92}, Wambsganss \& Paczy\`nski~\cite{WAM94}). Since the low redshift of the lens galaxy leads to a high relative transverse velocity between the observer, the lens, and the source, the lensed images of \obj\, are continuously flickering due to microlensing produced by the stars in the lens galaxy, on smaller timescales than in any other lens. The microlensing affecting the images of \obj\, leads to large variations in amplitude (Udalski et al.~\cite{UDA06}), reaching up $>$ 1 mag, often accompanied by chromatic microlensing of the quasar continuum (Wambsganss \& Paczy\,nski~\cite{WAM91}) used to study the accretion disc temperature profile (e.g. Kochanek~\cite{KOC04}, Anguita et al.~\cite{ANG08}, Eigenbrod et al.~\cite{EIG08}={\it Paper II}). Several studies have shown that the broad emission lines are also significantly affected by microlensing (Metcalf et al.~\cite{MET04}, Wayth et al.~\cite{WAY05}, Paper I). Based on the microlensing of \ion{C}{III]} observed at one epoch in 2002, Wayth et al. (\cite{WAY05}) derived a most likely size of the region emitting this line of 0.06 $h^{1/2}$ pc. The work presented here aims at improving the measurement of the BLR size by using monitoring data instead of a single epoch measurement and at constraining the structure of the BLR. The structure of the paper is the following. In Sect.~\ref{sec:obs}, we briefly summarise the data we used and explain the three different methods we applied to analyse the spectra. In Sect.~\ref{sec:phenomenology}, we apply these techniques to our data to understand how microlensing deforms the emission lines and to measure the flux ratios in several portions of the \ion{C}{IV} and \ion{C}{III]} emission lines. In Sect.~\ref{sec:simulation}, we explain the microlensing simulations we developed to derive the size of the broad emission lines and continuum emission region. In Sect.~\ref{sec:results} we present our results: measurement of the size of the BLR and of the continuum region, comparison of our BLR size with reverberation mapping, estimate of the black hole mass, and finally constraints on the BLR structure. Finally, we summarise our main findings in Sect.~\ref{sec:conclusions}.

\label{sec:conclusions} For the first time, we have derived the size of the broad line region in \obj\, based on spectrophotometric monitoring data and consisting of 39 different epochs obtained between Oct 2004 and Dec 2007. To reach this goal, we measured differential lightcurves between images A \& D for the \ion{C}{IV} and \ion{C}{III]} broad emission lines and compared them to microlensing simulations. This led to determining the half-light radius of the \ion{C}{IV} emitting region: $R_{\ion{C}{IV}} \sim $ 66 lt-days (19 lt-days $< R_{\ion{C}{IV}} <$ 176 lt-days at 68.3\% confidence), in very good agreement with the $R_{BLR}-L$ relation derived from reverberation mapping (Kaspi et al.~\cite{KAS07}). The size we derived for \ion{C}{III]} is $R_{\ion{C}{III]}} \sim 49$ lt-days (14 lt-days $< R_{\ion{C}{III]}} <$ 154 lt-days at 68.3\% confidence), compatible with the size estimated for the region emitting \ion{C}{IV}. Thanks to the variable amount of microlensing observed within a given emission line, we can also derive information on the structure of the broad line region. Differential lightcurves obtained for various velocity slices of the \ion{C}{IV} and \ion{C}{III]} lines show that the wings of the lines are more microlensed than the core, indicating that the former arise in a more compact region. This finding is confirmed by two other techniques. The first one demonstrates that a broad and single-peaked fraction of the emission line is microlensed, while a narrower fraction is unaffected by microlensing. This technique also suggests a slightly different structure for the \ion{C}{IV} and \ion{C}{III]} emission regions, with the narrow \ion{C}{IV} emission blueshifted with respect to the systemic redshift. The second technique assumes that the emission lines can be decomposed into a sum of Gaussian components. The \ion{C}{IV} is separated into a narrow (FWHM $\sim$ 2600\,km\,s$^{-1}$) and a broad (FWHM $\sim$ 6300\,km\,s$^{-1}$) component. In order to reproduce the small differences between the \ion{C}{IV} and \ion{C}{III]} lines, we need three Gaussian profiles (FWHM $\sim$ 1550, 3400, 8550\,km\,s$^{-1}$) to reproduce the latter. The lightcurves derived for these components are essentially flat and show that microlensing is more important when the FWHM increases. Although the individual Gaussian line components do not necessarily isolate individual emission regions, we find that these lightcurves are compatible with the continuum lightcurves, provided only the size of the emission region is modified. This allows us to derive a half-light radius for the regions emitting these components, as well as for their size relative to the continuum. The radii are consistent with the virial hypothesis and a radius that varies as FWHM$^{-2}$. Using the virial theorem, we derived a black hole mass $M_{BH} \sim $2.0$\times 10^8 M_{\sun}$ (0.5$\times 10^8 M_{\sun} < M_{BH} < $ 5.4$\times 10^8 M_{\sun}$ at 68.3\% confidence). Our analysis supports the findings by other authors (Brotherton et al.~\cite{BRO94a}, Marziani et al.~\cite{MAR10}) that the regions emitting the \ion{C}{IV} and \ion{C}{III]} lines are composed of at least two spatially distinct components, one emitting the narrow core of the line and another, more compact, emitting a broadest component. The broad (resp. narrow) component of the \ion{C}{IV} and \ion{C}{III]} lines do not have exactly the same profiles. The flux ratio \ion{C}{III]}/\ion{C}{IV} is very similar when measured in the broad and in the narrow components of the lines. This suggests that the ionisation parameter U is nearly the same in the two regions, a surprising result since the narrow components are found to arise in regions that are several times larger than the broad components. Other lines observed in our spectra seem to arise from a least two components: \ion{Mg}{II}\,$\lambda 2798$, \ion{Al}{II}\,$\lambda 1671$, \ion{He}{II}\,$\lambda 1640$, \ion{Al}{III}\,$\lambda 1857$. The situation is different for \ion{Si}{III]}\,$\lambda 1892$, which does not show emission from a broad component, which is indicative of a smaller electronic density $n_e$ in the region emitting the broadest part of the emission lines. On the other hand, we detect broad microlensed \ion{Fe}{II+III} but no ``extended'' emission. This suggests that Fe$_{UV}$ is produced in the inner part of the BLR or in a very compact region. Obtaining spectrophotometric monitoring data in the near-infrared where Balmer lines are detectable would be very useful for constraining photoionisation models and comparing the microlensing signal in Balmer (e.g. \ion{H}{$\beta$}), high ionisation (e.g. \ion{C}{IV}), and Fe$_{opt}$ lines. We demonstrated that the spectrophotometric monitoring of microlensing in a lensed quasar is a powerful technique for probing the inner regions of quasars, measuring the size of the broad line region, and infering its structure. More work is still needed to take full advantage of the method, but our results are very promising. Several improvements are possible to increase the accuracy of our size measurements and better characterise the BLR structure. First, microlensing simulations reproducing the signal in more than 2 lensed images should allow one to narrow the final probability distribution on the source size. Second, the implementation of spectral fitting using Markov-Chain Monte-Carlo should allow a more appropriate estimate of the error bars and more advanced modelling of the individual spectral components. Third, a fully coherent scheme should be developed to consistently model the emission line shape and the corresponding source intensity profile used in the simulation. These improvement will be the subject of a future work.

1012.2871

<BR /> Aims: We aim to use microlensing taking place in the lensed quasar QSO 2237 + 0305 to study the structure of the broad line region (BLR) and measure the size of the region emitting the C iv and C iii] lines. <BR /> Methods: Based on 39 spectrophotometric monitoring data points obtained between Oct. 2004 and Dec. 2007, we derived lightcurves for the C iv and C iii] emission lines. We used three different techniques to analyse the microlensing signal. Different components of the lines (narrow, broad, and very broad) were identified and studied. We built a library of the simulated microlensing lightcurves that reproduce the signal observed in the continuum and in the lines provided only the source size is changed. A Bayesian analysis scheme is then developed to derive the size of the various components of the BLR. <BR /> Results: 1. The half-light radius of the region emitting the C iv line is found to be R<SUB>C IV</SUB>} ∼ 66<SUP>+110</SUP><SUB>-46</SUB>} light-days = 0.06<SUP>+0.09<SUB>-0.04</SUB> pc = 1.7<SUP>+2.8<SUB>-1.1 × 10<SUP>17</SUP></SUB> cm (at 68.3% CI). Similar values are obtained for C iii]. Relative sizes of the carbon-line and V-band continuum emitting-regions are also derived with median values of R<SUP>line</SUP>/R<SUP>cont</SUP> in the range 4 to 29, depending on the FWHM of the line component. 2. The size of the C iv emitting region agrees with the radius-luminosity relationship derived from reverberation mapping. Using the virial theorem, we derive the mass of the black hole in QSO 2237 + 0305 to be M<SUB>BH</SUB> ~ 10<SUP>8.3 ± 0.3</SUP> M<SUB>⊙</SUB>. 3. We find that the C iv and C iii] lines are produced in at least 2 spatially distinct regions, the most compact one giving rise to the broadest component of the line. The broad and narrow line profiles are slightly different for C iv and C iii]. 4. Our analysis suggests a different structure for the C iv and Fe ii+iii emitting regions, with the latter produced in the inner part of the BLR or in a less extended emitting region than C iv. <P />Based on observations made with the ESO-VLT Unit Telescope # 2 Kueyen (Cerro Paranal, Chile; Proposals 073.B-0243(A&amp;B), 074.B-0270(A), 075.B-0350(A), 076.B-0197(A), 177.B-0615(A&amp;B), PI: F. Courbin). <P /></SUP></SUP>

1408249

2010ApOpt..49.6675H

1012

1012.0790_arXiv.txt

\label{Intro} The direct imaging and characterization of Earth-like planets, especially those capable of supporting life, is one of the outstanding goals of modern astrophysics and science in general. Currently, high-angular resolution astronomical imaging can be achieved interferometrically by combining the wavefronts from spatially separated telescopes~\cite{Labeyrie06}. Interferometric imaging and nulling require that the light from multiple apertures be combined. Although beam combining can be performed using bulk optics, integrated optic (IO) implementations offer a number of important advantages. These include spatial filtering, enhanced stability, on-chip fringe scanning, compactness and scalability. IO beam combiners for astronomical imaging were first proposed by Kern, Malbet, Schanen-Duport, and Benech in 1996~\cite{Kern96}. Using silicate-based glass IO devices, laboratory and on-sky stellar interferograms were demonstrated at astronomical H (1.5 $\mu$m -- 1.8 $\mu$m) and K (2.0 $\mu$m -- 2.4 $\mu$m) bands~\cite{BergerII,BergerIII,BergerIV,BergerV,BergerVI}. In many situations, especially those involving nulling and exoplanet search, operation in the infrared beyond 3 $\mu$m, where silicate-based glasses are not transparent, is required. On-sky interferometric measurements have been performed in the L band (3.0 $\mu$m -- 4.0 $\mu$m) using a guided-wave device, consisting of a two-beam fluoride glass fiber coupler, but this fiber-based technology is not easily scalable to multiple apertures~\cite{Mennesson99}. We have recently developed a prototype, single-mode, integrated optic, astronomical, beam combiner fabricated by titanium-indiffusion in an x-cut, lithium niobate (LiNbO$_3$) wafer~\cite{HKH09,HKHFiO,HKH2010}. The device operates in the 3.2 $\mu$m -- 3.8 $\mu$m spectral region, which lies in the L band, and has on-chip, electro-optically (EO) controlled, fringe scanning capabilities. Our results confirm that IO devices are well suited to perform the beam combining function for astronomical imaging. In the infrared wavelength region, a host star is normally a million times brighter than the planet orbiting it, which presents major difficulties when trying to image the planet directly. Nulling interferometry offers the possibility to overcome this problem by attenuating the stellar light, thus enhancing the visibility of the planet. Generally, achromatic phase shifting and broadband achromatic beam combining functions are required for deep nulling, with nulling depths of 10$^{-6}$ or better over the infrared spectrum. Although we can image the planet by nulling at a single wavelength, there are key reasons to use a broad wavelength band. First, there are several key biomarkers in the infrared spectrum from 6 $\mu$m to 18 $\mu$m. Second, the total integration time needed to detect a planet increases as the spectral bandwidth is reduced. One key component for achieving a broadband deep null is the broadband achromatic beam combiner. Generally, there are two basic types of planar, integrated optic beam combiners: reversed-Y combiners and directional couplers. The operation of a reversed-Y combiner is achromatic by symmetry. A reversed-Y combiner, however, when used as part of a nulling interferometer, suffers an inevitable 3 dB loss of signal. On the other hand, directional coupler type beam combiners can capture the entire signal, and hence are theoretically lossless provided that both interferometric outputs are recorded and subtracted. Unfortunately, the operation of the directional coupler is chromatic, and this wavelength dependence prevents the device from achieving deep broadband nulls. Various approaches have been suggested to mitigate chromaticity of 2 by 2 couplers and achieve a wavelength flattened response. These approaches include (i) asymmetric directional couplers, where the two constituent waveguides differ in width but are uniform along the direction of propagation~\cite{Takagi92}, (ii) Mach-Zehnder type interferometers that are wavelength-insensitive~\cite{Jinguji90}, and (iii) tapered velocity couplers, where the waveguide widths are varied along the propagation direction~\cite{Milton75}. The operation of the devices described by (i) and (ii) above depend on the interference between modes, while the tapered velocity mode coupler (iii) is meant to operate adiabatically, where there is little coupling between the local normal modes. The later device is also referred to as a mode-evolution coupler. Unfortunately these approaches do not yield an ultra-broadband achromatic response nor do they preserve the initial phase difference of the input beams in the interference term at the output. Lack of broadband chromaticity in either power splitting or phase can seriously degrade the deep nulls required for exoplanets search. In this paper, we propose and present a theoretical design of an achromatic, broadband, polarization-insensitive, mode-evolution, integrated optic, beam combiner suited for space-based nulling interferometry. The design is based on a system of three coupled waveguides along the lines of~\cite{Schneider00} and~\cite{Ishikawa07}. The proposed device can be realized using germanium-on-silicon or germanium air-bridge waveguide structures~\cite{HKH2010}. In Section~\ref{LNM}, the theory of normal modes of a three coupled waveguide system is presented. A design of an achromatic, polarization-insensitive, mode evolution beam combiner is proposed in Section~\ref{ProposedDesign}. In Section~\ref{CMTofLNM}, the derivation of the coupled mode equations that describe the mode coupling between the local normal modes is given, and a condition for adiabatic operation is presented. In Section~\ref{ABCDesign}, a numerical design example of the proposed achromatic beam combiner, based on candidate technologies for waveguide fabrication in mid-infrared wavelength region, is presented and the device performance is numerically evaluated. The results of the paper are summarized in Section~\ref{ABCDis}.

\label{ABCDis} We have presented the numerical design of a broadband, polarization insensitive, achromatic beam combiner for the operation in the astronomical N band based on Ge/Si heterostructure raised strip waveguides. The beam combining is intrinsically achromatic because of the symmetric arrangement of the three coupled waveguides. As opposed to a reversed-Y junction combiner which suffers a 3 dB loss, our device is theoretically lossless. Furthermore on-chip EO modulation is also possible by utilizing the free carrier plasma dispersion effect~\cite{Soref4} in silicon-based waveguides. Most importantly, the technology needed to actually fabricate the proposed design is quite promising and plausible. We believe the realization of such beam combiner will increase the possibility of using integrated optic combiners for space-based deep nulling interferometry.

1012.0790

Integrated optic beam combiners offer many advantages over conventional bulk optic implementations for astronomical imaging. To date, integrated optic beam combiners have only been demonstrated at operating wavelengths below 4 microns. Operation in mid-infrared wavelength region, however, is highly desirable. In this paper, a theoretical design technique based on three coupled waveguides is developed to achieve fully achromatic, broadband, polarization-insensitive, lossless beam combining. This design may make it possible to achieve the very deep broadband nulls needed for exoplanet searching.

12137815

2010idm..confE..63V

1012

1012.0273_arXiv.txt

Measurements of cosmic rays and gamma rays have a long tradition of delivering interesting insights into several areas of physics, e.g., particle physics, astrophysical objects, and the interstellar medium. Nowadays models are available that are in good agreement with the measured fluxes \cite{galprop}. Nevertheless, cosmic rays remain an interesting field to study new phenomena like dark matter, as yet unknown astrophysical objects, or even baryogenesis. \subsection{Cosmic rays and dark matter} This work will concentrate on the dark matter aspect. Even if the existence of dark matter seems to be proven \cite{darkmatter}, an exciting question remains concerning its nature. This is particularly interesting as the latest results of the positron flux measurements (PAMELA \cite{pamela}) and the combined positron and electron flux measurements (ATIC \cite{atic} and Fermi LAT \cite{fermi}) may be utilized to constrain dark matter properties. The general idea behind these interpretations is briefly explained in the following: It was shown that dark matter particle candidates do not exist within the standard model of particle physics, but well-motivated theories beyond this model, like supersymmetry \cite{neutralino} and universal extra dimensions \cite{lkp}, contain viable candidates. It is now believed that these particles are Majorana particles and are therefore able to self-annihilate. Whatever the exact physical processes in the underlying theories for these annihilation processes, it is assumed that standard model particles are among the final products, which then in turn contribute to the total flux of cosmic rays. The actual shape of these additional contributions are then used to constrain the parameters of the dark matter model under study. As the calculated dark matter annihilation fluxes are generally small, it appears to be very difficult or even impossible to look for deviations in the most abundant cosmic-ray species like protons and alphas. Therefore people concentrate on electrons, positrons, and antiprotons, which have (much) smaller fluxes compared to the primary species. The PAMELA, ATIC, and Fermi LAT positron and electron results show deviations from the predicted fluxes calculated within the frameworks of the available cosmic-ray models. These deviations could be interpreted as additional fluxes coming from dark matter annihilations. Despite the success of the aforementioned experiments, a few drawbacks exist for a reliable dark matter interpretation using positrons and electrons: The observed deviations are relatively small and might be also interpreted in more standard astronomical ways. Furthermore, the predicted dark matter annihilation fluxes tend to be even smaller than the fluxes needed to explain the deviations and often require the introduction of some kind of boosting mechanism \cite{boost}. This makes it even harder to disentangle the different contributions. \subsection{Antideuterons as a signature for dark matter} \begin{figure} \centerline{\includegraphics[width=0.93\linewidth]{dbar_flux_and_ratio.pdf}}\caption{\label{f-dbar_flux}\textbf{\textit{Left)}} Antideuteron fluxes from secondary interactions of cosmic rays (blue) with the interstellar medium and from dark matter annihilations (red) \cite{antideuteron,antideuteroncui}. \textbf{\textit{Right)}} Ratio of (anti)proton fluxes to antideuteron fluxes with and without an extra contribution from dark matter annihilations \cite{galprop,antideuteron,antideuteroncui}.} \end{figure} A promising channel in the field of indirect dark matter detection is the measurement of antideuterons in cosmic rays \cite{antideuteron}. As cosmic antideuterons have never been measured, only calculations exist. The left hand side of Fig.~\ref{f-dbar_flux} shows the predicted antideuteron flux from interactions of primary cosmic rays with the interstellar medium \cite{antideuteron} and the flux from dark matter annihilation in a generic model \cite{antideuteroncui}. It is interesting to note that the dark-matter-induced flux exceeds the secondary flux below 1-2\,GeV without using any boosting mechanism. What makes the antideuteron detection now very challenging becomes immediately evident by looking at the ratios of (anti)proton to antideuteron fluxes (Fig.~\ref{f-dbar_flux}, right). Below 1\,GeV, protons (antiprotons) are about $10^{10}$ to $10^{11}$ ($10^5$ to $10^6$) more abundant than antideuterons. Therefore any attempt to measure reliably cosmic antideuterons needs an exceptionally strong particle identification.

The measurement of the low-energy antideuteron flux is a promising way to search for dark matter indirectly. The GAPS experiment is specifically designed to perform this task by stopping antideuterons and forming and detecting exotic atoms. It is planned to carry out (ultra-)long duration balloon flights from the South Pole starting from 2014. A prototype experiment is currently under construction, and a flight is scheduled for the summer of 2011 from Taiki, Japan. The hardware, software, and simulations for this flight are currently under development.

1012.0273

The GAPS experiment is foreseen to carry out a dark matter search using a novel detection approach to detect low-energy cosmic-ray antideuterons. The theoretically predicted antideuteron flux resulting from secondary interactions of primary cosmic rays with the interstellar medium is very low. So far not a single cosmic antideuteron has been detected by any experiment, but well-motivated theories beyond the standard model of particle physics, e.g., supersymmetry or universal extra dimensions, contain viable dark matter candidates, which could led to a significant enhancement of the antideuteron flux due to self-annihilation of the dark matter particles.This flux contribution is believed to be especially large at small energies, which leads to a high discovery potential for GAPS. GAPS is designed to achieve its goals via a series of ultra-long duration balloon flights at high altitude in Antarctica, starting in 2014. The detector itself will consist of 13 planes of Si(Li) solid state detectors and a time of flight system. The low-energy antideuterons (&lt; 0.3 GeV/n) will be slowed down in the Si(Li) material, replace a shell electron, and form an excited exotic atom. The atom will be deexcited by characteristic x-ray transitions and will end its life by forming an annihilation pion star. This unique event structure will allow for nearly background free detection. To prove the performance of the different detector components at stratospheric altitudes, a prototype flight will be conducted in 2011 from Taiki, Japan.

4519287

2011PhRvD..83e3006G

1012

1012.0267_arXiv.txt

1012.0267

We analyze the LSND, KARMEN, and MiniBooNE data on short-baseline ν¯<SUB>μ</SUB>→ν¯<SUB>e</SUB> oscillations and the data on short-baseline ν¯<SUB>e</SUB> disappearance obtained in the Bugey-3 and CHOOZ reactor experiments in the framework of 3+1 antineutrino mixing, taking into account the MINOS observation of long-baseline ν¯<SUB>μ</SUB> disappearance and the KamLAND observation of very-long-baseline ν¯<SUB>e</SUB> disappearance. We show that the fit of the data implies that the short-baseline disappearance of ν¯<SUB>μ</SUB> is relatively large. We obtain a prediction of an effective amplitude sin⁡<SUP>2</SUP>2ϑ<SUB>μμ</SUB>≳0.1 for short-baseline ν¯<SUB>μ</SUB> disappearance generated by 0.2≲Δm<SUP>2</SUP>≲1eV<SUP>2</SUP>, which could be measured in future experiments.

12131231

2010arXiv1012.0784R

1012

1012.0784_arXiv.txt

\label{sec:intro} \subsection{A factor of 2} The early universe, at least from Big Bang Nucleosynthesis at one second onwards, is well described by a model which is exactly homogeneous and isotropic (up to linear perturbations) and contains only ordinary matter, and where the relation between matter and spacetime geometry is given by ordinary general relativity. Here ordinary matter means that the pressure is non-negative, and ordinary general relativity is based on the four-dimensional Einstein-Hilbert action. Such a model also works well when applied to the late universe, except that the distance and the expansion rate are underpredicted by a factor of two. The observed distance to the last scattering surface at redshift 1090 is a factor of 1.4--1.7 longer than predicted by the spatially flat homogeneous and isotropic model dominated by pressureless matter, keeping the Hubble parameter today fixed (and assuming a power-law spectrum of primordial perturbations) \cite{Vonlanthen:2010}. Observations of type Ia supernovae and large-scale structure are consistent with this cosmic microwave background (CMB) measurement, and they show that the discrepancy arises at redshifts of order one and smaller, when the universe is about ten billion years old. In the homogeneous and isotropic Friedmann-Robertson-Walker (FRW) models, the explanation for the longer distances is simple: the expansion of the universe has accelerated, so objects have receded further away. Most cosmological observations probe distances, but there are also some measurements of the expansion rate. Galaxy ages \cite{ages} and the radial baryon acoustic oscillation signal \cite{BAOradial} have been used to measure the expansion rate as a function of redshift, and the value of the Hubble parameter today is known with some accuracy \cite{Hubble}. The expansion rate observations agree well with the distance observations and support the interpretation of faster expansion being the cause of the longer distances. The measured Hubble parameter is about a factor of 2 larger than expected compared to the matter density, $\Omega_\mathrm{m0}\equiv8\pi\GN\rho_\mathrm{m0}/(3H_0^2)\approx0.25$ \cite{Peebles:2004}, or a factor of 1.2--1.5 larger if compared to the age of the universe, $H_0t_0\approx$ 0.8--1 instead of $H_0t_0=2/3$ \cite{Krauss:2003}. As the observations are beyond reasonable doubt, at least one of the three assumptions of the theoretical model is wrong. Either there is exotic matter with negative pressure (dark energy), general relativity does not hold on cosmological scales, or the homogeneous and isotropic approximation is not valid at late times. Apart from observations of the expansion rate and the distance scale, there is no evidence for negative-pressure matter or modifications of general relativity. For example, such matter has not been detected in the laboratory, nor have deviations from general relativity been observed in the Solar system\footnote{Apart from possibly the Pioneer anomaly and the flyby anomaly \cite{solar}.}. Likewise, no objects have been seen to accelerate away from each other. All of the relevant observations involve averages over large volumes or integrals over large distances. The situation is different from that of dark matter, for which there is evidence from several kinds of observations on various scales and eras, including local physics, such as the CMB peak structure, large-scale structure, rotation curves of galaxies, the motions of galaxies and gas in clusters, gravitational lensing and so on \cite{Roos:2010}. This is the reason why constructing alternatives to dark matter requires resort to baroque models, if it is possible at all \cite{teves}. In contrast, the various observations usually interpreted as indicating dark energy or modified gravity are all different tracers of the same quantity: longer distances and faster expansion. An effect which leads to faster expansion and correspondingly increases the distances can explain all of these observations. (Indeed, the only effect of the favorite dark energy candidate, vacuum energy, is to change the expansion rate.) Unlike the assumptions of ordinary matter and gravity, the assumption of only linear deviations from homogeneity and isotropy is known to be violated at late times due to the formation of non-linear structures, a process which has a preferred time of about ten billion years. Before concluding that new physics is needed to explain the observations, we should therefore study the possibility that the factor two failure of the predictions of the simple homogeneous and isotropic models is due to the known breakdown of homogeneity and isotropy \cite{Buchert:2000, Tatekawa:2001, Wetterich:2001, Schwarz:2002, Rasanen, Kolb:2004}. \subsection{Our clumpy universe} It is important to distinguish between {\it exact} homogeneity and isotropy and {\it statistical} homogeneity and isotropy. Exact homogeneity and isotropy means that the space has a local symmetry: all points and all directions are equivalent. Statistical homogeneity and isotropy simply means that if we consider a box anywhere in the universe, the mean quantities in the box do not depend on its location, orientation or size, provided that it is larger than the homogeneity scale. The early universe is nearly exactly homogeneous and isotropic, in two ways. First, the amplitude of the perturbations around homogeneity and isotropy is small. Second, the distribution of the perturbations is statistically homogeneous and isotropic. At late times, when density perturbations become non-linear, the universe is no longer locally near homogeneity and isotropy, and there are deviations of order unity in quantities such as the local expansion rate. However, the distribution of the non-linear regions remains statistically homogeneous and isotropic on large scales. The homogeneity scale appears to be around 100 Mpc today \cite{Hogg:2005} (though see \cite{inhom}). Due to the statistical symmetry, the average expansion rate evaluated inside each box is equal (up to statistical fluctuations), but this does not mean that it would be the same as in a completely smooth spacetime, because there are structures in the box. The feature that the average evolution of a clumpy space is not the same as the evolution of a smooth space is called {\it backreaction} \cite{Ellis:2005, Rasanen:2006b, Buchert:2007}. We can say that time evolution and averaging do not commute: if we smooth a clumpy distribution and calculate the time evolution of the smooth quantities with the Einstein equation, the result is not the same as if we evolved the full clumpy distribution and took the average at the end. Put simply, FRW models describe universes which are exactly homogeneous and isotropic. They do not describe universes which are only statistically homogeneous and isotropic. The effect of clumpiness on the average was first discussed in detail by George Ellis in 1983 under the name {\it fitting problem} \cite{fitting}. Clumpiness affects the expansion of the universe, the way light propagates in the universe and the relationship between the two. The possibility that the late time deviations from the simple homogeneous and isotropic models would be explained in terms of these changes due to structure formation can be called the {\it backreaction conjecture}. In section 2 we go through the basics of how structures affect the expansion rate. An increased average expansion rate due clumpiness may be a bit unfamiliar concept, so in section 3 we explain what average acceleration means physically using a simple toy model, and in section 4 we go on to discuss a semi-realistic model for the universe where we can get some numbers out. We also briefly mention the relation of the backreaction problem to the fascinating question of the Newtonian limit of general relativity. In section 5 we discuss the relation of the average expansion rate to light propagation. The order is a bit backwards, as it is the observed redshifts and distances which are the important quantities, but it is perhaps easier to start from the expansion rate. We conclude in section 6 with a summary.

Observations of the universe at late times are inconsistent with homogeneous and isotropic FRW models which have ordinary matter and gravity. The problem is usually addressed by adding exotic matter or modifying general relativity. However, non-linear structures also influence the expansion rate: this is an effect which is present in reality but missing in FRW models. The Buchert equations which do include the effect of structures show that it is possible for the average expansion of a clumpy dust universe to accelerate, and there are toy models which demonstrate this. The physical explanation is simple: faster expanding regions increase their fraction of the volume more rapidly, so the average expansion rate increases. In a semi-realistic model, the correct timescale of about $10^5\teq\sim$ 10 billion years and the right order of magnitude for the change of the expansion rate emerge from the physics of structure formation without new parameters. In Newtonian cosmology, backreaction is necessarily small for a statistically homogeneous and isotropic distribution, but this is not the case in general relativity. Therefore, if backreaction has a significant effect on the expansion rate in the real universe, this is due to non-Newtonian aspects of general relativity. Even if backreaction is important, the relation between the redshift and the average expansion rate is the same as in FRW models, if the distribution of structures is statistically homogeneous and isotropic and evolves slowly. In contrast, the relation between the average expansion rate and the angular diameter distance is different from the FRW case. This relation is a unique backreaction prediction which makes it possible to distinguish the effect of non-linear structures from FRW dark energy or modified gravity models. The present estimates of the effect of structure formation on the average expansion rate cannot be trusted beyond an order of magnitude, and it is possible that a careful study will reveal cancellations which lead to a negligible effect. The relation between the average expansion rate and light propagation should also be studied more rigorously, and the difference between general relativity and Newtonian gravity in the cosmological non-linear regime remains to be fully understood. There is much work to be done before we can say whether or not the backreaction conjecture that the failure of ordinary homogeneous and isotropic models at late times is due to the breakdown of homogeneity and isotropy is correct. Until this effect has been quantified, we do not know whether new physics is needed to explain the observations, or if they can be understood in terms of a complex realisation of the physics we already know.

1012.0784

The predictions of homogeneous and isotropic cosmological models with ordinary matter and gravity are off by a factor of two in the late universe. One possible explanation is the known breakdown of homogeneity and isotropy due to the formation of non-linear structures. We review how inhomogeneities affect the average expansion rate and can lead to acceleration, and consider a semi-realistic model where the observed timescale of ten billion years emerges from structure formation. We also discuss the relation between the average expansion rate and observed quantities.

12307034

2012SCPMA..55..354L

1012

1012.5959_arXiv.txt

1012.5959

Brightest Cluster Galaxies (BCGs) are mostly elliptical galaxies and very rarely have prominent star formations. We found that five out of 8,812 BCGs are E+A (i.e., post-starburst) galaxies, having an H δ absorption line with an equivalent width &gt;2.5 Å and no distinct emission lines in [O II] and H α <SUP>-</SUP>. The E+A features we identified from the BCGs for the first time are not as significant as those in general galaxies, indicating that historically the star formations were not very violent.

12213660

2011MNRAS.413..223M

1012

1012.0098_arXiv.txt

\label{sec:Intro} Cepheids are pulsating stars with periods between one day and about one hundred days. Although there was originally some confusion, see the review by \citet{Fernie-1969}, it is now believed that Cepheids can be grouped into two distinct classes: classical Cepheids and type II Cepheids (T2Cs, hereafter). The former are intermediate mass stars (4--10~$M_\odot$), while the latters are lower-mass stars ($\sim 1~M_\odot$) belonging to disc and halo populations (\citealt{Wallerstein-2002}; \citealt{Sandage-2006}). T2Cs are conventionally divided into three period groups, BL~Her stars (henceforth BL) at short periods, W~Vir stars (WV) at intermediate periods and RV~Tau stars (RV) at the longest periods. According to Gingold~(\citeyear{Gingold-1976}, \citeyear{Gingold-1985}), the BLs are stars evolving from the horizontal branch to the Asymptotic Giant Branch (AGB); the WVs are stars that along the AGB cross the Cepheid instability strip due to excursions towards higher effective temperatures; the RVs are stars that after their AGB phase are moving towards the white dwarf cooling sequence. However, some later evolutionary tracks (e.g.~\citealt{Pietrinferni-2006}) do not show the excursions, and the precise evolution state of T2Cs remains unclear. We follow the division by \citet{Soszynski-2008b} and adopt thresholds of 4~d and 20~d to divide the three groups. The RV stars tend to show alternating deep and shallow minima (and this is often taken as a defining characteristic), but the single period of the RVs will be used in this paper. In addition to these three groups, \citet{Soszynski-2008b} established a new group of T2Cs, peculiar W~Vir (pW) stars. They have distinctive light curves and tend to be brighter than normal T2Cs of the same period. It is suggested that many, perhaps all, of them are binaries, but their nature remains uncertain. The bright Galactic T2C $\kappa$~Pav appears to be a pW star (Matsunaga, Feast \& Menzies, 2009, hereafter M09). \citeauthor{Matsunaga-2006}~(\citeyear{Matsunaga-2006}, hereafter M06) discovered a tight PLR of T2Cs in globular clusters based on infrared photometry for 46 T2Cs. \citet{Feast-2008} used pulsation parallaxes of nearby T2Cs to calibrate this cluster PLR and to discuss the distances of the Large Magellanic Cloud (LMC) and the Galactic Centre. The Optical Gravitational Lensing Experiment (OGLE-III) have significantly increased samples of T2Cs in the Magellanic Clouds. \citet{Soszynski-2008b} found 197 T2Cs in the LMC. M09 investigated their near-infrared (near-IR) nature and confirmed that BL and WV stars follow a tight PLR like that of T2Cs in globular clusters, whereas RV and pW stars show a large scatter and are systematically brighter than the PLR for the BL and WVs. The calibration by \citet{Feast-2008} then leads to a distance modulus of the LMC of $18.46\pm 0.05$~mag. These investigations have established the T2Cs as a promising distance indicator. The aim of the present paper is to compare the T2C PLRs in the Small Magellanic Cloud (SMC), LMC and globular clusters. This is important for a further study of these relations and their possible dependence on metallicity and other factors. The OGLE-III survey discovered 43 T2Cs in the SMC (\citealt{Soszynski-2010b}, S10 hereafter). In the following, we first collect near-IR $JHK$ magnitudes of SMC T2Cs (Section~\ref{sec:Photometry}) and investigate their PLRs (Section~\ref{sec:PLR}). The reddening-free PLR in $VI$ is also discussed. In parallel with our work, \citeauthor{Ciechanowska-2010}~(\citeyear{Ciechanowska-2010}, C10) have obtained independent $J$ and $K$ for a subset of the SMC T2Cs and their data is combined with ours in some of the discussions. We find evidence that the BL and WV stars need to be discussed independently. In Section~\ref{sec:delmu} the apparent differences in distance moduli of the LMC and SMC are derived for these stars and compared with the differences obtained from other objects. The data used for classical Cepheids are discussed in an Appendix. In addition, the colours and period distributions of T2Cs in various systems are compared and contrasted in Sections~\ref{sec:evolution}. Section~\ref{sec:Conclusion} summarizes the present work.

\label{sec:Conclusion} In this paper PLRs are derived for T2Cs in the SMC. Evidence is presented that T2C PLR slopes differ in three different systems (globular clusters, LMC, SMC) when BL and WV stars are used in combined solutions. Treating the BL and WV stars separately, it was found that the difference in distance moduli between the SMC and LMC derived from WVs alone agrees closely with that obtained from other distance indicators. This implies that the absolute magnitudes of the WVs are, within the uncertainties, free of metallicity effects, which is in agreement with the results derived from globular clusters. On the other hand the SMC-LMC difference is smaller for the BLs and suggests that their absolute magnitudes are not immune from population effects. The relative frequencies of WV to BL stars also varies from system to system.

1012.0098

Period-luminosity relations (PLRs) of type II Cepheids (T2Cs) in the Small Magellanic Cloud (SMC) are derived based on OGLE-III, IRSF/SIRIUS and other data, and these are compared with results for the Large Magellanic Cloud (LMC) and Galactic globular clusters. Evidence is found for a change of the PLR slopes from system to system. Treating the longer period T2Cs (W Vir stars) separately gives an SMC-LMC modulus difference of 0.39 ± 0.05 mag without any metallicity corrections being applied. This agrees well with the difference in moduli based on different distance indicators, in particular the PLRs of classical Cepheids. The shorter period T2Cs (BL Her stars) give a smaller SMC-LMC difference suggesting that their absolute magnitudes might be affected either by metallicity or by age effects. It is shown that the frequency distribution of T2C periods also changes from system to system.

12213640

2011MNRAS.414.3458W

1012

1012.2865_arXiv.txt

Radiation from stars and black holes strongly affects their surroundings and plays a crucial role in topics such as stellar atmospheres, the interstellar medium (ISM), star formation, galaxy formation, supernovae (SNe) and cosmology. It is a well-studied problem \citep[e.g.][]{Mathews65, Rybicki, Mihalas84, Yorke86}; however, its treatment in multi-dimensional calculations is difficult because of the dependence on seven variables -- three spatial, two angular, frequency, and time. The non-local nature of the thermal and hydrodynamical response to radiation sources further adds to the difficulty. Depending on the problem of interest some simplifying assumptions may be made. An important case was considered by \citet{Stroemgren39} for an ultraviolet (UV) radiation source photo-ionising a static uniform neutral medium. When recombinations balance photo-ionisations, the radius of a so-called \hii region, \begin{equation} \label{eqn:Rstr} R_s = \left(\frac{3\dot{N}_\gamma}{4\pi \alpha_{\rm B} n_{\rm H}^2}\right)^{1/3}, \end{equation} where $\dot{N}_\gamma$ is the ionising photon luminosity, $\alpha_{\rm B}$ is the recombination rate, and $n_H$ is the ambient hydrogen number density. Furthermore he found that the delineation between the neutral and ionised medium to be approximately the mean free path of the ionising radiation. His seminal work was expanded upon by \citet{Spitzer48, Spitzer49, Spitzer54} and \citet{Spitzer50}, who showed that the ionising radiation heated the medium to $T \sim 10^4$ K. If the density is equal on both sides of the ionisation front, then this over-pressurised region would expand and drive a shock outwards \citep[e.g.][]{Oort54, Schatzman55}. These early works provided the basis for the modern topic of radiation hydrodynamics of the ISM. A decade later, the first radiation hydrodynamical numerical models of \hii regions in spherical symmetry and plane-parallel ionisation fronts were developed \citep[e.g.][]{Mathews65, Lasker66, Hjellming66}. They described the expansion of the ionisation front and the evolution of its associated shock wave that carries most of the gas away from the source. At the same time, theoretical models of ionisation fronts matured and were classified by \citet{Kahn54} and \citet{Axford61} as either R-type (rare) or D-type (dense). In R-type fronts, the ionised gas density is higher than the neutral gas density, and in D-type fronts, the opposite is true. R-type fronts travel supersonically with respect to the neutral gas, whereas D-type fronts are subsonic. Furthermore ``weak'' and ``strong'' R-type fronts move supersonically and subsonically with respect to the ionised gas, respectively. The same terminology conversely applies to D-type fronts. ``Critical'' fronts are defined as moving exactly at the sound speed. These works established the evolutionary track of an expanding \hii illuminated by a massive star in a uniform medium: \begin{enumerate} \item Weak R-type: When the star (gradually) starts to shine, the ionisation front will move supersonically through the ambient medium. The gas is heated and ionised, but otherwise left undisturbed. This stage continues until $r \sim 0.02R_s$. \item Critical R-type: As the ionisation front moves outwards, it begins to slow because of the geometric dilution of the radiation. It becomes a critical R-type front, which is equivalent to an isothermal shock in the neutral gas. \item Strong and weak D-type: The front continues to slow, becoming a strong D-type front, and then a critical D-type front. From this point forward, it is moving subsonically with respect to the ionised gas, i.e. a weak D-type front. Thus sound waves can travel across the ionisation front and form a shock. The ionisation front detaches from the shock, putting the shock ahead of the ionisation front. \item Expansion phase: After the shock forms, the \hii region starts to expand, lowering the interior density and thus the recombination rates. This increases the number of photons available for ionising the gas. The sphere expands until it reaches pressure equilibrium with the ambient medium at $r \sim 5R_s$. \end{enumerate} In the 1970's and 1980's, algorithmic and computational advances allowed numerical models to be expanded to two dimensions, mainly using axi-symmetric to simplify the problem \citep[e.g.][]{Bodenheimer79, Sandford82, Yorke83}. One topic that was studied extensively were champagne flows. Here the source is embedded in an overdense region, and the \hii region escapes from this region in one direction. The interface between the ambient and dense medium was usually set up to be a constant pressure boundary. When the ionisation front passes this boundary, the dense, ionised gas is orders of magnitude out of pressure equilibrium as the temperatures on both sides of initial boundary are within a factor of a few. In response, the gas is accelerated outwards in this direction, creating a fan-shaped outflow. Only in the past 15 years, computational resources have become large enough, along with further algorithmic advances, to cope with the requirements of three-dimensional calculations. There are two popular methods to solve the radiative transfer equation in three-dimensions: \begin{itemize} \item Moment methods: The angular moments of the radiation field can describe its angular structure, which are related to the energy energy, flux, and radiation pressure \citep{Auer70, Norman98}. These have been implemented in conjunction with short characteristics \citep[2D]{Stone92_RHD}, with long characteristics \citep{Finlator09}, with a variable Eddington tensor in the optically thin limit \citep[OTVET;][]{Gnedin01_OTVET, Petkova09}, and with an M1 closure relation \citep{Gonzalez07, Aubert08}. Moment methods have the advantage of being fast and independent of the number of radiation sources. However, they are diffusive and result in incorrect shadows in some situations. \item Ray tracing: Radiation can be propagated along rays that extend through the computational grid \citep[e.g.][]{Razoumov99, Abel99_RT, Ciardi01, Sokasian01, Whalen06, Rijkhorst06, Mellema06, Alvarez06, Trac07, Krumholz07_ART, Paardekooper10} or particle set \citep[e.g.][]{Susa06, Johnson07, Pawlik08, Pawlik10, Altay08, Hasegawa09}. In general, these methods are very accurate but computationally expensive because the radiation field must be well sampled by the rays with respect to the spatial resolution of the grid or particles. \end{itemize} Until the mid-2000's the vast majority of the three-dimensional calculations were performed with static density distributions. One example is calculating cosmological reionisation by post-processing of density fields from N-body simulations \citep{Ciardi01, Sokasian01, McQuinn07, Iliev06, Iliev07}. Any hydrodynamical response to the radiation field was thus ignored. Several radiative transfer codes were compared in four purely radiative transfer tests in \citet[hereafter RT06]{RT06}. Only recently has the radiative transfer equation been coupled to the hydrodynamics in three-dimensions \citep[e.g.][]{Krumholz07_FLD}. In the second comparison paper \citep[hereafter RT09]{Iliev09}, results from these radiation hydrodynamics codes were compared. Even more rare are ones that couple it with magneto-hydrodynamics \citep[e.g.][]{Krumholz07_ART}. The tests in RT06 and RT09 were kept relatively simple to ease the comparison. In this paper, we present our implementation, \moray, of adaptive ray tracing \citep{Abel02_RT} in the cosmological hydrodynamics adaptive mesh refinement (AMR) code, \enzo~\citep{BryanNorman1997, OShea2004}. The radiation field is coupled to the hydrodynamics solver at small time-scales, enabling it to study radiation hydrodynamical problems. We have used this code to investigate the growth of an \hii region from a 100\Ms~Population III (Pop III) star \citep{Abel07}, the early stages of reionisation from Pop III stars \citep{Wise08_Reion}, radiative feedback on the formation of high redshift dwarf galaxies \citep{Wise08_Gal, Wise10_Gal}, ultraviolet radiation escape fractions from dwarf galaxies before reionisation \citep{Wise09}, negative radiative feedback from accreting Pop III seed black holes \citep{Alvarez09}, and radiative feedback in accreting supermassive black holes \citep[][in prep.]{Kim11}. We have included \moray~in the latest public release of \enzo\footnote{http://enzo.googlecode.com}, and it is also coupled with the newly added MHD solver in \enzo~\citep{Wang09}. We have structured this paper as follows. In Section 2, we describe the mathematical connections between adaptive ray tracing and the radiative transfer equation. Furthermore, we detail how physics other than photo-ionisation and photo-heating are included. We then derive a geometric correction factor to any ray tracing method to improve accuracy. We end the section by describing a new computational technique to approximate an optically-thin radiation field with ray tracing and multiple sources. In Section 3, we cover the details of our radiation hydrodynamics implementation in \enzo, specifically (1) the ray tracing algorithms, (2) coupling with the hydrodynamics solver, (3) several methods to calculate the radiative transfer timestep, and (4) our parallelisation strategy. We present our results from the RT06 radiative transfer tests in Section 4. Afterwards in Section 5, we show the results from the RT09 radiation hydrodynamics tests. In Section 6, we expand on these tests to include more dynamical and complex setups to demonstrate the flexibility and high fidelity of \moray. Section 7 gives the results from spatial, angular, frequency, and temporal resolution tests. In Section 8, we illustrate the improvements from the geometric correction factor and our optically-thin approximation. We also show the effects of X-ray radiation and radiation pressure in this section. Finally in Section 9, we demonstrate the parallel scalability of \moray. Last Section 10 summarises our method and results.

In this paper, we have presented our implementation, \moray, of adaptive ray tracing \citep{Abel02_RT} and its coupling to the hydrodynamics in the cosmology AMR code \enzo, making it a fully functional radiation hydrodynamics code. As this method is photon conserving, accurate solutions are possible with coarse spatial resolution. A new geometric correction factor to ray tracing on a Cartesian grid was described, and it is general to any implementation. We have exhaustively tested the code to problems with known analytical solutions and the problems presented in the RT06 and RT09 radiative transfer comparison papers. Additionally we have tested our code with more dynamical problems -- champagne flows, Rayleigh-Taylor instabilities, photo-evaporation of a blastwave, beamed radiation, a time-varying source, and an \hii region with MHD -- to demonstrate the flexibility and fidelity of \moray. Because production simulations may not have the resolution afforded in these test problems, we have tested the dependence on spatial, angular, frequency, and temporal resolution. It provides accurate solutions even at low resolution, except for the large constant timesteps. However, we have described two methods to determine the radiative transfer timestep that are based on the variations in specific intensity or changes in neutral fraction inside the ionisation front. Both methods give very accurate results and provide the largest timestep to obtain an accurate solution, ultimately leading to higher computational efficiency. On the same topic, we have described a method to calculate the radiation field in the optically-thin limit with ray tracing. Being a ray tracing code, it scales with the number of radiation sources; nevertheless, it scales well to O($10^3$) processors for problems with $\sim10^9$ computational cells and $\sim10^4$ sources, such as reionisation calculations. We have also shown that the code shows good strong scaling in AMR calculations, given a large enough problem. The combination of AMR and adaptive ray tracing allows for high-resolution and high-dynamical range problems, e.g. present-day star formation, molecular cloud resolving cosmological galaxy formation, and \hii regions of Population III stars. Furthermore, we have included Lyman-Werner absorption, secondary ionisations from X-ray radiation, Compton heating from photon scattering, and radiation pressure into the code, which extends the reach of \moray~to study AGN feedback, stellar winds, and local star formation. Coupling the radiative transfer with MHD further broadens the applicability of our code. The full implementation is included in the latest public version of \enzo\footnote{http://enzo.googlecode.com}, providing the community with a full-featured radiation hydrodynamics AMR code.

1012.2865

We describe a photon-conserving radiative transfer algorithm, using a spatially-adaptive ray-tracing scheme, and its parallel implementation into the adaptive mesh refinement cosmological hydrodynamics code ENZO. By coupling the solver with the energy equation and non-equilibrium chemistry network, our radiation hydrodynamics framework can be utilized to study a broad range of astrophysical problems, such as stellar and black hole feedback. Inaccuracies can arise from large time-steps and poor sampling; therefore, we devised an adaptive time-stepping scheme and a fast approximation of the optically-thin radiation field with multiple sources. We test the method with several radiative transfer and radiation hydrodynamics tests that are given in Iliev et al. We further test our method with more dynamical situations, for example, the propagation of an ionization front through a Rayleigh-Taylor instability, time-varying luminosities and collimated radiation. The test suite also includes an expanding H II region in a magnetized medium, utilizing the newly implemented magnetohydrodynamics module in ENZO. This method linearly scales with the number of point sources and number of grid cells. Our implementation is scalable to 512 processors on distributed memory machines and can include the radiation pressure and secondary ionizations from X-ray radiation. It is included in the newest public release of ENZO.

5273514

2011ASSP...27..135M

1012

1012.3114_arXiv.txt

It is well known that the cluster environment affects the physical properties of galaxies, as there is a segregation with projected radius of morphology \citep{Dressler80}, color (e.g. [\citealp{Balogh+04_ApJL}]), luminosity (e.g. [\citealp{ABM98}]) and spectral indices such as the equivalent width of [OII] or H$_\alpha$ (e.g. [\citealp{Balogh+04_MN,Haines+06}]). At the same time, there are indications of velocity segregation of luminosity \citep{Biviano+92} and of star formation efficiency: emission-line galaxies tend to span a wider dispersion of velocities than galaxies without such lines \citep{Biviano+97}. In fact, it was long known that spiral galaxies in clusters span a wider distribution of velocities than ellipticals and S0s (e.g. [\citealp{MD77}]). However, those spirals without emission lines and that are not morphologically disturbed span the same velocity distribution as the early type galaxies \citep{Moss06}. Relative to passive galaxies (without H$_\alpha$ in emission), the trend for higher velocity dispersions of galaxies with emission (H$_\alpha$) lines found by \citep{Biviano+97} is reversed outside the virial radius \citep{RGKD05}. In the present work \citep{MMR11}, we take advantage of the large statistics of the Sloan Digital Sky Survey (SDSS) to better quantify the velocity modulation of the radial segregation of the diagnostics of star formation efficiency. We then model the SDSS observed fractions of Galaxies with Ongoing or Recent Efficient Star Formation (GORES), with the help of a cosmological simulation.

1012.3114

Using H<SUB>δ</SUB> and D<SUB> n </SUB>4000 as tracers of recent or ongoing efficient star formation, we analyze the fraction of SDSS galaxies with recent or ongoing efficient star formation (GORES) in the vicinity of 268 clusters. We confirm the well-known segregation of star formation, and using Abel deprojection, we find that the fraction of GORES increases linearly with physical radius and then saturates. Moreover, we find that the fraction of GORES is modulated by the absolute line-of-sight velocity (ALOSV): at all projected radii, higher fractions of GORES are found in higher ALOSV galaxies. We model this velocity modulation of GORES fraction using the particles in a hydrodynamical cosmological simulation, which we classify into virialized, infalling and backsplash according to their position in radial phase space at z = 0. Our simplest model, where the GORES fraction is only a function of class does not produce an adequate fit to our observed GORES fraction in projected phase space. On the other hand, assuming that in each class the fraction of GORES rises linearly and then saturates, we are able to find well-fitting 3D models of the fractions of GORES. In our best-fitting models, in comparison with 18% in the virial cone and 13% in the virial sphere, GORES respectively account for 34% and 19% of the infalling and backsplash galaxies, and as much as 11% of the virialized galaxies, possibly as a result of tidally induced star formation from galaxy-galaxy interactions. At the virial radius, the fraction of GORES of the backsplash population is much closer to that of the virialized population than to that of the infalling galaxies. This suggests that the quenching of efficient star formation is nearly complete in a single passage through the cluster.

12104189

2010ApJ...721..318O

1012

1012.3578_arXiv.txt

According to the model of AGNs, cold material close to the central black hole forms an accretion disc, matter inside which due to the dissipative forces transports inwards causing the accretion disc to heat up. Such a hot material in turn can inverse-Compton scatter photons up to X-ray energies \citep{blan}. From high energy astronomical sources a special interest deserve blazar type AGNs the standard model of which implies the presence of the supermassive black hole, surrounded by the accretion disc and ejecting twin relativistic jets. The observationally evident broadband emission spectrum of blazars is made of two components: the low energy (from radio to optical) domain attributed to the synchrotron emission and the high energy (from $X$-rays to $\gamma$-rays) part formed by either the inverse-Compton mechanism \citep{blan} or the curvature radiation \citep{g96,tg05}. A recent investigation of the parsec scale jets is very important. \cite{giro} argued that the high energy and radio emissions are strongly correlated. Our model on the other hand, as we will see, automatically provides a connection of radiation in high energy and radio domains. Magnetospheres of AGNs have strong magnetic fields, therefore, synchrotron cooling timescales are relatively short, leading to efficient energy losses. This in turn creates appropriate conditions for particles to transit to their ground Landau state. When this happens, relativistic electrons will move only along magnetic field lines without emitting in the synchrotron regime. Despite the very strong magnetic field, which prevents a continuous process of the synchrotron emission, there is a possibility to overcome the dissipative factors and maintain the radiation mechanism. \cite{machus1} have studied the cyclotron instability of two-component electron-positron plasma for the pulsar, NP 0532. It was found that the instability arises near the light cylinder surface (a hypothetical zone, where the linear velocity of rigid rotation equals exactly the speed of light) leading to a certain distribution of particles by pitch angles and the consequent synchrotron radiation. \cite{lomin} considered the magnetospheres of the pulsar NP 0532 and the Crab nebula, studying the generation of waves from optical to gamma-ray domains. The similar approach was presented by \cite{malmach} where the QLD was applied to the radio pulsars. The authors found that the transverse momenta of relativistic particles induced by the cyclotron instability caused the stable non-zero pitch angle distribution maintained by means of the QLD. Analyzing the data obtained from MAGIC Cherenkov telescope between 2007 October and 2008 February \citep{magic} we found that, the observed coincidence of signals in the optical and $\gamma$-ray domains are easily explained by the QLD process, which leads to the increase of the pitch angles, making the synchrotron process feasible \citep{difus,difus1}. In the magnetospheres of AGNs the magnetic fields are of the order of $10^4G$ \citep{paradig}, close to the supermassive black hole, and $100G-300G$ close to the light cylinder surface. Therefore, the aforementioned QLD mechanism could be of great importance for AGNs as well. For this purpose by considering the cyclotron instability excited in the radio domain, in \citep{difus3} we studied the quasi-linear interaction of proper modes of AGN magnetospheric plasmas with the resonant plasma particles investigating the QLD in the context of producing the soft and hard X-ray emission from AGNs. Under favorable conditions this mechanism could also be efficient for explaining the MeV-GeV energy synchrotron emission, strongly connected either with the submilimeter radio band, or with the infrared emission induced by the cyclotron instability. This will be the subject of the present paper, which is organized as follows. In Section 2 we describe our model, in Sect. 3 we apply the mechanism of QLD to AGNs and in Sect. 4 we summarize our results.

\label{sec:summary} In this section we apply the model of the QLD to the light cylinder lengthscales of typical AGNs. For explaining the high energy radiation, it is strongly believed that AGN magnetospheres are consist of highly relativistic electrons. This fact sets another problem - how to accelerate these particles to such high energies? In general, there are several mechanisms, which may account for the efficient acceleration of electrons. Indeed, as is shown in a series of papers, Fermi-type acceleration process \citep{cw99}, re-acceleration of electron-positron pairs as a feedback mechanism \citep{ghis} and centrifugal acceleration \citep{mr94,osm7,ra08,osm10} may provide very high Lorentz factors of the order of $\gamma_b\sim 10^{5-9}$. Therefore, in the framework of the paper an existence of such particles is assumed to be as a given fact. If we suppose an isotropic distribution of relativistic electrons, one can estimate the synchrotron cooling timescale \citep{difus3} \begin{equation}\label{ts} t_{cool}\approx 5\times 10^{-3}\times\left(\frac{10^2G}{B}\right)^2\times\left(\frac{10^8}{\gamma}\right)s. \end{equation} The value of the magnetic induction is given by \citep{difus3} \begin{equation} \label{b} B_{lc}\approx 260\times\left(\frac{L}{10^{45}erg/s}\right)^{1/2}\times\left(\frac{\Omega}{3\times 10^{-5}s^{-1}}\right)G, \end{equation} where $L$ is the bolometric luminosity of the AGN and $r_{lc} = c/\Omega$ is the light cylinder radius (a hypothetical zone, where the linear velocity of rigid rotation exactly equals the speed of light). $\Omega$ is the magnetic field lines' angular velocity of rotation, normalized to the value $3\times 10^{-5}s^{-1}$ \citep{belv}. Throughout the paper we assume that the magnetic field is robust enough to maintain the frozen-in condition in the magnetosphere of the AGN. Indeed, as one can see, for the typical magnetospheric parameters, $\gamma_b\sim 10^8$, $n_b\sim 10 cm^{-3}$, the following condition $B_{lc}^2/8\pi>\gamma_b mn_bc^2$ is satisfied, which means that the plasma particles will be forced to follow the rigidly rotating field lines. During such a motion, especially on the light cylinder lengthscales, the electrons will undergo the centrifugal force accelerating them up to very high Lorentz factors $\sim 10^{8-9}$ \citep{osm7,ra08} It is clear from Eq. (\ref{ts}) that for a certain class of physical parameters the synchrotron cooling timescale is of the order of $5\times 10^{-4}s$. On the other hand, the kinematic timescale of the system, $t_{kin}\sim r_{lc}/c\sim 3\times 10^4s$ is by many orders of magnitude bigger than $t_{cool}$, which in turn means that without the quasi-linear diffusion, particles very soon would stop emitting in the synchrotron regime, after transiting to their ground Landau level. \cite{kmm} showed that the anomalous Doppler effect generates the cyclotron waves with the frequency \citep{difus3} $$\omega\approx 6.8\times 10^9\times\left(\frac{\gamma_p}{100}\right)^4\times\left(\frac{10^8} {\gamma_b}\right)^2\times$$ \begin{equation} \label{om} \times\left(\frac{B}{100G}\right)^3 \times\left(\frac{10cm^{-3}}{n_b}\right)Hz, \end{equation} leading to the process of the quasi-linear diffusion, which, despite the efficient dissipative factors, creates the pitch angles. To demonstrate the present model, we consider an AGN with the bolometric luminosity $L = 10^{45}erg/s$. Let us examine the following parameters $\Omega = 3\times 10^{-5}rad/s$, $\gamma_b = 10^8$ and $n_b = 10 cm^{-3}$. Since the particles are distributed by the pitch angles [see Eq. (\ref{chi})], for analyzing the synchrotron emission it is reasonable to estimate a mean value of $\psi$ \begin{equation}\label{pitch} \bar{\psi} = \frac{\int_{0}^{\infty}\psi f(\psi)d\psi}{\int_{0}^{\infty}f(\psi)d\psi} \approx \frac{0.5}{\sqrt[4]{A}}. \end{equation} Then one can show from Eq. (\ref{pitch}) that for the aforementioned parameters the pitch angle is of the order of $8\times 10^{-3}rad$, therefore, relativistic electrons will inevitably emit photons with energies \citep{Lightman} \begin{equation} \label{eps} \epsilon_{eV}\approx 1.2\times 10^{-8}B\gamma^2\sin\psi. \end{equation} After substituting the value of $\overline{\psi}$ in Eq. (\ref{eps}), we see that the synchrotron emission generates radiation in the MeV domain. The quasi-linear diffusion works if the cyclotron modes are excited, therefore, it is essential to estimate the timescale of the corresponding instability ($t_{ins}$) and compare it with the kinematic timescale of the system. According to the work of \cite{kmm} the growth rate of the instability is given by \begin{equation}\label{inc1} \Gamma = \pi \frac{\omega_b^2}{\omega\gamma_p} \;\;\;\ if \;\;\;\ \frac{1}{2}\frac{u_x^2}{c^2}\ll\delta \end{equation} and \begin{equation}\label{inc2} \Gamma = \pi \frac{\omega_b^2}{2\omega\gamma_p}\frac{u_x^2}{\delta\cdot c^2} \;\;\;\ if \;\;\;\ \frac{1}{2}\frac{u_x^2}{c^2}\gg\delta, \end{equation} where $\omega_b\equiv\sqrt{4\pi n_b e^2/m}$ is the plasma frequency of beam electrons. It is easy to show that for $n_b = 10cm^{-3}$, $\gamma_p = 200$ (see Fig. \ref{fig1}), $V_{_{\parallel}}\sim c$ and $\rho\sim R_g$, one obtains $u_x^2/(2c^2)\ll 1$, implying that the increment of the instability is given by Eq. (\ref{inc1}). The cyclotron resonance makes sense if $t_{ins}/t_{kin}<1$, then, by taking into account the definition of the kinematic timescale, $r_{lc}/c$ and the instability timescale, $1/\Gamma$, one can show that the aforementioned condition reduces to \begin{equation} \label{cond}3.5\times10^{-3}\times\left(\frac{\gamma_p}{100}\right)^5\times\left(\frac{10^8} {\gamma_b}\right)^2\times\left(\frac{10cm^{-3}}{n_b}\right)^2<1. \end{equation} As is clear from Eq. (\ref{cond}), the condition is very sensitive to the Lorentz factor of the plasma components, and for relatively higher values of $\gamma_p$ the condition will violate. The upper limit of $\gamma_p$, when the condition is still valid is of the order of $300$ for $\gamma_b\sim 10^8$ and $n_b\sim 10cm^{-3}$. For studying the efficiency of the QLD, we examine the following set of the parameters $L = 10^{45}erg/s$, $\Omega = 3\times 10^{-5}rad/s$, $\gamma_p = 200$ and $n_b = \{5;10;15\}cm^{-3}$ and the results are demonstrated in Fig.\ref{fig1} where we show the behaviour of $\epsilon_{_{MeV}}$ versus $\gamma_b$. From the plots is clear that $\epsilon_{_{MeV}}$ is a continuously increasing function of the beam Lorentz factor, which is a natural result of the fact that more energetic particles produce more energetic photons. The behaviour of $\epsilon_{_{MeV}}$ versus $n_b$ is different, more dense beam electrons produce photons with lower energies. This can be seen from Eqs. (\ref{A},\ref{pitch}): $\bar{\psi}\sim\sqrt[4]{D_{_{\perp\perp}}}$, which by combining with $D_{_{\perp\perp}}\sim n_b^3$ [see Eqs. (\ref{dif},\ref{ek2})] confirms the dependence $\epsilon_{_{MeV}}(n_b)$. According to the results demonstrated in the figure, relativistic electrons with Lorentz factors $\gamma_b=\{1-2\}\times 10^8$ may provide the high energy radiation in the MeV-GeV domain. Since the generation of the synchrotron emission strongly depends on the cyclotron waves, we also investigate the behaviour of $\epsilon_{_{MeV}}$ versus $\omega$. Figure \ref{fig2} shows the function $\epsilon_{_{MeV}}(\omega)$ for several values of $n_b$. The set of parameters is the same as in the previous figure. As is clear from the plots, the $\{200-1200\}$MeV radiation is strongly connected with the submillimeter ($\sim [0.3-3]\times 10^{12}Hz$) and low infrared ($\sim [3-3.8]\times 10^{12}Hz$) emission. Another important parameter, the physical system depends on, is the Lorentz factor of the plasma component. Therefore, it is reasonable to demonstrate the function $\epsilon_{_{MeV}}(\omega)$ for different values of $\gamma_p$. These results are shown in Fig. \ref{fig3}, where the set of parameters is $L = 10^{45}erg/s$, $\Omega = 3\times 10^{-5}rad/s$, $\gamma_p = \{200;250;300\}$ and $n_b = \{5;10;15\}cm^{-3}$. As we see, the excited infrared domain extends up to $\sim 10^{13}Hz$, which is strongly connected with high energy emission ($100MeV$). As is clear from the figure, higher values of $\gamma_p$ correspond to lower synchrotron energies. Indeed, by taking into account the relation $\bar{\psi}\sim\sqrt[4]{D_{_{\perp\perp}}}$ combined with Eqs. (\ref{dif},\ref{ek2}) one can see that $\bar{\psi}\sim 1/\gamma_p^2$. As is clear from the results, the quasi-linear diffusion together with the cyclotron instability may guarantee production of high energy radiation in the MeV-GeV domains strongly connected with submillimeter/infrared emission. The major difference in results from our previous work is that in \citep{difus3} we studied physical conditions leading to excitation of $X$-rays connected with the relatively low frequency radio band (KHz-MHz), whereas in the present paper both energies (produced as by synchrotron as by cyclotron mechanisms respectively) are much higher. This investigation sets another problem: since the QLD is a feasible mechanism providing the aforementioned high energies, one of the important next steps could be testing of MeV-GeV AGNs exhibiting an efficient submillimeter/infrared radiation and see if the strong correlation is observationally evident. This in turn, could be a certain test for estimating the AGN magnetospheric parameters, such as the density and the Lorentz factors of plasma component and beam electrons. A particular future objective is to investigate theoretically the radiative signatures of both high and low energy emissions respectively. This work we are going to perform sooner or later. The main aspects of the present work can be summarized as follows: \begin{enumerate} \item Mechanisms producing strongly connected high and low energy radiation was studied by taking into account the QLD in the AGN magnetospheres. Considering a physical regime different from that of \citep{difus3}, we investigate the efficiency of the QLD in a region close to the light cylinder surface. \item For the considered physical parameters, it has been shown that the cyclotron instability appears for relatively low frequency range, producing radiation in the submillimeter/infrared domains. On the other hand, despite the short cooling timescales, the effect of diffusion on particles recreates the pitch angles and produces the high energy radiation in the MeV-GeV bands. \item The problem was studied versus three major magnetospheric parameters: the beam and the plasma component's Lorentz factors and the beam electrons' density. It was shown that the photon energy, $\epsilon_{_{MeV}}$, is a continuously increasing function of the beam Lorentz factor and the beam density. Contrary to this, by increasing the plasma component Lorentz factor, the corresponding photon energy decreases. \end{enumerate}

1012.3578

For active galactic nuclei (AGNs), we study the role of the mechanism of quasi-linear diffusion (QLD) in producing the high-energy emission in the MeV-GeV domains strongly connected with the submillimeter/infrared radiation. Considering the kinetic equation governing the stationary regime of the QLD, we investigate the feedback of the diffusion on electrons. We show that this process leads to the distribution of particles by pitch angles, implying that the synchrotron mechanism is no longer prevented by energy losses. Examining a reasonable interval of physical parameters, we show that it is possible to produce MeV-GeV γ-rays that are strongly correlated with submillimeter/infrared bands.

3826094

2012MNRAS.421.1136A

1012

1012.4996_arXiv.txt

Clusters of galaxies contain large reservoirs of hot, ionized gas. This plasma, although invisible in the optical waveband, can be observed in both X-ray and microwave bands of the electromagnetic spectrum through thermal Bremsstrahlung radiation and its scattering of the cosmic microwave background (CMB) respectively. This inverse Compton scattering results in a decrement in the intensity of CMB photons in the direction of the cluster at frequencies $< 218$~GHz, and is known as the Sunyaev--Zel'dovich (SZ) effect \citep{1970CoASP...2...66S, 1999PhR...310...97B, 2002ARA&A..40..643C}. To describe the full spectral behaviour of the SZ effect, one needs to consider three main components. These include the thermal SZ effect caused by thermal (random) motion of scattering electrons, including thermal weakly relativistic electrons, the kinematic SZ effect caused by peculiar velocity of the cluster with respect to Hubble flow, and relativistic effects caused by presence of the energetic nonthermal electrons in the hot plasma of the cluster that are responsible for synchrotron emission of radio halos or relics. However, since the last two processes have significantly smaller effects on the overall spectral distortion at cm wavelengths, we only consider the thermal SZ effect in this paper. Moreover, we ignore the effects of weakly relativistic thermal electrons, which are negligible at cm wavelengths. A main science driver for studying clusters through their thermal SZ signal arises from the fact that SZ surface brightness is independent of redshift. This provides us with a powerful opportunity to study galaxy clusters out to high redshift. However, estimating the physical properties of the clusters depends strongly on the model assumptions. In this paper we aim to show how employing different parameterisations for a cluster model affects the constraints on cluster properties. These tasks are conveniently carried out through Bayesian inference using a highly efficient parameter space sampling method: nested sampling (Skilling 2004). This sampling method is employed using the package \textsc{Multinest} (Feroz \& Hobson 2008; Feroz, Hobson \& Bridges 2009).\textsc{ Multinest} explores the high dimensional parameter space and calculates both the probability distribution of cluster parameters and the Bayesian evidence. This algorithm is employed to analyse real multi-frequency SZ observations made by the Arcminute Microkelvin Imager (AMI), (AMI Consortium: Zwart et~al.\ 2008). The rest of the paper is organised as follows. In Section~2, we describe the AMI telescope. In Section~3, we discuss Bayesian inference. Section~4 gives details of how we model interferometric SZ data. In Section 5, we describe the modelling of the SZ signal using both the isothermal $\beta$-model and an "entropy"- GNFW pressure model. Section 6 outlines the assumptions needed to estimate cluster physical parameters and describes how different parameterisations introduce different constraints and biases in the resulting marginalised posterior probability distributions. In Section 7, we describe how to generate a simulated SZ cluster in a consistent manner for both models, and in Section 8, we present our results. Finally, Section 9 summarises our conclusions.

We have studied two parameterised models, the traditional isothermal $\beta$-model and the ``entropy''-GNFW pressure model, to analyse the SZ effect from galaxy clusters and extract their physical parameters using AMI SA simulated data. In our analysis we have described the current assumptions made on the dynamical state of the ICM including spherical geometry, hydrostatic equilibrium and the virial mass-temperature relation. In particular we have shown how different parameterisations which relate the thermodynamical quantities describing the ICM to the cluster global properties via these assumptions lead to biases on the cluster physical parameters within a particular model. In this context, we first generated a simulated cluster using the isothermal $\beta$-model observed with the AMI SA and used these simulated data to study three different parameterisations in deriving the cluster physical parameters. We showed that in generating AMI simulated data, it is extremely important to select the model parameters describing the SZ signal in a way that leads to the consistent cluster parameter inferences upon using the three different parameterisation methods. We found that each parameterisation introduces different constraints and biases in the posterior probability distribution of the inferred cluster parameters which arise from the way we implement assumptions about the cluster structure and its composition. The biases in the posterior probability distributions of the cluster parameters are more pronounced in parameterisations I and II, as the results depend strongly on the relatively unconstrained cluster model shape parameters: $r_{\rm c}$ and $\beta$. However, the biases introduced by the choice of priors are even worse in parameterisation I, in which the gas temperature is assumed to be an independent free parameter. This, along with the assumption of isothermality, causes the priors to dominate in extracting the cluster physical parameters regardless the type of prior chosen for the gas temperature (AMI Consortium: Rodr\'{i}guez-Gonz\'{a}lvez et~al.\ 2011 and AMI Consortium: Zwart et~al.\ 2010). The cluster physical parameters estimated using parameterisation I depend strongly on the model parameters. Although it can constrain the cluster position and its $M_{\rm g}(r_ {\rm 200})$, it fails to recover the true input values of most of the simulated cluster properties. For example the inferred values for mass and temperature at $r_{\rm200}$ are $M_{\rm T}(r_{\rm 200})=(6.43\pm 5.43)\times 10^{15}\,\rm{M_\odot}$ and $T_{\rm g}(r_{\rm 200})= (10.61 \pm 5.28)\,\rm{keV}$ whereas the corresponding input values of simulated cluster are: $M_{\rm T}(r_{\rm 200})=5.83\times 10^{14}\,\rm{M_\odot} $ and $T_{\rm g}(r_ {\rm 200})=5\, \rm{keV} $. In terms of the application to the real data, we have noticed similar biases in the results of our analysis of 7 clusters using this parameterisation (AMI Consortium: Zwart et~al. 2010). In order to improve our analysis methodology in parameterisations II and III, the correlation between the cluster total mass and its gas temperature is taken into account. In parameterisation II we relate $M_{\rm T}(r_ {\rm 200})$ and $T_{\rm g}(r_{\rm 200})$ using the hydrostatic equilibrium whereas in parameterisation III we use virial mass-temperature relationship. It should be noted that the derived $T_{\rm g}(r_{\rm 200})$ in parameterisation II is the gas temperature at the overdensity radius $r_{\rm 200}$ which is then assumed to be constant throughout the cluster. In parameterisation III, however, $T_{\rm g}(r_{\rm 200})$ is the mean gas temperature internal to radius $r_{\rm 200}$ and is assumed to be constant. We notice that analysing the same simulated data set using parameterisation II can constrain the 1-D posterior distribution of the cluster physical parameters better than parameterisation I such that $M_{\rm T}(r_{\rm 200})= (6.8 \pm 2.1)\times 10^{14}\,\rm{M_ \odot}$ and $T_{\rm g}(r_{\rm 200})= (3.0 \pm 1.2)\,\rm{keV}$. Since parameterisation II uses the full parametric hydrostatic equilibrium, the temperature estimate depends on $r_{\rm c}$ and $\beta$ and is therefore biased low. These results were also confirmed in our analysis of the bullet like cluster A2146 (AMI Consortium: Rodr\'{i}guez-Gonz \'{a}lvezet~al.\ 2011). Relating the cluster total mass and its temperature via virial theorem in parameterisation III leads to less bias in cluster physical parameters compared to the other two parameterisations as it is less model dependent: $M_{\rm T}(r_ {\rm 200})= (4.68 \pm 1.56)\times 10^{14}\,\rm{M_\odot}$ and $T_{\rm g}(r_{\rm 200})= (4.3 \pm 0.9)\,\rm{keV}$. A detailed comparison between our different parameterisations both using simulated data and on the bullet like cluster A2146 (AMI Consortium: Rodr\'{i}guez-Gonz\'{a}lvez et~al. \ 2011) found that parameterisation III can give more reliable results for cluster physical properties as it is less dependent on model parameters. Parameterisation II also gives convincing estimates for the cluster total mass and its gas content although its temperature estimate is poorly justified, as it depends strongly on the model parameters. Moreover, young or disturbed clusters are unlikely to be well-described by hydrostatic equilibrium. We therefore used parameterisation III as our adopted analysis methodology in our follow-up studies of the real clusters including the joint SZ and weak lensing analysis of six clusters (AMI Consortium: Hurley-Walker et~al.\ 2011) and the analysis of LoCuss cluster sample (AMI Consortium: Rodr\'{i}guez-Gonz \'{a}lvez et~al.\ 2011; AMI Consortium: Shimwell et~al.\ 2011). In order to make sure that our results are not biased by one realisation of primordial CMB, we have studied $\rm {100}$ CMB realisations for the three parameterisations. The 1-D marginalised posterior probability distributions of $M_{\rm T}(r_{\rm 200})$ and $T_{\rm g}(r_{\rm 200})$ are shown in Figs. 21, 22 and 23 for each parameterisation. The solid blue line represents the true value corresponding to the simulated cluster and the dashed red line shows the mean value of the distributions. Table 13 also presents the numerical results of this analysis. \begin{table} \caption{The results of $100$ CMB realisations for the three parameterisations assuming h=0.7.} \begin{tabular}{@{}ccc@{} } \hline parameterisation & $M_{\rm T}(r_{\rm 200})\,\rm{M_\odot}$ & $T_{\rm g}(r_{\rm 200})\,\rm {keV}$ \\\hline I & $(6.18 \pm 5.23)\times 10^{15}$ & $11.18 \pm 5.16$ \\ II & $(8.067 \pm 2.61)\times 10^{14}$ & $3.94 \pm 1.67$ \\ III & $(5.94 \pm 2.26)\times 10^{14}$ & $4.97 \pm 1.21$ \\ \hline \end{tabular} \end{table} Comparing the 1D posterior distributions along with the mean values of the $M_{\rm T}(r_ {\rm 200})$ estimates in the three parameterisations for these $100$ realisations show that parameterisation I can hardly constrain the simulated cluster properties and recover the input true values. Parameterisation II can constrain the cluster total mass, however, the gas temperature estimate is biased low as it depends on unconstrained model shape parameters. On the other hand, parameterisation III can indeed constrain both cluster mass and its gas temperature and the results are unbiased. \begin{figure} \includegraphics[width=80mm]{jointplotpar1.ps} \caption{$100$ realisations of 1-D marginalised posterior probability distributions of $M_{\rm T}(r_{\rm 200})$ and $T_{\rm g}(r_{\rm 200})$ using isothermal $/beta$-model--parameterisation I. The solid blue line represents the true value corresponding to the simulated cluster and the dashed red line shows the mean value of the distributions.} \end{figure} \begin{figure} \includegraphics[width=80mm]{jointplotpar2.ps} \caption{$100$ realisations of 1-D marginalised posterior probability distributions of $M_{\rm T}(r_{\rm 200})$ and $T_{\rm g}(r_{\rm 200})$ using isothermal $\beta$-model-- parameterisation II. The solid blue line represents the true value corresponding to the simulated cluster and the dashed red line shows the mean value of the distributions.} \medskip \includegraphics[width=80mm]{jointplotpar3.ps} \caption{$100$ realisations of 1-D marginalised posterior probability distributions of $M_{\rm T}(r_{\rm 200})$ and $T_{\rm g}(r_{\rm 200})$ using isothermal $\beta$-model-- parameterisation III. The solid blue line represents the true value corresponding to the simulated cluster and the dashed red line shows the mean value of the distributions.} \end{figure} In order to remove the assumption of isothermality which is of course a poor assumption both within the cluster inner region and at the large radii and to improve our analysis model for the cluster ICM which can be fitted accurately throughout the cluster, we also studied the SZ effect using ``entropy''-GNFW pressure model. This model assumes a 3-D $\beta$-model like radial profile describing the entropy in the ICM as well as the GNFW profile for the plasma pressure. This choice is reasonable as the entropy is a conserved quantity and describes the structure of the ICM while the pressure is related to the dark matter component of the cluster. Moreover, among all the thermodynamical quantities describing the ICM, entropy and pressure show more self- similar distribution in the outskirts of the cluster. The combination of these two profiles then allows us to relate the SZ observable properties to the cluster physical parameters such as its total mass. This model also allows the electron pressure and its number density profiles to have different distributions leading to a 3-D radial temperature profile. In this context we simulated a second cluster using an entropy-GNFW pressure profile with the same physical parameters and thermal noise as the first cluster at $r_{\rm 200}$. We then analysed the second simulated cluster using "entropy"-GNFW pressure model with different parameterisations. In this model temperature is no longer isothermal so that we can not use parameterisation I where a single temperature is assumed as an independent input parameter. The results of our analysis using parameterisation II and III show that while the characteristic scaling radius describing the GNFW pressure profile is constrained, the shape parameters defining the entropy profile remain unconstrained. Moreover, all the cluster physical parameters lie within $1 \sigma$ errorbars from the corresponding true values of the simulated cluster in the two parameterisations. However, parameterisation III provides tighter constrains in 1-D marginalised posterior distribution of the temperature and the overall results are less model dependent so that it can be reliably used in the analysis of galaxy clusters in particular when the assumption of hydrostatic equilibrium breaks (e.g. in disturbed clusters and clusters that are going through merging). We conclude that using the ``entropy''-GNFW pressure model overcomes the limitations of the isothermal $\beta$-model in fitting cluster parameters over a broad radial extent. However, AMI simulated data do not strongly prefer one model over the other. We investigated this conclusion further by fitting both GNFW pressure profile and isothermal $\beta$-model to a simulated cluster with $\theta_{\rm 500}=2.5 \arcmin$ and $Y_{\rm 500}=2.5 \times 10^{-3}(\rm{arcmin})^{2}$. The result is shown in Fig. 24 with blue dashed line representing the fit using the isothermal $\beta$-model and the red representing the fit using GNFW pressure profile. However, we aim to compare these two models in our future studies using the real data. \begin{figure} \includegraphics[width=80mm]{GNFW_v_BETA.ps} \caption{The SZ flux amplitude versus AMI SA observing baseline for a cluster with $\theta_{\rm 500}=2.5\arcmin$ and $Y_{\rm 500}=2.5 \times 10^{-3}(\rm{arcmin})^{2}$. Blue dashed line represents the fit using the isothermal $\beta$-model and the red represents the fit using GNFW pressure profile.} \end{figure}

1012.4996

Most Sunyaev-Zel'dovich (SZ) and X-ray analyses of galaxy clusters try to constrain the cluster total mass (M<SUB>T</SUB>(r)) and/or gas mass (M<SUB>g</SUB>(r)) using parametrized models derived from both simulations and imaging observations, and assumptions of spherical symmetry and hydrostatic equilibrium. By numerically exploring the probability distributions of the cluster parameters given the simulated interferometric SZ data in the context of Bayesian methods, and assuming a β-model for the electron number density n<SUB>e</SUB>(r) described by two shape parameters β and r<SUB>c</SUB>, we investigate the capability of this model and analysis to return the simulated cluster input quantities via three parametrizations. In parametrization I we assume that the gas temperature is an independent free parameter and assume hydrostatic equilibrium, spherical geometry and an ideal gas equation of state. We find that parametrization I can hardly constrain the cluster parameters and fails to recover the true values of the simulated cluster. In particular it overestimates M<SUB>T</SUB>(r<SUB>200</SUB>) and T<SUB>g</SUB>(r<SUB>200</SUB>) (M<SUB>T</SUB>(r<SUB>200</SUB>) = (6.43 ± 5.43) × 10<SUP>15</SUP> M<SUB>⊙</SUB> and T<SUB>g</SUB>(r<SUB>200</SUB>) = (10.61 ± 5.28) keV) compared to the corresponding values of the simulated cluster (M<SUB>T</SUB>(r<SUB>200</SUB>) = 5.83 × 10<SUP>14</SUP> M<SUB>⊙</SUB> and T<SUB>g</SUB>(r<SUB>200</SUB>) = 5 keV). We then investigate parametrizations II and III in which f<SUB>g</SUB>(r<SUB>200</SUB>) replaces temperature as a main variable; we do this because f<SUB>g</SUB> may vary significantly less from cluster to cluster than temperature. In parametrization II we relate M<SUB>T</SUB>(r<SUB>200</SUB>) and T<SUB>g</SUB> assuming hydrostatic equilibrium. We find that parametrization II can constrain the cluster physical parameters but the temperature estimate is biased low (M<SUB>T</SUB>(r<SUB>200</SUB>) = (6.8 ± 2.1) × 10<SUP>14</SUP> M<SUB>⊙</SUB> and T<SUB>g</SUB>(r<SUB>200</SUB>) = (3.0 ± 1.2) keV). In parametrization III, the virial theorem (plus the assumption that all the kinetic energy of the cluster is the internal energy of the gas) replaces the hydrostatic equilibrium assumption because we consider it more robust both in theory and in practice. We find that parametrization III results in unbiased estimates of the cluster properties (M<SUB>T</SUB>(r<SUB>200</SUB>) = (4.68 ± 1.56) × 10<SUP>14</SUP> M<SUB>⊙</SUB> and T<SUB>g</SUB>(r<SUB>200</SUB>) = (4.3 ± 0.9) keV). We generate a second simulated cluster using a generalized Navarro-Frenk-White pressure profile and analyse it with an entropy-based model to take into account the temperature gradient in our analysis and improve the cluster gas density distribution. This model also constrains the cluster physical parameters and the results show a radial decline in the gas temperature as expected. The mean cluster total mass estimates are also within 1σ from the simulated cluster true values: M<SUB>T</SUB>(r<SUB>200</SUB>) = (5.9 ± 3.4) × 10<SUP>14</SUP> M<SUB>⊙</SUB> and T<SUB>g</SUB>(r<SUB>200</SUB>) = (7.4 ± 2.6) keV using parametrization II, and M<SUB>T</SUB>(r<SUB>200</SUB>) = (8.0 ± 5.6) × 10<SUP>14</SUP> M<SUB>⊙</SUB> and T<SUB>g</SUB>(r<SUB>200</SUB>) = (5.98 ± 2.43) keV using parametrization III. However, we find that for at least interferometric SZ analysis in practice at the present time, there is no differences in the Arcminute Microkelvin Imager (AMI) visibilities between the two models. This may of course change as the instruments improve.

12157934

2010PhRvD..82l3010B

1012

1012.5748_arXiv.txt

Antikaon ($K^-$ meson) condensation in dense baryonic matter formed in heavy ion collisions as well as in neutron stars was first proposed by Kaplan and Nelson \cite{Kap}. There was lots of interest in the study of antikaon condensation in neutron stars \cite{Bro,Tho,Ell,Lee,Pra97,Gle99,Kno,Sch,Pal,Bani1,Bani2,Bani3,Bani4,Pons,Bani5} after their work. A first-order phase transition from hadronic matter to antikaon condensed matter was investigated in several cases using relativistic field theoretical models \cite{Gle99,Bani1,Pons,Bani5}. The phase transition was either studied using Maxwell construction or governed by Gibbs' rules for phase equilibrium along with global baryon number and charge conservation \cite{Gle92}. In those cases the focus was on the equation of state (EoS) and neutron star structure as well as critical temperature of antikaon condensation. A first-order phase transition may proceed through the nucleation of droplets of the new phase. The formation of droplets of exotic matter such as antikaon condensed matter and quark matter, could be possible in neutron stars when the protoneutron star cools down to a temperature of $\sim$ 10 MeV and is deleptonised \cite{Nor,San,Sat,Hei,Bom,Min}. This nucleation process could be due to quantum and thermal nucleation mechanisms \cite{Bom}. The thermal nucleation of antikaon condensed matter was already studied using the homogeneous nucleation theory of Langer \cite{Nor,San,Lan}. Recently it has been shown in the context of nucleation of quark matter in hot and neutrino-free neutron stars that the thermal nucleation is more efficient than the quantum nucleation process at higher temperatures \cite{Bom}. The homogeneous nucleation theory of Langer \cite{Lan,Tur} is applicable close to a first-order phase transition. The hadronic matter becomes metastable near the phase transition point when there is a sudden change in state variables. In this case thermal and quantum fluctuations play important roles in the metastable phase. Small ranged and localised fluctuations in state variables of the metastable hadronic matter might lead to the nucleation of droplets of stable antikaon condensed matter. Those droplets of the stable phase having radii larger than a critical radius will survive and grow. A droplet may grow beyond the critical size if the latent heat is transported from the surface of the droplet into the metastable phase. It was argued that the heat transportation could be achieved through the thermal dissipation and viscous damping \cite{Tur,Las,Raj}. The influence of thermal conductivity and shear viscosity on the thermal nucleation time was studied in a first-order phase transition from hadronic to quark matter \cite{Bom,Las}. However there is no such calculation for the thermal nucleation of antikaon condensed matter. We are motivated to study the effect of shear viscosity on the thermal nucleation rate of droplets of antikaon condensed matter in this work. Shear viscosity of pure neutron and neutron star matter has been calculated by several groups \cite{Flo1,Flo2,Cut,Ben,Yak,Glam}. Recently we have investigated the shear viscosity in antikaon condensed matter \cite{Bani6}. Though we considered a first-order antikaon condensation, we did not take into account the nucleation process. We performed the calculation of shear viscosity in neutron star matter using the EoS derived from relativistic field theoretical models. Now the question is how this shear viscosity of nucleonic phase may impact the nucleation of antikaon condensed phase. The earlier calculation of the nucleation of quark matter adopted a parametrised form of the shear viscosity \cite{Bom}. In this calculation we adopt the same EoS for the computation of thermal nucleation time of the antikaon condensed phase and shear viscosity of nuclear matter. We organise the paper in the following way. We describe models for homogeneous nucleation, shear viscosity and EoS in Sec. II. Results of this calculation are discussed in Sec. III. Sec. IV gives the summary and conclusions.

We have investigated the first-order phase transition from nuclear matter to antikaon condensed matter through the thermal nucleation of a critical droplet of anikaon condensed matter. Our main focus in this calculation is the role of shear viscosity in the prefactor and its consequences on the thermal nucleation rate. We adopt the same EoS derived in the relativistic mean field model for the calculation of shear viscosity and thermal nucleation rate. In this connection, we have constructed critical droplets of antikaon condensed matter above the critical density for different values of surface tension. Droplet radii increase with increasing surface tension. We obtain thermal nucleation time as a function of temperature for a set of values of surface tension and find that the thermal nucleation time is strongly dependent on the surface tension. Our results show that thermal nucleation of a critical antikaon droplet could be possible during the cooling stage of a hot neutron star after neutrinos are emitted for a lower value of surface tension $\sigma <$ 20 MeV fm$^{-2}$. Further, a comparison of our results with that of the $T^4$ approximation shows that the $T^4$ approximation overestimates our results of thermal nucleation time. This comparison highlights the importance of shear viscosity in our calculation. We have considered the thermal nucleation in a nonrotating neutron star and neglected the bulk viscosity in the prefactor. However neutron stars rotate very fast at birth. They emit gravitational waves and become r-mode unstable \cite{Chat}. The bulk viscosity plays an important role in damping the r-mode instability. The bulk viscosity might dominate over the shear viscosity above 10 MeV \cite{Chat}. It would be interesting to investigate the effect of bulk viscosity on the thermal nucleation time in a rotating neutron star along with that of shear viscosity.

1012.5748

We investigate a first-order phase transition from hadronic matter to antikaon condensed matter during the cooling stage of protoneutron stars. The phase transition proceeds through the thermal nucleation of antikaon condensed matter. In this connection we study the effect of shear viscosity on the thermal nucleation rate of droplets of antikaon condensed matter. Here we adopt the same equation of state for the calculation of shear viscosity and thermal nucleation time. We compute the shear viscosity of neutron star matter composed of neutrons, protons, electrons, and muons using the relativistic mean field model. The prefactor in the nucleation rate, which includes the shear viscosity, is enhanced by several orders of magnitude compared with the T<SUP>4</SUP> approximation of earlier calculations. Consequently the thermal nucleation time in the T<SUP>4</SUP> approximation overestimates our result. Further, the thermal nucleation of an antikaon droplet might be possible in our case for surface tension smaller than 20MeVfm<SUP>-2</SUP>.

12207299

2011JCAP...04..017E

1012

1012.2732_arXiv.txt

\label{section:conclude} We have found and remedied a common problem in the literature \cite{toobigtooearly,CayonGordonSilk} that occurs when estimating $\fnl$ from massive high redshift clusters. We found that the value of $\fnl$ necessary to make the existence of these massive clusters certain (i.e. to give 100\% confidence that they would be detected in the given surveys) is $\fnl \gtrsim 1000$, which is significantly larger than what has been claimed (e.g.~$\fnl > 550$ in \cite{toobigtooearly}). We have also demonstrated that this discrepancy is explained by the fact that \cite{toobigtooearly} and \cite{CayonGordonSilk} use a non-Gaussian mass function, derived from a Gaussian mass function that is not valid in the whole range of $\ln\sigma^{-1}$ needed for the computation. Most importantly the Gaussian mass function does not scale correctly outside of this range. We have rectified this oversight by using a Gaussian mass function by Tinker et al.~\cite{Tinkeretal} that has the proper asymptotic behaviour. We also found that below a certain critical $\fnl$ value, the results in both refs.~\cite{toobigtooearly,CayonGordonSilk} can be trusted with confidence. However, this value differs for each references and depends on an arbitrary cutoff of the mass integral in equation (\ref{eq:defExp}). Our 95\% confidence lower bound of $\fnl\gtrsim410$ agrees reasonably well with reference \cite{toobigtooearly} because it occurs below this critical $\fnl$ value for their cutoff. We then considered the case where the non-Gaussian part of the primordial spectrum is dominated by $\gnl$. Such a situation can easily arise e.g.~in curvaton models \cite{Curvaton1, Curvaton2, Curvaton3}. We estimated $\gnl > 2.0 \times 10^6$. We thus demonstrated, that within current observational limits, $\gnl$ appears to have more potential to explain the observed excess of high-redshift massive clusters. Non-Gaussianity is not the only potential explanation for the tension caused by the existence of these high redshift clusters. A systematic error that consistently over-estimated the masses would have a similar effect. In \cite{toobigtooearly} this possibility was considered. To minimise the effects of systematic errors, the mass measurements used were always taken to be those whose quoted errors were consistent with the smallest mass value. We use the same mass estimates. For many clusters, this involved comparing mass estimates from SZ effect measurements, X-ray measurements and weak lensing measurements. Any systematic errors would need to be present in all three measurement methods. It was also found in \cite{toobigtooearly} that all the masses would need to be systematically over-estimated by 1.5 $\sigma$ in \emph{each} measurement technique to make this ensemble of clusters fully consistent with $\fnl=0$. A different expansion history is another potential explanation for the existence of these clusters \cite{Alam:2010tt,Bhattacharya:2010wy,Xia:2009ys}. A modified equation of state for dark energy that resulted in an earlier onset of the accelerated expansion would suppress structure growth at smaller redshifts/later times. This would have the effect of causing the linear growth function $D(z)$ to drop more slowly from low redshifts to high redshifts. The net result is more structure and thus more, large mass, clusters at high redshifts than what would be expected in $\Lambda$CDM. The estimates of $\fnl$ quoted here and in \cite{toobigtooearly} and \cite{CayonGordonSilk} are far outside the observational limits set by WMAP. In \cite{toobigtooearly} and \cite{CayonGordonSilk} they suggested remedying this problem by introducing running of $\fnl$. While this would explain the apparent discrepancy of the magnitude of $\fnl$ over different scales, it is also likely to introduce problems with the actual spectral index of the perturbations, since usually if $\fnl$ acquires running, so does the magnitude of the perturbations. (For discussions on running non-linearity parameters, see \cite{Scale1, Scale2, Scale3, Scale4, Scale5}.) We have demonstrated that instead of introducing non-zero $n_{\fnl}$, the abundance of the massive clusters can be explained by introducing $\gnl$ almost within the current observational bounds. We also like to draw attention to a potentially important observational result: there appears to be an overabundance of large voids (see \cite{Voids1,Voids2} and references therein), and while large $\gnl$ makes both heavy clusters and large voids more probable, a positive $\fnl$ actually makes the voids \emph{less} likely, increasing the tension with observations even more \cite{D'Amico:2010kh}. In general it seems that a large value of $\gnl$ is in much better agreement with all observations (CMB, clusters, halo bias, voids), than large values of $\fnl$. With better measurements and N-body simulations a slightly smaller value of $\gnl$ could perhaps accommodate the abundance of the heavy clusters. This stresses the importance of extracting limits from the Planck microwave temperature map not only on $\fnl$ but also on $\gnl$.

1012.2732

There are observations of at least 14 high-redshift massive galaxy clusters, which have an extremely small probability with a purely Gaussian initial curvature perturbation. Here we revisit the estimation of the contribution of non-Gaussianities to the cluster mass function and point out serious problems that have resulted from the application of the mass function out of the range of its validity. We remedy the situation and show that the values of f<SUB>NL</SUB> previously claimed to completely reconcile (i.e. at ~ 100% confidence) the existence of the clusters with ΛCDM are unphysically small. However, for WMAP cosmology and at 95% confidence, we arrive at the limit f<SUB>NL</SUB>gtrsim411, which is similar to previous estimates. We also explore the possibility of a large g<SUB>NL</SUB> as the reason for the observed excess of the massive galaxy clusters. This scenario, g<SUB>NL</SUB> &gt; 2 × 10<SUP>6</SUP>, appears to be in more agreement with CMB and LSS limits for the non-Gaussianity parameters and could also provide an explanation for the overabundance of large voids in the early universe.

12163136

2011A&A...527A.122H

1012

1012.5324_arXiv.txt

\label{introduction} The bulk of interstellar matter is found in regions of low-to-moderate opacity to UV and visible light where stellar radiation governs the chemical and thermal state of the gas. These photodissociation or photon-dominated regions \cite[PDRs, for a review see][]{hollenbach99} are responsible for reprocessing much of the energy output from stars, reemitting this energy in the infrared-millimeter wavelengths, including a rich mixture of gas lines (e.g., emission in fine-structure, rotational, and rovibrational lines). They are privileged objects for studying the chemical and physical processes of the interstellar medium (ISM). The motivation of this study is to test our understanding of the excitation of H$_2$ in a regime of space parameters that has been poorly studied. Through its ability to observe pure rotational lines of H$_2$ towards a number of bright PDRs, the Infrared Space Observatory (ISO) has given important information on the local microphysics of interstellar gas. One conclusion of these studies is that the H$_2$ line intensities and the gas temperature, derived from the first rotational levels of H$_2$, are higher than model calculations \cite[]{bertoldi97,draine99a,thi99,kemper99,li2002,habart2003a}. Changes in the description of heating and cooling processes or of the H$_2$ formation rate have been proposed to account for this discrepancy \cite[]{habart2004}. With H$_2$ UV absorption line observations performed by the Far Ultraviolet Spectroscopic Explorer (FUSE), the H$_2$ formation and excitation processes in diffuse and translucent clouds have also been reconsidered. In particular, UV data show significant amounts of rotationally excited H$_2$, for which UV photons could not be the unique heating source \cite[]{gry2002,nehme2008}. The H$_2$ excitation in the diffuse interstellar medium may be tracing the dissipation of interstellar turbulence in C-shocks \cite[]{flower98} or interstellar vortices \cite[]{falgarone2005}. There has been very little exploration of the physics of PDRs with moderate FUV fields and densities, an intermediate regime between the diffuse clouds and the bright PDRs. The sensitivity of the infrared spectrograph (IRS) onboard the Spitzer Space Telescope provides a unique opportunity to observe H$_2$ pure rotational line emission in low-brightness sources. As part of the SPECPDR\footnote{See http://www.cesr.fr/$\sim$joblin/SPECPDR public/SPECPDR.html} program dedicated to the study of PDRs with Spitzer, we used this capability to study the H$_2$ line emission in regions with $\chi \sim 5-10^3$ times the Solar Neighborhood Far-UV interstellar radiation field \footnote{For a discussion on the definitions of the mean interstellar radiation field used in the literature on PDRs see Appendix B of \cite{allen2004} or Appendix C of \cite{lepetit2006}.} as given by \cite{draine78}. Our sample consists of well-known PDRs, so that they are good test cases for models. These observations are analyzed in combination with previous ISO and ground-based data of rotational and rovibrational H$_2$ lines when available, in order to help constrain the relative roles of ultraviolet pumping and collisions in establishing the level populations. The paper is organized as follows. In Sect. \ref{sample}, we present our PDR sample. In Sect. \ref{observation}, we present the IRS spectrometer observations. In Sect. \ref{model_presentation}, we briefly describe the PDR model used to analyze the data, and in Sect. \ref{comparison_observation_model} we compare the observed H$_2$ line emission to the PDR model predictions. In Sect. \ref{h2_diffus}, we compare our results with UV absorption measurements in diffuse clouds. In Sect. \ref{origin}, we discuss several possibilities of explaining our results. In Sect. \ref{galaxies}, in light of our PDR observations, we discuss H$_2$ line infrared emission measurements in galaxies and along Galactic lines of sight. Our conclusions are summarized in Sect. \ref{conclusion}.

\label{conclusion} Thanks to Spitzer spectroscopic observations, we were able to detect H$_2$ pure rotational lines emission in PDRs with modest FUV fields and densities. The low/moderate $\chi$ regions studied here are very widespread in galaxies and may contain a large fraction of the molecular gas. However, this intermediate regime between the diffuse cloud and the bright PDRs has been poorly studied. To analyze the H$_2$ line emission observations, we used an updated version of the Meudon PDR code. Our results allow strong constraints to be placed upon the H$_2$ excitation in the interstellar medium of galaxies. The main results from this work can be summarized as follows. The IRS wavelength coverage allows us to detect several strong H$_2$ pure rotational lines from 0-0 S(0) to S(3) at 28.2, 17.03, 12.29, and 9.66 $\mu$m, the aromatic band features at 6.2, 7.7, 8.6, 11.3 $\mu$m, the dust mid-IR continuum emission, and the fine structure lines of ionized gas [NeII] at 12.8 $\mu$m, [SIII] at 18.7 and 33.4 $\mu$m and [SiII] at 34.9 $\mu$m. The observed mid-IR spectra are typical of PDRs. A single temperature cannot describe the full set of observed H$_2$ line intensities. A combination of at least two H$_2$ gas components, with one cool/warm ($\sim$100-300K) and another warm/hot ($\sim$300-700K) with much lower column densities (a few percent of the first component) is required. The ortho-to-para ratios derived are about $\sim$1 for the first component and about $\sim$3 for the second warmer component. The non-equilibrium behavior has already been noted in previous ISO observations of PDRs \cite[]{moutou99,fuente99,habart2003a}. By comparing the observations with the PDR model predictions, we find that the model can account for the first low H$_2$ rotational line (e.g., 0-0 S(0)-S(1)) probing the bulk of the gas at moderate temperature, as well as the ro-vibrational line (e.g., 1-0 S(1) observed with ground-based telescopes) probing the UV-pumped gas. In contrast, the model underestimates the excited rotational lines (e.g., 0-0 S(2)-S(3)) by large factors ($\ge$10 in some objects). In the lowest excited PDR, the discrepancy between the model and the data starts even from the rotational $J=3$ level (e.g., 0-0 S(1)). This discrepancy between observations and the PDR model predictions has the same order of magnitude as reported for diffuse interstellar clouds. However, in PDRs the power radiated per H atom in the rotational excited levels is more than one order of magnitude larger than in diffuse clouds. The discrepancy between the data and the PDR model could indicate that, in the models, the column density of H$_2$ is too low in the outer zones where H$_2$ is excited (by inelastic collisions or UV pumping), or, alternatively, that a small fraction of the H$_2$ gas is hotter than predicted by models. We discuss whether an enhancement in the H$_2$ formation rate, or a local increase in photoelectric heating, as proposed for brighter PDRs in former ISO studies, may also work in low-excitation PDRs. An enhancement in the H$_2$ formation rate improves the comparison with observations, but the models still fall short. Further work is needed to quantify the impact of the evolution of very small dust particles across PDRs on the gas energetics. Out-of-equilibrium processes or mechanical heating (by weak shocks or dissipation of turbulence as proposed for diffuse interstellar clouds) are alternative possibilities. % Although we cannot decide at this point, we emphasize the need for further development of the models. Progress is also expected from PDR spectroscopy of additional cooling lines and data on the gas kinematics, to be obtained from the ground and Herschel. In particular, observations of gas cooling lines of species such as C$^+$ at high spectral resolution with the Heterodyne Instrument for the Far Infrared (HIFI) will provide missing information on the gas velocity within the PDR layer where CO is photodissociated. The combination of these data with models % should help constrain the relative radiative and dynamical influence of stars on the physical conditions within PDRs. In the longer term, the Mid-InfraRed Instrument (MIRI) on James Webb Space Telescope (5-28 $\mu$m, diffraction-limited resolution of $\sim$0.2'' at 5 $\mu$m) will provide maps of the H$_2$ rotational line emission, as well as of the small dust emission with unprecedented angular resolution. \begin{appendix}

1012.5324

<BR /> Aims: We present spectroscopic observations obtained with the infrared Spitzer Space Telescope, which provide insight into the H<SUB>2</SUB> physics and gas energetics in photodissociation regions (PDRs) of low to moderate far-ultraviolet (FUV) fields and densities. <BR /> Methods: We analyze data on six well known Galactic PDRs (L1721, California, N7023E, Horsehead, rho Oph, N2023N), sampling a poorly explored range of excitation conditions (χ ~ 5-10<SUP>3</SUP>), relevant to the bulk of molecular clouds in galaxies. Spitzer observations of H<SUB>2</SUB> rotational lines are complemented with H<SUB>2</SUB> data, including ro-vibrational line measurements, obtained with ground-based telescopes and ISO, to constrain the relative contributions of ultraviolet pumping and collisions to the H<SUB>2</SUB> excitation. The data analysis is supported by model calculations with the Meudon PDR code. <BR /> Results: The observed column densities of rotationally excited H<SUB>2</SUB> are observed to be much higher than PDR model predictions. In the lowest excitation PDRs, the discrepancy between the model and the data is about one order of magnitude for rotational levels J ≥ 3. We discuss whether an enhancement in the H<SUB>2</SUB> formation rate or a local increase in photoelectric heating, as proposed for brighter PDRs in former ISO studies, may improve the data-model comparison. We find that an enhancement in the H<SUB>2</SUB> formation rates reduces the discrepancy, but the models still fall short of the data. <BR /> Conclusions: This large disagreement suggests that our understanding of the formation and excitation of H<SUB>2</SUB> and/or of PDRs energetics is still incomplete. We discuss several explanations, which could be further tested using the Herschel Space Telescope. <P />Appendix is only available in electronic form at <A href="http://www.aanda.org">http://www.aanda.org</A>

12132696

2010arXiv1012.5438L

1012

1012.5438_arXiv.txt

The relativistic contribution to the rate of precession of perihelion of Mercury is calculated accurately using general relativity \cite{einstein0,einstein1,einstein2,einstein3,schwarzschild,droste}. However, the problem is commonly discussed in undergraduate and graduate classical mechanics textbooks, without introduction of an entirely new, metric theory of gravity. One approach \cite{goldstein2,jose,peters,phipps1,phipps2} is to define a Lagrangian that is consistent with both Newtonian gravity and the momentum-velocity relation of special relativity. The resulting equation of motion is solved perturbatively, and an approximate rate of precession of perihelion of Mercury is extracted. This approach is satisfying in that all steps are proved, and a brief introduction to special relativity is included. On the other hand, one must be content with an approximate rate of precession that is about one-sixth the correct value. Another approach \cite{goldstein,TM1,barger,fowles,hand} is that of a mathematical exercise and history lesson. A modification to Newtonian gravity is given, without proof, resulting in an equation of motion that is the same as that derived from general relativity. The equation of motion is then solved using appropriate approximations, and the correct rate of precession of perihelion of Mercury is extracted. Both approaches provide an opportunity for students of classical mechanics to understand that relativity is responsible for a small contribution to perihelic precession and to calculate that contribution. We present a review of the approach using only special relativity, followed by an alternative solution of the equation of motion derived from Lagrange's equations. An approximate rate of perihelic precession is derived that agrees with established calculations. This effect arises as one of several small corrections to Kepler's orbits, including reduced radius of circular orbit and increased eccentricity. The method of solution makes use of coordinate transformations and the correspondence principle, rather than the standard perturbative techniques, and is approachable by undergraduate physics majors. A relativistic particle of mass $m$ orbiting a central mass $M$ is commonly described by the Lagrangian \cite{goldstein2,jose,peters,phipps1,phipps2,TM2,barger2,potgieter,DesEri,HuangLin,SonMas} \begin{gather} L = -mc^2\gamma^{-1} - U(r), \label{eq_lagrangian} \end{gather} where: $\gamma^{-1} \equiv \sqrt{1 - v^2/c^2}$; $v^2 = \dot{r}^2 + r^2 \dot{\theta}^2$; and $U(r) = -GMm/r$. ($G$ is Newton's universal gravitational constant, and $c$ is the speed of light in vacuum.) The equations of motion follow from Lagrange's equations, \begin{equation} \frac{\rmd}{\rmd t} \frac{\partial L}{\partial\dot{q}_i} - \frac{\partial L}{\partial q_i} = 0, \end{equation} for each of $\{ q_i \} = \{ \theta,r \}$, where $\dot{q_i} \equiv \rmd q_i/\rmd t$. The results are: \begin{equation} \dfrac{\rmd}{\rmd t} [\gamma r^2 \dot{\theta}] = 0, \label{eq_AngMom} \end{equation} which implies that $\ell \equiv \gamma r^2\dot\theta = \mbox{constant}$; and \begin{equation} \gamma \ddot{r} + \dot{\gamma} \dot{r} + \frac{GM}{r^2} - \gamma r \dot{\theta}^2 = 0. \label{eq_EOM} \end{equation} The first of these [Eq.~(\ref{eq_AngMom})] is the relativistic analogue to the Newtonian equation for the conservation of angular momentum per unit mass, and is used to eliminate $\dot{\theta}$ in Eq.~(\ref{eq_EOM}), \begin{equation} \gamma r \dot{\theta}^2 = \frac{\ell^2}{\gamma r^3}. \label{eq_part1} \end{equation} Time is eliminated by successive applications of the chain rule, together with the conserved angular momentum \cite{goldstein4,jose2,TM4,fowles2,hamill}: \begin{equation} \dot{r} = - \frac{\ell}{\gamma} \frac{\mathrm{d}}{\mathrm{d} \theta} \frac{1}{r}; \label{eq_r_dot} \end{equation} and, therefore, \begin{equation} \gamma \ddot{r} = - \dot{\gamma} \dot{r} - \frac{\ell^2}{\gamma r^2} \frac{\mathrm{d^2}}{\mathrm{d} \theta^2} \frac{1}{r}. \label{eq_part2} \end{equation} Substituting Eqs.~(\ref{eq_part1})~and~(\ref{eq_part2}) into the equation of motion Eq.~(\ref{eq_EOM}) results in \begin{equation} \ell^2\frac{\mathrm{d}^2}{\mathrm{d} \theta^2} \frac{1}{r} - GM\gamma + \frac{\ell^2}{r} = 0. \label{eq_EOM2} \end{equation} We anticipate a solution of Eq.~(\ref{eq_EOM2}) that is near Keplerian and introduce the radius of a circular orbit for a nonrelativistic particle with the same angular momentum, $r_\text{c} \equiv \ell^2/GM $. The result is \begin{equation} \frac{\mathrm{d}^2}{\mathrm{d} \theta^2} \frac{r_\text{c}}{r} + \frac{r_\text{c}}{r} = 1 + \lambda, \label{eq_SR} \end{equation} where $\lambda\equiv \gamma-1$ is a velocity-dependent correction to Newtonian orbits due to special relativity. The conic-sections of Newtonian mechanics \cite{goldstein3,jose3,TM3,fowles3,smiths,hamill2} are recovered by setting $\lambda = 0$ $(c \rightarrow \infty)$: \begin{equation} \frac{\mathrm{d}^2}{\mathrm{d} \theta^2} \frac{r_\text{c}}{r} + \frac{r_\text{c}}{r} = 1, \label{eq_Newton0} \end{equation} which implies that \begin{equation} \frac{r_\text{c}}{r} = 1 + e\cos{\theta}, \label{eq_Newton} \end{equation} where $e$ is the eccentricity.

The present approach to incorporating special relativity into the Kepler problem results in an approximate orbit equation [Eq.~(\ref{eq_class_rel})] that has the same form as that derived from general relativity in this limit [Eq.~(\ref{eq_gen_rel})] and is easily compared to that describing Kepler's orbits [Eq.~(\ref{eq_Newton0})]. This orbit equation clearly describes three corrections to a Keplerian orbit due to special relativity: precession of perihelion; reduced radius of circular orbit; and increased eccentricity. The predicted rate of precession of perihelion of Mercury is identical to established calculations using only special relativity. Each of these corrections is exactly one-sixth of the corresponding correction described by general relativity in the Keplerian limit. This derivation of an approximate orbit equation is complementary to existing calculations of the rate of precession of perihelion of Mercury using only special relativity. The central-mass problem is described by a Lagrangian that is consistent with both Newtonian gravity and the momentum-velocity relation of special relativity. However, coordinate transformations and the correspondence principle are used to solve the equations of motion resulting from Lagrange's equations, rather than the standard perturbative methods. The resulting closed-form, approximate orbit equation exhibits several characteristics of relativistic orbits at once, but is limited to describing small relativistic corrections to approximately Newtonian, near-circular orbits. This orbit equation, derived using only special relativity, provides a qualitative description of corrections to Keplerian orbits due to general relativity. Exact solutions to the special relativistic Kepler problem require a thorough understanding of special relativistic mechanics \cite{rindler2,synge} and are, therefore, inaccessible to most undergraduate physics majors. The present approach and method of solution is understandable to nonspecialists, including undergraduate physics majors whom have not had a course dedicated to relativity. \appendix

1012.5438

Kepler's orbits with corrections due to Special Relativity are explored using the Lagrangian formalism. A very simple model includes only relativistic kinetic energy by defining a Lagrangian that is consistent with both the relativistic momentum of Special Relativity and Newtonian gravity. The corresponding equations of motion are solved in a Keplerian limit, resulting in an approximate relativistic orbit equation that has the same form as that derived from General Relativity in the same limit and clearly describes three characteristics of relativistic Keplerian orbits: precession of perihelion; reduced radius of circular orbit; and increased eccentricity. The prediction for the rate of precession of perihelion is in agreement with established calculations using only Special Relativity. All three characteristics are qualitatively correct, though suppressed when compared to more accurate general-relativistic calculations. This model is improved upon by including relativistic gravitational potential energy. The resulting approximate relativistic orbit equation has the same form and symmetry as that derived using the very simple model, and more accurately describes characteristics of relativistic orbits. For example, the prediction for the rate of precession of perihelion of Mercury is one-third that derived from General Relativity. These Lagrangian formulations of the special-relativistic Kepler problem are equivalent to the familiar vector calculus formulations. In this Keplerian limit, these models are supposed to be physical based on the likeness of the equations of motion to those derived using General Relativity. The derivation of this orbit equation is approachable by undergraduate physics majors and nonspecialists whom have not had a course dedicated to relativity.

12167462

2011ApJ...727L..37T

1012

1012.5112_arXiv.txt

Intensity oscillations with a period of three to thirty minutes have been frequently observed in polar plumes \citep[e.g.,][]{Ofman1997,DeForest1998,Ofman1999,Banerjee2010} and active region loops \citep[e.g.,][]{Berghmans1999,DeMoortel2000,DeMoortel2002,Robbrecht2001,Marsh2003,King2003,McEwan2006,Marsh2009,Stenborg2011}. In spectroscopic studies, small-amplitude oscillations (usually a few percent of the background emission) have been found in line intensities \citep{Banerjee2009} that are often accompanied by small fluctuations (at most a couple of km/s) in the Doppler velocities \citep{Wang2009a,Wang2009b,Kitagawa2010,Mariska2010}. These quasi-periodic disturbances usually show propagating speeds of 50-200~km/s and are almost interpreted as slow-mode magneto-acoustic waves propagating into the corona along the magnetic field without exception. Recently, both imaging and spectroscopic observations have revealed that upflows with velocities of 50-150~km/s are prevalent in active regions \citep{Sakao2007,Hara2008,DePontieu2009,McIntosh2009a,He2010,Guo2010,Peter2010,Bryans2010}, quiet Sun \citep{McIntosh2009b} and coronal holes \citep{DePontieu2009,McIntosh2010a,McIntosh2010b}. These upflows appear as weak upward propagating disturbances in coronal images, and in spectroscopic observations they are identified as significant blue-wing asymmetries in emission profiles of spectral lines formed at transition-region and coronal temperatures. These faint upflows, with a life time of 50-150 seconds, are believed to be associated with type-II spicules or rapid blue-shifted events observed in the chromosphere \citep{DePontieu2009,Rouppe2009}. They are suggested to provide hot plasmas into the corona and may thus play an important role in coronal heating process \citep{DePontieu2009,McIntosh2009b,Peter2010,DePontieu2010,Hansteen2010}. These rapid upflows often recur at the same location on time scales of three to fifteen minutes \citep[e.g.,][]{McIntosh2009a, McIntosh2009b}, and would naturally cause quasi-periodic low-contrast oscillations in coronal images \citep[also see][]{Xia2005}. Thus, the discovery of these rapid quasi-periodic upflows challenges the universal wave interpretation of coronal oscillations. \cite{DePontieu2010} analyzed time series data that were previously studied by \cite{Wang2009b} to illustrate the presence of slow-mode waves, and found that the intensity and Doppler shift oscillations are {\em also} accompanied by oscillations in the line width and excess emission of the blue wing. They concluded that while the ``flows vs waves'' picture was not unambiguously resolved, the presence of these multi-moment in-phase oscillatory signatures is consistent with propagating quasi-periodic upflows causing the observed signature. In this paper, we present new results derived from a timeseries data set obtained by the EUV Imaging Spectrometer \citep[EIS,][]{Culhane2007} onboard Hinode. We show that coherent oscillatory behaviors in intensity, Doppler shift, line width, and blueward asymmetry are clearly present almost everywhere at the root of fan-like structures in the boundary of an active region. We demonstrate that these oscillation signatures are caused by repetitive high-speed upflows, which can be directly identified in the image sequence simultaneously obtained by the X-Ray Telescope \citep[XRT,][]{Golub2007} onboard Hinode.

Applying a single Gaussian fit to the line profiles, it has been found that coronal emission lines usually show blue shift of the order of 30~km/s at boundaries of some ARs \citep{Marsch2004,Marsch2008,Harra2008,DelZanna2008,Doschek2008,Murray2010}. These blue shifts were thought to be genesis of the slow solar wind \citep{Sakao2007,Harra2008,Doschek2008}. However, from Figure~\ref{fig.2} we can see that the Fe~{\sc{xii}}~195.12\AA{} and Fe~{\sc{xiii}}~202.04\AA{} profiles at the root of the fan-like structure (y=-83$^{\prime\prime}$$\sim$-43$^{\prime\prime}$) in the AR boundary are actually very asymmetric with prominent excess emission in blue wings, confirming the results of \cite{McIntosh2009a}. Although the Fe~{\sc{xii}}~195.12\AA{} line is potentially blended, we place great confidence in its strong blue wing asymmetry here since the potential blends are sitting in the red wing of the profile. Such a result suggests the presence of continuous fast-moving upflows (around 100~km/s) and that the emission consists of multiple components. The centroid of the line profile derived by using a single Gaussian fit only reflects the ensemble velocity of the multiple emission components. And different components may have different velocities, which naturally broadens the profile. The lack of blue asymmetry above y=-43$^{\prime\prime}$ seems to suggest that rapid upflows occur mainly at the root of the fan-like structures. Perhaps the larger field inclination with height along the loops (thus larger angle between the flow direction and the line of sight) makes it difficult to resolve the upflow signatures. While the profiles of Fe~{\sc{xiii}}~202.04\AA{} are noisy there, the Fe~{\sc{xii}}~195.12\AA{} line profiles reveal clear prominent red wing asymmetries, which might be caused by the blends of the line (see above). The lower part of the slit (y$\leq$-83$^{\prime\prime}$) is dominated by loop structures in the AR core, where weak red wing asymmetries are present in profiles of both lines. It might be related to the complexity of the emission and magnetic structures. The asymmetries are more prominent for the Fe~{\sc{xii}}~195.12\AA{} line, which should be related to the blends. Examining the XRT movie associated with Fig.~\pref{fig.1}, we can see plasma moving outward rapidly along loop structures in the fan. These continuous upflows have previously been identified by \cite{Sakao2007} having an average speed of $\sim$100~km/s. We believe that these rapid upflows are responsible for the blueward asymmetry in coronal emission line profiles, consistent with the results of \citet{McIntosh2009a}. An isolated strong upflow event (indicated by the arrow in Figure~\ref{fig.1} and Figure~\ref{fig.2}) was clearly visible in both the imaging and spectroscopic observations. As the rapidly moving plasma crossed the slit at around 19:40, we immediately observed a significant enhancement of the emission in the blue wing of the emission line. This enhanced blue wing asymmetry resulted in an enhancement in the line intensity, Doppler shift, and non-thermal width derived by a single Gaussian fit. This we, believe is an isolated example of the process occurring frequently at the roots of the fan structure and is highly unlikely to be the result of wave passage. Although the cadence of the XRT observation is much lower than that of EIS, and the plasma sampled by the XRT Ti-Poly filter has a higher temperature than the formation temperature of Fe~{\sc{xii}}, from Figure~\ref{fig.3} we see a good correspondence in the temporal evolutions of the XRT and coronal line intensities. Figure~\ref{fig.3} shows that the evolutionary patterns of the de-trended intensity, Doppler shift, and non-thermal width are highly similar (the correlation coefficient between each pair of line parameters is around 0.6). The R-B pattern, to some extent, is also similar to those of the other four. In Figure~\ref{fig.4} we find that patches of large non-thermal width (contours) often coincide with those of large blue asymmetry (darker colors). The correlation coefficient between non-thermal width and R-B is 0.40 for Fe~{\sc{xii}}~195.12\AA{} and 0.29 for Fe~{\sc{xiii}}~202.04\AA{}. The relatively low value for the latter might be partly caused by the much lower signal to noise ratio of the Fe~{\sc{xiii}}~202.04\AA{} line profiles. The coherent behaviors revealed in Figure~\ref{fig.3} and Figure~\ref{fig.4} strongly suggest that continuous upflows with quasi-periodic enhancement of the flow intensity are responsible for the quasi-periodic enhancement of the line intensity, Doppler shift, and non-thermal width determined from a single Gaussian fit. We note that the correlations at several instances are not obvious or even not present. This is likely to be caused by the poor spectral resolution and high photon noise of the EIS instrument \citep{DePontieu2010}. As an example, Figure~\ref{fig.5} shows the timeseries for Fe~{\sc{xii}}~195.12\AA{} and Fe~{\sc{xiii}}~202.04\AA{} at y=-54$^{\prime\prime}$$\sim$-50$^{\prime\prime}$. Correlated changes in intensity, Doppler shift, non-thermal line width, and R-B are clearly present. It is also clear that the intensity ratio between the secondary and primary Gaussian components generally varies with R-B, suggesting that the fast-moving plasma is resolved by our guided double Gaussian fit. Note that the less-than-ideal correlations at some instances are actually the result of the poor spectral resolution and photon noise of the EIS instrument \citep{DePontieu2010}. The relative velocity of the secondary component is rather stable at $\sim$ 100~km/s, except for several instants when the R-B values are relatively small. As we mentioned in the introduction, quasi-periodic intensity oscillations have been almost universally interpreted as waves. However, from Figures~\ref{fig.3}-\ref{fig.5}, it is clear that repetitive flows can also produce oscillatory signatures, as demonstrated previously by \cite{DePontieu2010}. Recently, \cite{Verwichte2010} presented a slow wave model to argue that the wave interpretation is still valid for the observed quasi-periodic intensity perturbations. However, the Figure~3 of their paper shows a frequency-doubling of the line width oscillation compared to the intensity and Doppler shift oscillations, which is {\em not} observed by EIS in our observation. Instead, all of these single Gaussian parameters show a reasonable correlation over several hours. Moreover, in the wave scenario the R-B values often can be positive and negative over an oscillation period \--- this is also {\em not} the case in our observation. From our Figures~\ref{fig.2} and~\ref{fig.5} we can see that in the fan root region the R-B values almost remain the same sign in the entire 175-minute observation period, indicating the presence of continuous blueward emission from upflows. In conclusion, we find that coherent oscillatory behaviors in intensity, Doppler shift, line width, and blueward asymmetry are clearly present almost everywhere at the root of the fan-like structures in the boundary of an active region. With coordinated imaging observation, we conclude that the quasi-periodicities we observed are more likely to be caused by quasi-periodic high-speed upflows. There is no doubt that both waves and flows are present on the Sun. We emphasize that it is difficult to distinguish between upflows and waves only through the intensity evolution. Spectroscopic observations reveal more information, and a combination of imaging and spectroscopic observations is critical for the correct interpretation. So far we have found two observations where flows seem to be a better interpretation for the quasi-periodicity. With more detailed spectroscopic observations we are sure that we can find more evidences of flows.

1012.5112

Quasi-periodic propagating disturbances are frequently observed in coronal intensity image sequences. These disturbances have historically been interpreted as being the signature of slow-mode magnetoacoustic waves propagating into the corona. The detailed analysis of Hinode EUV Imaging Spectrometer (EIS) timeseries observations of an active region (known to contain propagating disturbances) shows strongly correlated, quasi-periodic, oscillations in intensity, Doppler shift, and line width. No frequency doubling is visible in the latter. The enhancements in the moments of the line profile are generally accompanied by a faint, quasi-periodically occurring, excess emission at ~100 km s<SUP>-1</SUP> in the blue wing of coronal emission lines. The correspondence of quasi-periodic excess wing emission and the moments of the line profile indicates that repetitive high-velocity upflows are responsible for the oscillatory behavior observed. Furthermore, we show that the same quasi-periodic upflows can be directly identified in a simultaneous image sequence obtained by the Hinode X-Ray Telescope. These results are consistent with the recent assertion of De Pontieu &amp; McIntosh that the wave interpretation of the data is not unique. Indeed, given that several instances are seen to propagate along the direction of the EIS slit that clearly shows in-phase, quasi-periodic variations of intensity, velocity, width (without frequency doubling), and blue wing enhanced emission, this data set would appear to provide a compelling example that upflows are more likely to be the main cause of the quasi-periodicities observed here, as such correspondences are hard to reconcile in the wave paradigm.

12201896

2011ASPC..448..391B

1012

1012.1883_arXiv.txt

Initially dismissed as potential habitats, planets orbiting M dwarfs have lately seen renewed interest \citep{Tarter2007,Scalo2007}. With lower luminosities, their ``habitable zones'' (HZ), the range of orbits in which an Earth-like planet could support surface water \citep{Kasting1993}, are significantly closer-in than for solar-type stars. This proximity leads to new perils such as increased susceptibility to stellar activity and stronger tidal effects. The discovery of exoplanets spurred more detailed and careful evaluation of these phenomena, and many researchers now argue that they are not as dangerous for life as previously feared. On the contrary, these planets may be ideal laboratories to test models of geophysics, atmospheric dynamics, celestial mechanics, photochemistry, aeronomy, and ultimately habitability. Without analogs in our Solar System or adequate remote sensing capabilities, the surface properties of M dwarf planets can only be considered theoretically. Nonetheless, progress has been made in several areas, such as modeling of planetary interiors \citep{Sotin2007,ONeillLenardic2007}, atmospheric mass loss \citep{Yelle2004,Segura2010}, atmospheric dynamics \citep{Joshi2003,HengVogt2010}, and tidal effects \citep{Jackson2008c,Heller_sub}. This chapter is a multidisciplinary study of the potential habitability of M star planets, but with an astrophysical bias. A full treatment would require far more space than permitted by this format. For a more comprehensive analysis of M dwarf planet habitability (and habitability in general), see \citet{Tarter2007}, \citet{Scalo2007}, and $\S$ 7. This chapter is organized as follows. First we examine the current detection limits of terrestrial planets due to stellar variability. Second, we explore the different chemical reactions in planetary atmospheres due to different stellar spectral energy distributions. Next we examine tidal effects. Fourth, we consider the possible existence of ``habitable evaporated cores'' of giant planets. Fifth, we explore the magnetic fields of terrestrial planets. Finally, we report key interdisciplinary findings from a recent NASA Astrobiology Institute workshop titled ``Revisiting the Habitable Zone.''

Less than one month after the Cool Stars XVI meeting, \citet{Vogt2010} reported the RV detection of a potentially rocky planet ($m_p \ge 3 \textrm{M}_\oplus$) orbiting in the HZ of the M3 star Gl 581 in a nearly circular orbit. If confirmed, this planet is the first discovered near the middle of the HZ of a main sequence star, and, as expected, that star is an M dwarf. So how does this planet measure up in terms of potential habitability? Vogt et al. note that the host star is extremely quiescent, to the point that they cannot really detect any jitter, and hence stellar activity is not currently an issue for the planet. However, the star is at least a few Gyr old \citep{Bonfils2005}, hence we cannot exclude the possibility that fatal harm was done to a potential biosphere in the past. The planet is tidally locked, or in a spin-orbit resonance \citep{Heller_sub}, but we do not know its rotation rate. The planet's mass is large enough that it can sustain tectonic activity for 10 Gyr, assuming it formed with a similar ratio of radiogenic isotopes as the Earth (tidal heating is minimal, even for $e=0.2$). Therefore, from the RV data, we can not discern any major issue impeding habitability. Unfortunately, though, remote detection of its biosphere is not possible for the foreseeable future, as at 0.15 AU from its host star, reflection spectra will be unavailable from any currently planned space mission. In spite of early skepticism, planets orbiting M dwarfs can be inhabited. Although many issues have been identified, such as tidal locking and atmospheric removal, more careful modeling has shown that these phenomena are not so deadly. While improvements in our knowledge via modeling will continue, transmission spectra of planets transiting M dwarfs will provide, at least for the next decade, our only observational means to directly assess the habitability of planets orbiting cool stars.

1012.1883

Terrestrial planets are more likely to be detected if they orbit M dwarfs due to the favorable planet/star size and mass ratios. However, M dwarf habitable zones are significantly closer to the star than the one around our Sun, which leads to different requirements for planetary habitability and its detection. We review 1) the current limits to detection, 2) the role of M dwarf spectral energy distributions on atmospheric chemistry, 3) tidal effects, stressing that tidal locking is not synonymous with synchronous rotation, 4) the role of atmospheric mass loss and propose that some habitable worlds may be the volatile-rich, evaporated cores of giant planets, and 5) the role of planetary rotation and magnetic field generation, emphasizing that slow rotation does not preclude strong magnetic fields and their shielding of the surface from stellar activity. Finally we present preliminary findings of the NASA Astrobiology Institute's workshop "Revisiting the Habitable Zone." We assess the recently-announced planet Gl 581 g and find no obvious barriers to habitability. We conclude that no known phenomenon completely precludes the habitability of terrestrial planets orbiting cool stars.

12168040

2011ApJ...729....1X

1012

1012.1598_arXiv.txt

Soft Gamma-Repeaters (SGRs) are thought to be a type of magnetars: isolated neutron stars with much higher magnetic field than normal pulsars. The characteristic strength of magnetic field is believed to be about $10^{14}-10^{15}$ G, inferred from the observational measurements of the pulse period and its derivative assuming a dipole field. Recurrent X-ray and soft gamma-ray flares have been seen from SGRs, with a typical luminosity up to $10^{42}$ ergs/s since their discovery in 1979 (for an observational review, see Mereghetti 2008). Several major observational properties of the flares in SGRs can be explained by the most popular model in which the flares are produced by Alfv\'{e}n wave which propagates in the magnetosphere of the neutron star, due to the magnetic field diffusion associated with small cracks in the neutron star crust \citep{tho95}. Besides these recurrent flares, several giant flares with much higher luminosity have been observed. The luminosity during giant flares could reach more than a million times the Eddington luminosity of a common neutron star, i.e. $10^{44}-10^{46}$ erg/s. As an example, for the most extreme giant flare observed from SGR 1806-20 in 2004, the peak luminosity was about $2\times10^{46}$ erg/s \citep{hur05,mer05,pal05,ter05}. Energy of giant flares was thought to have a magnetospheric origin other than a neutron star crust origin. In the above magnetar model, giant flares are caused by a fireball of pair-dominated plasma that expands at relativistic speed. The unwinding internal magnetic field will build up the magnetic energy outside the neutron star slowly by pushing out inside electric currents into the magnetosphere. When the magnetosphere currents reach a point at which the magnetosphere becomes dynamical unstable, the twist field outside the star associated with dissipation will have a sudden relaxation and the magnetic field will have a reconfiguration, consequently the destructive magnetic instability may produce the neutron star crust fracturing in large scale, which can lead to an extreme outburst \citep{tho95,lyu06}. The total energy inferred from the magnetic strength of magnetar is of about $10^{47}$ erg, which is significantly more than energy contained in common neutron stars in the form of rotational kinetic energy \citep{bra06}. During a giant flare a large portion of the energy contained in magnetar would be lost, so the giant flares are very rarely seen. Up to now only three events have been detected: the 1979 flare from SGR 0526-66, the 1998 flare from SGR 1900+14, and the 2004 flare from SGR 1806-20. Interestingly, it was suggested before the 2004 Giant flare that magnetars would have a supergiant flare produced by a dynamic overturning instability. These type of flare should be 100 times less frequent than normal giant flares but could be observed within about 10 times the distance because of its much higher energy release \citep{eic02}. The 2004 giant flare was just like the predicted events. Giant flares have some common properties. In the X-ray waveband the light curves of giant flares consist of a short hard initial spike followed by a long softer pulsating tail. The peak luminosities of the 1979 flare and the 1998 flare were about $10^{44}$ erg/s \citep{maz79,fen81,hur99,fer99,fer01} while during the 2004 flare it reached to $2\times10^{46}$ erg/s \citep{hur05,mer05,pal05,ter05}. The rise time is typically smaller than a few milliseconds during which the luminosities of giant flares increase rapidly. After the initial spike there is the pulsating tail lasting about hundreds of seconds, which is thought to originate from the cooling of the ejected plasma which remains confined in a trapped thermal fireball near the stellar surface due to the strong magnetic field in magnetar model \citep{tho95}. For the three known giant flares, the pulsating tails always show a strong periodic modulation at the neutron star spin period and have complex pulse profiles. They are usually not single-peaked and show strong evolution with time. During the 1998 flare from SGR 1900+14, the pulse profile of the pulsating tail showed four peaks and a dip which were almost evenly spaced at about 1.03 s interval in each period of 5.16 s \citep{fer01}, and the four-peak structure evolved to being nearly sinusoidal with time. The multi-peak features are thought to be described as collimated X-ray jets from trapped fireballs, in comparison to the 1979 giant flare from SGR 0526-66 during which only two sub-pulses have been observed. It was suggested that the number of sub-pulses detected in each period of these two sources could be due to different number of X-ray jets \citep{tho01}. For the 1998 giant flare the phases of these four sub-pulse peaks were recognized stable, which may suggest that the X-ray jets are tied to surface features on the neutron star. The four-peaked pattern is thought to be direct evidence of multipolar structure of the neutron star's magnetic field \citep{fer01}. Investigations in the quiescent emission of SGR 1900+14 revealed that before the 1998 giant flare the pulse profile of the source always showed complex multi-structure \citep{woo01}, and once it evolved to sinusoidal during the flare it keeped unchanged for more than 1.5 years. It is suggested that if the multipeak profile indicates multipolar field structure and the simple sinusoidal profile indicates dipolar field geometry, during the time of the 1998 giant flare the magnetic field, both inside and outside the neutron star, underwent a global reconfiguration, while during the quiescent stage the magnetic field was relative stable \citep{woo01}. The evolution of pulse profile was also observed during the 2004 flare from the SGR 1806-20. It showed an opposite trend compared with the 1998 flare; the pulse profile evolved from a simple sinusoidal to a complex multi-peaked structure \citep{woo06}. The phases of these multi-peak are generally thought stable, indicating that there was no substantially field change during the tail \citep{pal05,mer05,boggs07}, even though we see obvious phase shifts of the sub-pulses on short time scales (see our study below). In the quiescent X-ray light curve prelude to the 2004 flare, the pulse profile was nearly simple sinusoidal, it became more complex after the flare. Considering the pulse profile changes during the flare, it was recognized that the pulse profile changes in the quiescent light curve were flare-induced, consistent with SGR 1900+14 \citep{woo07}. Besides the two SGRs from which evolving pulse profile has been observed, during magnetar flares from magnetar candidates, for example CXOU J164710.2-455216, the pulse profile was also observed to evolve. During the observations of 4.3 day prior to and 1.5 day subsequent to two remarkable events that were detected with Swift on 2006 September 21 (a 20-ms burst and a glitch), the pulse profile was found to change from single-peaked to three distinct peaks \citep{mun07}. For this type of source with lower X-ray energy release, a large-scale rearrangement of the magnetic field was thought unreasonable. It is suggested that the pulse shape was governed by the distribution of currents within the magnetosphere, and the profile change should be due to the change in the distribution of currents in the magnetosphere \citep{mun07}. In the following study we focus on the pulse profile during the pulsating tail of the 2004 flare form SGR 1806-20. In our previous study of SGR 1806-20 giant flare using {\it GoodXenon} mode data of the RXTE/PCA {\it only}, we have decomposed this pulse profile into four sub-pulses. We modeled the pulse profile with four Gaussian functions and a constant background flux level, and investigated the X-ray waveform evolution of each sub-pulses, excluding the first 2-3 cycles because of data gaps in the {\it GoodXenon} mode data \citep{xin09}. In that study we found the phases of the sub-pulses were not fixed during the tail. The phase of one sub-pulse was decreasing while the phase of another sub-pulse was increasing with time -- an apparent anti-correlation between the phase drifts of the two sub-pulses. This suggests that the phases of the sub-pulses moved although the multi-peak waveform structure remained unchanged. Modeling the sub-pulses as multi-peak Gaussian gives good description of the overall phase distribution of the photons in each sub-pulses, but not accurate and sensitive in determination of the phases of the sub-pulse peaks. Furthermore, the data gap in the {\it GoodXenon} mode data recorded at the beginning of the Giant flare because of telemetry saturation prevents us from studying the sub-pulse phase evolution during the first few cycles of the giant flare. In this study, we intend to measure the phases of the sub-pulse flux peaks, which show the sub-pulse phase shifts more obviously. We determined the sub-pulse phase by 7-point Gaussian fit to each sub-pulse peak with {\it Standard 1} mode data instead of {\it GoodXenon} data. This data mode does not have data gaps and uniquely helps determination of the sub-pulse phases during the first few cycles. We found strong evidence of systematic phase shifts of those sub-pulse peaks, in agreement with our previous results but further helped us measure the characteristic timescales of the phase evolution of the sub-pulses. This timescale might indicate the timescale for the magnetic field to settle on the neutron star after a reconfiguration, which has been suggested to associate with the giant energy release during the flare.

We have determined the phases of the sub-pulses in the X-ray waveform during the 2004 giant flare and studied their evolutions. There existed phase drifts of these sub-pulses up to 0.06 in less than 100 seconds. The phases of the sub-pulses became stable after a characteristic timescale of tens of seconds. This indicates that the magnetic instability responsible for the magnetic field reconfiguration damps out on the above timescale and the magnetic field became stable after such a characteristic time scale. Two sub-pulses separated by about 0.5 phase, namely the 1st sub-pulse and the 4th sub-pulse in our study, show opposite evolution trend. The apparently correlated evolution trend is more obvious during the early stage of the giant flare. If the two sub-pulses correspond to the emission regions of two opposite magnetic poles or of the same magnetic pole, the phase evolution and the opposite phase trend indicate the two emission regions are probably tied together during the giant flare. The $\sim0.5$ phase separation of the two sub-pulses and the potentially physically correlated phase drifts then likely provide evidence that we can see at least two poles of the neutron star. This provides an additional evidence for a multipolar magnetic field geometry other than that discussed in Feroci et al (2001) on the four sub-pulse profile. The initial spike of a giant flare has a sharp rise which is typically smaller than a few milliseconds. For the 1979 event the rise time was smaller than 2 ms. For the 1998 event it was smaller than 4 ms \citep{maz99}. During the 2004 giant flare from SGR 1806-20 the rise time was much smaller than ever observed, which took only about 0.3 ms \citep{pal05}. The steep rise time of giant flare is thought associated with the Alfv\'{e}n crossing time of the inner magnetosphere, which is as short as $30 \mu s$ estimated by $R_{NS}/c$, where $R_{NS}$ is the radius of the neutron star and $c$ is the light speed \citep{lyu06,tan07}. During the steep rise the twist field outside the neutron star suffers a sudden relaxation, which leads to a rapid release of energy stored in the magnetosphere \citep{lyu06,bra06}. Besides the initial steep rise there was another timescale noted in previous studies, i.e. the intermediate rise time to the peak. The intermediate rise time to the peak during the giant flare from SGR 1900+14 was 3.1 ms, which was shorter than the 9.4 ms observed during the 2004 flare from SGR 1806-20. This timescale is thought to be limited by the propagation of a fracture, so it could be used to estimate the fracture size during giant flares \citep{tan07}. For all of the three giant flares observed, the duration of the initial spike was about hundreds of milliseconds: 0.25 s for the 1979 giant flare from SGR 0526-66; 0.35 s for the 1998 giant flare from SGR 1900+14; and 0.5 s for the 2004 giant flare from SGR 1806-20 \citep{mer051}. It may reflect the Alfv\'{e}n wave crossing time of the star which is in the range between 0.2 s and 0.5 s estimated by $R_{NS}/V_{A,NS}$, where $V_{A,NS}$ is the Alfv\'{e}n wave crossing speed in the star \citep{tho95}. The e-folding decay time of the giant flare was thought to describe the cooling time of the trapped fireball, which was formed by energy released in the initial spike trapped in the neutron star magnetosphere by high magnetic field \citep{tho95,pal05,mer08}. The emission from trapped fireball in higher energy band seemed to cool a little faster than that in lower energy band, which leads to a softening of the dominated emission during the later part of the pulsating tail. The characteristic time scales of the phase drifts of the sub-pulses are a little smaller, but comparable to the decay time scale of the giant flares. We therefore infer that the magnetic field reconfiguration settle earlier than the cooling of the trapped fireball in the magnetar model. In this work we have found the phases of the sub-pulses were obviously drifting during the first ten cycles or so and gradually became stable. The characteristic timescale of phase drifts is in the range 37 s -- 84 s. This may tell us the instability of the magnetic field during the giant flare damped in tens of seconds before the field configuration became stable. Thus this time scale may be used to constrain the magnetic fields and the mechanism for the energy release of the giant flare. \begin{figure} \includegraphics[width=6in]{f1.eps} \caption{Top: the (2-60 keV) RXTE/PCA light curve during the 2004 outburst of SGR 1806-20 generated from the Standard-1 mode data. The time resolution of the light curve is 0.125 s. Bottom: evolution of the average rates of each cycles during the pulsating tail.} \end{figure} \begin{figure} \includegraphics[width=6in]{f2.eps} \caption{Pulse profile evolution during the first 35 cycles of the 2004 giant flare obtained from {\it Standard 1} data. The non-pulsed background trend is removed from the data. Time sequence increases from bottom to top. Count rate corresponding to each cycle was shifted by 1500 counts/s one by one. The circle marks the peak of the major sub-pulse during the first period, consistent with the sub-pulse \#3 identified in subsequent cycles.} \end{figure} \begin{figure} \includegraphics[width=6in]{f3.eps} \caption{An example of the model fit to the waveform in which each of the three major sub-pulse peaks were fitted with a Gaussian (seven-point fits). The best-fit Gaussians were over-plotted as dashed lines. The central phases of the Gaussians gives the phases of the sub-pulses. A constant level at the mean count rate of the cycle has been added.} \end{figure} \begin{figure} \includegraphics[width=6in]{f4.eps} \caption{Phase evolution of the sub-pulses \#1, \#2, and \#4 measured by seven-points Gaussian fits to the sub-pulse peaks.} \end{figure} \begin{figure} \includegraphics[width=6in]{f5.eps} \caption{Phase evolution of the sub-pulses modeled with exponential phase drift. Notice that the phases of sub-pulse \#3 was obtained from a decomposition of the entire waveform into 4 Gaussian components and the phases of the \#1 and \#4 were obtained from 7-points Gaussian fits. The phase drifts should be considered as those in relative to the sub-pulse \#2 -- the major sub-pulse.} \end{figure} \begin{figure} \includegraphics[width=6in]{f6.eps} \caption{The model fit to data around sub-pulse 1 considering the left shoulder of sub-pulse 2 of the 3rd period. } \end{figure} \begin{figure} \includegraphics[width=6in]{f7.eps} \caption{The model fit to data around sub-pulse 4 considering the right shoulder of sub-pulse 2 of the 3rd period. } \end{figure} \begin{figure} \includegraphics[width=6in]{f8.eps} \caption{Phase evolution of the sub-pulses modeled with exponential phase drift. Different from Figure 5, the phases of the sub-pulses 1 and 4 were obtained from fitting with the sloping sides of sub-pulse 2 to a Gaussian function plus a linear slope. Notice that data with reduced-$\chi^2$ larger than 10 were excluded in the plot. Around period No. 17--23 the model does not fit sub-pulse 4 well. } \end{figure}

1012.1598

We analyzed the observations of SGR 1806-20 performed with the Rossi X-ray Timing Explorer during its 2004 giant flare. We studied the phase evolution of the sub-pulses identified in the X-ray waveform and found that the sub-pulses varied in phase with time and then gradually settled, which might indicate drifts of the emission regions relative to the neutron star surface, or changes in the local emission geometry before the magnetic field became stable. The characteristic e-folding timescale of the phase drifts measured starting about 15 s following the initial flux spike is in the range between 37 s and 84 s. This leads to the first measurements of the characteristic timescale for the magnetic field of the neutron star to settle after a field reconfiguration during the giant flare.

12213847

2011MNRAS.414..860K

1012

1012.4718_arXiv.txt

\label{sec_intro} Since the realisation that the Planetary Nebula luminosity Function (PNLF; Jacoby, Ciardullo \& Ford 1988) can be used as a standard candle, it has been firmly established as a reliable extra-galactic distance indicator for galaxies out to the Virgo (18.0$\pm$1.2~Mpc; \citealp{2001A&A...375..770F}) and Fornax (18.6$\pm$0.6~Mpc; \citealp{1999ApJ...515...29M}) clusters and beyond \citep{JCF90, FJP07, G07,C10}. Agreement with traditional population-I distance indicators such as Cepheids and SN~Ia \citep{C03} provides us with a rare link with which we can compare these independent methods between very different stellar populations of different ages and metallicities. However, despite the robustness of the PNLF as a distance indicator, the underlying detailed physics responsible is not well understood. The strength of the PNLF distance technique resides in the fit of the PNLF function having a definitive absolute magnitude bright-end cut-off at M$_{\mathrm{[OIII]}}=-4.47^{+0.02}_{-0.03}$ \citep{C02}, regardless of galaxy type and age, though there does appear to be a small dependence on metallicity \citep{D92}. \citet{C05} outlined this fundamental inconsistency. The PN at the bright-end of the PNLF have central star luminosities of $\sim$6000~L$_{\sun}$, which requires the central stars to be $>$0.6~M$_{\sun}$ \citep{2008ApJ...681..325M,2004A&A...423..995M} and corresponds to a mass on the main-sequence of $>$2~M$_{\sun}$ \citep{W00}. Stars with such main-sequence masses, and consequent short lifetimes ($\sim$1-2~Gyr), are not expected to be present in elliptical galaxies whose populations are typically $\sim$10~Gyr old and where there is no evidence for recent star formation. However, the bright PN that have evolved from such stars are nonetheless detected in these populations. This conundrum has been a major obstacle in our interpretation and understanding of the apparently fixed nature of the bright-end of the PNLF across galaxy types for nearly 20~yrs. To address the problem we urgently need to deconstruct the PNLF in fine detail. \\ To evaluate the varied characteristics of individual PN constituting the PNLF, we require a self-contained system at a known distance whose PN population is sufficiently nearby to permit investigation into individual PN morphologies, their abundances, expansion velocities, central star temperatures and consequently their masses. Such detailed analyses can be achieved in relatively nearby galaxies such as the LMC \citep{2010arXiv1002.3410R} and SMC \citep{2002AJ....123..269J}, but being metal-poor and comprised of intermediate-age populations they cannot represent valid proxies for an old elliptical galaxy and cannot be used to address the problem outlined above. \\ Fortunately, the bulge of our Galaxy, being relatively nearby ($\sim$8~Kpc), has the potential to represent such a system. There is mounting evidence that our Galactic bulge formed within a very short time-scale \citep{rich05,zoccali06,fulbright07} about $\sim10$~Gyr ago \citep{ortolani95,feltzing00,Z03}. A rapid star formation history in the early universe indicates that early-type spiral bulges undergo a comparable formation mechanism to elliptical galaxies \citep{peletier99,falcon02}. We can therefore exploit its proximity and population age as a proxy for an elliptical galaxy amendable to detailed study. This work has the potential to determine whether the PNLF bright-end is comprised of PN resulting from old, population-II stars which have found some peculiar path to enhanced luminosity, e.g. through binarity \citep{C05,M09a}, or if it is in fact dominated by younger, higher mass, bipolar nebulae mainly of Type-I as defined by \citet{KB94} and \citet{TPP97}. These are thought to derive from more massive progenitors which suffer third dredge-up and undergo hot-bottom burning, subsequently becoming nitrogen and helium enriched. Such detailed analysis of an old stellar population is currently, and for the forseeable future, impossible to do anywhere else. The next nearest old population is the bulge of M31 and in order to obtain spectra of comparable quality/depth to what exists for Galactic bulge PN requires instrumentation beyond what is currently available. Hence, this work will provide a significantly improved understanding of the PN population in all elliptical galaxies as well as in spiral bulges. \\ Our ability to study such samples has been recently enhanced thanks to the highly sensitive AAO/UKST SuperCosmos $\halpha$ Survey of the southern Galactic plane (SHS; \cite{P05}) which enabled doubling the number of known bulge PN in this region as reported in the Macquarie/AAO/Strasbourg $\halpha$ PN catalogues of \citet{P06} and \citet{M08}, dubbed MASH and MASH-II. This combined sample is more representative of the bulge PN population spanning a wider evolutionary range and allows us here to construct a new $\othreec$ Galactic bulge PNLF of unprecedented coverage. This will enable detailed study of the various sub-sets of PN (whether Type-I chemically or of certain morphologies) within the PNLF, whether they are confined to certain regions within the PNLF and, importantly, whether one population sub-set is solely responsible for constituting the bright-end. These analyses will be presented in \citet{K10b}, hereafter Paper II, while this paper describes the process for defining robust $\othreec$ flux measurements. \\

In this paper, we have increased the number of PN with directly measured $\othreec$ fluxes in the $10\arcdeg\times10\arcdeg$ toward the Galactic bulge region by a factor of eight, with the addition of 387 previously unobserved PN. We have also re-observed 48 PN with directly measured $\othreec$ fluxes from the literature. Our new fluxes account for $\sim$80\% of known PN towards the Galactic bulge. We have provided additional fluxes for five PN in the peripheral regions, and for a special PN outside the bulge region (PHR1315-6555) that is the subject of another paper \citep{P10}. This brings the total number of PN for which we have provided measured fluxes for to 441. \\ Photometric errors were derived via consideration of the S/N measurement, the velocities of the PN, the transmission profile of the filter curve and those due to the variation in the sensitivity factor calculated from the standard stars used. We found that these errors under-estimated the true error after comparison between duplicate measurements of a PN, so we increased the errors accordingly. We have compared our fluxes to all previously published $\othreec$, $\othreeb$, $\halpha$ and $\hbeta$ fluxes, and found there to be no statistically significant systematic offset between the most trusted datasets and our new flux measurements to within the errors. Angular diameter measurements for all PN observed are also included, along with preliminary estimates of the dereddened fluxes based on the current best available spectra. \\ These data therefore provide accurate fluxes and angular diameters for the largest sample of PN in the $10\arcdeg\times10\arcdeg$ region toward the Galactic bulge ever collated. Deriving from the same high quality, uniform, high resolution \linebreak MOSAIC-II imaging, precise angular diameter measurements and greater insight into the fine morphological detail of each PN can be made. Combination of these factors will allow for the most detailed and accurate construction of the bulge PNLF to date. This will be the subject of Paper II.

1012.4718

We present ? fluxes and angular diametres for 435 Planetary Nebulae (PNe) in the central 10°× 10° region towards the Galactic bulge. Our sample is taken from the new discoveries of the MASH PN surveys as well as previously known PN. This sample accounts for 80 per cent of known PN in this region. Fluxes and diametres are measured from narrow-band imaging with the MOSAIC-II camera on the 4-m Blanco Telescope at the Cerro-Tololo Inter-American Observatory. This is the largest (∼60 deg<SUP>2</SUP>), uniform ? survey of the inner Galactic bulge ever undertaken. 104 of the objects have measured ?, ?, Hα or Hβ fluxes from the literature, which we use to undertake a detailed comparison to demonstrate the integrity of our new fluxes. Our independent measurements are in excellent agreement with the very best literature sources over two orders of magnitude, while maintaining good consistency over five orders of magnitude. The excellent resolution and sensitivity of our data allows not only for a robust set of homogenous PN fluxes, but provides greater detail into their intricate, otherwise undetermined ? morphologies. These new, extensive measurements significantly increase the sample of reliable ? fluxes for Galactic bulge PN making it a valuable resource and a prelude to the construction of our new Galactic bulge PN luminosity function (Paper II).

2965719

2011AJ....141...71F

1012

1012.4232_arXiv.txt

Very low--mass stars and brown dwarfs (VLMs and BDs; M $\lesssim$ 0.10 M$_{\sun}$) are among the most populous constituents of the Galaxy yet their origins remain controversial. Whether they form in a manner similar to higher mass stars or require additional or completely different processes is currently under debate (see \citealt{2007prpl.conf..427B}; \citealt{2007prpl.conf..459W}; \citealt{2007prpl.conf..443L} and references there-in). One important characteristic that indicates a difference in formation is their binary frequency and multiplicity statistics (separation, mass ratio, eccentricities, etc). Contrary to higher mass stars, VLM binaries are found tightly bound ($\rho$ $<$ 20 AU), in small frequency (10-20$\%$) and with high mass ratios (typically q$\sim$1 -- e.g. \citealt{2003AJ....126.1526B}, \citealt{2003ApJ...587..407C}, \citealt{2003ApJ...586..512B}, \citealt{2007ApJ...671.2074A}, \citealt{2008AJ....135..580R}). Over the past decade, theoretical and observational studies have converged on several competing mechanisms to explain the formation of VLM stars and brown dwarfs. Among the most prominent are (1) magnetoturbulent or gravoturbulent fragmentation (\citealt{2004ApJ...617..559P}; \citealt{2003MNRAS.339..577B}; \citealt{2004A&A...414..633G,2004A&A...423..169G,2006A&A...452..487G}); (2) ejection of a protostellar embryo from the natal core (\citealt{2001AJ....122..432R}; \citealt{2005MNRAS.356.1201B}); or (3) disk fragmentation (\citealt{2006A&A...458..817W}; \citealt{2009arXiv0911.3662S}). Characteristics of multiple systems predicted by each have been used to differentiate which mechanism is most probable for the population. Early versions of the ejection model were favored in part because they predicted that brown dwarf binaries should have separations no greater than 10 AU (\citealt{2002MNRAS.332L..65B}). This was supported through early observational studies. However recent findings of widely separated VLM systems (typically $\rho$ $>$ 100 AU) are contradictory and shed doubt on the viability of the ejection scenario (\citealt{2004ApJ...614..398L}; \citealt{2009ApJ...697..824A}; \citealt{2007ApJ...660.1492C}; \citealt{2009ApJ...691.1265L}; \citealt{2005A&A...440L..55B}; \citealt{2007ApJ...667..520C}; \citealt{2007ApJ...659L..49A}; \citealt{2009arXiv0903.3251R}). Updated simulations have successfully created a handful of wide brown dwarf binaries (\citealt{2005MNRAS.356.1201B,2009MNRAS.392..590B}) but the growing number in the field indicates a fraction too high to be explained by just this mechanism. The existence and frequency of wide substellar pairs may not discount one mechanism versus another; rather multiple formation mechanisms may be at work in the field. Recent work has shown that brown dwarfs widely separated from nearby stars have a higher frequency of also being tight multiples (\citealt{2005AJ....129.2849B}; \citealt{2009AJ....137....1F}). This has been explored in the disk fragmentation models of \citet{2009arXiv0911.3662S} where similar high frequencies of brown dwarf multiples were produced when widely separated from a more massive companion. In this article, we report the identification of another widely separated brown dwarf binary (or triple; see Burgasser et al 2010) companion, the NLTT 20346 and 2MASS J0850+1057 comoving system. In Section 2, we review the discovery and observational data used to characterize the system. In Section 3, we discuss the astrometric details of the system including a contaminating artifact that likely skewed previous parallax measurements for 2MASS J0850+1057. In Section 4, we analyze the details of the system including the activity and masses of each component. In Section 5, we discuss its importance among the population of VLM wide companion systems. Results are summarized in Section 6.

We have added NLTT 20346/2MASS J0850+1057 to the growing list of VLM (M$_{tot}<$0.2~M$_{\sun}$) multiples widely separated from a more massive companion. Table ~\ref{system} is a selected compilation of these widely-separated multiples and demonstrates a range in both mass and separation of components. Within this list, NLTT 20346/2MASS J0850+1057 has a significantly lower binding energy and is one of the lowest total mass triple, quadruple, and/or quintuple systems known. Figure ~\ref{fig:be2} shows the binding energy versus total mass of systems gathered from the literature\footnote{Stellar companions were gathered from the catalogs of \citealt{1991A&A...248..485D}, \citealt{1992ApJ...396..178F}, and \citealt{1997A&AS..124...75T}; and young UCD companion systems from \citealt{2005ApJ...633..452K, 2006ApJ...649..306K}, \citealt{2007ApJ...663..394K}, \citealt{2009ApJ...691.1265L}, and \citet{2009ApJ...697..824A}. Details on the field UCD systems were gathered from the Very Low Mass Binary Archive ($http://vlmbinaries.org$; see \citealt{2007prpl.conf..427B} and references therein.)} and demonstrates that this new system falls well below the binding energy limitation set by known tight low mass (M$_{tot}$ $<$ 0.2 M$_{\sun}$) multiples (see \citealt{2007ApJ...660.1492C,2003ApJ...587..407C} and \citealt{2003ApJ...586..512B}). For comparison, \citet{1998MNRAS.299..955P} use the virial theorem to obtain the gravitational binding energy of the Pleiades cluster. Their derived value ($\sim$ 5.3 x 10$^{38}$ erg) is smaller than the estimated binding energy for the NLTT 20346/2MASS J0850+1057 system (and smaller than the lowest binding energy systems found to date), although it is unclear how long this association will remain bound. An interesting investigation would be to search for stars in the vicinity of NLTT 20346/2MASS J0850+1057 as the hierarchical nature of this system combined with the common kinematics, potential youth, and wide separation between the M dwarf and brown dwarf systems might be indicative of its own moving group. % \begin{figure}[htbp] \begin{center} \includegraphics[width=1.0\hsize]{Binding_Energy2.eps} \end{center} \caption{A plot of the binding energy vs. total mass. Objects marked with filled circles are tight low--mass systems (typically M$_{tot}$ $<$ 0.2M$_{\sun}$ and $\rho<$20 AU). Wide systems ($\rho>$100 AU) containing a UCD companion are marked as five point stars and those highlighted in red also contain a VLM binary (hence at least a triple). Those marked as squares are systems containing a tight or widely separated UCD with an age $<$ 500 Myr. Objects marked by open circles come from stellar companion catalogs. The NLTT 20346/2MASSJ0850+1057 system is the enlarged five-point star. The minimum binding energy corresponding to tight very low mass systems is labeled (see \citealt{2007ApJ...660.1492C,2003ApJ...587..407C} and \citealt{2003ApJ...586..512B}). } \label{fig:be2} \end{figure} \begin{deluxetable*}{llllllllllllllllllllllll} \tabletypesize{\footnotesize} \label{tab:tab1} \tablecaption{Details on Triple/Quadruple Systems Containing at Least Two Ultracool Dwarfs \label{system}} \tablewidth{0pt} \tiny \tablehead{ \colhead{Name} & \colhead{SpT} & \colhead{SpT} & \colhead{$\rho$} & \colhead{$\rho$} & \colhead{$\rho$} & \colhead{$\rho$} & \colhead{M}& \colhead{M} & \colhead{M }& \colhead{q\tablenotemark{a}} & \colhead{BE} & \colhead{Ref}\\ & & & \colhead{($\arcsec$)} & \colhead{(AU)}& \colhead{($\arcsec$)} & \colhead{(AU)}& \colhead{($M_{\sun}$)} & \colhead{($M_{\sun}$)} & \colhead{($M_{\sun}$)}& & \colhead{10$^{41}$ Erg}\\ & \colhead{\tiny{Primary}} & \colhead{\tiny{Secondary}}& \colhead{\tiny {star-UCD}} & \colhead{\tiny {star-UCD}}& \colhead{\tiny{UCD-UCD}}& \colhead{\tiny{UCD-UCD}} & \colhead{\tiny{Primary}}& \colhead{\tiny{Secondary}}& \colhead{\tiny{Total}}& & \\ \colhead{(1)} & \colhead{(2)} & \colhead{(3)} & \colhead{(4)} & \colhead{(5)} & \colhead{(6)} & \colhead{(7)} & \colhead{(8)} & \colhead{(9)} & \colhead{(10)} & \colhead{(11)} & \colhead{(12)} & \colhead{(13)} \\} \startdata NLTT 20346 & M5+M6 & L7+L6.5 & 248 & 7700 & 0.16 & 4.7 & 0.290 & 0.070 & 0.360 & 0.24 & 0.37 & 1A,1B\\ \hline G 171-58 & F8 & L4+L4 & 218 & 9202 & 0.33 & 10.2 & 1.150 & 0.095 & 1.245 & 0.08 & 1.7 & 2\\ HD 221356 & F8 & M8+L3 & 452 & 11900 & 0.57 & 15 & 1.020 & 0.160 & 1.180 & 0.16 & 1.9 & 3\\ G 124-62 & dM4.5e & L1+L1 & 44 & 1496 & 0.42 & 14.3 & 0.210 & 0.144 & 0.354 & 0.69 & 2.8 & 4\\ eps Ind & K5 & T1+T6 & 402 & 1460 & 0.73 & 2.6 & 0.670 & 0.072 & 0.742 & 0.11 & 4.6 & 5\\ Gl 417 & G0+G0 & L4.5+L6 & 90 & 2000 & 0.07 & 1.5 & 0.940 & 0.143 & 1.083 & 0.15 & 9.4 & 6A,6B\\ LP 213-67 & M6.5 & M8+L0 & 14 & 230 & 0.12 & 2.8 & 0.100 & 0.176 & 0.276 & 1.76 & 11 & 7A,7B\\ GJ 1001 & M4 & L4.5+L4.5 & 19 & 180 & 0.09 & 1 & 0.250 & 0.136 & 0.386 & 0.54 & 26 & 8A,8B,8C\\ Gl337 & G8+K1 & L8+L8/T & 43 & 880 & 0.53 & 10.9 & 1.740 & 0.110 & 1.850 & 0.06 & 30 & 9A,9B\\ HD65216 & G5 & M7+L2 & 7 & 253 & 0.17 & 5.9 & 0.940 & 0.167 & 1.107 & 0.18 & 87 & 10\\ GJ569 & M2.5 & M9.0+M9.0 & 5 & 50 & 0.10 & 0.9 & 0.350 & 0.126 & 0.476 & 0.36 & 123 & 11\\ Kelu-1 & L0.5/T7.5 & L3p & -- & -- & 0.29 & 5.4 & 0.125 & 0.055 & 0.180 & 0.44 & 178 & 12A,12B,12C\\ LHS1070 & M5.5 & M8.5+M9.0 & 1 & 4 & 0.45 & 3.4 & 0.210 & 0.138 & 0.348 & 0.66 & 1134 & 13A,13B\\ HD130948 & G2 & L4+L4 & 0.4 & 7 & 0.13 & 2.4 & 1.030 & 0.109 & 1.139 & 0.11 & 2178 & 14A,14B\\ \enddata \tablerefs{1A=This Paper 1B=Burgasser et al (2010) 2=\citet{2010AJ....139..176F} 3=\citet{2007ApJ...667..520C} 4=\citet{2005A&A...440..967S} 5=\citet{2003A&A...398L..29S} 6A=\citet{2001AJ....121.3235K} 6B=\citet{2003AJ....126.1526B} 7A=\citet{2000MNRAS.311..385G} 7B=\citet{2003ApJ...587..407C} 8A=\citet{2004AJ....128.1733G} 8B=\citet{1999ApJ...519..802K} 8C=\citet{1999Sci...283.1718M} 9A=\citet{2001AJ....122.1989W} 9B=\citet{2005AJ....129.2849B} 10=\citet{2007MNRAS.378.1328M} 11=\citet{2000ApJ...529L..37M} 12A=\citet{1997ApJ...491L.107R} 12B=\citet{2005ApJ...634..616L} 12C=\citet{2008arXiv0811.0556S} 13A=\citet{1994A&A...291L..47L} 13B=\citet{2001A&A...367..183L} 14A=\cite{2002ApJ...567L.133P} 14B=\citet{2008arXiv0807.2450D}} \tablenotetext{a}{q represents the mass ratio $\frac{M_{secondary}}{M_{primary}}$} \end{deluxetable*} Figure ~\ref{fig:be} shows the M$_{tot}$ vs. separation for the same companion systems. The studies of \citet{2001AJ....121..489R} and \citet{2003ApJ...586..512B} suggested an empirical limit for the stability of VLM multiples based on objects that were known at the time of the respective works. Neither cut-off seems appropriate for NLTT 20346/2MASS J0850+1057 or the collection of slightly more massive widely separated systems plotted on Figure ~\ref{fig:be} (see discussions in \citealt{2010AJ....139..176F} and \citealt{2010AJ....139.2566D}). \citet{2010AJ....139.2566D} found a log-normal limitation on the separation of binaries in their catalog of 1342 wide ($>$ 500 AU) low-mass (majority had M$_{tot}$ $>$ 0.3) systems and this cut-off does encompass NLTT 20346/2MASS J0850AB and all low mass objects on Figure ~\ref{fig:be}. \begin{figure}[htbp] \begin{center} \includegraphics[width=1.0\hsize]{Binding_Energy1.eps} \end{center} \caption{A plot of the total mass vs. separation. Symbols are described in Figure ~\ref{fig:be2}. We have over-plotted the \citet{2001AJ....121..489R} curve (center) which distinguished the cut-off for the formation of wide stellar pairs as well as the \citet{2003ApJ...586..512B} line (far right--which is specific for M$_{tot}<$0.2 M$_{\sun}$ field systems) and the log-normal limitation found in \citet{2010AJ....139.2566D} (far left--for systems with M$_{tot}$ $>$ 0.3 M$_{\sun}$).} \label{fig:be} \end{figure} The stability of binary systems will be a function of age and total mass (see \citealt{1987ApJ...312..367W}). While NLTT 20346/2MASS J0850AB is fit by the cut-off limitation of higher mass companions, it is significantly different from other known slightly less massive higher order multiples so we investigate one possible mechanism for its formation. Recent work by \citet{2009MNRAS.392..413S} has been successful in accounting for widely separated VLM binaries using simulations of gravitational fragmentation of massive extended disks. In their smoothed particle hydrodynamic (SPH) simulations, a system with M$_{DISK}$=0.7M$_{\sun}$, R$_{DISK}$=400AU, M$_{star}$=0.7 M$_{\sun}$ is evolved for up to 20 kyr followed by an N-body dynamical evolution for up to 200 Kyr. After 12 simulations, 96 stars are formed with brown dwarf or low--mass secondaries and among those companions, 9 are tight VLM multiples. The characteristics of the triple systems were listed in \citet{2009MNRAS.392..413S} and one was found to have a total secondary mass of 0.148 M$_{\sun}$, a close brown dwarf-brown dwarf binary separation of $\sim$ 1 AU and a wide separation from the central star of 7700 AU. This is slightly more massive but highly analogous to our proposed system indicating that gravitational fragmentation could account for the existence of NLTT 20346/2MASS J0850+1057. We found no analogs created using different formation mechanisms therefore comparisons with the predictions from ejection, etc. are not possible at this time. The simulations of \citet{2009arXiv0911.3662S} showed a significant population of low mass (M$_{secondary}$ $<$ 100 M$_{J}$) companions at distances out to 10,000 AU. This is in agreement with recent studies that have revealed a growing number of ultracool dwarfs with separations approaching and in some cases exceeding 10,000 AU from a companion (see Table 1 in \citealt{2009AJ....137....1F} and references there-in; \citealt{2010AJ....139.2566D};\citealt{2010MNRAS.404.1817Z}). The caveat in this comparison is that the observed sample covers primarily field age objects (1-5 Gyr) while the \citet{2009arXiv0911.3662S} simulations stop after only 200 kyr of dynamical evolution. The existence of a population of older, widely separated systems suggests that dynamical interactions into field ages does not disrupt all systems out to 10,000 AU. \citet{2009MNRAS.392..413S} find that the nine low-mass binaries formed through gravitational fragmentation that remain bound to their parent star have high eccentricities ($<$e$_{BIN}$$>$ of 0.7 $\pm$ 0.2). They postulate that the dynamical interactions which form them, and/or subsequent dynamical interactions with other stars that have condensed out of the disk, account for the increased eccentricity of the orbit. The masses of the primary stars in that work are not reported although most are described as Sun-like. We can compare the eccentricities of known widely separated VLM binaries with the simulation predictions although the primary masses are most likely smaller. Three of the VLM binary systems from Table ~\ref{system}, GJ 569B, HD 130948B, and 2MASSJ 0850+1057, have dynamical mass measurements and eccentricities measured (\citealt{2008arXiv0807.2450D}; \citealt{2010arXiv1001.4800K}). The mean eccentricity for these three systems is 0.4 $\pm$ 0.2, falling within 1$\sigma$ of the simulations. Further information on the eccentricity of VLM binaries widely separated from a more massive component will enhance our understanding of dynamical interactions of low-mass hierarchical systems.

1012.4232

We report our discovery of NLTT 20346 as an M5+M6 companion system to the tight binary (or triple) L dwarf 2MASS J0850359+105716. This nearby (~31 pc), widely separated (~7700 AU) quadruple system was identified through a cross-match of proper motion catalogs. Follow-up imaging and spectroscopy of NLTT 20346 revealed it to be a magnetically active M5+M6 binary with components separated by ~2'' (50-80 AU). Optical spectroscopy of the components shows only moderate Hα emission corresponding to a statistical age of ~5-7 Gyr for both M dwarfs. However, NLTT 20346 is associated with the XMM-Newton source J085018.9+105644, and based on X-ray activity the age of NLTT 20346 is between 250 and 450 Myr. Strong Li absorption in the optical spectrum of 2MASS J0850+1057 indicates an upper age limit of 0.8-1.5 Gyr, favoring the younger age for the primary. Using evolutionary models in combination with an adopted system age of 0.25-1.5 Gyr indicates a total mass for 2MASS J0850+1057 of 0.07 ± 0.02 M <SUB>sun</SUB>, if it is a binary. NLTT 20346/2MASS J0850+1057 joins a growing list of hierarchical systems containing brown dwarf binaries and is among the lowest binding energy associations found in the field. Formation simulations via gravitational fragmentation of massive extended disks have successfully produced a specific analog to this system. <P />This paper includes data gathered with the 6.5 m Magellan Telescopes located at Las Campanas Observatory, Chile.