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1609 | 1609.06602_arXiv.txt | Young and rapidly rotating stars are known for intense, dynamo generated magnetic fields. Spectropolarimetric observations of those stars in precisely aged clusters are key input for gyrochronology and magnetochronology. We use ZDI maps of several young K-type stars of similar mass and radius but with various ages and rotational periods, to perform 3D numerical MHD simulations of their coronae and follow the evolution of their magnetic properties with age. Those simulations yield the coronal structure as well as the instant torque exerted by the magnetized, rotating wind on the star. As stars get older, we find that the angular momentum loss decreases with $\Omega_{\star}^3$, which is the reason for the convergence on the Skumanich law. For the youngest stars of our sample, the angular momentum loss show signs of saturation around $8 \Omega_{\odot}$, which is a common value used in spin evolution models for K-type stars. We compare these results to semi-analytical models and existing braking laws. We observe a complex wind speed distribution for the youngest stars with slow, intermediate and fast wind components, which are the result of the interaction with intense and non axisymmetric magnetic fields. Consequently, in our simulations, the stellar wind structure in the equatorial plane of young stars varies significantly from a solar configuration, delivering insight about the past of the solar system interplanetary medium. | \label{intro} Among all the stellar properties, the characteristics of solar-like stars' winds are probably the most difficult to probe. Emissions throughout the electromagnetic spectrum unveil some of the properties of the photosphere and the coronae of stars, and internal structures can be probed with asteroseismology. Winds, on the contrary, produce very few detectable signals, although they are likely to exist in all stars possessing a hot corona, as supersonic outflows are the only stable way to balance coronal pressure with the near zero interstellar medium pressure \citep{Parker1958,Velli1994}. \citet{LinskyWood1996} have shown that absorption by neutral hydrogen at the astropause could be detected in Ly{$\alpha$} spectra around astrospheres of nearby solar-like stars, unraveling properties of the stellar wind shocking against the interstellar medium. A growing sample of solar-type stars with positive detection for stellar winds led to a relationship between X-ray fluxes originating from coronal loops and mass loss rates \citep{Wood2002}. The ``strength" of stellar winds, the mass loss rate $\dot{M}$, has consequently been related to the magnetic activity of the star. \citet{Wood2005a} have obtained the relation: $\dot{M} \propto F_X^{1.34 \pm 0.18}$, for $F_X \leq 10^6$ ergs cm$^{-2}$ s$^{-1}$, where $F_X$ is the X-ray flux. Beyond this value, weaker mass loss rates are observed, suggesting a saturation effect, that is below the usual $F_X$ saturation value \citep{Randich2000,Pizzolato2003,Gudel2004}. In parallel, the development of Zeeman Doppler Imaging (ZDI) \citep{Semel1989,DonatiBrown1997,PiskunovKochukhov2002} has opened the study of surface magnetic fields for cool stars. Large scale magnetic fields have been detected in the whole mass range that is thought to correspond to the existence of a convective envelope ($0.1 M_{\odot} - 1.4 M_{\odot}$). The study of the geometrical and topological properties of the field in the context of stellar evolution is still in progress \citep{DonatiLandstreet2009,See2015} and raises theoretical questions about their generation through dynamo processes in convective envelopes \citep[see][and references therein]{Brun2015}. Nonetheless, the magnetic field amplitude of these stars has been shown to be a growing function of the rotation rate \citep{Noyes1984,BrandenburgSaar2000,Vidotto2014b}. This is necessary to explain the rotational braking of cool main sequence dwarfs, as evolutionary models need the wind to carry angular momentum at a rate proportional to $\Omega_{\star}^3$ \citep{Kawaler1988,Bouvier1997,Matt2015} all along the main sequence. However, recent studies suggest that the wind braking could stop or strongly decay for evolved stars, around a solar Rossby number $Ro \approx 2.5$ \citep{vanSaders2016}, perhaps because of a change in magnetic topology \citep{Reville2015a,Garraffo2015a}. Hence wind, magnetism, and rotation are likely to evolve coherently through the whole life of solar-like stars. After \citet{Schatzman1962} understood that a magnetized outflow would carry away stellar angular momentum, \citet{WeberDavis1967} demonstrated that this loss is proportional to the Alfv\'en radius squared. Several studies have followed to try to estimate the Alfv\'en radius from stellar parameters and thus give scaling laws for the angular momentum loss \citep{Mestel1968,Kawaler1988}. The latest braking laws have been successfully introduced in stellar evolution models \citep{Matt2012,GalletBouvier2013}, and we recently demonstrated that the magnetic topology could be included in those formulations through a simple scalar parameter, the open magnetic flux \citep{Reville2015a,Reville2015b}. Most studies \citep[see, e.g.,][]{Matt2012,Reville2015a} have been made in two dimensions with axisymmetric configurations \citep[see][for a 3D study of angular momentum loss with idealized magnetic field topologies]{Cohen2014,Garraffo2015a}, and were not able to capture the structure of complex magnetic fields observed by ZDI. 3D MHD simulations are now taking into account this complexity \citep{Cohen2011,Vidotto2014a,doNascimento2016,AlvaradoGomez2016a,AlvaradoGomez2016b} to derive a self-consistent coronal structure. The complex structure of the corona is needed to study the interaction between stars and close-in planets, which has been shown to be very sensitive to 3D effects \citep{Strugarek2015}. Yet, to our knowledge, the influence of realistic magnetic fields on the long time variation of the wind properties has not been studied. This work proposes to include observed, realistic magnetic fields in the context of stellar evolution. We used spectropolarimetric observations of the surface fields of solar-like stars to constrain 3D MHD simulations of stellar winds. The stars of our sample share similar properties except their rotational periods and their ages, which range from 25 Myr to 4.5 Gyr. We developed a coherent framework to take into account the evolution of the coronal properties with time, inspired by X-ray flux observations, spin evolution models and theoretical, ab initio models \citep[see][]{HolzwarthJardine2007,CranmerSaar2011,Suzuki2013b}. We confirm that the evolution with age of global properties of the wind, such as the mass and angular momentum loss, follows simple prescriptions in agreement with the spin evolution models. These prescriptions can be recovered by the semi-analytical model we developed in \citet{Reville2015b}, except for the saturation of angular momentum that appears only in our simulations. Also, the three dimensional structure of the young stars' winds shows interesting features, such as a trimodal speed distribution, that we explain through various interactions with the intense magnetic field. We show that superradial expansion is a key factor to explain the fastest wind components of young stars' wind. We also observe regions of fast wind encountering slower streams in the equatorial plane, the so-called Corotating Interactions Regions (CIRs), that could be more common in the wind of young stars. This paper is organized as follows: the ingredients of our model are described throughout Section \ref{sec:model}. In Subsection \ref{subsec:zdi}, we describe the observations used to constrain the surface magnetic fields of our simulations. Subsection \ref{subsec:evol} describes our prescriptions for the coronal properties and Subsection \ref{subsec:num} our numerical setup. The results are presented in two parts, Section \ref{sec:glob} where we look at global properties such as the angular momentum and mass loss over time, and Section \ref{sec:vel} where the tridimensional structure of the wind is detailed, with a special focus on the velocity distribution. We summarize and reflect upon our findings in Section \ref{sec:ccl}. | \label{sec:ccl} Our study addresses 3D simulations constrained by spectropolarimetric observations of magnetic fields in the context of stellar evolution, along the main sequence. Our findings can be divided into two parts. First, considering the global and integrated properties, the mass and angular momentum loss follow simple prescriptions thanks to an appropriate modeling of the evolution of the temperature and the coronal density with the rotation rate. An angular momentum proportional to the $\Omega_{\star}^3$ can be obtained using our prescriptions with observed magnetic field amplitudes, which is required to observe the convergence of spin rates on the Skumanich law. Our 3D simulations follow the braking law we derived in \citet{Reville2015a} if the $K_3$ constant is reduced by $15\%$. This can be understood because our simulations are made with a larger (and different for each case) coronal temperature compared to this previous work. Hence, the wind described is here faster, and for a given magnetic field strength, the Alfv\'en radius is closer to the star. However, the fit shows little deviation, and one constant $K_3$ in the scaling law is able to describe the whole range of temperature of our sample. A more general braking law should quantify the influence of the temperature on $K_3$, since the variation remain limited in this study. Also, the use of a fully 3D geometry could be involved in the variation of $K_3$. Nonetheless, because this formulation expresses the dependence of the Alfv\'en radius on a magnetization parameter that includes the thermodynamics of the wind, it is likely to be valid for a wide range of magnetic fields, rotation rates, coronal temperatures and densities, if one allows a small dependence of $K_3$ on the temperature. With this adapted formulation, the semi-analytical we developed in \citet{Reville2015b} can be applied. The estimation method of the open flux of the simulation with a potential extrapolation that was tested on 2.5D configurations is perfectly operational with 3D non-axisymmetric fields. Our semi-analytical model is consequently able to estimate closely the evolution of angular momentum with the rotational period, as long as the mass loss rate of the simulations does not deviate too much from the spherically symmetric value used in the semi-analytical model. This deviation grows as the stars rotate faster and possess more intense magnetic fields. Large coronal loops are able to confine more plasma, and the mass loss seems to plateau for $\Omega_{\star} > 8 \Omega_{\odot}$. This behavior has consequences on the angular momentum loss that shows signs of saturation beyond this rotation rate, whose value is coherent with the saturation value used in rotation evolution models for K-type stars \citep{GalletBouvier2015}. Although the saturation of angular momentum loss is often associated with the saturation observed in the X-ray fluxes, the precise process behind this saturation remains unknown. Some works have suggested a stochastic change of the dynamo process generating the magnetic field could be involved in a topology switch from small scales to large scales that turns on the $\Omega_{\star}^3$ braking law \citep{Barnes2003,Brown2014,Garraffo2015b}. However, no such transition is observed in our sample, as all our stars possess a strong dipolar field. In our simulations, the AML saturation seems to be due to the confinement of the outflow in large coronal loops that reduces the mass loss, which can be associated with the dependence of the wind braking on the filling factor \citep[see][]{CranmerSaar2011,GalletBouvier2013,GalletBouvier2015}. These results need, however, to be confirmed by more simulations of fast rotators and are likely to be highly dependent on the prescriptions we used for the coronal base densities and temperature. The mass loss rate of young stars in our study, although up to $6$ to $9$ times the solar one \footnote{In table \ref{table2}, we have considered an upper value of the solar mass loss rate $\dot{M} = 3 \times 10^{-14} M_{\odot}/$yr.}, does not reach the highest values derived in \citet{Wood2005a}. A much higher dependence on rotation, for either the temperature or the coronal density would have been necessary to observe $100 \dot{M}_{\odot}$ values in our simulations. Our semi analytical model could be used to study the influence of different prescriptions on the variation of $\dot{J}$. Change in the exponents of the evolution laws, or the solar initial values, could be tested. A more physical description of the stellar wind acceleration, driven for example by Alfv\'en waves turbulence \citep[see][]{Suzuki2006,MatsumotoSuzuki2012,Suzuki2013a,Sokolov2013,Oran2013,Lionello2014} can help to understand how coronal parameters evolve with age. Future works will be dedicated to including such processes in our simulations. To solve this issue, more observations are also critical. In the work of \citet{Wood2005a} \citep[see also][]{WoodRev2004}, the analysis of the Ly$\alpha$ absorption spectra is coupled with numerical simulations of the terminal shock, where the wind speed is kept constant at the slow solar wind value around $400$ km/s. Our study, and this is the second part of our findings, brings more accurate constraints on the wind velocity amplitudes and distribution for solar-like stars that could be used to improve those calculations. Indeed, several studies have now shown that the wind speeds are likely to increase for young stars, mostly because of higher coronal temperature and magneto-centrifugal acceleration \citep{WashShib1993,HolzwarthJardine2007,Matt2012,Suzuki2013b,Reville2015a,Reville2015b}. We show that the speed distribution of young and active stars follows a trimodal structure due to the interaction with the magnetic field. Moreover, the one-dimensional magneto-centrifugal wind solution \citep{WeberDavis1967,Sakurai1985} is at best a low estimate of the fastest component of the wind, as other magnetic processes are able to accelerate the wind. For instance, we observe in our simulations fast streams in the vicinity of the star caused by latitudinal or longitudinal superradial expansion of flux tubes due to the fast rotation and the non-axisymmetry of the magnetic field. As slow and fast winds exist around solar-like stars, they can interact in the equatorial plane, creating Corotating Interaction Regions, which could be more common for younger stars and could be forming shocks closer to the star (the usual distance observed in the solar system is between $1$ and $5$ AU). This could have consequences on the energetics of expanding stellar winds of young stars and on exoplanetary space weather. The current sheet of young stars is also strongly corrugated when the dipole is inclined and polarity variations occur closer to one another, which has important consequences on the interaction with close-in planets, especially when they are within the Alfv\'en surface \citep{Strugarek2015}. This is, however, a very preliminary study and a more focused work needs to be done in that sense, improving, for instance, the numerical scheme and shock absorbing method. | 16 | 9 | 1609.06602 |
1609 | 1609.06328_arXiv.txt | The radio and gamma-ray pulsar \psr\ was recently found to be in a decades-long orbit with the Be star \mttwo, with the pulsar moving rapidly towards periastron. This binary shares many similar characteristics with the previously unique binary system \psrb/LS~2883. Here, we describe radio, X-ray, and optical monitoring of \psr/\mttwo. Our extended orbital phase coverage in radio, supplemented with \Fermi\ gamma-ray data, allows us to update and refine the orbital period to 45--50~yr and time of periastron passage to 2017 November. We analyse archival and recent \Chandra\ and \Swift\ observations and show that \psr/\mttwo\ is now brighter in X-rays by a factor of $\sim 70$ since 2002 and $\sim 20$ since 2010. While the pulsar is still far from periastron, this increase in X-rays is possibly due to collisions between pulsar and Be star winds. Optical observations of the H$\alpha$ emission line of the Be star suggest that the size of its circumstellar disc may be varying by $\sim 2$ over timescales as short as 1--2~months. Multiwavelength monitoring of \psr/\mttwo\ will continue through periastron passage, and the system should present an interesting test case and comparison to \psrb/LS~2883. | \label{sec:intro} Recent radio observations of the 143~ms pulsar \psr\ find it to be part of a very eccentric, long orbital period binary system, with the pulsar expected to reach periastron in late 2017 with its high-mass, Be star companion \citep{lyneetal15}. The pulsar was discovered by \Fermi\ \citep{abdoetal09} and is associated with TeV source TeV~J2032+4130 \citep{camiloetal09}. These characteristics make the \psr\ binary system very similar to the previously unique pulsar system \psrb/LS~2883 (see, e.g. \citealt{dubus13}, for review), and X-ray results reported here support this similarity. Subsequent to its discovery in gamma rays, \psr\ was detected in radio by \citet{camiloetal09}, who also reanalysed a 49~ks \Chandra\ observation taken in 2004 of the field of the Cygnus OB2 association (at a distance $d=1.33\pm0.06\mbox{ kpc}$; \citealt{kiminkietal15}), to which the pulsar likely belongs. They confirm an X-ray source at the position of the radio source, which also corresponds to optical source 213 of \citet{masseythompson91}, a B0~V star (hereafter \mttwo). \mttwo\ has a mass of either $14.5\,M_{\sun}$ \citep{wrightetal15} or $17.5\,M_{\sun}$ \citep{kiminkietal07} and bolometric luminosity $L_{\rm bol}=1.51\times 10^4\,L_{\sun}=5.79\times 10^{37}\mbox{ erg s$^{-1}$}$ \citep{wrightetal15}, in broad agreement with values appropriate to its spectral and luminosity classification, i.e. $15.0\pm2.8\,M_{\sun}$ and $16100\pm130\,L_{\sun}$ \citep{hohleetal10}. An optical spectrum shows \mttwo\ has a H$\alpha$ equivalent width (\EW) of $-12.6$~\AA\ \citep{camiloetal09}, which is typical of Be stars and due to a circumstellar disc surrounding the star. \citet{camiloetal09} fit a power-law to the \Chandra\ spectrum and find an unabsorbed 0.5--10~keV X-ray luminosity $L_{\mathrm X}\approx 6\times 10^{30}\mbox{ erg s$^{-1}$}(d/\mbox{1.3 kpc})^2$. This X-ray luminosity is compatible with that of either Be stars (see, e.g. \citealt{berghoferetal97,gagneetal11,nazeetal14}) or pulsars of age $\sim 10^5\mbox{ yr}$ (see, e.g. \citealt{yakovlevpethick04,pageetal06,potekhinetal15}), where the pulsar age is taken to be its characteristic spin-down age. More recent analysis of timing observations of \psr\ reveals that its timing noise can be removed by considering a timing model in which the pulsar is in an eccentric (with eccentricity $\epsilon>0.94$), decades-long orbit \citep{lyneetal15}. Because of the long orbital period, radio measurements up to that time only cover about 20\% of the orbit, and previous observations (at all wavelengths) of the pulsar/Be-star binary system have been when the pulsar is on the apastron-side of the orbit. Fig.~\ref{fig:orbit} shows a schematic diagram of the system. Radio and gamma-ray telescopes continue to monitor the pulsar as it moves towards periastron. These observations will refine the orbital parameters, including eccentricity and mass function, and could, along with VLBI measurements, directly determine distance and orbital inclination \citep{lyneetal15}. \begin{figure} \includegraphics[width=0.9\columnwidth,angle=270,trim={0 4.5cm 0 3.8cm},clip]{f1.ps} \caption{Schematic diagram illustrating the approximate orbital motion of \psr\ and its Be-star companion \mttwo\ about their common centre of mass. The orbit shown is that of model~2 (see Table~\ref{tab:spinpars}), which has an orbital period $P_{\rm b}= 17000$~d and is projected on to the plane containing the line of sight and major axis of the orbit. The inclination $i$ of the plane of the orbit to the plane of the sky is assumed to be 60$^\circ$. The circles mark 200~d intervals and indicate time from predicted epoch of periastron, MJD~58069. The pulsar moves counter-clockwise in the diagram and has been approaching the Be star since discovery in late 2008/early 2009. The thick line shows the portion of the pulsar orbit covered by radio observations reported here, MJD~54689--57538. Note that orbital velocity is proportional to the separation between the circles, with a 1000-light-second separation indicating a velocity of about 18~km~s$^{-1}$. The small ellipse near the origin shows the orbit of the Be star, assuming that it has a mass of $15\,M_{\sun}$ and that the pulsar has a mass of $1.35\,M_{\sun}$.} \label{fig:orbit} \end{figure} Many of the above characteristics are typical of Be X-ray binary systems, albeit with orbital periods of $\lesssim 1\mbox{ yr}$, which can shine at up to Eddington luminosities ($10^{38}\mbox{ erg s$^{-1}$}$) when the neutron star nears periastron and accretes matter from the circumstellar disc (of size a few times the Be star radius, and larger for isolated Be stars; \citealt{klusetal14,reigetal16}) of its Be star companion \citep{reig11}. In this work, we are concerned with X-ray (as well as radio and optical) emission when the pulsar is far from periastron and not accreting from the circumstellar disc of the Be star, which might occur near periastron. Thus the X-ray luminosity is expected to be much lower, as indeed we find, as well as a brightening that seems in accord with the well-studied 3.4~yr orbital period gamma-ray binary that contains \psrb. There are several previous \Chandra\ observations of Cygnus~OB2 (see Table~\ref{tab:obsx}), and although significantly off-axis, some of these contain the pulsar/Be-star binary system in the field of view. As discussed above, \citet{camiloetal09} describe results for the 2004 observation (ObsID~4501), while \citet{rauwetal15} describe results for the 2010 observations (ObsID~10944, 10945, 10951, 10962). In Section~\ref{sec:radio}, we report results of an updated timing solution based on monitoring at radio wavelengths, supplemented with \Fermi\ gamma-ray data. We (re)analyse all \Chandra\ and \Swift\ X-ray data, as well as our 2016 4.9~ks \Chandra\ Target of Opportunity observation, and report our results in Section~\ref{sec:xray}. In Section~\ref{sec:optical}, we report on recent optical measurements of the H$\alpha$ \EW\ and their implication for the size of the Be star circumstellar disc. In Section~\ref{sec:discuss}, we summarize and briefly discuss our findings, including a few comparisons with \psrb/LS~2883. | \label{sec:discuss} In this work, we described recent multiwavelength observations of the high-energy binary system containing the 143~ms radio pulsar \psr\ and the B0~Ve companion star \mttwo. The orbit is very eccentric ($0.94<\epsilon<0.99$) and large ($P_{\rm orb}\approx 45-50$~yr), and the pulsar is accelerating rapidly towards periastron passage in 2017 November. We updated orbital parameters of the system obtained via radio monitoring, combined with \Fermi\ observations, of the pulsar. Archival and recent \Chandra\ and \Swift\ observations show that the \psr/\mttwo\ system has brightened significantly in X-rays, especially within the last year as it approaches periastron. We can understand the current and past behavior of the \psr/\mttwo\ system by comparing it to the high-energy binary pulsar system \psrb/LS~2883, since the two systems have many similarities. \psrb\ is a 47.76~ms radio pulsar in an eccentric ($\epsilon\approx 0.87$) 3.4~yr orbit with LS~2883, which is a 09.5~Ve star \citep{johnstonetal92,negueruelaetal11}. \psrb/LS~2883 has been observed across the electromagnetic spectrum during each periastron passage since its discovery (see, e.g. \citealt{chernyakovaetal15}, for discussion of the most recent passage in 2014), as well as around apastron (see, e.g. \citealt{hirayamaetal99}). The resulting studies show that its multiwavelength emission can be understood as shock interaction between the relativistic wind emitted by the pulsar and the circumstellar disc and wind of the companion star \citep{tavaniarons97}. However, in attempting to extrapolate the observed behavior of \psrb\ to that of \psr, it is important to note that the binary separation at apastron is $\sim 11$~AU in the case of \psrb, while \psr\ has only been observed up to this point at a binary separation $>10$~AU. At this large distance, the wind from \mttwo\ is possibly tenuous, and its collision with the pulsar wind is possibly weak. Also LS~2883 is four times more luminous and a different stellar type \citep{negueruelaetal11}, so its wind properties are likely different from that of \mttwo. For \Chandra\ observations taken in 2002, 2004, and around 2010, \psr\ was quite distant from \mttwo, at a binary separation of $\gtrsim 30$~AU (see Fig.~\ref{fig:orbit}). Thus the observed X-ray emission [with $L_{\rm X}\approx(0.05-0.2)\times 10^{32}\mbox{ erg s$^{-1}$} (d/1.3\mbox{ kpc})^2$] could be due to what is effectively an isolated Be star or an isolated young pulsar. In the former case, bright X-ray emission from powerful wind shocks of O stars is quite common, and spectra are often fit with an optically thin thermal plasma model (as performed in Section~\ref{sec:chandra}); the source of X-ray emission from stars of later stellar types is less certain, with a transition around early B stars, like \mttwo, that have $L_{\rm X}\sim 10^{30}-10^{32}\mbox{ erg s$^{-1}$}$ \citep{berghoferetal97,gagneetal11,nazeetal14}. Using the relation between X-ray and bolometric luminosities, $\log L_{\rm X}/L_{\rm bol}\approx -7.2$, found for O and bright B stars (although there is large dispersion at the luminosity of stars similar to \mttwo; \citealt{rauwetal15}), we find $L_{\rm X}=4\times 10^{30}\mbox{ erg s$^{-1}$}$, which matches the X-ray luminosity of \mttwo\ in 2002. For isolated pulsars, X-ray radiation can have non-thermal and thermal contributions. Non-thermal emission can be generated by a relativistic wind, which produces a ratio between X-ray luminosity to rotational energy loss of $L_{\rm X}/\dot{E}\lesssim 10^{-3}$ and a spectrum that is best-fit by a power-law with $\Gamma\approx 1-3$ \citep{becker09}. For \psr, $\dot{E}=1.5\times 10^{35}\mbox{ erg s$^{-1}$}$, and the measured power law is $\Gamma\approx 1.5-2.5$ (Table~\ref{tab:spectrafit}). Thus a pulsar wind can easily be the source of observed X-rays. Meanwhile, thermal emission for intermediate age ($\sim 10^5$~yr) neutron stars gives $L_{\rm X}\sim 10^{31}-10^{33}\mbox{ erg s$^{-1}$}$ (see, e.g. \citealt{potekhinetal15}). The more recent observations since late 2015 show significant brightening in X-rays (see Fig.~\ref{fig:j2032lc}), with $L_{\rm X}\approx(0.6-2.9)\times10^{32}\mbox{ erg s$^{-1}$} (d/1.3\mbox{ kpc})^2$. Although these luminosities are somewhat lower than the X-ray luminosity seen for \psrb\ at apastron ($L_{\rm X}\approx 5\times 10^{32}\mbox{ erg s$^{-1}$}$; \citealt{hirayamaetal99,uchiyamaetal09}), the brightening of \psr\ suggests that the pulsar has entered the regime where the pulsar wind is interacting strongly with the Be star wind. X-ray spectral studies of \psrb\ find that most observations made by \Chandra, \textit{Suzaku}, \Swift, and \textit{XMM-Newton} can be fit with a power law model (see, e.g. \citealt{tavaniarons97,chernyakovaetal09,chernyakovaetal14}). \citet{chernyakovaetal06,chernyakovaetal09} and \citet{uchiyamaetal09} show that the value of $\Gamma$ varies with orbital phase and that $\Gamma$ changes from 1.8 around apastron to 1.2 right before the pulsar enters the circumstellar disc of LS~2883 at periastron, a decline that is similar to what is found from our limited spectra of \psr\ (see Table~\ref{tab:spectrafit}). The \psr/\mttwo\ system will continue to be monitored across the electromagnetic spectrum as the pulsar approaches periastron. When it is near periastron, the system may brighten even more if the pulsar accretes from the circumstellar disc of the Be star. \psr/\mttwo\ will thus serve as an invaluable tool for comparing and contrasting to the very well-studied and previously unique gamma-ray binary \psrb/LS~2883. | 16 | 9 | 1609.06328 |
1609 | 1609.04788_arXiv.txt | We conduct an analysis of the Planck 2015 data that is complete in reionization observables from the large angle polarization $E$-mode spectrum in the redshift range $6 < z < 30$. Based on 5 principal components, all of which are constrained by the data, this single analysis can be used to infer constraints on any model for reionization in the same range; we develop an effective likelihood approach for applying these constraints to models. By allowing for an arbitrary ionization history, this technique tests the robustness of inferences on the total optical depth from the usual steplike transition assumption, which is important for the interpretation of many other cosmological parameters such as the dark energy and neutrino mass. The Planck 2015 data not only allow a high redshift $z>15$ component to the optical depth but prefer it at the $2\sigma$ level. This preference is associated with excess power in the multipole range $10 \lesssim \ell \lesssim 20$ and may indicate high redshift ionization sources or unaccounted for systematics and foregrounds in the 2015 data. | \label{sec:intro} The epoch of reionization of the Universe remains one of the least well-understood aspects of the standard model of cosmology (see e.g.~\cite{2016ASSL..423.....M}). Yet its impact on the interpretation of its fundamental properties is comparatively large in this era of precision cosmology. In addition to the intrinsic astrophysical interest in ionization sources, reionization uncertainties impact cosmic microwave background (CMB) inferences on the initial power spectrum and hence also cosmic acceleration through the growth of structure \cite{Hu:2003pt}. In the future it will be one of the leading sources of error in the interpretation of neutrino mass measurements from gravitational lensing \cite{Smith:2006nk,Allison:2015qca}, the study of large scale anomalies in the CMB \cite{Mortonson:2009xk,Mortonson:2009qv}, and the inflationary consistency relation \cite{Mortonson:2007tb}. The standard approach to parametrizing the impact of reionization on the CMB is with the total Thomson optical depth through the reionization epoch. Although it is indeed the total optical depth that is important for the interpretation of most other aspects of cosmology, it is usually assumed that reionization occurs promptly in a steplike transition. Interestingly, the central value of the inferred optical depth from this approach has been steadily drifting downwards from its first detection in WMAP1 \cite{Kogut:2003et} to the current but still proprietary Planck 2016 High Frequency Instrument (HFI) results \cite{Aghanim:2016yuo,Adam:2016hgk}. Relaxing this sharp transition assumption can in principle raise the optical depth inference from the CMB as well as change its implications for sources of high redshift ionization (e.g.~\cite{2012ApJ...756L..16A}). In particular, the angular scale of the peak and the width of the reionization bump in the $E$-mode CMB polarization power spectrum carries coarse grained information on the redshift dependence of the ionization history. There is an alternate, model independent approach introduced in Ref.~\cite{Hu:2003gh} that fully addresses these concerns. The impact of {\it any} ionization history on the large angle CMB polarization spectrum can be completely characterized by a handful of reionization parameters constructed from the principal components (PCs) of the ionization history with respect to the $E$-mode power spectrum. This approach has the advantage over redshift binning alternatives of being observationally complete without introducing numerous highly correlated parameters \cite{Lewis:2006ym}. Conversely, it does not provide an accurate, local reconstruction for visualizing the ionization history itself. This approach was implemented and tested on WMAP3 \cite{Mortonson:2007hq} and WMAP5 \cite{Mortonson:2008rx} power spectra which showed that those data allowed for, but did not particularly favor, contributions to the optical depth from high redshift. It was adopted in the Planck 2013 analysis but exclusively to test the impact of marginalizing the ionization history on inflationary parameters rather than drawing inferences on reionization itself \cite{Planck:2013jfk}. The impact on massive neutrinos and gravitational wave inferences was also examined in Ref.~\cite{Dai:2015dwa}. In this work, we analyze the public Planck 2015 data, including the Low Frequency Instrument (LFI) large angle polarization power spectrum, using the observationally complete PC basis. In addition we further develop the technique as a method to probe reionization itself. This development is timely as the technique should come to its full fruition with the upcoming final release of Planck data which will be the definitive result on large angle polarization for years to come. We demonstrate that the Planck 2015 data already have more information on the ionization history than just the total optical depth and provide effective likelihood tools for interpreting this information in any given model for reionization within the redshift range analyzed. Tested here, these techniques can be straightforwardly applied to the final release when it becomes available. We begin with a review of the approach itself in Sec. \ref{sec:PC}. In Sec. \ref{sec:MCMC}, we analyze the Planck 2015 data and explore the origin of the additional information on the high redshift ionization history. We develop and test an effective likelihood approach for utilizing our analysis to constrain the parameters of any given model of reionization from $6 < z <30$ in Sec. \ref{sec:likelihood}. We discuss these results in Sec. \ref{sec:discussion}. | \label{sec:discussion} By analyzing the Planck 2015 data with an observationally complete PC basis for the ionization history, we show that it allows and even favors high redshift, $z\gtrsim 15$, optical depth at the $\sim 2\sigma$ level. The standard analysis which includes just the total optical depth and assumes a sharp steplike transition excludes this possibility by prior assumption of form rather than because it is required by the data. The same is true for power-law models that additionally vary the duration of reionization. While a $2\sigma$ result amongst the 5 PC parameters is not on its own surprising, it originates from the first and best constrained component and hence has consequences for the total optical depth. The total optical depth is important for understanding a host of other cosmological parameters from the amplitude of the current matter power spectrum $\sigma_8$ to the inferences from CMB lensing. At the very least, this analysis highlights the need for a complete treatment of CMB reionization observables to guarantee a robust interpretation of the optical depth. This preference for extra high redshift optical depth mainly originates from the large angle polarization spectrum in the Planck 2015 data and appears related to excess power in the multipole range $10 \lesssim \ell \lesssim 20$. It is only slightly weakened by marginalizing gravitational lensing information in the temperature power spectrum, which is known to favor a higher optical depth, but more significantly changed by replacing the Planck LFI with WMAP9 polarization data. While excess polarization power in this range favors additional sources of high redshift ionization such as population III stars or dark matter annihilation it could also indicate contamination from systematics and foregrounds. The latter have been significantly improved in the as yet proprietary Planck 2016 intermediate results. These results indicate that the low redshift end of the optical depth as tested by steplike models, or equivalently the low $\ell$ polarization power, is both better measured and lower than the central value in the Planck 2015 data \cite{Aghanim:2016yuo}. On the other hand, these results exacerbate the tension with gravitational lensing in the shape of the temperature power spectrum which probes the total optical depth. It will be interesting to see if this complete analysis still prefers an additional high redshift component in the final Planck release. Regardless of the outcome of resolving the mild tension between steplike reionization scenarios and the Planck 2015 data, the complete PC approach developed here is useful because with a single analysis one can infer constraints on the parameters of any reionization model within the specified redshift range, here $6 < z <30$, but easily extensible to any desired range. What has presented an obstacle for this approach in the past is the lack of tools for converting posterior parameter constraints on the PCs to parameter constraints on models and so be able to combine them with other sources of reionization information. For example, the ionization history can also be tested in the CMB through the kinetic Sunyaev-Zel'dovich effect from temperature fluctuations beyond the damping scale, but in a manner that is highly model dependent (e.g. \cite{Mortonson:2010mi,Zahn:2011vp,Battaglia:2012im,Adam:2016hgk}). Towards this end, we have developed and tested an effective likelihood code for inferring constraints on any given ionization history provided by a model. This approach should be especially useful in constraining models where small high redshift contributions to the optical depth need to be separated from the total. \medskip \noindent {\it Acknowledgments}: We thank Austin Joyce, Adam Lidz, and Pavel Motloch for useful discussions. W.H. thanks the Aspen Center for Physics, which is supported by National Science Foundation Grant PHY-1066293, where part of this work was completed. C.H. and W.H. were supported by NASA ATP NNX15AK22G, U.S.~Dept.\ of Energy Contract No. DE-FG02-13ER41958, and the Kavli Institute for Cosmological Physics at the University of Chicago through Grants No. NSF PHY-0114422 and No. NSF PHY-0551142. Computing resources were provided by the University of Chicago Research Computing Center. V.M. was supported in part by the Charles E.~Kaufman Foundation, a supporting organization of the Pittsburgh Foundation. \vfill | 16 | 9 | 1609.04788 |
1609 | 1609.01352_arXiv.txt | Many extremely low-mass (ELM) white-dwarf (WD) stars are currently being found in the field of the Milky Way. Some of these stars exhibit long-period nonradial $g$-mode pulsations, and constitute the class of ELMV pulsating WDs. In addition, several low-mass pre-WDs, which could be precursors of ELM WDs, have been observed to show short-period photometric variations likely due to nonradial $p$ modes and radial modes. They could constitute a new class of pulsating low-mass pre-WD stars, the pre-ELMV stars. Here, we present the recent results of a thorough theoretical study of the nonadiabatic pulsation properties of low-mass He-core WDs and pre-WDs on the basis of fully evolutionary models representative of these stars. | An increasing number of low-mass WDs, including ELM WDs ($M_{\star}\sim 0.18-0.20 M_{\sun}$, H-rich atmospheres), are currently being detected through the ELM survey (see Brown et al. 2016 and references therein). These WD stars, which likely harbor cores made of He, are thought to be the result of strong mass-loss events at the red giant branch stage of low-mass stars in binary systems before the He flash onset that, in this way, is avoided (Althaus et al. 2013; Istrate et al. 2016). The increasing interest in ELM WDs has lead to the discovery of long-period ($\Pi \sim 1000-6300$ s) $g$-mode pulsations in some of them (ELMVs). The existence of ELMV stars ($7000 \lesssim T_{\rm eff} \lesssim 10\,000$ K and $6 \lesssim \log g \lesssim 7$; red circles in Fig. \ref{Fig1}) constitutes an unprecedented opportunity for probing their subsurface layers and ultimately to place constraints on the currently accepted formation scenarios by means of asteroseismology (Winget \& Kepler 2008; Fontaine \& Brassard 2008; Althaus et al. 2010). Apart from ELMVs, short-period ($\Pi \sim 300-800$ s) pulsations in five objects that are probably precursors of ELM WDs have been detected in the few last years. These stars have typically $8000 \lesssim T_{\rm eff} \lesssim 13\,000$ K and $4 \lesssim \log g \lesssim 5.5$ (green circles in Fig. \ref{Fig1}) and show a surface composition made of H and He. They are called pre-ELMV stars and constitute a new class of pulsating stars. Also, the discovery of long-period ($\Pi \sim 1600-4700$ s) pulsations in three additional objects located at the same region of the HR diagram has been reported. These stars are emphasized with black squares surrounding the green circles in Fig. \ref{Fig1}. The nature of these objects is uncertain and could be identified as pre-ELM stars as well as SX Phe and/or $\delta$ Scuti pulsating stars. In Table \ref{tabla1} we include an updated compilation of the effective temperature, gravity and range of observed periods for all the known pre-ELMV and ELMV stars. \articlefigure{Corsico_A_Fig1.eps}{Fig1}{The location of the known ELMVs (red circles) and pre-ELMVs (light green circles) along with the other several classes of pulsating WD stars (dots of different colors) in the $\log T_{\rm eff} - \log g$ plane. The three stars emphasized with squares surrounding the light green circles can be identified as pre-ELMV stars as well as SX Phe and/or $\delta$ Scuti stars. In parenthesis we include the number of known members of each class. Two post-VLTP (Very Late Thermal Pulse) evolutionary tracks for H-deficient WDs and two evolutionary tracks for low-mass He-core WDs are plotted for reference. Dashed lines indicate the theoretical blue edge for the different classes of pulsating WDs.} \begin{table}[!ht] \caption{Stellar parameters and observed period range of the known pre-ELMV (upper half of the table) and ELMV (lower half of the table) stars. For ELMVs, the $T_{\rm eff}$ and $\log g$ values are computed with 1D model atmospheres after 3D corrections.} \smallskip \begin{center} {\small \begin{tabular}{lccccr} \tableline \noalign{\smallskip} Star & $T_{\rm eff}$ & $\log g$ & $M_{\star}$ & Period range & Ref.\\ & [K] & [cgs] & [$M_{\sun}$] & [s] & \\ \noalign{\smallskip} \tableline \noalign{\smallskip} SDSS J115734.46+054645.6 & $11\,870\pm260$ & $4.81\pm0.13$ & $0.186$ & $364$ & (3) \\ SDSS J075610.71+670424.7 & $11\,640\pm250$ & $4.90\pm0.14$ & $0.181$ & $521-587$ & (3) \\ WASP J024743.37$-$251549.2 & $11\,380\pm400$ & $4.576\pm0.011$ & $0.186$ & $380-420$ & (1) \\ SDSS J114155.56+385003.0 & $11\,290\pm210$ & $4.94\pm0.10$ & $0.177$ & $325-368$ & (3) \\ KIC 9164561(*) & $10\,650\pm200$ & $4.86\pm0.04$ & $0.213$ & $3018-4668$ & (5) \\ WASP J162842.31+101416.7 & $9200\pm600$ & $4.49\pm0.05$ & $0.135$ & $668-755$ & (2) \\ SDSS J173001.94+070600.25(*) & $7972\pm200$ & $4.25\pm0.5$ & $0.171$ & $3367$ & (4) \\ SDSS J145847.02+070754.46(*) & $7925\pm200$ & $4.25\pm0.5$ & $0.171$ & $1634-3279$ & (4) \\ \tableline \noalign{\smallskip} SDSS J222859.93+362359.6 & $7890\pm120$ & $5.78\pm0.08$ & $0.142$ & $3254-6235$ & (6) \\ SDSS J161431.28+191219.4 & $8700\pm170$ & $6.32\pm0.13$ & $0.172$ & $1184-1263$ & (6) \\ PSR J173853.96+033310.8 & $8910\pm150$ & $6.30\pm0.10$ & $0.172$ & $1788-3057$ & (8) \\ SDSS J161831.69+385415.15 & $8965\pm120$ & $6.54\pm0.14$ & $0.179$ & $2543-6125$ & (9) \\ SDSS J184037.78+642312.3 & $9120\pm140$ & $6.34\pm0.05$ & $0.177$ & $2094-4890$ & (10) \\ SDSS J111215.82+111745.0 & $9240\pm140$ & $6.17\pm0.06$ & $0.169$ & $108-2855$ & (7) \\ SDSS J151826.68+065813.2 & $9650\pm140$ & $6.68\pm0.05$ & $0.197$ & $1335-3848$ & (7) \\ \noalign{\smallskip} \tableline \end{tabular} } {\footnotesize References: (*) Not secure identification as pre-ELM WD (see text); (1) Maxted et al. (2013); (2) Maxted et al. (2014); (3) Gianninas et al. (2016); (4) Corti et al. (2016); (5) Zhang et al. (2016); (6) Hermes et al. (2013b); (7) Hermes et al. (2013a); (8) Kilic et al. (2015); (9) Bell et al. (2015); (10) Hermes et al. (2012) } \label{tabla1} \end{center} \end{table} | The origin and basic nature of pulsations exhibited by pre-ELMVs and ELMVs have been established (Steinfadt et al. 2010; C\'orsico et al. 2012; Jeffery \& Saio 2013; Van Grootel et al. 2013; C\'orsico \& Althaus et al. 2014ab, 2016; C\'orsico et al. 2016; Gianninas et al. 2016). The next step is to start exploiting the period spectra of these stars with asteroseismological analysis. Asteroseismology of low-mass He-core WDs will provide crucial information about the internal structure and evolutionary status of these stars, allowing us to place constraints on the binary evolutionary processes involved in their formation. | 16 | 9 | 1609.01352 |
1609 | 1609.01955_arXiv.txt | Using one-dimensional hybrid expanding box model we investigate properties of the solar wind in the outer heliosphere. We assume a proton-electron plasma with a strictly transverse ambient magnetic field and, beside the expansion, we take into account influence of a continuous injection of cold pick-up protons through the charge-exchange process between the solar wind protons and hydrogen of interstellar origin. The injected cold pick-up protons form a ring distribution function that rapidly becomes unstable and generate Alfv\'en cyclotron waves. The Alfv\'en cyclotron waves scatter pick-up protons to a spherical shell distribution function that thickens over that time owing to the expansion-driven cooling. The Alfv\'en cyclotron waves heat solar wind protons in the perpendicular direction (with respect to the ambient magnetic field) through the cyclotron resonance. At later times, the Alfv\'en cyclotron waves become parametrically unstable and the generated ion acoustic waves heat protons in the parallel direction through the Landau resonance. The resulting heating of the solar wind protons is efficient on the expansion time scale. | Properties of the solar wind plasma in the outer heliosphere are importantly influenced by inelastic collisions between ions and neutrals of interstellar origin \citep{burlal96,zankal09,rist09}. The solar wind ions and interstellar neutrals interact through the charge-exchange process. The neutrals (mostly hydrogen) become ionized and the interaction between the newly created, pick-up ions and the ambient solar wind plasma leads to a deceleration of the solar wind and has likely an important effect on ion thermal energetics \citep{richal95,wangal00b,richal08}. The nonthermal pick-up ions play likely an important role at the termination shock \citep{burlal08,deckal08}. The actual form of the pick-up ion distribution function just upstream from the termination shock importantly influences the shock properties and ion acceleration \citep{gide10,masc14,yangal15}. A pick-up ion dominated distant solar wind can be to some extend described using a macroscopic multi-fluid approach \citep{zankal14}. Global multi-fluid simulations show that the interaction between solar wind and interstellar neutrals has an important effect on the global properties of the outer heliosphere \cite[e.g.,][]{usmaal12}. However, the highly nonthermal newly born pick-up ion population leads naturally to plasma instabilities where strongly kinetic processes such as the cyclotron and Landau resonances are expected \citep{gary93}. Depending on the orientation of the background magnetic field, the pick-up have different distribution functions \citep{zaca00}. A ring distribution function is generated by injection perpendicular to the ambient solar wind. Such a distribution has an effective perpendicular temperature anisotropy and may generate Alfv\'en cyclotron waves (or mirror instabilities, etc.); the generated waves scatter the ring pick-up protons to a spherical shell-like velocity distribution function \citep{leip87,wiza94} and may heat directly the solar wind protons through the cyclotron instability \citep{grayal96,richal96}. Injection parallel to the ambient magnetic field leads to several beam-type instabilities \citep{daga98}. For a general orientation of the ambient magnetic field a ring-beam distribution function is generated and both the temperature anisotropy and the differential velocity could be a source of free energy for kinetic instabilities \citep{vahe15}. The pick-up ion driven instabilities are active in the solar wind as indicated by in situ observation of the generated waves \citep{joycal10,cannal14a,cannal14b,aggaal16}. The linear and nonlinear studies of pick-up ion driven instabilities usually assume a homogeneous plasma that is at odds with ubiquitous turbulent fluctuations observed in the distant solar wind \citep{fratal16}. In situ observations of enhanced wave activity \citep{joycal10,cannal14a,cannal14b,aggaal16} indicate that the pick-up ion driven instabilities are active in the solar wind. This is supported by direct hybrid simulations of \cite{hellal15} showing that kinetic instabilities may coexist with turbulence. The wave activity generated by pick-up ions may be a local source of turbulence as assumed in phenomenological transport models of turbulence \citep{zankal96,mattal99,isenal03,smital06,isenal10,adhial15}; the enhanced level of turbulence leads to enhanced cascade and heating rates. However, connection between turbulence and kinetic instabilities driven by non-thermal particle velocity distribution functions is far from being understood. In this paper we investigate long-time evolution of the expanding solar wind plasma under the effect of a continuous injection of pick-up protons. Section~\ref{heb} describes the simulation method, the hybrid expanding box model. Section~\ref{results} presents the main simulation results. In section~\ref{discussion} we discuss the simulation results. | \label{discussion} We investigated effects of continuous injection of pick-up protons in the distant solar wind using 1-D hybrid expanding box model. We assumed an ideal 1-D homogeneous system parallel to the ambient magnetic field directed along the transverse direction with respect to the radial direction. We assumed a slowly expanding ($t_e=10^4\omega_{c\mathrm{p}0}^{-1}$) in the two transverse directions and a slow continuous injection of cold pick-up protons due to the charge-exchange process with the solar wind protons ($t_{cx}=100 t_e$). The injection of pick-up protons leads to the formation of a ring velocity distribution distribution. This distribution becomes rapidly unstable and generates AIC waves that scatter through the cyclotron resonance the pick-up protons which consequently form a shell velocity distribution distribution. The AIC waves also scatter the solar wind protons and heat them in the perpendicular direction and cool them in the parallel one. The continuous injection of pick-up protons keep the instability active, and, at later times, the AIC waves grow slowly (secularly, about linearly in time). The AIC waves become eventually unstable with respect to a parametric instability that leads to formation of a secondary, shorter wavelength population of AIC waves and compressible ion-acoustic waves. The ion-acoustic waves interact with solar wind protons through the Landau resonance. The combined effect of the AIC and ion-acoustic waves lead to an efficient proton parallel and perpendicular heating. The pick-up proton shell distribution thickens during the evolution due to diffusion owing to the AIC waves and due to the expansion driven cooling; eventually, a power-law distribution is expected with a cut off near the injection velocity. The pick-up proton generated AIC waves and the ion-acoustic waves generated by the parametric instabilities may be partly responsible for the enhanced level of density fluctuations observed in the outer heliosphere \citep{bellal05,zankal12}. In the presented simulation we initialized the solar wind protons with $\beta_\mathrm{p}=0.2$. To test the sensitivity of the simulation results with respect to $\beta_\mathrm{p}$ we performed additional simulations with $\beta_\mathrm{p}=0.1$ and 1. For higher solar wind proton beta the cyclotron perpendicular heating becomes less efficient but overall behavior remain the same. For higher $\beta_\mathrm{p}$ one expects that the mirror instability would become important \citep{garyal97} and, also, properties of the parametric instabilities depend on $\beta_\mathrm{p}$ \citep{holl94}. Therefore, the present results are relevant for $\beta_\mathrm{p}\lesssim 1$. The present model is in many respects simplified, only one dimension is considered, the solar wind velocity is assumed to be constant and is about ten time faster than in the reality. Also, the system is assumed homogeneous, no pre-existing fluctuations/turbulence is assumed, Coulomb collisions and electron impact ionization are neglected, etc. However, the model self-consistently describes the kinetic plasma behavior in the expanding solar wind where pick-up protons are continuously injected. The one-dimensional geometry strongly reduces the available physics, only parallel propagating modes are allowed. In a low beta plasma the ring is expected to generate cyclotron waves primarily along the magnetic field which justifies the 1-D geometry but other instabilities (such as the mirror instability) may appear at oblique angles and may modify the nonlinear behavior. Also the 1-D geometry tend to leads to larger amplitude fluctuations the are prone to parametric instabilities. At 2-D and 3-D the effect of parametric instabilities will be likely reduced. The numerical resolution of the code could be source of other problems; our choice of the spatial grid, the box size, and the time step guaranties a good resolution of the AIC waves. While the used number of particles per cell is substantial (note that at $t=t_e$ there are about $10^3$ particles per cell for pickup protons) the resulting numerical noise may lead to enhanced scattering of the ring pick-up protons. \citep{floral16}. This is likely a minor problem since the continuous injection of pickup ions tends to keep the system unstable and rapidly generates fluctuations well above the numerical noise level. The presence of large scale fluctuations/turbulence will modify the initial local pick-up ion distribution function and the resulting instabilities \citep{zaca00}. For injection at nearly parallel angles with respect to the magnetic field beam-type instabilities are expected \citep{gary93}. These instabilities have different nonlinear properties but quite generally one expects formation of partial spherical shells \cite[cf.,][]{wiza94,mattal15}. Coulomb collisions are expected to be weak in the solar wind but may lead at the expansion time scale to scattering of pick-up ions \cite[via interaction with the solar wind ions, cf.,][]{tracal15,hell16} reducing thus the source of free energy for instabilities. Electron impact ionization is expected to have generally a subdominant effect with respect to the charge-exchange process in the solar wind \cite[but it is likely important in the heliosheath, cf.,][]{deckal08}. This process would have a similar effect as the charge exchange (i.e., generation of pick-up ions) but without the reduction of the solar wind ion density. We expect the overall effect of electron impact ionization would comparable to that of charge exchange. The connection between instabilities driven by the pick-up ions and turbulence remains an important open question. The turbulent fluctuations (or another wave activity) already present in the solar wind would scatter the pick-up ions and possibly reduce the source of free energy for instabilities. On the other hand, instabilities driven by pick-up ion ring-beam velocity distribution function generate typically waves at short, ion scales, often at nearly parallel angles with respect to the ambient magnetic field. Coupling between the strongly oblique turbulent fluctuations and quasi-parallel waves at ion scales is likely weak \cite[while waves generated at strongly oblique angles seem to participate in the cascade, cf.,][]{hellal15}. There are many open problems that will be subject of future work. In concluding, the present work shows that (i) collisionless plasma with energetically important population of pick-up ions generally needs a fully kinetic treatment, (ii) pick-up ion generated waves are able to directly quite efficiently heat the solar wind protons, (iii) the distribution function of pick-up ions at later times/larger distances has a wide spread of velocities/energies owing to scattering of the initial distribution on the generated waves and to the expansion-driven cooling, (iv) the hot neutrals of the solar wind origin have a complex velocity distribution function since they are neutralized at different distances and the solar wind ion temperature and the bulk velocity vary substantially over the time/distance. | 16 | 9 | 1609.01955 |
1609 | 1609.04425_arXiv.txt | {We recently discovered that the active galactic nucleus (AGN) of Mrk~1018 has changed optical type again after 30 years as a type 1 AGN. Here we combine \textit{Chandra}, \textit{NuStar}, \textit{Swift}, \textit{Hubble Space Telescope} and ground-based observations to explore the cause of this change. The 2--10\,keV flux declines by a factor of $\sim$8 between 2010 and 2016. We show with our X-ray observation that this is not caused by varying neutral hydrogen absorption along the line-of-sight up to the Compton-thick level. The optical-UV spectral energy distributions are well fit with a standard geometrically thin optically thick accretion disc model that seems to obey the expected $L\sim T^4$ relation. It confirms that a decline in accretion disc luminosity is the primary origin for the type change. We detect a new narrow-line absorber in \Lya\ blue-shifted by $\sim$700\,\kms\ with respect to the systemic velocity of the galaxy. This new \Lya\ absorber could be evidence for the onset of an outflow or a companion black hole with associated gas that could be related to the accretion rate change. However, the low column density of the absorber means that it is not the direct cause for Mrk 1018's changing-look nature.} | Active galactic nuclei (AGN) and some X-ray binaries and are thought to be powered by accretion of material onto a black hole (BH). They commonly show significant variability at optical-to-X-ray wavelengths on short timescales; this can be well described by noise processes \citep[e.g.][]{Nandra:1997,McHardy:2004,Mushotzky:2011}. The variability timescale is expected to scale with the mass of the BH and is therefore longest for super-massive BHs (SMBH) and can reach up to several hundred years \citep[e.g.][]{McHardy:2006}. It is therefore difficult or even impossible to directly measure the long-term high-amplitude fluctuations of AGN over the required timescales. Dramatic changes in the soft X-ray brightness or the strength of broad Balmer lines emitted in the broad line region (BLR) have been reported in some AGN. Examples of such a ''changing-look'' AGN caused by absorbing clouds passing in front of the nucleus and/or variable reflection components are NGC~4151 \citep[e.g.][and references therein]{Puccetti:2007}, NGC1365 \citep[e.g.][]{Risaliti:2009} and NGC~4051 \citep{Guainazzi:1998}, but these cloud events can be more common \citep{Markowitz:2014}. Prominent examples of AGN with appearing BLR are Mrk~1018 \citep{Cohen:1986}, NGC~1097 \citep{Storchi-Bergmann:1993}, and NGC~2617 \citep{Shappee:2014}, and examples with disappearing BLR are NGC 7603 \citep{Tohline:1976}, Mrk~590 \citep{Denney:2014}, SDSS~J0159+0033 \citep{LaMassa:2015}, and SDSS~J1011+5442 \citep{Runnoe:2016}. These events can either be explained by flares from tidal disruption events (TDEs) that are due to accretion of a star \citep[e.g.][]{Komossa:1999,Halpern:2004,Merloni:2015}, or intrinsic changes in the accretion disc flow depending on their light curves. \citet{McElroy:2016} (hereafter Paper I) reported the surprising discovery that \object{Mrk 1018} ($z=0.035$), which turned from a type 1.9 to a bright type 1 AGN around 1984, has changed back to a type 1.9 nucleus after about 30 years. The optical continuum brightness dropped by an order of magnitude between 2010 and early 2016. While we discuss in Paper I that a TDE probably cannot explain the variability of Mrk~1018, several other options including a cloud event still appeared possible. In this Letter, we present follow-up Director's Discretionary Time (DDT) and archival X-ray and UV spectroscopic data that show that the changing classification is driven by accretion rate changes and not by an obscuration event. | Based on follow-up X-ray observations we rule out an obscuring-cloud event as the cause of the change of type again after 30 years, as discovered in McElroy et al. (2016). All observations, in particular the optical-UV SED, are consistent with a declining accretion rate of a geometrically thin, optically thick accretion disc. Based on the appearance of a new NAL in \Lya, we speculate whether the onset of an outflow or a putative binary SMBH system is driving instabilities in the accretion disc that cause the decline in luminosity. However, the NAL might also be completely unrelated to the accretion disc changes. Continuous monitoring from the radio to X-rays is needed to further constrain the nature of the dramatic changes at the heart of the nucleus. | 16 | 9 | 1609.04425 |
1609 | 1609.02203_arXiv.txt | We report the first hard X-ray observations with {\it NuSTAR} of the BL Lac type blazar PKS~2155-304, augmented with soft X-ray data from XMM-{\it Newton} and $\gamma$-ray data from the {\it Fermi} Large Area Telescope, obtained in April 2013 when the source was in a very low flux state. A joint {\it NuSTAR} and XMM spectrum, covering the energy range 0.5 - 60 keV, is best described by a model consisting of a log-parabola component with curvature $\beta = 0.3^{+0.2}_{-0.1}$ and a (local) photon index $3.04\pm 0.15$ at photon energy of $2\;{\rm keV}$, and a hard power-law tail with photon index $2.2\pm 0.4$. The hard X-ray tail can be smoothly joined to the quasi-simultaneous $\gamma$-ray spectrum by a synchrotron self-Compton component produced by an electron distribution with index $p = 2.2$. Assuming that the power-law electron distribution extends down to $\gamma_{\rm min} = 1$ and that there is one proton per electron, an unrealistically high total jet power of $L_p \sim 10^{47}\;{\rm erg\,s^{-1}}$ is inferred. This can be reduced by two orders of magnitude either by considering a significant presence of electron-positron pairs with lepton-to-proton ratio $n_{\rm e+e-}/n_{\rm p} \sim 30$, or by introducing an additional, low-energy break in the electron energy distribution at the electron Lorentz factor $\gamma_{\rm br1} \sim 100$. In either case, the jet composition is expected to be strongly matter-dominated. | PKS~2155-304 is one of the most extensively studied BL Lac objects. It is a strong emitter of electromagnetic radiation in all observable bands, from radio to very high energy (VHE) $\gamma$ rays. Its $E \times F(E)$ broad-band spectrum reveals two prominent peaks located respectively in the far UV/soft X-ray band, and in the multi-GeV part of the high energy $\gamma$-ray band. As such, PKS~2155-304 belongs to the class of jet-dominated active galaxies with the jet pointing close to our line of sight - known as blazars - and, specifically, to a sub-class known as high-energy peaked BL Lac objects, or HBLs (see, e.g., \citealt{Pad95}). The two-peak spectral energy distribution (SED) of HBL blazars is generally (and most successfully) interpreted in the context of leptonic synchrotron self-Compton (SSC) models \citep[e.g.,][]{Ghi98}, where the low-energy component is presumably due to synchrotron emission, while the high energy component is due to inverse Compton scattering by the same electrons that produce the synchrotron peak. The optical spectra of the HBL blazars are generally devoid of emission lines even in the low jet flux states, implying a rather weak isotropic radiation field associated with the accretion. In such objects, it is generally believed that the dominant population of ``seed'' photons (as seen in the co-moving frame of the relativistic jet) are the synchrotron photons produced within the jet. From an observational standpoint, in HBL-type blazars perhaps the least is known about the lowest-energy part of the inverse-Compton peak. This is primarily due to the limited sensitivity of instruments in the relevant energy range, from $\sim 20$ keV to $\sim 100$ MeV. In particular, the onset of the high-energy peak contains important information about the lowest-energy electrons in the jet, which, in the context of any emission model, are most plentiful, and thus are a sensitive probe of the total content of particles in the jet. Notably, this low-energy end of the electron population cannot be reliably studied in the synchrotron component, since at low energies, the synchrotron emission is likely self-absorbed. Fortunately, the successful launch of the {\it NuSTAR} mission, sensitive in the 3 - 79 keV energy range, opened a new window for sensitive searches for the low-energy ``tail'' of the electron distribution in the inverse Compton component. In this paper, we report {\it NuSTAR} observations of PKS~2155-304, one of the brightest and also most luminous HBL blazars. This object, at $z=0.116$, has been known as a bright X-ray emitter since its discovery by HEAO-1 A3 \citep{Sch79}. Subsequent X-ray observations consistently show soft X-ray spectra, with photon index $\Gamma > 2.5$ in the 2 - 10 keV band \citep[e.g.,][]{Sem93,Bri94,Ede95,Urr97,Zha99,Kat00,Tan01,Bhag14}. Rapid variability on hourly time scales in the X-ray and optical bands is common; see \citep{Zha99,Ede01,Tan01,Kat00}. PKS~2155-304 is a known bright VHE $\gamma$-ray source \citep{Cha99,Aha05} and is highly variable on timescales down to $\sim$ minutes in the VHE $\gamma$ rays \citep{Aha07}. For the most recent multi-band observations involving {\it Fermi}-LAT and VHE observatories, see \cite{Aha09} or \cite{Che15}. {\it NuSTAR} observed PKS~2155-304 multiple times in 2013, as a part of multi-frequency monitoring with ground-based observatories, spanning radio through VHE bands. Here, we focus on the X-ray spectroscopy afforded by the first observation, conducted strictly simultaneously with XMM-{\it Newton}, in April 2013 for cross-calibration purposes. The joint {\it NuSTAR} and XMM-{\it Newton} spectrum reveals spectral complexity, and specifically, a soft spectrum in the 2 -- 10 keV range, hardening at the high-energy part of its bandpass. While a similar hard spectral ``tail'' was previously measured in the spectrum of this object by HEAO-1 \citep{Urr82} as well as by {\it Beppo}-SAX \citep{Gio98}, this was done with less sensitive, non-imaging instruments; the sensitive {\it NuSTAR} observation allows us to reliably confirm its presence, and characterize the spectrum in more detail. With relatively simple modeling of the broad-band SED in the context of SSC models, we are able to draw inferences about the distribution of radiating particles over a broad range of energies. Unless otherwise specified, we adopt the concordance cosmology, $\Omega_{\rm M} = 0.3, \Omega_{\Lambda} = 0.7$, and $H_{\rm 0} = 70$ km\ s$^{-1}$\ Mpc$^{-1}$. | PKS~2155-304 displayed a relatively low state during the first {\it NuSTAR} observations of the source in April 2013, with the measured 2 -- 10 keV X-ray flux of only $\sim 1.1 \times 10^{-11}$ erg cm$^{-2}$ s$^{-1}$, roughly three times lower than the lowest X-ray flux in August-September 2008, reported by \cite{Aha09}. {\it NuSTAR} data reveal a steep ($\Gamma \sim 3$) spectrum below $\sim 10$ keV, hardening to $\Gamma \sim 2$ above $\sim 10$ keV. When fitted with strictly simultaneous XMM-{\it Newton} data, the soft component is best-fit as a log-parabolic model, and the hard tail is even more significant. It is naturally expected that such spectral hardening as we detect in the combined {\it NuSTAR} and XMM-{\it Newton} data would be more easily detectable when the source is in a state of a relatively low soft X-ray flux. This is because the soft X-ray and VHE $\gamma$-ray variability in HBL BL Lacs is generally more rapid and has larger amplitude than that at lower energy of the respective peaks. This is partially due to more rapid energy losses with increasing particle energy. Therefore, the chance of detecting the presumably less variable onset (low-energy end) of the inverse Compton component is actually {\sl greater} when the high-energy tail of the synchrotron peak is weak, and does not dilute the Compton component. Indeed, our data taken in an extremely low-flux state reveal such a component. An application of the SSC model allows us to estimate the particle content in the jet. If we assume one proton per electron, then the total power of the jet is dominated by two orders of magnitude by particles, amounting to $L_p \sim 10^{47}\;{\rm erg\,s^{-1}}$. This would require a very large amount of power to be delivered via accretion, and would imply accretion at a highly super-Eddington rate. This, in turn, is unlikely given the absence of any quasi-thermal spectral components one would expect to be present in the optical/UV spectra of this source. Therefore, we consider a more plausible scenario, where the jet contains significantly more than one lepton per proton, meaning that by number, the jet is dominated by electron-positron pairs. This allows the reduction of the required jet power by two orders of magnitude, bringing it to more realistic values. The required jet power can also be reduced by introducing an additional break in the electron injection spectrum, e.g., with $\gamma_{\rm br1} \sim 100$ and $p_1 = 1$. In either case explored here, the total power of the jet is dominated by particles rather than by magnetic fields. In summary, while the presence of electron-positron pairs was previously postulated in relativistic jets of FSRQs (see \citealt{Sik00}), the new constraint from {\it NuSTAR} on the low-energy part of the electron distribution suggests that copious pairs may be present in jets associated with the lineless, HBL-type blazars. | 16 | 9 | 1609.02203 |
1609 | 1609.02942_arXiv.txt | {We report on the first focused hard X-ray view of the absorbed supergiant system \igr\ performed with the {\it NuSTAR} observatory. The pulsations are clearly detected with a period of $P_{\rm spin}=139.866(1)$ s and a pulse fraction of about 50-60\% at energies from 3 to 80 keV. The source demonstrates an approximately constant X-ray luminosity on a time scale of more than dozen years with an average spin-down rate of $\dot P\simeq6\times10^{-10}$ s s$^{-1}$. This behaviour of the pulsar can be explained in terms of the wind accretion model in the settling regime. The detailed spectral analysis at energies above 10 keV was performed for the first time and revealed a possible cyclotron absorption feature at energy $\sim 23$~keV. This energy corresponds to the magnetic field $B\simeq3\times10^{12}$ G at the surface of the neutron star, which is typical for X-ray pulsars.} | \igr\ was discovered by the {\it INTEGRAL} observatory during deep observations of the Galactic Center region as a faint persistent hard X-ray source \citep{2004AstL...30..382R}. Follow-up observations performed with the {\it XMM-Newton} observatory in soft X-rays allowed for an improvement in the determination of the source position, enabling near-infrared observations with the {\it NTT}/ESO telescope and establishing the nature of its companion as a supergiant star of B1Ib type \citep{2010A&A...510A..61T} at a distance of 12.4 kpc. \citet{2011A&A...532A.124M} classified the optical star as B0-B1 I, which is broadly in agreement with the above mentioned type. Using archival data from the {\it BeppoSAX} observatory, \citet{2003ApJ...596L..63A} revealed that the serendipitous source SAX\,J1802.7-2017 is spatially associated with \igr. These authors showed that this source is an X-ray pulsar with a spin period of 139.6 s, residing in a binary system with an orbital period of 4.6 days. Later, the orbital parameters of the system were improved by \citet{2005A&A...439..255H} and \citet{2015A&A...577A.130F}. The average spectrum of \igr, measured in a wide energy band \citep{2005A&A...430..997L}, demonstrates a cutoff at high energies, which is typical for X-ray pulsars \citep[see, e.g.,][]{2005AstL...31..729F}. In addition, a relatively high absorption value, $N_{\rm H}=6.8\times10^{22}$ cm$^{-2}$, was revealed at low energies \citep{2005A&A...439..255H}, and the source was classified as an absorbed binary system with a supergiant companion \citep{2015A&ARv..23....2W}. No search for a cyclotron absorption line and no pulse phase resolved spectroscopic study for \igr\ has been performed to date. In this paper we report on results of observations of \igr\ collected with the {\it NuSTAR} observatory and {\it Swift}/XRT telescope in Aug 2015. The main purpose of these observations was to characterize the broadband spectrum with high accuracy and to search for a cyclotron absorption line. | In this paper, we report results of the {\it NuSTAR} observations of the absorbed supergiant system \igr. They can be summarized as follows: 1) the system demonstrates approximately constant X-ray luminosity on a time scale of more than a dozen years; 2) during this time interval, the pulsar spun down with a rate of $\dot P\simeq6\times10^{-10}$ s s$^{-1}$; 3) the possible presence of a cyclotron absorption line at $\sim 23$~keV is found in the source X-ray spectrum. The observed increase of the rotation period of the neutron star in \igr\ can be explained in terms of the wind accretion model in a settling regime \citep{2012MNRAS.420..216S,2014EPJWC..6402001S}, which can be realized for slowly rotating magnetized neutron stars at X-ray luminosities below $\sim 4\times 10^{36}$~erg s$^{-1}$. The observed stable X-ray luminosity suggests that the X-ray pulsar has reached its equilibrium period, which, in this model, depends on the binary orbital period $P_b$, the stellar wind velocity $v_w$, the neutron star magnetic moment $\mu$ and the mass accretion rate $\dot M$: \begin{equation} \label{e:Peq} P_{eq}\simeq 940[\mathrm{s}]\mu_{30}^{12/11}\left(\frac{P_b}{10\mathrm{d}}\right) \dot M_{16}^{-4/11}v_8^4\,, \end{equation} \noindent where the characteristic parameters are $\mu_{30}\equiv \mu/10^{30}[\mathrm{G\,cm}^3]$, $\dot M_{16}\equiv \dot M/10^{16}[\mathrm{g\,s}^{-1}]$, $v_8\equiv v_w/10^8[\mathrm{cm\,s}^{-1}]$. A very strong dependence on the wind velocity suggests that it is more reliable to estimate the wind velocity by inverting this formula. From the observed X-ray luminosity (see Table 1) we find $\dot M_{16}\simeq 3$, and assuming that the absorption feature is the cyclotron line, we find the neutron star dipole magnetic moment $\mu_{30}\simeq 2$. Therefore, assuming that the observed pulsar period is close to its equilibrium value, $P_{eq}=P_{spin}=139.6$~s, we can estimate the wind velocity $v_8\sim 0.7$. We stress that this estimate depends only very weakly on parameters ($\dot M$ and $\mu$ and other model-dependent numerical coefficients). Note that this rather low value is very close to the wind velocity measured in the prototypical persistent HMXB with OB supergiant Vela X-1 \citep{2016A&A...591A..26G}, suggesting similarity between the two sources. In the model of settling quasi-spherical accretion, the neutron star close to equilibrium can exhibit either spin-up or spin-down, depending on whether the actual $\dot M$ is above or below the critical mass accretion rate, $\dot M_{eq}$ (see Eq. (2) in \citealt{2014EPJWC..6402002P}). For the parameters of \igr, we find $\dot M_{eq}\simeq 3.6\times 10^{16}$~g s$^{-1}$, i.e. indeed spin-down of the neutron star is possible in this system. The observed value of the negative torque acting on the neutron star in \igr\, $\dot\omega\simeq 2\times 10^{-13}$~rad~s$^{-2}$, does not exceed the maximum possible negative torque, $\omega_{sd,max}\simeq 10^{-12}$~rad~s$^{-2}$, for the parameters of \igr\ (see Eq. (6) in \citealt{2014EPJWC..6402002P}). Thus, the observed steady-state spin-down is consistent with expectations from the settling accretion model. As discussed in Section~\ref{sec:spec}, the spectral analysis revealed the possible presence of a cyclotron absorption line in the spectrum of \igr\ at energies of $\sim23-24$ keV. An additional independent test for the existence of this feature comes from the timing properties of the source. Namely, the pulse profile and pulsed fraction dependencies on the energy have prominent features near the same energy. Particularly, the relative intensity of two peaks in the profile changes around 22-24 keV (see Fig.\,\ref{fig:2dpprof}). Such behaviour was shown to be typical for another well known X-ray pulsar V\,0332+53 with a well established cyclotron feature \citep{2006MNRAS.371...19T}. The observed non-monotonic dependence of the pulsed fraction on energy (see Fig.\,\ref{fig:ppfr}) is also typical for pulsars with cyclotron lines \citep{2009AstL...35..433L,2009A&A...498..825F}. These additional observational facts indirectly support the presence of the cyclotron absorption line in the \igr\ spectrum despite of its low significance for the {\sc powerlaw*highcut} continuum model. The measured X-ray luminosity in \igr, $L_x\simeq 3\times 10^{36}$ \lum, suggests that accretion onto the neutron star occurs in the subcritical regime where the radiation plays a secondary role \citep[see, e.g.][]{1976MNRAS.175..395B,2015MNRAS.447.1847M}. In this case, the accretion flow is decelerated in a collisionless shock formed at some height above the neutron star surface \citep{1982ApJ...257..733L,2004AstL...30..309B}. The formation of a cyclotron line downstream of the shock occurs in the resonance layer in the inhomogeneous magnetic field, so the line profile can be more complicated than the simple Doppler-broadened line \citep{1996ASSL..204.....Z}. For example, depending on the geometry, the line may have a flat bottom or show emission wings. The complex shape of the residuals shown in Fig.\,\ref{fig:avspec}, which are obtained assuming a Gaussian form of the line, may suggest a complicated line profile or may be evidence for the presence of complicated magnetic field structure near the surface of the neutron star, as discussed, for example, in \cite{2012MNRAS.420..720M}. Moreover, as the optical thickness of the flow at the resonant energies is very high even at low mass accretion rate, the cyclotron line can be also formed in the accretion channel above the shock region \citep{2015MNRAS.454.2714M}. In this case, one would expect the cyclotron absorption feature to be redshifted relative to the actual cyclotron energy. This interesting result should be studied in the future with deeper observations. | 16 | 9 | 1609.02942 |
1609 | 1609.02459_arXiv.txt | We conduct a linear stability calculation of an ideal Keplerian flow on which a sinusoidal zonal flow is imposed. The analysis uses the shearing sheet model and is carried out both in isothermal and adiabatic conditions, with and without self-gravity (SG). In the non-SG regime a structure in the potential vorticity (PV) leads to a non-axisymmetric Kelvin-Helmholtz (KH) instability; in the short-wavelength limit its growth rate agrees with the incompressible calculation by \citet{Lithwick2007}, which only considers perturbations elongated in the streamwise direction. The instability's strength is analysed as a function of the structure's properties, and zonal flows are found to be stable if their wavelength is $\gtrsim 8H$, where $H$ is the disc's scale height, regardless of the value of the adiabatic index $\gamma$. The non-axisymmetric KH instability can operate in Rayleigh-stable conditions, and it therefore represents the limiting factor to the structure's properties. Introducing SG triggers a second non-axisymmetric instability, which is found to be located around a PV maximum, while the KH instability is linked to a PV minimum, as expected. In the adiabatic regime, the same gravitational instability is detected even when the structure is present only in the entropy (not in the PV) and the instability spreads to weaker SG conditions as the entropy structure's amplitude is increased. This eventually yields a non-axisymmetric instability in the non-SG regime, albeit of weak strength, localised around an entropy maximum. | Zonal flows are azimuthally symmetric shear flows exhibiting a strip-like alternating pattern which are observed in a wide range of settings. Possibly the most famous natural example is represented by Jupiter's belts, whose origin is still controversial; recent numerical simulations have confirmed and reproduced the presence of zonal flows in both gas and ice giant planets by modelling convective turbulence \citep[e.g.][]{Sunetal1993, HeimpelAurnou2007}. Zonal flows also occur in laboratory plasma experiments, where they are generated by non-linear transfer of energy from drift waves between small and large scales \citep{Diamondetal2005}. More recently, they have been encountered in 2D numerical simulations of shearing sheet models of accretion discs in magnetohydrodynamic (MHD) turbulent regimes \citep[e.g.][]{Johansenetal2009, Simonetal2012, KunzLesur2013, BaiStone2014}. They are seen to become more prominent and possess larger amplitudes with increasing box size, with their wavelength usually taking the largest possible value in the radial direction (i.e. $k_x = \frac{2\pi}{L_x}$, $k_x$ and $L_x$ being the wavenumber and the box size in the $x$ direction) \citep{Johansenetal2009, Simonetal2012}. The zonal flows are also observed to have longer lifetimes when considering larger box sizes \citep{BaiStone2014}, except when very small boxes are considered \citep{Johansenetal2009}, and they do not appear to be dependent on the initial conditions applied to the system, with \citet{Simonetal2012} initialising their simulations with either two flux tubes or a uniform toroidal field and seeing no difference in the wavenumber of the structure nor a significant discrepancy in its appearance time. Using a 2D shearing sheet model of a hydrodynamical disc, \citet{Lithwick2007, Lithwick2009} showed that zonal flows in an incompressible, inviscid fluid can be unstable to the formation of vortices \citep[see also][]{Gill1965, LernerKnobloch1988} via the Kelvin-Helmholtz (KH) instability. The instability is however often missed by numerical simulations as it operates at small values of the azimuthal wavenumber $k_y$, implying that the computational domain must be substantially elongated in the $y$ direction. Some 2D numerical simulations have shown that coherent vortices associated with zonal flows \citep[e.g.][]{UmurhanRegev2004, JohnsonGammie2005} appear to be long-lived; what's more, as \citet{UmurhanRegev2004} point out, accretion discs are characterised by Reynolds numbers $Re = VL/\nu$ ($V$ and $L$ being the characteristic velocity and length scales in the disc and $\nu$ the kinematic viscosity) that are orders of magnitude larger than what it is possible to implement in a simulation, which would potentially make vortices' lifetimes in discs extremely long. The vortices are also seen to produce an outward transport of angular momentum, necessary for accretion onto the central object to occur; this could thefore potentially prove to be an alternative to magnetohydrodynamic (MHD) turbulence in discs -- or regions of discs -- which are not sufficiently ionised for MHD turbulence to take place, as long as a suitable source for the zonal flow is pinpointed. Lastly, zonal flows can play an important role in the dynamics of planetary formation in protoplanetary discs. A key obstacle in the early formation of protoplanets is the `metre-size barrier', representing the difficult chances of survival and growth of metre-sized bodies during the period of fast inward migration they experience due to gas drag \citep{Weidenschillingold1977}. The presence of density and velocity fluctuations in the disc, such as zonal flows, alters the drag force exerted on the material by the gas, potentially slowing down -- or halting altogether -- the inward drift of metre-sized bodies \citep{Whipple1972, KlahrLin2001, HaghighipourBoss2003, FromangNelson2005, Katoetal2009}. This would create a local enhancement in their density, promoting both the growth of these bodies through coagulation \citep{Weidenschilling1997, DullemondDominik2005, Braueretal2008a, Braueretal2008b} and the triggering of gravitational instability, potentially leading to fragmentation \citep{YoudinShu2002, Johansenetal2006, Johansenetal2007}. Owing to the widespread interest in zonal flows mentioned above, the aim of this paper is to perform a linear analysis on the stability of a Keplerian accretion disc when an axisymmetric structure is imposed on the system; such analysis will be conducted using the compressible shearing sheet (2D) approximation in both isothermal and adiabatic conditions, also taking into account the effects of the disc's self-gravity (SG). The paper is arranged as follows: in Section~\ref{sec:model} we present the basic equations of the shearing sheet model employed to tackle the problem, and derive the equations governing both the axisymmetric structures and their non-axisymmetric disturbances. We also describe the method by which they are solved numerically. In Section~\ref{sec:results} we present results on instabilities in both isothermal and adiabatic cases and compare these with relevant literature. The paper closes in Section~\ref{sec:conclusions} with the conclusions drawn from the results. | \label{sec:conclusions} We carried out a linear stability calculation of a steady Keplerian flow on which an axisymmetric structure has been imposed both in the absence and presence of self-gravity. For the isothermal case with no self-gravity, the growth rates of the non-axisymmetric instability induced by the presence of the axisymmetric structure in the potential vorticity $\zeta$ were compared to the incompressible calculation carried out by \citet{Lithwick2007}. The growth rates from the two calculations agree well for structures and disturbances of short wavelength, although a discrepancy appears as the azimuthal wavenumber increases due to a simplification in the model employed by \citet{Lithwick2007}. Thanks to the level of agreement reached, the non-axisymmetric instability found was pinpointed to be of the Kelvin-Helmholtz type. The strength of the Kelvin-Helmholtz (KH) instability was investigated as a function of the structure's amplitude and its wavelength. Zonal flows were found to be stable to non-axisymmetric instabilities, even for large amplitudes, as long as their wavelength $\gtrsim 8H$, with $H$ being the disc scale height; the result was found to be independent of the adiabatic index $\gamma$. Zonal flows having wavelength smaller than that critical value can undergo a KH instability, which would break the structure up into vortices. On the other hand, long-lived zonal flows observed in magnetorotational instability (MRI) simulations, can be explained either by the structure possessing a sufficiently long (and hence stable) wavelength, or by the box employed being too small in the azimuthal direction, preventing the capture of the instability. In the latter case, the results presented in this paper suggest that the aspect ratio of the box should be $L_y/L_x \gtrsim 2-2.5$ in order to successfully capture the KH instability. % On the other hand, if the zonal flow wavelength is less than $\sim2H$, a small zonal flow amplitude is enough to trigger a KH instability. The Rayleigh criterion was also looked at as a potential way of constraining the size of the zonal flow, although its resulting axisymmetric instability needs a 3D model to operate. It was found that it is possible for the KH instability to operate in a Rayleigh stable regime, meaning that the length-scale and the amplitude of any zonal flow would be limited by the non-axisymmetric instability. As highlighted by \citet{Lithwick2007, Lithwick2009} the KH instability observed leads to the formation of vortices in the disc, while no direct widening of the zonal flow is seen; it is however possible that the non-linear dynamics of the vortices induced by the instability might indirectly cause the formation of a wider, stable zonal flow configuration. Self-gravity was subsequently introduced and a second, distinct type of non-axisymmetric instability was detected. This newly pinpointed instability was of a gravitational nature and was found to be linked -- by considering the correlation integral of the wave's energy with the potential vorticity -- to maxima in the potential vorticity, while the previously detected KH instability is associated with minima in $\zeta$; this agrees with results in the literature, among which are \citet{LinPapaloizou2011, LovelaceHohlfeld2013}. This was confirmed by looking at the $x$--profile of the energy for modes associated with each instability. A brief analysis on the amount of compressibility and vorticity associated with each instability was also carried out, and this in general confirmed the initial expectation that the KH and gravitational instabilities should present an excess of vorticity and compressibility, respectively. However in the long $2\pi/k_y$ regime, both instabilities showed a predominance of vorticity, underlining the importance of the potential vorticity's role in this problem. An ideal adiabatic case was also considered to analyse the claim by \citet{Lovelaceetal1999} that a non-self-gravitating linear Rossby wave instability can be triggered by the presence of a local maximum in the entropy profile but not in the potential vorticity, which is usually a requirement for the development of a Rossby wave instability. A gravitational non-axisymmetric instability was detected, and its activation point shifted down to weaker gravity conditions as the amplitude of the entropy structure was increased. Eventually an instability was detected for no self gravity, therefore confirming the presence of such an instability as claimed by \citet{Lovelaceetal1999}. It was however unclear whether this instability, which was found to be localised near an entropy maximum using the Poisson summation formula method, was of a different type from the one obtained for non-zero self-gravity. A study into the exact nature of this instability was beyond the scope of this paper, but an analysis of its strength as a function of the zonal flow's amplitude and wavelength was carried out; the result was qualitatively similar to the KH instability in the isothermal case, except for the entropy-induced instability being weaker by an order of magnitude. It could also be triggered for a smaller range of amplitude and wavenumber values. In this case the flow is stable to axisymmetric disturbances according to the Rayleigh criterion for all values of the structure's amplitude and wavelength, so once again the properties of the zonal flow are limited by the non-axisymmetric instability. This work is based on an ideal model of the problem at hand, where the imposed zonal flow is assumed to be in a steady state with diffusive and -- in the adiabatic case -- cooling effects being neglected. While this ideal situation allows us to consider and analyse the dynamics of the flow without having to untangle them from secondary effects, it obviously represents an unrealistc scenario in the case of a real life accretion disc. For that reason, the next step is to expand this work by including effects such as those mentioned above and analysing their effects upon the evolution and stability of the system. | 16 | 9 | 1609.02459 |
1609 | 1609.06170_arXiv.txt | Coherent jets containing most of the kinetic energy of the flow are a common feature in observations of atmospheric turbulence at planetary scale. In the gaseous planets these jets are embedded in a field of incoherent turbulence on scales small relative to the jet scale. Large scale coherent waves are sometimes observed to coexist with the coherent jets and the incoherent turbulence with a prominent example of this phenomenon being the distortion of Saturn's North Polar Jet (NPJ) into a distinct hexagonal form. Observations of this large scale jet-wave-turbulence coexistence regime raises the question of identifying the mechanism responsible for forming and maintaining this turbulent state. The coherent planetary scale jet component of the turbulence arises and is maintained by interaction with the incoherent small-scale turbulence component. It follows that theoretical understanding of the dynamics of the jet-wave-turbulence coexistence regime can be facilitated by employing a statistical state dynamics (SSD) model in which the interaction between coherent and incoherent components is explicitly represented. In this work, a two-layer beta-plane SSD model closed at second order is used to develop a theory that accounts for the structure and dynamics of the NPJ. An asymptotic analysis is performed of the SSD equilibrium in the weak jet damping limit that predicts a universal jet structure in agreement with observations of the NPJ. This asymptotic theory also predicts the wavenumber (six) of the prominent jet perturbation. Analysis of the jet-wave-turbulence regime dynamics using this SSD model reveals that jet formation is controlled by the effective value of $\beta$ and the required value of this parameter for correspondence with observation is obtained. As this is a robust prediction it is taken as an indirect observation of a deep poleward sloping stable layer beneath the NPJ. The slope required is obtained from observations of the magnitude of the zonal wind component of the NPJ. The amplitude of the wave six perturbation then allows identification of the effective turbulence excitation maintaining this combined structure. The observed jet structure is then predicted by the theory as is the wave six disturbance. The wave six perturbation, which is identified as the least stable mode of the equilibrated jet, is shown to be primarily responsible for equilibrating the jet with the observed structure and amplitude. | Coherent structures emergent from small scale turbulence are often observed in planetary atmospheres with the zonal jets of the gaseous planets being familiar examples~\citep{Ingersoll-90, Vasavada-and-Showman-05, Sanchez-etal-2000,Galperin-etal-2014}. While this phenomenon of spontaneous large scale jet organization from small scale turbulence has been extensively investigated in both observational and theoretical studies~\citep{Kraichnan-1967,Rhines-1975,Williams-79a, Williams-03,Panetta-93, Nozawa-and-Yoden-97, Huang-Robinson-98, Lee-05, Manfroi-Young-99, Vallis-Maltrud-93, Cho-Polvani-1996, Read-etal-2004,Showman-2007, Scott-Polvani-2008, Galperin-etal-04} the physical mechanism underlying it remains controversial. The prominence of jets in planetary turbulence is in part due to the jet being a nonlinear stationary solution of the dynamics in the limit of vanishing dissipation and therefore not disrupted by nonlinear advection on the time scale of the large scale shear. However, its being a stationary solution is insufficient by itself to serve as an explanation for the observed jets for three reasons. First, strong jets typically assume a characteristic structure for a given set of system parameters, while any zonally symmetric flow is a fixed point of the inviscid dynamics. Second, nonlinear stationary states lack a mechanism of maintenance against dissipation and so can not explain the fact that the observed jets, which are not maintained by coherent external forcing such as by an imposed pressure gradient, persist much longer than the dissipation time scale. Third, planetary jets commonly appear to be unstable; for example, the north polar jet (NPJ) of Saturn robustly satisfies the Rayleigh--Kuo necessary condition for barotropic instability in a dissipationless stationary flow~\citep{Antunano-etal-2015, Sanchez-etal-2014}, and barotropic instability of this jet has been verified by eigenanalysis~\citep{Barbosa-etal-2010}. The aforementioned considerations imply that a comprehensive theory for the existence of large scale jets in the atmospheres of the gaseous planets and in particular Saturn's NPJ must provide a mechanism for the formation and maintenance of the jet from incoherent turbulence, the particular structure assumed by the jet and its stability. In addition to these the case of the NPJ also requires that the theory account for the prominent coherent wave six perturbation that distorts the jet into a distinct hexagonal form. The primary mechanism by which the large scale jets of the gaseous planets are maintained is upgradient momentum flux resulting from straining of the perturbation field by the mean jet shear which produces a spectrally nonlocal interaction between the small-scale perturbation field and the large-scale jet. This mechanism has been verified in observational studies of both the Jovian and Saturnian atmospheres~\citep{Ingersoll-etal-2004,Salyk-etal-2006, Delgenio-etal-2007}, as well as in numerical simulations~\citep{Nozawa-and-Yoden-97,Huang-Robinson-98,Showman-2007} and in laboratory experiments~\citep{Wordsworth-etal-2008}. This upgradient momentum transfer mechanism has been found to maintain mean jets both in barotropic forced dissipative models~\citep{Huang-Robinson-98, Farrell-Ioannou-2003-structural} and in baroclinic free turbulence models~\citep{Panetta-93, Farrell-Ioannou-2009-closure} and can be traced to the interaction of the perturbation field with the mean shear \citep{Farrell-Ioannou-1993-unbd, Huang-Robinson-98, Bakas-Ioannou-2013-jas,Srinivasan-Young-2014}. Excitation of the observed small scale forced turbulence in the case of both the Jovian and Saturnian jets is believed to be of convective origin ~\citep{Vasavada-etal-2006, Showman-2007,Gierasch-etal-2000, Sanchez-etal-2000,Porco-etal-2003,Delgenio-etal-2007,Showman-2007}. For our purposes it suffices to maintain the observed amplitude of small scale field of turbulence. We choose to maintain this turbulent field in the simplest manner though introducing a stochastic excitation. The structure of the stochastic excitations is not important so long as it maintains the observed amplitude of turbulence given that the anisotropy of the turbulence is induced by the mean shear of the jet. In this work Saturn's NPJ is studied using the statistical state dynamics (SSD) of a two-layer baroclinic model, specifically a closure at second order in its cumulant expansion (cf.~Ref.~\citep{Hopf-1952}). The implementation of SSD used is referred to as the stochastic structural stability theory (S3T) system~\citep{Farrell-Ioannou-2003-structural}. In S3T the nonlinear terms in the perturbation equation for the second cumulant involves the third cumulant which is parameterized by a stochastic excitation rather than being explicitly calculated while the nonlinear interaction of the perturbations with the mean jet are fully retained. For this reason the S3T system may be described as quasi-linear (QL) in accord with the fact that quasilinearity is a general attribute of second order closures \citep{Herring-1963}. S3T has been applied previously to the problem of jet formation in barotropic turbulence~\citep{Farrell-Ioannou-2007-structure, Srinivasan-Young-2012,Tobias-Marston-2013,Parker-Krommes-2013-generation,Bakas-Ioannou-2013-prl,Constantinou-etal-2014} to jet dynamics in the shallow water equations \citep{Farrell-Ioannou-2009-equatorial} and to jet dynamics in baroclinic turbulence~\citep{Farrell-Ioannou-2008-baroclinic, Farrell-Ioannou-2009-closure}. The S3T system employs an equivalently infinite ensemble in the dynamical equation for the second cumulant and as a result provides an autonomous and fluctuation-free dynamics for the statistical mean turbulent state which greatly facilitates analytical study\footnote{ Formal justification of the second order S3T closure has been given by Bouchet et al. \cite{Bouchet-etal-2013, Bouchet-etal-2014} who show that for $\alpha = \lambda \tau\ll1$ to leading order in $\alpha$ the statical dynamics asymptotically approaches the dynamics of the S3T second-order closure with mean flow $O(1/\alpha)$ larger than the perturbation field. The dimensionless parameter $\alpha$ is the product of $\tau= L_y/(U_{\rm max}-U_{\rm min})$, the shear time of the jet, and, $\lambda$, the inverse of the time scale on which the large scale flow evolves, which is inversely proportional to the square root of the energy injection rate, $\varepsilon$. However, these formal results are too conservative. For example, it has been demonstrated that S3T theory is predictive of the bifurcation of the statistically homogeneous state of barotropic beta plane turbulence to the jet state, in which arbitrarily small jets emerge, and that these predictions are valid at the bifurcation point, which is the regime of $\alpha \gg1$ \citep{Constantinou-etal-2014}. Similar validity of perturbative structure instability in the $\alpha>>1$ regime has been demonstrated for the case of 3-D Couette flow turbulence \citep{Farrell-Ioannou-2016-bifur}. In retrospect we understand that the fundamental underlying reason for the robust validity of the S3T dynamics has a physical basis in that it captures the mechanism determining the statistical state of shear turbulence which is interaction between the mean flow and the perturbations supporting it by means of quadratic fluxes which are accurately obtained from the second cumulant. The fact that this interaction is contained in the closure at second order is consistent with the success of this closure in capturing the qualitative and in many examples also the quantitative behavior of turbulent equilibria in shear flow.}. When applying S3T to the study of zonal jets it is useful to equate the ensemble mean and zonal mean by appeal to the ergodic hypothesis. A two-layer model is employed in order to provide the possibility for baroclinic and barotropic dynamics both for the jet itself and for the perturbations that are involved in the equilibration dynamics. One reason this is important is that a barotropic, equivalent barotropic or shallow water model with the observed Rossby radius would not allow the problem freedom to adopt barotropic dynamics corresponding to formation of deep jets. In the event we find that the statistical equilibrium jets are either barotropic or close to it so that the Rossby radius is not a relevant parameter \citep{Showman-2007}. We find that jet formation is tightly controlled by the effective vorticity gradient, $\beta$. As this is a robust requirement of the dynamics, the observed jet structure is taken as an indirect observation of this parameter. Saturn's NPJ is similar in structure and amplitude to strong midlatitude jets on the gaseous planets such as Jupiter's $24^\circ\;\rm N$ jet while the planetary value of $\beta_{sat}(74^\circ)=1.6\times 10^{-12}~\rm m^{-1}\,s^{-1}$ at the latitude of the NPJ is too weak to stabilize a jet with the observed amplitude ($98.7~\rm m\,s^{-1}$), which poses a dynamical dilemma~\cite{Antunano-etal-2015}. Theory and observation can be brought into correspondence by inferring a deep strongly statically stable layer beneath the jet giving rise to the equivalent of a topographic $\beta$ effect. The $\beta$ used in the model is then the dyamical superposition of the effects of both the planetary and the topographic components. With this inferred effective $\beta$ the observed jet structure accords with the theory. While the first cumulant provides the structure of the jet, the second determines the planetary scale wave disturbance superposed on the jet. With the inferred value of $\beta$ and an incoherent turbulence excitation level consistent with observation this wave is found to have wavenumber six and the amplitude required to produce the observed hexagonal shape of the NPJ. The role of this wave in the dynamics is to equilibrate the jet with the observed velocity structure and amplitude while providing the pathway for dissipation of the energy that the jet is continuously extracting from the small scale turbulence. Previously advanced explanations for the prominent wave six perturbation to Saturns's NPJ are that it arises as the surface expression of an upward propagating Rossby wave the origin of which is attributed to a wave six corrugation of an inferred deep lower layer \citep{ Sanchez-etal-2014} and that the wave six results from nonlinear equilibration of a linear instability of the jet \citep{Morales-etal-2011,Morales-etal-2015}. However, the equilibrated wave six instability predicts closed vortices which are not seen in observations of the NPJ. The nonlinear S3T equilibrium obtained satisfies the Rayleigh--Kuo necessary condition for barotropic jet instability in both the prograde and retrograde jet and significant interaction between the jet and the modes associated with both these vorticity gradient structures is seen. Although the Rayleigh--Kuo criterion is not sufficient to ensure instability of a barotropic jet, experience has shown that, absent careful contrivance of the velocity profile, satisfaction of this necessary condition coincides with modal instability. Therefore, finding this criterion satisfied absent an instability directs attention toward the mechanism responsible for the implied careful contrivance. In the case of the NPJ this mechanism is shown in this work to be continuous feedback regulation between the coherent jet and the incoherent turbulence that adjusts the jet to marginal stability under conditions of sufficiently strong forcing by the small scale incoherent components to produce an unstable jet profile. The widely debated enigma of the stability of the zonal jets of the gaseous planets and in particular the stability of the NPJ in the face of observed strong vorticity gradient sign reversals is in this way resolved by the jets having been adjusted to (in most cases marginal) stability by perturbation-mean flow interaction between the first and second cumulants of the S3T dynamics. This mechanism of regulation to marginal stability by feedback between the first and second cumulant is familiar as the agent underlying establishment of turbulent equilibria in Rayleigh-Benard convection \citep{Malkus-1956,Herring-1963} and in establishing the baroclinic adjustment state in baroclinic turbulence \citep{Stone-1978,Farrell-Ioannou-2008-baroclinic,Farrell-Ioannou-2009-closure}. This mechanism of equilibration also has implications for the problem of identifying how energy transferred upscale from the excited small scales to large scales in geostrophic turbulence is dissipated as is required to maintain statistical equilibrium. Ekman damping associated with no slip boundaries is not available in the absence of solid boundaries and there is negligible diffusive damping at the jet scale. In fact, in the case of the NPJ, the eddy fluxes, including the eddy damping, are explicitly calculated for the SSD equilibrium state and these are found to be dominantly upgradient and therefore in toto are responsible for maintaining rather than dissipating the jets. In model studies hypodiffusion is commonly used to allow establishment of a statistically steady state \citep{Danilov-04,Scott-Polvani-2007}. While hypodiffusion is often employed without physical justification it can be related to radiative damping of baroclinic structures \citep{Scott-Polvani-2007}. However, we find that the dynamics of jet formation result in primarily barotropic jet structure so that thermal damping is not relevant. Instead, in the case of Saturn's NPJ we identify the physical dissipation mechanism responsible for equilibrating the jet to be energy transfer directly from the coherent jet to a wave six structure followed by dissipation of the energy by this wave. We note that in this planetary scale turbulence regime both the upscale energy transfer maintaining the jet as well as the downscale energy transfer to the wave six mode regulating its amplitude occur directly between remote scales and in neither case do these involve a turbulent cascade. | Large-scale coherent structures such as jets, meandering jets are characteristic features of turbulence in planetary atmospheres. While conservation of energy and enstrophy in inviscid 2D turbulence predicts spectral evolution leading to concentration of energy at large scales, these considerations cannot predict the phase of the spectral components and therefore can neither address the central question of the organization of the energy into jets with specific structure nor the existence of the coherent component of the planetary scale waves. In order to study structure formation additional aspects of the turbulence dynamics beyond conservation principles must be incorporated in the analysis. SSD models have been developed to study turbulence dynamics and specifically to solve for turbulent state equilibria consisting of coexisting coherent mean structures and incoherent turbulent components which together constitute the complete state of the turbulence at second order. In this work a second order SSD of a two-layer baroclinic model was used to study the jet-wave-turbulence coexistence regime in Saturn's NPJ. This second order SSD model, referred to as the S3T model, is closed by a stochastic parameterization that accounts for both the neglected nonlinear dynamics of the perturbations from the zonal mean as well as the excitation maintaining the turbulence. The equation for the zonal mean retains its interaction through Reynolds stress with the perturbations. In this model a jet forms as an instability and grows at first exponentially eventually equilibrating at finite amplitude. Exploiting the simplicity of the asymptotic regime in which the jet is undamped makes it possible to obtain a universal jet structure and jet amplitude scaling. Given that the associated jet structure and its amplitude scaling is robust in the SSD model we conclude that the observed structure of the NPJ can only be maintained as an equilibrium state with a value of $\beta$ greater than the planetary value. This requirement implies existence of a topographic beta effect with a specific predicted value. Incorporating the implied poleward decreasing stable layer depth into the model results in the model producing the observed jet structure. In the model a stable retrograde mode of the Rossby wave spectrum with wavenumber six becomes neutrally stable as the jet amplitude increases under Reynolds stress forcing by the small scale turbulence increased and by inducing strong non-normal interaction with the jet this wave six arrests its growth via perturbation Reynolds stresses. This composite structure equilibrates in the form of a hexagonal jet in agreement with the NPJ observations. Among the correlates of this theory is the predicted existence of the observed robust vorticity gradient reversals in both the prograde and retrograde jets as well as the location and structure of these reversals. | 16 | 9 | 1609.06170 |
1609 | 1609.08671_arXiv.txt | {Supernova remnants exhibit shock fronts (shells) that can accelerate charged particles up to very high energies. In the past decade, measurements of a handful of shell-type supernova remnants in very high-energy gamma rays have provided unique insights into the acceleration process. Among those objects, \rxj~(also known as \gthree) has the largest surface brightness, allowing us in the past to perform the most comprehensive study of morphology and spatially resolved spectra of any such very high-energy gamma-ray source. Here we present extensive new \hess\ measurements of \rxj, almost doubling the observation time compared to our previous publication. Combined with new improved analysis tools, the previous sensitivity is more than doubled. The \hess\ angular resolution of $0.048^\circ$ ($0.036^\circ$ above 2~TeV) is unprecedented in gamma-ray astronomy and probes physical scales of 0.8 (0.6) parsec at the remnant's location. The new \hess\ image of \rxj\ allows us to reveal clear morphological differences between X-rays and gamma rays. In particular, for the outer edge of the brightest shell region, we find the first ever indication for particles in the process of leaving the acceleration shock region. By studying the broadband energy spectrum, we furthermore extract properties of the parent particle populations, providing new input to the discussion of the leptonic or hadronic nature of the gamma-ray emission mechanism.} | \label{sec:intro} Highly energetic particles with energies up to $10^{20}$ electron volts (eV, 1~eV~=~$1.6\times 10^{-19}$~J) hit the atmosphere of the Earth from outer space. These cosmic rays (CRs) are an important part of the energy budget of the interstellar medium. In our Galaxy the CR energy density is as large as the energy density of thermal gas or magnetic fields, yet the exact connection and interaction between these different components is poorly understood~\citep{2015ARA&A..53..199G}. Among the measured properties of CRs is the energy spectrum measured at Earth, which extends over many orders of magnitude. At least up to a few times $10^{15}$~eV, these particles are likely of Galactic origin -- there must be objects in the Milky Way that accelerate charged particles to these energies. The composition of Galactic CRs is also known~\citep{Agashe:2014kda}: at GeV to TeV energies, they are dominantly protons. Alpha particles and heavier ions make up only a small fraction of CRs. Electrons, positrons, gamma rays and neutrinos contribute less than 1\%. Establishing the Galactic sources of charged CRs is one of the main science drivers of gamma-ray astronomy. The standard paradigm is that young supernova remnants (SNRs), expanding shock waves following supernova explosions, are these accelerators of high-energy Galactic CRs~\citep[for a review, see for example][]{2013A&ARv..21...70B}. Such events can sustain the energy flux needed to power Galactic CRs. In addition, there exists a theoretical model of an acceleration process at these shock fronts, known as diffusive shock acceleration (DSA)~\citep{1977DoSSR.234.1306K,1977ICRC...11..132A,1978MNRAS.182..147B,1978ApJ...221L..29B}, which provides a good explanation of the multiwavelength data of young SNRs. In the past decade, a number of young SNRs have been established by gamma-ray observations as accelerators reaching particle energies up to at least a few hundred TeV. It is difficult to achieve unequivocal proof, however, that these accelerated particles are protons, which emit gamma rays via the inelastic production of neutral pions, and not electrons, which could emit very high-energy (VHE; Energies $E > 100$\,GeV) gamma rays via inverse Compton (IC) scattering of ambient lower energy photons. For old SNRs, for which the highest energy particles are believed to have already escaped the accelerator volume, the presence of protons has been established in at least five cases. For W28, the correlation of TeV gamma rays with nearby molecular clouds suggests the presence of protons~\citep{2008A&A...481..401A}. At lower GeV energies, four SNRs (IC\,443, W44, W49B, and W51C) have recently been proven to be proton accelerators by the detection of the characteristic pion bump in the \emph{Fermi} Large Area Telescope (\fermi) data~\citep{FermiPion,2016ApJ...816..100J,W49BForth}. At higher energies and for young SNRs, this unequivocal proof remains to be delivered. It can ultimately be found at gamma-ray energies exceeding about 50~TeV, where electrons suffer from the so-called Klein-Nishina suppression~(see for example \citet{2013APh....43...71A} for a review), or via the detection of neutrinos from charged pion decays produced in collisions of accelerated CRs with ambient gas at or near the accelerator. \rxj~(also known as \gthree) is the best-studied young gamma-ray SNR~\citep{Hess1713a,Hess1713b,Hess1713c,FermiRXJ}. It was discovered in the \rosat\ all-sky survey~\citep{Pfeffermann} and has an estimated distance of 1\,kpc~\citep{2003PASJ...55L..61F}. It is a prominent and well-studied example of a class of X-ray bright and radio dim~\citep{Lazendic} shell-type SNRs\footnote{Another example with very similar properties is RX~J0852.0$-$4622 (Vela Junior); see \citet{VelaJrForth}.}. The X-ray emission of \rxj\ is completely dominated by a non-thermal component~\citep{Koyama,Slane,CassamXMM,UchiyamaChandra,2008ApJ...685..988T}, and in fact, the first evidence for thermal X-ray line emission was reported only recently~\citep{2015ApJ...814...29K}. Despite the past deep \hess\ exposure and detailed spectral and morphological studies, the origin of the gamma-ray emission (leptonic, hadronic, or a mix of both) is not clearly established. All scenarios have been shown to reproduce the spectral data under certain assumptions~(as discussed by \citet{2014MNRAS.445L..70G} and references therein). In addition to such broadband modelling of the emission spectra of \rxj, correlation studies of the interstellar gas with X-ray and gamma-ray emission are argued to show evidence for hadronic gamma-ray emission~\citep{2012ApJ...746...82F}. We present here new, deeper, \hess\ observations, analysed with our most advanced reconstruction techniques yielding additional performance improvements. After a detailed presentation of the new \hess\ data analysis results and multiwavelength studies, we update the discussion about the origin of the gamma-ray emission. \begin{figure*} \centering \begin{subfigure}{0.49\textwidth} \includegraphics[width=\textwidth]{Figs/ExcessONRegion.pdf} \end{subfigure} \begin{subfigure}{0.49\textwidth} \includegraphics[width=\textwidth]{Figs/Excess2TeV.pdf} \end{subfigure} \caption{\hess\ gamma-ray excess count images of \rxj, corrected for the reconstruction acceptance. On the left, the image is made from all events above the analysis energy threshold of 250~GeV. On the right, an additional energy requirement of $E>2$\,TeV is applied to improve the angular resolution. Both images are smoothed with a two-dimensional Gaussian of width $0.03^{\circ}$, i.e.\ smaller than the 68\% containment radius of the PSF of the two images ($0.048^{\circ}$ and $0.036^{\circ}$, respectively). The PSFs are indicated by the white circles in the bottom left corner of the images. The linear colour scale is in units of excess counts per area, integrated in a circle of radius $0.03^{\circ}$, and adapted to the width of the Gaussian function used for the image smoothing.} \label{fig:hess-maps} \end{figure*} | \label{sec:summary} The new \hess\ measurement of \rxj\ reaches unprecedented precision and sensitivity for this source. With an angular resolution of $0.048^\circ$ (2.9 arcminutes) above gamma-ray energies of 250 GeV, and $0.036^\circ$ (2.2 arcminutes) above energies of 2 TeV, the new \hess\ map is the most precise image of any cosmic gamma-ray source at these energies. The energy spectrum of the entire SNR confirms our previous measurements at better statistical precision and is most compatible with a power law with an exponential cut-off, both a linear power-law model at gamma-ray energies of 12.9\,TeV and a quadratic model at 16.5\,TeV. A spatially resolved spectral analysis is performed in a regular grid of 29 small rectangular boxes of $0.18^\circ$ (10.8 arcminutes) side lengths, confirming our previous finding of the lack of spectral shape variation across the SNR. The broadband emission spectra of \rxj\ from various regions are fit with present age parent particle spectra in both a hadronic and leptonic scenario, using \emph{Suzaku} X-ray and \hess\ gamma-ray data. From the resolved spectra in the 29 small boxes in the leptonic scenario, we derive magnetic field, energy cut-off, and particle index maps of the SNR. For the latter parameter, we do the same for the hadronic scenario. The leptonic and hadronic parent particle spectra of the entire remnant are also derived without further detailed assumptions about the acceleration process. These particle spectra reveal that the \fermi\ and \hess\ gamma-ray data require a two-component power-law with a break at 1-3\,TeV, challenging our standard ideas about diffusive particle acceleration in shocks. In either leptonic or hadronic scenarios, approaches more involved than one or two zone models are needed to explain such a spectral shape. Neither of the two scenarios (leptonic or hadronic), or a mix of both, can currently be concluded to explain the data unambiguously. Either better gamma-ray measurements with the future CTA, with much improved angular resolution and much higher energy coverage, or high sensitivity VHE neutrino measurements will eventually settle this case for \rxj. Comparing the gamma-ray to the \xmm\ X-ray image of \rxj, we find significant differences between these two energy regimes. As concluded before by~\citet{2008ApJ...685..988T}, the bright X-ray hotspots in the western part of the shell appear relatively brighter than the \hess\ gamma-ray data. The most exciting new finding of our analysis is that in some regions of \rxj\ the SNR is larger in gamma rays than it is in X-rays -- the gamma-ray shell emission extends radially farther out than the X-ray shell emission in these regions. We interpret this as VHE particles leaking out of the actual shock acceleration region -- we either see the shock precursor or particles escaping the shock region. Such signs of escaping particles are a longstanding prediction of DSA, and we find the first such observational evidence with our current measurement. | 16 | 9 | 1609.08671 |
1609 | 1609.08501_arXiv.txt | {}% {A direct search of $\gamma$-ray emission centered on multi-frequency selected candidates is a valuable complementary approach to the standard search adopted in current $\gamma$-ray Fermi-LAT catalogs. Our sources are part of the 2WHSP sample that was assembled with the aim of providing targets for Imaging Atmospheric Cherenkov Telescopes (IACT). A likelihood analysis based on their known position enabled us to detect 150 $\gamma$-ray excess signals that have not yet been reported in previous $\gamma$-ray catalogs (1FGL, 2FGL, 3FGL). By identifying new sources, we solve a fraction of the extragalactic isotropic $\gamma$-ray background (IGRB) composition, improving the description of the $\gamma$-ray sky. } % {We perform data reduction with the Fermi Science Tools using positions from 400 high synchrotron peaked (HSP) blazars as seeds of tentative $\gamma$-ray sources; none of them have counterparts from previous 1FGL, 2FGL and 3FGL catalogs. Our candidates are part of the 2WHSP sample (currently the largest set of HSP blazars). We focus on HSPs characterised by bright synchrotron component with peak flux $\nu f_{(\nu)} \geq 10^{-12.1}$ ergs/cm$^{2}$/s, testing the hypothesis of having a $\gamma$-ray source in correspondence to the WHSP positions. Our likelihood analysis considers the 0.3-500 GeV energy band, integrating over 7.2 yrs of Fermi-LAT observation and making use of the Pass 8 data release. } {From the 400 candidates tested, a total of 150 2WHSPs showed excess $\gamma$-ray signature: 85 high-significance detections with test statistic (TS)$>$25, and 65 lower-significance detections with TS between 10 to 25. We assume a power law spectrum in the 0.3-500 GeV band and list in Table \ref{tableFermi1} the spectrum parameters describing all 150 new $\gamma$-ray sources. We study the $\gamma$-ray photon spectral index distribution, the likelihood of detection according to the synchrotron peak brightness (figure of merit parameter), and plot the measured $\gamma$-ray LogN-LogS of HSP blazars, also discussing the portion of the IGRB that has been resolved by the present work. We also report on four cases where we could resolve source confusion and find counterparts for unassociated 3FGL sources with the help of high-energy TS maps together with multi-frequency data. The 150 new $\gamma$-ray sources are named with the acronym 1BIGB for the first version of the Brazil ICRANet Gamma-ray Blazar catalog, in reference to the cooperation agreement supporting this work. } {} | Catalogs of $\gamma$-ray sources currently compiled by the Fermi-LAT team are based on $\gamma$-ray data only, and their standard detection method is blind with respect to information coming from other wavelengths. This approach is clean and unbiased with respect to any class of potential $\gamma$-ray emitters. However, there are populations of astrophysical objects that are now known to emit $\gamma$-rays, and the knowledge of their position in the sky can be used to facilitate the detection and identification of new $\gamma$-ray sources. Based on this principal, we select a sample of candidates to be used as seeds for a direct search of $\gamma$-ray signatures using likelihood analysis with the Fermi Science Tools. Blazars are the most abundant $\gamma$-ray sources in the latest Fermi-LAT 3FGL catalog, being 1147 (660 BL Lacs and 487 Flat Spectrum Radio Quasars - FSRQ) of the total 3034 \citep{3FGL}. Even so, one third of the known blazars from 5BZcat\footnote{The 5BZcat \citep{5BZcat} is a large sample of 3 561 identified blazars. Multi-frequency data for the 5BZcat is available at http://www.asdc.asi.it/bzcat with a direct link to the SED-builder tool.} are not confirmed as $\gamma$-ray emitters. Probably many of them are faint $\gamma$-ray sources that are hard to identify by automatic search methods only based on Fermi-LAT data. The blazar population has been extensively studied by means of a multi-frequency approach considering dedicated databases on radio, microwave, infra-red (IR), optical, ultra-violet (UV), and X-ray, since they are characterised by radiation emission extending along the whole electromagnetic spectrum, up to TeV energies. A particular family of extreme sources with the synchrotron component peaking at frequencies $\nu_{peak}$ larger than 10$^{15}$ Hz is classified as a high synchrotron peak blazar \citep[HSP,][]{padgio95,abdo10} and is the dominant population associated with extragalactic very high-energy \citep[VHE: E$>$0.1 TeV,][]{TeVAstronomy} sources in the 2nd Catalog of Hard Fermi-LAT Sources \citep[2FHL,][]{2FHL}. Therefore, HSPs constitute a key population for the detection of point-like $\gamma$-ray sources within Fermi-LAT data. A large sample of HSP blazars was recently assembled using a multi-frequency selection procedure that exploits the unique features of their spectral energy distribution (SED). This sample is known as the 1WHSP catalog \citep{1WHSP} and was built using a primary source-selection based on IR colours \citep[following][]{WiseBlazars}, later demanding all potential candidates to have a radio, IR and X-ray counterpart. The sources had to satisfy broadband spectral slope criteria (from radio to X-rays) that were fine-tuned to match the SED of typical HSP blazars. In addition, their multi-frequency SEDs were inspected individually using the SED-builder tool (http://www.asdc.asi.it) fitting the synchrotron component with a third degree polynomial to determine the $\nu_{peak}$ parameter, only keeping cases with $\nu_{peak}>10^{15}$ Hz. The catalog name ``WHSP" stands for WISE High Synchrotron Peak blazars, since all sources have an IR counterpart from the WISE mission \citep{WISE}, which defines their positions. The 1WHSP catalog includes 992 objects at Galactic latitude $|b|>20^{\circ}$. A total of 299 1WHSPs have a confirmed $\gamma$-ray counterpart in 1FGL, 2FGL and 3FGL \citep{1WHSP}, but many HSPs with bright synchrotron peak are still not detected/confirmed in the $\gamma$-ray band. Given the importance of finding new HSP blazars, an extension of the 1WHSP sample \citep[the 2WHSP,][]{2WHSP} has ben assembled. It considers sources located at latitudes as low as $|b|=10^{\circ}$ with a total of $\approx$1 693 sources, 439 of which have counterparts within the error circles from the 3FGL catalog. The 2WHSP sample avoids the selection based on IR colours that was used as a primary step for the 1WHSP catalog. This brings an overall improvement in completeness\footnote{ \label{b10} Also to improve the completeness of the final sample, known HSP sources at $|b|<$10$^{\circ}$ were incorporate in the 2WHSP catalog.}, since some HSP blazars were out of the 1WHSP sample owing to the contamination of IR colours by the elliptical-galaxy thermal emission. Compared to the 1WHSP, the 2WHSP sample incorporates extra X-ray catalogs like Einstein IPC, IPC slew and Chandra \citep{Harris1993,Elvis1992,Evans2010} as well as updated versions from 3XMM-DR5 and XMM-slew catalogs \citep{Rosen2015,Saxton2008}. In addition, Swift-XRT alone performed a series of $\sim$160 new X-ray observation of WHSP sources (enabling us to better estimate synchrotron peak parameters) and an extensive study of X-ray extended sources helped to avoid contamination with spurious objects (more details are given by \cite{2WHSP}). The catalogs are available at: www.asdc.asi.it/1whsp or /2whsp; where multi-frequency SEDs can be quickly built using open access online tools. Since the 2WHSP catalog supersedes the 1WHSP (with improved selection and better estimate of synchrotron peak parameters), from now on we only refer to the 2WHSP sample. | The 2WHSP catalog was built to select promising VHE candidates for the present and future generation of Cherenkov Telescope Arrays, therefore we have tested the efficiency of a direct search for $\gamma$-ray signatures associated with 2WHSP blazars, achieving significant results. We have detected 150 $\gamma$-ray excess signals out of 400 seed positions based on 2WHSP sources that had no counterpart in previous 1FGL, 2FGL, and 3FGL catalogs. A total of 85 sources were found with high-significance with TS$>$25, and we also report on 65 lower-significance detections with TS between 10 to 25. The 150 new $\gamma$-ray sources presented in Table \ref{tableFermi1} are named with acronym 1BIGB (first version of the Brazil ICRANet Gamma-ray Blazar catalog) which corresponds with the 2WHSP seed-positions used for our likelihood analysis. Clearly, the subsample of 2WHSP blazars that have not yet been detected by Fermi-LAT is a key representative population of faint $\gamma$-emitters, and we show how the new detections down to TS$>$10 level can probe the faint-end of the flux-distribution (see Fig. \ref{GammaFlux} and \ref{hist-flux}). As discussed in Sect. \ref{faintdetec}, a $\gamma$-ray source-search based on the seed positions from HSP blazars can be used to unveil faint HE sources down to TS=10 without compromising the $\gamma$-ray sample with spurious detections. Our current work enabled us to associate a relevant fraction of the IGRB to a population of faint $\gamma$-ray emitters that had been previously unresolved. Moreover, we show the increasing relevance of faint-HSPs for the IGRB composition with respect to energy (see Table \ref{IGRB}), specially for E$>$10 GeV, reaching 6-8$\%$ in the 100-200 GeV band. Motivated by this first assessment, we plan to perform a complete $\gamma$-ray analysis of the 2WHSP sample, down to the lowest fluxes, and probably extend the search to other blazar families with potential to improve the $\gamma$-ray description of lower-significance $\gamma$-ray blazars, also helping to constrain the origins of the extragalactic diffuse $\gamma$-ray background. We have worked out the possibility of solving source confusion when considering multi-frequency data for identifying potential $\gamma$-ray emitters in a certain ROI, and building energy dependent TS maps to help disentangle hard-steep components from confused sources. We also addressed cases of unassociated 3FGL sources by studying high-energy TS maps to evaluate possible counterparts. This could be a key for solving cases of unassociated $\gamma$-ray sources (just as discussed in Sect. \ref{un1}, \ref{un2} and \ref{un3}) showing that we can improve the $\gamma$-ray signature localization based on currently available databases. Certainly, it is interesting to evaluate if this kind of approach could be applied systematically as a complementary refinement for the building of upcoming $\gamma$-ray catalogs. | 16 | 9 | 1609.08501 |
1609 | 1609.01420_arXiv.txt | Visible-light observations of Coronal Mass Ejections (CMEs) performed with coronagraphs and heliospheric imagers (in primis on board the SOHO and STEREO missions) have offered so far the best way to study the kinematics and geometrical structure of these fundamental events. Nevertheless, it has been widely demonstrated that only combination of multi-wavelength data (including X-ray spectra, EUV images, EUV-UV spectra, and radio dynamic spectra) can provide complete information on the plasma temperature and density distributions, non-thermal motions, magnetic fields, and other physical parameters, for both CMEs and CME-related phenomena. In this work, we analyze three CMEs by combining simultaneous data acquired in the polarized visible light by the LASCO-C2 coronagraph and in the UV \hi\ \lya\ line (1216~\AA) by the UVCS spectrometer, in order to estimate the CME plasma electron density (using the polarization-ratio technique to infer the 3D structure of the CME) and temperature (from the comparison between the expected and measured \lya\ intensities) along the UVCS field of view. This analysis is primarily aimed at testing the diagnostic methods that will be applied to coronagraphic observations of CMEs delivered by the Metis instrument on board the next ESA-Solar Orbiter mission. We find that CME cores are usually associated with cooler plasma ($T \sim 10^6$~K), and that a significant increase of the electron temperatures is observed from the core to the front of the CME (where $T > 10^{6.3}$~K), which seems to be correlated, in all cases, with the morphological structure of the CME as derived from visible-light images. | During major solar eruptions (or Coronal Mass Ejections -- CMEs), huge bubbles of highly ionized plasma (often associated with shock waves and Solar Energetic Particles -- SEPs) expand into the interplanetary space, affecting a significant fraction of the whole heliosphere, and eventually propagating even out of the heliopause as interstellar shocks in the most dramatic cases, as recently reported with Voyager~1 observations \citep{gurnett2015}. Solar transients (flares, CMEs and prominence eruptions) have an impact on all the planetary objects, interacting with their magnetospheres or magnetospheric-like structures and inducing geomagnetic storms \citep[see review by][]{akasofu_2011}, as well as beautiful auroras on Earth and other planets \citep{hultqwist_2008}. These events likely played a role even in the development of life on Earth, for instance by modulating the rate of galactic cosmic rays impacting on the early Earth's atmosphere \citep[via the well-known ``Forbush decrease'' effect; see][]{lockwood1071} and on the atmospheric chemistry \citep{airapetian2016}. Moreover, the study of these events is very important also from a theoretical point of view, because in order to understand their origin and interplanetary evolution it is necessary to consider many different plasma physical processes (such as plasma instabilities, magnetic reconnections, wave-particle interactions, etc.), phenomena that are only partially understood. From the observational point of view, the study of solar eruptions can be performed using many different data delivered daily by both ground- and space-based observatories. Nevertheless, images of the solar disk can provide only information on the location of the source region, on the eruption start time and on the early expansion phases. Then, after the eruptions take off and leave the Sun, their subsequent evolution during the expansion and propagation phases can be followed only with two classes of instruments: space-based coronagraphs and heliospheric imagers. In fact, these are the only instruments covering the huge amount of space traveled by solar eruptions during their propagation from the Sun to the Earth and beyond. Without data provided by space-based coronagraphs and heliospheric imagers it would be simply impossible to characterize the real CME propagation angle and CME speed, and to investigate the physical processes occurring during their interplanetary expansion. It is well known today that, after the main impulsive acceleration phase occurring in the lower corona, solar eruptions are subject to many different processes affecting their evolution, which is never like a simple radial and self-similar expansion. During their early propagation phases, CMEs are often channelled by coronal streamers and/or deflected away from nearby coronal holes \citep[see][and refeences therein]{mostl2015}, or towards the interplanetary current sheet \citep[e.g.,][]{byrne2010,isavnin2014}; these interactions may modify their propagation directions up to 25$^\circ$ with respect to the location of the source region \citep{kay2013}, influencing the strength of the eventual impact on Earth. Significant rotations of CMEs around their propagation axis are also observed in many cases \citep[e.g.,][]{thompson2012,bemporad2011}, which change the orientation of the magnetic field of the associated magnetic cloud impacting on the Earth's magnetosphere, and in turn their capability to induce geomagnetic storms. Moreover, significant magnetic drag occurs during the interplanetary propagation of CMEs, leading to further accelerations or decelerations of the ejecta (depending on the expansion speed relative to the ambient solar wind) that affect the expected arrival times at Earth \citep[e.g.][]{iju2014,temmer2011}. Furthermore, when multiple eruptive events are ejected in sequence, CME-CME interactions may occur increasing their final geoeffectiveness \citep[see][]{far06,wu07}. All these phenomena make CME observations by coronagraphs and heliospheric imagers crucial for understanding these events and forecasting their impact on the Earth. Moreover, previous experience with visible-light (VL) coronagraphs shows that unique information can be derived only when data acquired at different wavelengths are combined together, i.e., not only images in the VL, but also X-ray spectra, EUV images, EUV-UV spectra, and radio dynamic spectra. In particular, the combination of observations acquired in the VL by different coronagraphs and in the UV spectral range by the UVCS spectrometer \citep{kohl1995} on SOHO, allowed to characterize the distribution of plasma temperatures (electron and ion) and their evolution inside the CME core and front, to study many CME-related phenomena, such as post-CME current sheets and CME-driven shocks, and to reveal the three-dimensional (3D) CME structure \citep[see][for an extensive review of these results]{koh06}. More recently, UVCS spectra have been combined for the first time with SOHO and STEREO VL images to perform the first stereoscopic and spectroscopic reconstruction of a CME \citep{susino14}, to derive the physical parameters of coronal plasma -- including the magnetic field -- across CME-driven shocks \citep[e.g.,][]{bemporad14,susino15}, and to derive kinetic temperature, gas pressure, and filling factor in erupting prominences \citep{heinzel16}. All these works demonstrate the importance of complementarity of VL and UV observations of solar eruptions. In the near future, combined VL and UV images will be provided by the Metis coronagraph \citep{antonucci12,fineschi12,romoli14} on board the ESA-Solar Orbiter mission, due to launch in October 2018. The Metis instrument will acquire the first-ever simultaneous observations of the solar corona in the polarized visible light (broad-band 580-640~nm) and in the UV (narrow-band around the \hi\ \lya\ 1216~\AA\ line), even if, since 2012, the instrument has lost its spectroscopic capabilities. Hence, the aim of the present work is to perform the first tests on the diagnostic capabilities for CMEs that will be possible with future Metis data. This test is performed here using available simultaneous observations of real CMEs acquired in the polarized visible light by the SOHO/LASCO-C2 coronagraph \citep{brueckner95} and in the UV by the UVCS spectrometer. Because the Metis coronagraph will not provide spectroscopic observations, in this work we focused only on the observed evolution of UV \lya\ intensities, thus simulating the information that will be provided by Metis. The paper is organized as follows: after a description of the selected events and datasets (Section 2), we describe the diagnostics (Section 3) we applied for the determination of CME electron density (Section 3.1) and electron temperature (Section 3.2), and then we discuss our results (Sections 4 and 5). | In this work we demonstrate how future observations of CMEs that will be provided by Metis coronagraph on-board Solar Orbiter will be analyzed to infer physical parameters of the plasma involved in the eruption and to derive, in particular, the electron temperature of the ejected gas. To this end, we first selected eruptive events observed at the same times and at the same coronal locations both in LASCO-C2 polarized-brightness images, and in the UV \hi\ \lya\ line by UVCS; it turns out that over the whole era of UVCS, only three events were sampled by both instruments. Then, we analyzed VL data and showed that they can be used to derive not only the electron column density and volumetric density inside the CME, but also the average location of the emitting plasma along the LOS. This in turn can be used, on one hand to better constrain the plasma electron densities, on the other hand, to measure the CME propagation direction with respect to the POS and to derive the unprojected CME speed at different latitudes (i.e., along the UCVS slit). The unprojected speeds can be converted into Doppler-dimming factors, the missing ingredient needed to combine VL with UV intensities. Given the electron densities and Doppler-dimming factors, the combination of VL and UV intensities provides an estimate of CME plasma electron temperatures. As we showed here, for the three events we selected the unprojected speeds were so small that the \lya\ emission is still dominated by the radiative component. Nevertheless, we expect that for major and faster CMEs the situation could be even reversed, with the \lya\ emission being dominated by the collisional component, in particular in the denser parts such as the CME cores. In this work we found that the CME cores are usually associated with cooler plasma, and that a significant rise of temperatures is observed moving from the core to the front of the CME. The determination of electron temperatures inside CMEs is of crucial importance. In fact, one of the main problems left open after the UVCS era is the real evolution of the CME thermal energy during their expansion. Different authors found that during the expansion additional heating sources need to be considered in order to reproduce the observed UV emission, with heating rates comparable \citep{akm01, murphy2011} or even larger \citep{lee2009, landi2010} than kinetic and potential energies carried by the CME. Furthermore, \citet{bem07} demonstrated with the only existing multi-slit study of a CME based on UVCS data, that the CME plasma temperature is increasing during the expansion, implying again the existence of an additional thermal-energy source. Same results have been recently confirmed by \citet{lee2015} using Hinode/XRT images. Nevertheless, a clear interpretation for the source of this additional thermal energy is missing so far. In this work we demonstrate how CME electron temperatures can be derived using VL images and UV \lya\ intensities. Nevertheless, being limited here to the one dimensional FOV of the UVCS slit and to the single time when both VL $pB$ and UV \lya\ emissions were observed, it was not possible to study neither the CME plasma temperature distribution within the whole CME bubble, nor its time evolution during the CME expansion. Full investigation of these aspects will be possible thanks to future data that will be provided by the Metis coronagraph on board the Solar Orbiter. In fact, sequences of VL and UV images that will be acquired at the same time, will allow to study both the thermal energy distribution within the CME bubble at a given time, and its evolution during the CME propagation. Further analysis on the kinematics and 3D structure of CMEs will be possible also thanks to the synergies between Metis and other next-generation coronagraphs, such as ASPIICS \citep[][]{renotte15} on board the ESA Proba-3 mission. Moreover, for events that will be observed during quadratures also by other spacecrafts (such as Solar Probe Plus), the combination of these information with \textit{in situ} measurements made close to the Sun will allow to tightly constrain the temporal evolution of thermal energy of the ejected plasma during its early interplanetary propagation, thus letting connections with the still open issue of \textit{in situ} detections of high-ionization states of heavy ions in interplanetary CMEs. | 16 | 9 | 1609.01420 |
1609 | 1609.03425_arXiv.txt | We analyze \Chandra\ X-ray images of a sample of 11 % quasars that are known to contain kiloparsec scale radio jets. The sample consists of five high-redshift ($z\geq3.6$) flat-spectrum radio quasars, and six intermediate redshift ($2.1<z<2.9$) quasars. The dataset includes four sources with integrated steep radio spectra and seven with flat radio spectra. A total of 25 radio jet features are present in this sample. We apply a Bayesian multi-scale image reconstruction method to detect and measure the X-ray emission from the jets. We compute deviations from a baseline model that does not include the jet, and compare observed X-ray images with those computed with simulated images where no jet features exist. This allows us to compute $p$-value upper bounds on the significance that an X-ray jet is detected in a pre-determined region of interest. We detected 12 of the features unambiguously, and an additional 6 marginally. We also find residual emission in the cores of 3 quasars and in the background of 1 quasar that suggest the existence of unresolved X-ray jets. The dependence of the X-ray to radio luminosity ratio on redshift is a potential diagnostic of the emission mechanism, since the inverse Compton scattering of cosmic microwave background photons (IC/CMB) is thought to be redshift dependent, whereas in synchrotron models no clear redshift dependence is expected. We find that the high-redshift jets have X-ray to radio flux ratios that are marginally inconsistent with those from lower redshifts, suggesting that either the X-ray emissions is due to the IC/CMB rather than the synchrotron process, or that high redshift jets are qualitatively different. | Jets in active galactic nuclei transfer the energy generated by the central supermassive black hole (SMBH) to large( $>100$ kpc) distances. The impact of jets on the environment contributes to the formation and evolution of structures in the early Universe \citep{croton2006}. The innermost jets (parsec-scales or smaller) of radio-loud quasars are highly relativistic and their observed radiation can be Doppler amplified when observed at small angles to the line of sight. These jets can be bright at high-energies and thus can provide interesting observational probes of the state of the SMBH activity \citep{begelman84}, but they remain spatially unresolved in X-rays and $\gamma$-rays. Large scale X-ray jets span distances out to hundreds of kiloparsecs away from the SMBH and encode the history of SMBH activity during the jet's lifetime (a few Myrs). Their X-ray emission can be resolved with the \Chandra\ X-ray Observatory. The number of such X-ray jets has significantly increased since the launch of \Chandra\ in 1999, but it is still relatively small in comparison to the number of known quasars. There are about 100 large scale X-ray jets detected to date and only a few of them have good quality X-ray morphology data \citep{mas11}. Though a direct connection between SMBH activity and the existence of kpc-scale jets is ambiguous, and the X-ray emission mechanism is not well understood, high-redshift ($z> 3$) jets could potentially establish the dominant energy environment in the early Universe. Such jets probe the physics of the earliest (first $\sim$2~Gyr of the Universe in the quasars studied) actively accreting SMBH systems and are also interesting for other reasons. For instance, the ambient medium in these high-redshift galaxies is probably higher than that of lower redshift galaxies \citep[e.g.,][]{dey06} and this may manifest itself in jets with different morphologies, with increased energy dissipation, or with the jets being slower in general than their lower-redshift counterparts. The X-ray radiation could be attributed to either synchrotron emission by highly relativistic electrons (Lorentz factors of $\gamma \sim 10^7-10^8$) in relatively strong magnetic fields, or inverse Compton scattering of the cosmic microwave background (IC/CMB) photons off the low energy ($\gamma \sim 10^3$) large-scale jet particles \cite[for a review see][]{har06}. In the simplest scenario, such models have diverging predictions at high redshift. Specifically, we expect a strong redshift dependence in the X-ray--to--radio- energy flux ratio, $\ratxr = \frac{F_{\rm x}}{\nu_{\rm r} f_{\rm r}}$, where the radio fluxes are given at the observed frequency, and the X-ray fluxes are modeled over the energy range $0.5-7$~keV \citep{mas11}. Typically, $\ratxr \propto U_{\rm CMB} \propto~(1+z)^{4}$ for IC/CMB, whereas in synchrotron models, we do not expect a strong dependence,\footnote{The energies of the synchrotron emitting electrons are different in the observed radio and X-ray spectra, $\gamma_{\rm r} \sim 10^3$ vs. $\gamma_{\rm x} \sim 10^7$ and these electrons may originate in the same population or two different populations. Therefore, there could be some weak redshift dependence in the synchrotron model.} $\ratxr \propto (1+z)^{0}$. Below we compare the predictions of these two models for the highest-redshift relativistic jets. Most \Chandra\ studies of quasar jets have so far targeted known arcsecond-scale radio jets \citep[e.g.,][]{sam04,mar11}, as most known examples are at $z$ $\stackrel{<}{{}_\sim}$2 \citep{bri84,liu02}. At the time our program began, there were two high-$z$ quasars with kpc-scale X-ray jets: GB~1508$+$5714 at $z=4.3$ \citep{sie03,yua03,che04} and 1745$+$624 at $z=3.9$ \citep{che06}. They were observed to have large $\ratxr$ values consistent with the IC/CMB model \citep{sch02,che04}, although the small number of high-$z$ detections precluded any definitive statements \citep{kat05,che06}. We have therefore obtained \Chandra\ X-ray observations of an additional four high-redshift ($z>3.6$; GB\,1508$+$5714 was previously analyzed) and six intermediate-redshift ($2 \leq z \leq 3$) quasars with known radio jets. The highest redshift X-ray and radio jet discovered in the sample studied (at $z=4.72$, in GB~1428+4214 (1428+422)) was presented and discussed in detail by \cite{che12}. New and archival arcsecond-resolution NRAO\footnote{ The National Radio Astronomy Observatory is operated by Associated Universities, Inc.\ under a cooperative agreement with the National Science Foundation.} imaging observations of these quasars are also presented. The small number of X-ray photon counts observed from jets relative to their corresponding quasar cores means that detecting X-ray jets is inherently challenging. Statistically, we must test the hypothesis that a baseline model of the quasar core and a flat background, without a jet, is insufficient to explain the observed data. We do this test using a multi-scale Bayesian method known as Low Count Image Reconstruction and Analysis\footnote{ LIRA is implemented as a package for the R statistical programming language (r-project.org) that is available for downloading and use at {\tt github.com/astrostat/LIRA}. } \citep[LIRA;][]{esch04, con07}. The algorithm models the residual as a multi-scale component, and generates a series of images that capture the emission that may be present in excess of the baseline model. We can then compute a $p$-value\footnote{ Formally, a $p$-value is the probability that the baseline null hypothesis can generate a value for the test statistic as large as that which is observed. In this case, it defines the likelihood that a given intensity can be obtained under the assumption that the baseline model is the truth. That is, when the $p$-value is small, the chances that the feature under consideration can be attributed to a fluctuation is small. This allows us to {\sl reject} the null hypothesis when this probability falls below a pre-defined threshold. Note, however, that it should never be interpreted as a measure of the probability that the alternate hypothesis is true, nor, if the null cannot be rejected, as a measure of the probability that the null hypothesis is true \citep{wl16}. } by generating a series of Monte Carlo simulations of images under the baseline model and fitting each of these simulated images using LIRA. \cite{stei15} (hereafter Paper~I) show how an upper-bound on the $p$-value can be computed with a small number of MCMC replicates. We are interested in detecting whether X-ray jets exist in regions where jets were previously observed in the radio band. In this paper, we will not consider X-ray detections without a corresponding radio emission (such a detection of a jet that was recently reported by \cite{simionescu2016}) when matching our results to the IC/CMB or synchrotron emission model. We run LIRA to detect jets in pre-defined regions of an X-ray image. Using the jets detected in X-rays we are able to observe how $\ratxr$ is dependent on redshift and whether it matches the predictions of the IC/CMB or synchrotron emission model. Section 2 describes the sample selection and initial processing of the X-ray and corresponding radio observations. Section 3 outlines how LIRA is used to find evidence that a jet exists in a region where one is observed in radio imaging. Section 4 elaborates on the results of the image analysis methods when applied to the new X-ray observations. Section 5 gives a final description of our results in context and we summarize our results in Section 6. | \subsection{X-ray Morphology} The posterior mean images $\tau_1 \Lambda_1$ representing deviations from the baseline model of the \Chandra\ data (panel (c) of Figure~\ref{fig:10307} and of Figures~\ref{fig:10308}--\ref{fig:2241} in Appendix B provide a view of the X-ray morphology with the quasar core removed. However, not all the structures seen in these images are significant, since many features could be attributed to statistical fluctuations. We have developed a method (Paper~I; see also Section 2 above) to assess the significance of the emission for well-defined ROIs. Note that these ROIs {\sl must be set prior to the analysis} and cannot be deduced from the LIRA output, since doing so would increase the false detection rate. We adopt regions based on the locations of the radio jet features, but note that the jet X-ray emission may not always be spatially coincidental with the radio emission \citep{sch02}, nor even have a radio counterpart \citep{jor04, simionescu2016}. Offsets between the radio and X-ray peaks in the jet features have been reported \cite[e.g.][]{sie2007, worrall2009}. For instance, in 1754$+$676 (ObsID 7872), the X-ray jet is not detected, but the image showing deviations from the baseline (see panel (c) of Figure~\ref{fig:7872}) suggests the existence of an emission feature between the quasar and the radio jet region. A longer \Chandra\ observation is necessary to confirm this emission. In another source 0805$+$046 (ObsID 10308), the image showing deviations from the baseline (panel (c) of Figure~\ref{fig:10308}) displays considerable emission outside the narrow radio jet, suggesting a more complex X-ray morphology. We used the complementary regions to assess the possibility of unexpected X-ray emission present outside the pre-defined regions. In all sources but 1508$+$5714, we do not find a strong indication that such emission is present, though the complementary regions cover a large area and thus statistical tests have relatively low power to detect smaller compact structures. Future studies of the X-ray morphology in the vicinity of this source is required for understanding the origin of this emission. The $p$-value upper bound test relies on the test regions being pre-defined. This is done in order to avoid the loss of power in the test that arises when multiple hypotheses are tested. We thus take the regions directly from the radio data and do not optimize the regions based on the X-ray data. This could result in the size of location of the regions to be slightly misaligned, reducing the significance of detection. For example, ROI 1 in source 1428$+$422 (ObsID 7874) is a marginal detection with an upper bound on the $p$-value of 0.010, but decreasing the region size from 91 to 77 pixels, an arguably better fit for this object in radio, results in an improvement in the $p$-value upper bound to 0.009, which crosses the threshold into a significant detection. Areas of deviation from the baseline may be difficult to detect in a large region or a region may not encompass all of the relevant area of the image, and the better the region fits an area with deviation from the baseline, the lower the nominal $p$-value is. But when large numbers of regions are tested, the $p$-value threshold must be reduced correspondingly in order to prevent false claims of detections due to fluctuations. For instance, if the $\xi$ in 20 regions (say) of radii stepping from 75 to 95 pixels are tested, the appropriate threshold of $\alpha$ must be reduced by a factor of 20, to $\alpha=0.0005$, to maintain the same level of significance. We emphasize that ROI selection must be consistent across the analysis, and {\sl must be defined before} applying the significance test. We also note that we apply the test to a total of 47 ROIs, so we expect at most 1 false positive amongst the claimed significant and marginal detections at the significance threshold of $\alpha<0.02$. As noted above, our current method of region selection depends on radio data. In the future, we plan to develop methods that are independent of the radio (or other wavebands) selections and autonomously generate regions that adaptively fit the deviations from the baseline in the LIRA output. Such a method is needed since the X-ray emission does not always follow the radio closely. Naturally, any such method will trade-off ROI optimization for the statistical power of the detection routine. \subsection{Redshift Dependence in Large-Scale X-ray/Radio Emission} The origin of X-ray jet emission is still under debate. An early hint at the advantage of studying high-redshift jets came from the $z$=4.3 quasar 1508$+$5714 (Siemiginowska et al.\ 2003; Yuan et al.\ 2003). This quasar has higher X-ray to radio luminosity ratio ($\ratxr>$100) than any of its lower-$z$ counterparts (Cheung 2004). This appears consistent with the $(1+z)^{4}$ amplification in the energy density of Cosmic Microwave Background (CMB): \begin{equation} \ratxr \propto u_{\rm CMB}/u_{\rm B} \propto (1+z)^{4}(\delta/B)^{2} \,, \end{equation} as expected under the IC/CMB model (e.g., Schwartz 2002). Figure~\ref{fig:fluxRatio} compares the energy flux ratio $\ratxr = [F_{\rm x}~{(0.5-7\rm keV)}]/[\nu_{\rm r} f_{\rm r}]$ of the detected and marginally detected jets across redshift from our \Chandra\ sample. We seek to establish whether or not the energy flux ratio varies with redshift. Figure~\ref{fig:density} shows the posterior distribution from LIRA of the energy flux ratio for each detected and marginally detected source (blue corresponds to low, and red to high redshift). It is visually apparent that there is a difference in the distributions of the sources with higher redshifts ($z>3$) and those with lower redshifts. In order to establish a statistical measure of the significance of this difference, we split the detected and marginally detected sources into two samples consisting of the 18 at low redshift ($z<3$; sample $L$) and the 3 at high redshift ($z>3$; sample $H$). We then use a hierarchical Gaussian model to examine whether low and high redshift quasars differ in terms of the mean and variance of their energy flux ratio. Appendix A describes a procedure for evaluating the posterior probability that the difference between the mean $\log_{10}$ energy flux ratio of the high and low redshift jets ($\mu_H - \mu_L$) is greater than zero. Figure~\ref{fig:mustat} shows the distribution of $\mu_H-\mu_L$ calculated from the posterior output. We find an empirical probability of $95\%$ that $\mu_H-\mu_L \geq 0$, which is at best marginal evidence that the observed difference cannot be due to a statistical fluctuation. Though highly suggestive, because of the small number of sources represented in this paper, and given the disproportionate numbers of jets in the two samples, there is insufficient evidence to conclude that the mean $\log_{10}$ energy flux ratio differs between two groups of jets. More observations at $z>3$ are required in order to obtain more reliable results. \begin{figure*}[ht] \centering \includegraphics[width=0.9\columnwidth]{fluxPlot_final.png} \caption{ Ratio of X-ray to radio flux $\ratxr$ vs redshift for the detected and marginally detected regions of interest of the jets. The circles are the energy flux ratios from jets detected in this study. The diamonds are estimated energy ratios from the marginally detected jets. The error bars form the $68\%$ interval from the LIRA iterations. } \label{fig:fluxRatio} \end{figure*} \begin{figure*}[ht] \centering \includegraphics[width=0.7\columnwidth]{fluxRatio_density.png} \caption{ The posterior distribution of the ratio of X-ray to radio fluxes, $\ratxr$, for each detected and marginally detected jet. The energy flux ratio is calculated at every iteration of LIRA. The color corresponds to redshift. } \label{fig:density} \end{figure*} \begin{figure*}[ht] \centering \includegraphics[width=0.8\columnwidth]{muPlot.png} \caption{ The difference in the mean X-ray to radio energy flux ratio between high- and medium-redshift quasars, $\mu_H - \mu_L$ at every LIRA iteration. $\mu_L$ and $\mu_H$ are the average log flux ratio across the lower-redshift ($2<z<3$) and higher-redshift ($z>3$) detected and marginally detected redshift jets. } \label{fig:mustat} \end{figure*} The difference between the jets in two redshift groups is interesting because it can also indicate that radio-loud quasars at $z>3$ are different from low redshift ones. \cite{volonteri2011} hypothesized that the jets at $z>3$ are systematically slower in comparison to the jets at $z<3$. If true, this could affect the energy flux ratio in the framework of the IC/CMB model, due to a strong dependence of this ratio on the jet Doppler factor. In this case, assuming that the comoving jet magnetic field is roughly the same at different redshifts, the observed increase in the energy flux ratio should be smaller than that expected from the $(1+z)^4$ scaling. However, the jet magnetization may evolve with redshift \citep[e.g.][]{singal2013} which complicates the redshift scaling even further. Our results indicate that the high redshift jets are different, but more observations are needed to study the origin of this difference. \subsection{Quasars at High Redshift} The X-ray emission of radio loud quasars observed with {\it Chandra} is unresolved and contained within $< 1.5\arcsec$ circular regions. This emission could be due to a mixture of at least three components: a hot corona directly related to the accretion process, a parsec scale jet, and an unresolved portion of the kpc-scale outflow emitting X-rays via IC/CMB. We measured a standard range of photon indices for the assumed power law model for quasar core spectra and found them consistent with either process. However, we detected relatively low values of photon indices in a few quasars, including two at the highest redshift (see Table~\ref{tbl:quasar2}). Lower values of the photon index are predicted if the jet dominates the X-ray emission. In this case the beamed jet would make the quasars to appear more luminous. We notice that such a trend is present in our small sample and five sources with $\Gamma_X < 1.6$ are more luminous, with an average 2-10~keV luminosity of $6.3 \pm 2.5 \times (10^{46}$erg~s$^{-1})$, than the other six with $\Gamma_X>1.6$ and the average luminosity of $7.9 \pm 2.5 \times (10^{45}$erg~s$^{-1})$. Formal correlation tests show that $\Gamma_X$ and ${\log}L_X$ are indeed correlated, with Pearson's correlation coefficient $\rho=0.61$ ($p=0.047$), and Kendall's $\tau=0.53$ ($p=0.024$). Such relation has been seen in analysis of large samples of radio-loud and radio-quiet quasars \citep{bechtold1994, young2009,lanzuisi2013}. \cite{bechtold1994} found a similar trend in a sample of radio loud quasars and argued that it could be caused by an increased absorption. On the other hand \cite{young2009} did not find a significant correlation in a radio-loud subsample of quasars observed with XMM-Newton. Future studies of large number of radio loud quasars in X-ray and radio band are necessary for understanding the presence and origin of this correlation. In the analysis of the {\it Chandra} images we assumed that the quasar emission is point-like. However, we detect signatures of non-point--like emission in the two highest redshift ($z>4$) quasars (1428+422 and 1508+5714) and in the one at $z=2.3$ (1834+612) (as evidenced by the fact that the core component Q is not fully accounted for in the baseline model, see Table~\ref{tbl:jets} and Figures~\ref{fig:10311},\ref{fig:7874},\ref{fig:2241}). This is unlikely to be due to uncertainties in the shape of the \Chandra\ PSF, since the residual core emission is not present in the multi-scale components for other sources. We suggest that these are due to non-negligible contributions from unresolved kpc-scale jets emitting IC/CMB. The effect of this residual component on our analysis is conservative, i.e., the imperfect modeling of the core tends to increase the strength of the baseline model and systematically dampen the added multi-scale component. Thus, our results indicate that the unresolved X-ray cores of radio loud high redshift quasars can contain significant contributions from kpc-scale jet emission. This result is in agreement with studies based on the optical-to-X-ray luminosity ratio \citep{saez2011,wu2013} and along the expectations from the IC/CMB model. Such a jet contribution can potentially bias population studies of quasars and needs to be taken into account in investigations of radio loud quasars at high redshift. | 16 | 9 | 1609.03425 |
1609 | 1609.04205_arXiv.txt | The Amsterdam-ASTRON Radio Transients Facility And Analysis Center (AARTFAAC) all sky monitor is a sensitive, real time transient detector based on the Low Frequency Array (LOFAR). It generates images of the low frequency radio sky with spatial resolution of 10s of arcmin, MHz bandwidths, and a time cadence of a few seconds, while simultaneously but independently observing with LOFAR. The image timeseries is then monitored for short and bright radio transients. On detection of a transient, a low latency trigger will be generated for LOFAR, which can interrupt its schedule to carry out follow-up observations of the trigger location at high sensitivity and resolutions. In this paper, we describe our heterogeneous, hierarchical design to manage the ~240 Gbps raw data rate, and large scale computing to produce real-time images with minimum latency. We discuss the implementation of the instrumentation, its performance, and scalability. | 16 | 9 | 1609.04205 |
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1609 | 1609.05729_arXiv.txt | We report on the discovery of a planetary companion candidate with a minimum mass $M\,\sin i=4.6\pm 1.0\,M_{\rm{Jupiter}}$ orbiting the K2$\,$III giant star HD~175370 (KIC~007940959). This star was a target in our program to search for planets around a sample of 95 giant stars observed with \textit{Kepler}. This detection was made possible using precise stellar radial velocity measurements of HD~175370 taken over five years and four months using the coud\'e echelle spectrograph of the 2-m Alfred Jensch Telescope and the fibre-fed echelle spectrograph HERMES of the 1.2-m Mercator Telescope. Our radial velocity measurements reveal a periodic ($349.5\pm 4.5$ days) variation with a semi-amplitude $K=133\pm 25\,\rm{m\,s^{-1}}$, superimposed on a long-term trend. A low-mass stellar companion with an orbital period of $\sim 88$ years in a highly eccentric orbit and a planet in a Keplerian orbit with an eccentricity $e=0.22$ are the most plausible explanation of the radial velocity variations. However, we cannot exclude the existence of stellar envelope pulsations as a cause for the low-amplitude radial velocity variations and only future continued monitoring of this system may answer this uncertainty. From \textit{Kepler} photometry we find that HD~175370 is most likely a low-mass red-giant branch or asymptotic-giant branch star. | Planets around K-giant stars may provide us with clues on the dependence of planet formation on stellar mass. The progenitors of K-giant stars are often intermediate-mass main-sequence A-F stars. While on the main sequence these stars are not amenable to precise radial velocity (RV) measurements. There is a paucity of stellar lines due to high effective temperatures and these are often broadened by rapid rotation. Therefore one cannot easily achieve the RV precision needed for the detection of planetary companions. On the other hand, when these stars evolve to giant stars they are cooler and have slower rotation rates and a RV accuracy of a few$\,\,\rm{m\,s^{-1}}$ can readily be achieved. The giant stars thus serve as proxies for planet searches around intermediate-mass (1.2--2 $M_{\odot}$) early-type main-sequence stars. Since the discovery of the first exoplanet around K-giant stars \citep{Hatzes93,Frink02}, over 90 exoplanets (3 per cent\footnote{The Extrasolar Planets Encyclopedia: http://exoplanet.eu} of the total in September 2016) have been discovered orbiting giant stars. These planet-hosting giant stars are on average more massive than planet-hosting main-sequence stars. \citet{Johnson10} showed, based on an empirical correlation, that the frequency of giant planets increases with stellar mass to about 14 per cent for A-type stars. This is consistent with theoretical predictions of \citet{Kennedy08}, who concluded that the probability that a given star has at least one gas giant increases linearly with stellar mass from 0.4 to 3 $M_\odot$. Statistical analysis of microlensing and transiting data reveals that cool Neptunes and super-Earths are even more common than Jupiter-mass planets \citep{Cassan12,Howard12}. Unlike for a main-sequence star where there is more or less a direct mapping between effective temperature and stellar mass, it is more problematic to determine the stellar mass of a giant star. The evolutionary tracks for stars covering a wide range of masses all converge to the similar region of the H-R diagram. One way to obtain the stellar mass is to rely on evolutionary tracks. However, they are not only model dependent, but they require accurate stellar parameters such as the effective temperature and heavy element abundance. \citet{Lloyd11} argued that the masses of giant stars in Doppler surveys were only in the range 1.0 -- 1.2 $M_\odot$ and thus were not intermediate-mass stars. Clearly, we cannot disentangle the effect of stellar mass on the observed planet properties if we cannot get a reliable measurement of the stellar mass. The stellar mass can be derived from solar-like oscillations. The first firm discovery of solar-like oscillations in a giant star was made using RVs by \citet{Frandsen02}. However, it was only recently that solar-like oscillations were unambiguously found in late-type giant stars, owing to space-based photometric observations by the \textit{CoRoT} \citep{deRidder09} and the \textit{Kepler} \citep{Gilliland10} missions. Solar-like p-mode oscillations of the same degree are equally spaced in frequency and the spacing is related to the square root of the mean stellar density ($\propto$$(M/R^{3})^{1/2}$), whereas the frequency of maximum oscillation power is $\propto$ $M/(R^2 \sqrt{T_{\rm{eff}}})$ \citep{Kjeldsen95}. From these empirical relations we can calculate both the stellar mass and radius in a more or less model independent way. The \textit{Kepler} Space Mission has been monitoring a sample of over 13,000 red-giant stars which can be used for asteroseismic studies. All stars show stellar oscillations that have been analysed to determine their fundamental stellar parameters \citep{Huber10,Stello13} and internal stellar structure \citep{Bedding11}. Furthermore, an estimate of the stellar age can be obtained using stellar models \citep{Lebreton14}. Giant stars observed with \textit{Kepler} represent a unique sample for planet searches as a planetary detection would mean that we can determine reliable stellar properties via asteroseismic analysis, characteristics not well known for many other planet-hosting giant stars. For this reason, we started a planet-search program among \textit{Kepler} asteroseismic-giant stars in 2010. We have distributed our targets over four different telescopes in order to maximize the detection and to minimize the impact of telescope resources at a single site. These telescopes include the 2-m telescope at Th\"uringer Landessternwarte Tautenburg (TLS), Germany (29 stars, since 2010), the 1.2-m Mercator telescope, La Palma, Spain (38 stars, since 2011), the 2.5-m Nordic Optical Telescope, La Palma, Spain (12 stars, since 2012) and the 2.7-m telescope at McDonald Observatory, Texas, US (33 stars, since 2012). In total our sample contains 95 giant stars, a statistically significant number given an expected detection rate of $\sim$15 per cent. We observe some targets in common at different sites, in order to check our measurements independently. Until now, three \textit{Kepler} giant stars are known to harbor planets discovered by detecting transits in the \textit{Kepler} light curves. \citet{Huber13} found two planets in the Kepler-56 system, while \citet{Lillo-Box14} confirmed the hot-Jupiter Kepler-91~b. \citet{Ciceri15}, \citet{Ortiz15} and \citet{Quinn15} discovered the warm Jupiter Kepler-432~b. An additional planetary candidate has also been reported to Kepler-432 \citep{Quinn15} and Kepler-56 \citep{Otor16} found via a long-term RV monitoring. Here, we report on the discovery of a giant planetary candidate orbiting the K-giant star HD~175370 (KIC~007940959) in the \textit{Kepler} field found via the RV method. | Our spectral analysis of HD~175370 corresponds to a K2$\,$III giant star. No carbon depletion or oxygen enhancement was detected in high-resolution spectra indicating that HD~175370 is a low-mass star, where hydrogen burning occurred in a radiative core, dominated by the proton-proton reactions so that the CNO cycle did not play a crucial role \citep{Charbonnel93}. Based on the \textsc{GBM} code \citep{hekker2013}, we deduce that HD~175370 is most likely a RGB or AGB star, and not a RC star. We conclude that the RV variations of HD~175370 are caused by a low-mass stellar companion with an orbital period of $\sim 88$ years in a highly eccentric orbit and a possible planetary companion with an orbital period $P=349.5\pm 4.5\jed{d}$, eccentricity $e=0.22\pm 0.10$, and a semi-amplitude $K=133\pm 25\jed{m\,s^{-1}}$. Our interpretation of the RV changes in terms of a planet is supported by the lack of variability in the spectral line bisectors, \textit{Kepler} light curve, and the lack of chromospheric activity corresponding to the orbital period of the planet. Furthermore, if 349.5 d were the rotation period of the star then this would result in a {\it maximum} projected stellar rotational velocity of 3.4 $\pm$ 0.5 $\rm{km\,s^{-1}}$ (RGB star) and 3.5 $\pm$ 0.5 $\rm{km\,s^{-1}}$ (AGB star), or nearly a factor of two smaller than our measured $v\,\sin i$. However, even the lack of stellar variability in stellar activity indicators corresponding to the orbital period of the planet do not prove its existence. One caveat is 42~Dra which has been claimed to host a planetary companion \citep{Doellinger09}. Our continued RV measurements of 42~Dra show that in the past three years the amplitude of the orbital motion decreased by a factor of four (Hatzes et al., in preparation). This casts serious doubts on the planet hypothesis for the RV variations. When the RV amplitude was high there were no variations in H$\alpha$, the spectral line bisectors, or in the \textit{Hipparcos} photometry. In this case the standard tools for planet confirmation - the same ones that we use for HD 175370 - failed. We suggest that measurements of such standard activity indicators may not be sufficient and that long-term monitoring is essential in order to confirm planets around K-giant stars. We note that we do not see any variations in the RV amplitude of the orbital motion due to HD~175370~b. The star 42~Dra showed coherent RV variations with a constant amplitude over three years. The RV variations in HD~175370 are constant for more than five years. We found two periods in the \textit{Kepler} light curves of 389.4~d and 370~d. Both of them are very close to the orbital period of the satellite, 372.5~d. We have shown that most of the \textit{Kepler} giant stars from our sample show the dominant period of $390\pm 20\jed{d}$, which agrees with the orbital period of the \textit{Kepler} satellite within the 1-$\sigma$ error bars. We therefore believe that finding a period in the \textit{Kepler} photometry which is close to the orbital period of the satellite is suspicious and should not be related to a real physical process unless one looks at all K giants in the field and detrends them systematically. Even though we cannot prove that the RV period is unrelated to the photometric periods, there is no evidence about the opposite. Continued RV monitoring of this system may help answer this uncertainty. A possible explanation of the RV changes observed for HD~175370 could be an unknown envelope pulsation causing photometric and spurious Keplerian-like RV variations. \citet{Jorissen16a,Jorissen16b} found among 13 low-metallicity giants three cases where small amplitude variations ($K$ ranging from 0.1 to 0.9 $\rm{km\,s^{-1}}$) with periods very close to 1 year are superimposed on a long-period Keplerian orbit. They do not give a conclusive answer on the origin of these variations, but leave an envelope pulsation as an option. In choosing a hypothesis based on a known phenomenon (exoplanets) as opposed to an unknown phenomenon (envelope oscillations) we chose the former until more evidence comes to light. In addition, we have shown that the photometric periods we have found in the \textit{Kepler} light curves of HD~175370 are suspicious of being due to the orbital period of the satellite. Oscillation modes dominate the short time-scale variations in the light curve and their amplitudes are at least ten times larger than the expected transit depth. It would therefore be difficult to detect any transit events in the light curve, which would constrain the orbital inclination. HD~175370 is one of the few close binary systems to host a giant planetary candidate and thus may be important for understanding planet formation in binary systems. To date, only about six binary systems with separations less than 25~AU have been found to host giant planets \citep{Morais08,Ramm09}. With a binary orbital separation of 22~AU HD~175370 adds to this small list. HD~175370 is unique in that it has the largest binary eccentricity ($e$ = 0.88) of these binary systems hosting planetary candidates. At closest approach the primary-secondary separation is only a factor of 2.7 larger than the planet-star semi-major axis. However, we fixed the orbital period and eccentricity of the binary orbit to values obtained from varying these parameters separately and searched for a minimum in the residuals, and therefore the errors of both parameters are expected to be large, which also implies a large error on the orbital separation. We estimate that uncertainties of the period and eccentricity are $\sim$4000 d and $\sim$0.02, respectively, which implies that at closest approach the primary-secondary separation could be a factor of 3.5 larger than the planet-star semi-major axis. A dynamical study may place additional constraints on the age and future evolution of this system. HD~175370~b is one of the first planetary candidates discovered around a \textit{Kepler} giant star via the RV method. Unlike for other evolved giant stars with extrasolar planets detected by ground-based RV surveys, for HD~175370 we have high-quality \textit{Kepler} photometry from which we can determine the stellar mass and radius in a more or less model independent way. The relative uncertainties of our derived stellar mass, stellar radius, and $\log g$ are 16, 14, and 1 per cent, respectively. HD~175370 is a relatively old star with an age of at least 5.4 Gyrs. More discoveries of planets around evolved giant stars can give us a better understanding of the mass dependence of planet formation and the evolution of planetary systems. | 16 | 9 | 1609.05729 |
1609 | 1609.02279_arXiv.txt | We present the Collection of Elemental Routines for Echelle Spectra (CERES). These routines were developed for the construction of automated pipelines for the reduction, extraction and analysis of spectra acquired with different instruments, allowing the obtention of homogeneous and standardised results. This modular code includes tools for handling the different steps of the processing: CCD image reductions, identification and tracing of the echelle orders, optimal and rectangular extraction, computation of the wavelength solution, estimation of radial velocities, and rough and fast estimation of the atmospheric parameters. Currently, CERES has been used to develop automated pipelines for thirteen different spectrographs, namely CORALIE, FEROS, HARPS, ESPaDOnS, FIES, PUCHEROS, FIDEOS, CAFE, DuPont/Echelle, Magellan/Mike, Keck/HIRES, Magellan/PFS and APO/ARCES, but the routines can be easily used in order to deal with data coming from other spectrographs. We show the high precision in radial velocity that CERES achieves for some of these instruments and we briefly summarize some results that have already been obtained using the CERES pipelines. | The possibility of obtaining at the same time high spectral resolution and wide spectral coverage has made echelle spectrographs a highly demanded type of instrument. Nowadays, most astronomical facilities count with at least one of these spectrographs \citep[see e.g.][]{vogt:1994,dekker:2000,noguchi:2002} and they are vastly used for a wide list of astronomical applications, like the study stellar atmospheres and the search of stellar and substellar companions by measuring radial velocity variations. In particular, the development of echelle spectrographs has significantly raised in the last couple of decades due to the high radial velocity precision that they can achieve with careful calibration, a capability that has been used for discovering $\approx$ 500 extrasolar planets \citep[e.g.][]{mayor:1995,giguere:2015}. One of the drawbacks of echelle spectrographs compared to typical spectrographs is the relative complexity demanded in the data reduction process, due to the fact that it contains several instrumental artefacts that need to be removed in order to extract a wavelength calibrated spectrum amenable for astrophysical analysis. The major complexity relies in the presence of multiple orders. These orders have quite different intensity levels due to wavelength dependent efficiency of the instrument which can obstruct their identification in some cases. Additionally, these orders have in general a significant curvature which has to be taken into account during the extraction and in some particular cases contiguous orders tend to overlap each other in the vertical direction which difficults a proper estimation of the scattered light. Moreover, echelle spectrographs can contain additional calibration fibres and image slicers that further complicate the processing of the data. An important fraction of current echelle spectrographs have their own reduction pipelines specifically designed for the properties of each instrument \citep[e.g.][]{bochanski:2009,mink:2011}, while in other particular cases there is no dedicated pipeline at all. This fact can produce some inconsistencies when spectra obtained from different instruments are used in the same analysis, in particular when the reduction steps include human intervention. Even in the case of working with data of a particular telescope, the automatisation of the data processing is desirable and is specially crucial when working on obtaining precision radial velocities at different epochs, because slight changes in the reduction steps can introduce significant systematic effects that propagate to the estimation of the Doppler shifts. There have been already some attempts to develop computational tools for the automated processing of data originated from different echelle spectrographs, but their use has not been extended for more than a couple of instruments so far. For example, \citet{mills:2003} presented the open source code ECHOMOP, which has been mostly used in the processing of data obtained from the Utrecht Echelle Spectrograph of the 4.2-m William Hershell Telescope, while \citet{sosnowaska:2015} presented the flexible reduction library for the ESPRESSO project which is also able to process data from HARPS and HARPS-N. Along the same line, the MIDAS system \citep{banse:1983} developed by ESO included a package designed to process echelle spectra \citep{ballester:1992}, which has been used to develop pipelines for most of ESO spectrographs. In this paper we present a new set of computational routines for developing fully automated reduction pipelines for data of echelle spectrographs. This modular code called CERES (\textbf{C}ollection of \textbf{E}lemental \textbf{R}outines for \textbf{E}chelle \textbf{S}pectra) is mostly written in Python, but contains also C and Fortran routines when speedy execution demands it. We have developed reduction pipelines for thirteen different spectrographs, and these recipes can be used as a guide for building pipelines for other instruments. The principal aim of the pipelines that we have developed is the handling of low signal to noise ratio data and the measurement of precision radial velocities in the context of extrasolar planets. In \S~\ref{structure} we describe the structure of the CERES echelle pipelines and the corresponding variations for each type of spectrograph. In \S~\ref{instruments} we list the instruments that are currently supported by CERES, while in \S~\ref{results} we discuss the performance of some of the CERES pipelines. Finally, in \S~\ref{concl} we summarise our work. | \label{sec:conclusions} We have presented CERES, a set of routines that allow the development of robust and fully automated pipelines for the reduction, processing and analysis of echelle spectra. We have constructed pipelines for thirteen different instruments with quite different specifications, reaching results that are almost as good or better than those of dedicated pipelines when available. In this regard, the CERES pipeline for the FEROS spectrograph stands out due to the high RV precision that can be achieved in the high signal-to-noise ration regimen ($\sigma_{RV}$=7.5 ms$^{-1}$), and also for its good behaviour for low signal-to-noise ratio data, which has allowed the discovery of several extrasolar planets orbiting stars even fainter than $V=14$. Moreover, we have developed reduction pipelines for instruments that do not have any dedicated pipeline, like the one for the DuPont spectrograph, for which we can achieve an RV precision of $\sigma_{RV}$=400 ms$^{-1}$ and complete automatization despite the lack of stability of the instrument. The recipes for these thirteen instruments can be used as guide for constructing pipelines for other instruments. In addition to reduced spectra, radial velocities and bisector spans, the CERES pipelines estimate rough and fast atmospheric parameters which are useful for quick target vetting at the telescope. There are several limitations of CERES that are being considered for future upgrades. For example, our extraction algorithms do not perform a previous rectification of the curvature of the echelle orders. The correction for the curvature can increase the resolution of the extracted spectrum, but it could introduce complications for the optimal extraction algorithm because it will require an interpolation between pixels. In addition, given that the principal applications of CERES have been focused on fibre-fed spectrograph, it does not contain a routine to correct for sky contamination, which can be useful in the case of slit spectrographs. Finally, an important upgrade for CERES will consist in developing automated routines for the computation of precise radial velocities using the iodine cell technique. The full CERES code, along with the pipelines for the instruments described in Section \ref{instruments} have been made publicly available\footnote{https://github.com/rabrahm/ceres}. \label{concl} | 16 | 9 | 1609.02279 |
1609 | 1609.07389_arXiv.txt | We present new MUSE observations of quasar field Q2131$-$1207 with a log \nhi=19.50$\pm$0.15 sub-DLA at z$_{\rm abs}$=0.42980. We detect four galaxies at a redshift consistent with that of the absorber where only one was known before this study. Two of these are star forming galaxies, while the ones further away from the quasar ($>$140 kpc) are passive galaxies. We report the metallicities of the HII regions of the closest objects (12+log(O/H)=8.98$\pm$0.02 and 8.32$\pm$0.16) to be higher or equivalent within the errors to the metallicity measured in absorption in the neutral phase of the gas (8.15$\pm$0.20). For the closest object, a detailed morpho-kinematic analysis indicates that it is an inclined large rotating disk with V$_{\rm max}$=200$\pm$3 km~s$^{-1}$. We measure the masses to be $M_{\rm dyn}$=7.4$\pm$0.4$\times$10$^{10}$ M$_{\odot}$ and M$_{\rm halo}$=2.9$\pm$0.2$\times$10$^{12}$ M$_{\odot}$. Some of the gas seen in absorption is likely to be co-rotating with the halo of that object, possibly due to a warped disk. The azimuthal angle between the quasar line of sight and the projected major axis of the galaxy on the sky is 12$\pm$1 degrees which indicates that some other fraction of the absorbing gas might be associated with accreting gas. This is further supported by the galaxy to gas metallicity difference. Based on the same arguments, we exclude outflows as a possibility to explain the gas in absorption. The four galaxies form a large structure (at least 200 kpc wide) consistent with a filament or a galaxy group so that a fraction of the absorption could be related to intra-group gas. | One of the key unknowns in the study of galaxy evolution is how galaxies acquire their gas and how they exchange this gas with their surroundings. Since gas, stars, and metals are intimately connected, gas flows affect the history of star formation and chemical enrichment in galaxies. Accretion is required to explain some of the basic observed properties of galaxies including the gas-phase metallicity \citep{erb06a}. Moreover, galaxies are believed to interact with the intergalactic medium (IGM) by pervading it with hydrogen ionising photons and by injecting heavy elements formed in stars and supernovae through supersonic galactic winds. Indeed, observations of the IGM indicate significant quantities of metals at all redshifts \citep{pettini03,ryanweber09,dodorico13,shull14a,becker15}. The presence of these metals is interpreted as a signature of strong galactic outflows in various models \citep{aguirre01,oppenheimer06}. Hydrodynamical simulations provide predictions of the physical properties of these gas flows \citep{keres05,brook11,keating15}. In recent years, much attention has been focused on the circumgalactic medium (CGM), a loosely defined term that describes the gas immediately surrounding galaxies over scales of $\sim 300$ kpc \citep{shull14b}. The CGM is at the heart of these physical processes. Therefore study of the CGM is crucial for understanding both the inflows of gas accreting into galaxies and the outflows carrying away the energy and metals generated inside galaxies. Outflows are commonly probed by the presence of interstellar absorption lines from cool gas blue-shifted by hundreds of km~s$^{-1}$ relative to the systemic velocities of the background galaxies \citep{shapley03,steidel10}. Strong \mgii\ absorbers in particular \citep{martin12,schroetter15} have been observed to extend out to 100 kpc along the galaxies' minor axes \citep{bordoloi11}, a fraction of which could be associated with galactic winds. Outflows are ubiquitous in galaxies at various redshifts \citep{pettini01,pettini02,cabanac08,heckman15}. Interestingly, the circumgalactic gas has also been probed in emission by \citet{steidel11} who stacked narrow-band images of z$\sim$2-3 galaxies and revealed diffuse Ly-$\alpha$ haloes extending to 80 kpc. More recently, the GTO (Guaranteed Time Observations) MUSE team used a 27-hr deep field to report Ly$\alpha$ haloes in individual emitters down to a limiting surface brightness of $\sim$10$^{-19}$ erg/s/cm$^2$/arcsec$^2$ with a scale length of a few kpc \citep{wisotzki16}. While observational evidence for outflows is growing, direct probes of infall are notoriously more difficult to gather. So far, the accretion of cool gaseous material has been directly observed only in the Milky Way in high velocity clouds in 21 cm emission at distances of 5--20 kpc \citep{lehner11, richter14}. At larger distances, nearby spirals exhibit both extraplanar HI clouds and morphological disturbances, which may be attributed to gas infall \citep{sancisi08}. However, the emission from these diffuse structures is difficult to map at high redshift. Nevertheless, detections suggestive of cool gas inflows have recently been reported in a few objects \citep{rubin11,martin12,bouche13,diamond16}. A powerful tool to study the CGM gas is offered by absorption lines in quasar spectra. We have initiated a novel technique to examine this gas in absorption against background sources whose lines of sight pass through the CGM of galaxies using 3D spectroscopy \citep{bouche07,peroux11a}. Over the past few years, we have demonstrated the power of this technique for studying the CGM using VLT/SINFONI by successfully detecting the galaxies responsible for strong \nhi\ absorbers at redshifts $z \sim 1$ and $z \sim 2$ \citep{peroux11b,peroux12,peroux13,peroux14,peroux16}. These detections have enabled us to map the kinematics, star formation rate, and metallicity of this emitting gas, and to estimate the dynamical masses of these galaxies. Out of our 6 detections for absorbers with known \nhi, we find evidence for the presence of outflows in two of them, while three are consistent with gas accretion. The remaining system at z$\sim$2 is poorly constrained \citep{peroux16}. Having demonstrated the power of 3D spectroscopy for study of the CGM at high-$z$ \citep{peroux16}, we now extend the technique at low-redshift with the MUSE optical spectrograph. Here, we present results from new observations of a sub-DLA at z$_{\rm abs}$=0.42980. The manuscript is organised as follows: Section 3 presents the ancillary and new observations of the absorber and the quasar field. Section 3 shows the analysis performed on the new MUSE and ancillary observations presented here. Finally, in section 4, we explore different scenarios to explain the gas seen in absorption in relation with the objects observed in the field. Throughout this paper we adopt an $H_{0}=70$~\kms~Mpc$^{-1}$, $\Omega_{\rm M}=0.3$, and $\Omega_{\rm \Lambda}=0.7$ cosmology. | The results presented in this work can be summarised as follows: \begin{itemize} \item We have obtained integral field spectroscopic observations with the VLT/MUSE instrument of a $\sim 1$\,arcmin field centred on the quasar Q2131$-$1207. The sightline to this quasar intersects a sub-DLA absorber at $z_{\rm abs} = 0.42980$, characterised by neutral hydrogen column density $\log [N$(H\,\textsc{i})]/cm$^{-2} = 19.50\pm0.15$ and abundance ${\rm [X/H]} = -0.54 \pm 0.18$. Approximately 0.1\% of the absorbing gas is in molecular form. \item We identify four galaxies at redshifts consistent with that of the sub-DLA absorber, where only one (the brightest of the four) was known previously. The two galaxies closest to the quasar sight-line (52 and 61 kpc) exhibit signs of on-going star formation; while the other two, which are further away at projected distances of $> 140$\,kpc, are passive galaxies. We report the metallicities of the HII regions of the closest of these objects (12+log(O/H)=8.98$\pm$0.02 and 8.32$\pm$0.16) which are to be compared with the metallicity measured in absorption in the neutral phase of the gas (12+log(O/H)=8.15$\pm$0.20). For galaxy "a", we derive a flat metallicity gradient of +0.01$\pm$0.03 dex kpc$^{\rm -1}$. \item For the brightest object, dubbed galaxy "a", a detailed morphological analysis indicates that the object is a extended galaxy showing indications of sub-structure on scales of 1". The kinematical modelling shows that it is an inclined (sin $i$=0.87$\pm$0.01) large rotating disk with V$_{\rm max}$=200$\pm$3 km~s$^{-1}$. We measure the dynamical mass of the object to be $M_{\rm dyn}$=7.4$\pm$0.4$\times$10$^{10}$ M$_{\odot}$ while the halo mass is M$_{\rm halo}$=2.9$\pm$0.2$\times$10$^{12}$ M$_{\odot}$. \item Some of the gas seen in absorption is likely to be co-rotating with galaxy "a", possibly due to a warped disk. In addition, we measure an azimuthal angle of 12$\pm$1 degrees which may suggest that some fraction of the absorption may arise in gas being accreted \citep{peroux16}. This is further supported by the galaxy to gas metallicity difference, where the metallicity of the gas is found to be lower than the one from the galaxy. We exclude outflows as a possibility to explain the gas in absorption but speculate that some of it may be related to intra-group gas or filament based on indications of a large structure (at least 200 kpc wide) formed by the four galaxies. \item Our new observations of the field indicate that the closest object to the quasar line-of-sight (a dwarf galaxy at $z_{\rm gal} = 0.74674$) is unrelated to the absorber. This chance alignment once more calls for caution when associating absorbers with galaxies without spectroscopic information. \end{itemize} Finally, it is clear from the work presented here that further MUSE observations of low-redshift (z$<$0.8) absorbers in quasar fields hold great potential for constraining the properties of the circumgalactic medium of galaxies. | 16 | 9 | 1609.07389 |
1609 | 1609.00083_arXiv.txt | Modelling of gravitational waves from binary black hole inspiral has played an important role in the recent observations of such signals. The late-stage ringdown phase of the gravitational waveform is often associated with the null particle orbit (``light ring'') of the black hole spacetime. With simple models we show that this link between the light ring and spacetime ringing is based more on the history of specific models than on an actual constraining relationship. We also show, in particular, that a better understanding of the dissociation of the two may be relevant to the astrophysically interesting case of rotating (Kerr) black holes. | 16 | 9 | 1609.00083 |
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1609 | 1609.00193_arXiv.txt | It has been observed that photons in the prompt emission of some gamma-ray bursts (GRBs) are highly polarized. The high polarization is used by some authors to give a strict constraint on the Lorentz invariance violation (LIV). If the Lorentz invariance is broken, the polarization vector of a photon may rotate during its propagation. The rotation angle of polarization vector depends on both the photon energy and the distance of source. It is believed that if high polarization is observed, then the relative rotation angle (denoted by $\alpha$) of polarization vector of the highest energy photon with respect to that of the lowest energy photon should be no more than $\pi/2$. Otherwise, the net polarization will be severely suppressed, thus couldn't be as high as what was actually observed. In this paper, we will give a detailed calculation on the evolution of GRB polarization arising from LIV effect duration the propagation. It is shown that the polarization degree rapidly decrease as $\alpha$ increases, and reaches a local minimum at $\alpha\approx \pi$, then increases until $\alpha\approx 3\pi/2$, after that decreases again until $\alpha \approx 2\pi$, etc. The polarization degree as a function of $\alpha$ oscillates with a quasi-period $T\approx \pi$, while the oscillating amplitude gradually decreases to zero. Moreover, we find that a considerable amount (more than $60\%$ of the initial polarization) of polarization degree can be conserved when $\alpha\approx \pi/2$. The polarization observation in a higher and wider energy band, a softer photon spectrum, and a higher redshift GRB is favorable in order to tightly constrain LIV effect. | \label{sec:introduction} Lorentz invariance is one of the foundations of Einstein's special relativity. It has been tested to a high accuracy using both the laboratory and cosmos experiments. In some quantum gravity theories \citep{Kostelecky:1989,Gambini:1999,Amelino-Camelia:2002,Myers:2003}, however, Lorentz invariance may be broken. In such a case, the propagation of light in vacuum exhibits a nontrivial dispersion relation compared to that in the special relativity. One of the most extensively discussed dispersion relation with Lorentz invariance violation (LIV) is $E_{\pm}^2=p^2\pm 2\xi p^3/M_{\rm pl}$, where $M_{\rm pl}$ is the Planck energy, and $\xi$ is a dimensionless parameter. According to this dispersion relation, the group velocity of light in vacuum, $v_g=\partial E/\partial p$, is no longer a constant, but is energy dependent. Therefore, two photons with different energies emitted simultaneously from a cosmological source will have a slightly time delay when they arrive the earth. Such a time delay should be detectable if the source is far enough away from the earth. Gamma-ray bursts (GRBs) provide an effect tool to test Lorentz invariance. As one of the most energetic explosions in the universe, GRBs are detectable out to redshift $z\approx 10$. In fact, GRBs have been widely used to constrain LIV effect \citep{Ellis:2006,Jacob:2008,Abdo:2009,Chang:2012,Nemiroff:2012fk,Zhang:2014wpb,Vasileiou:2015}. The value of $\xi$ constrained in this way is usually in the order of unity. The strictest limit on the LIV energy scale from GRB 090510 is $E_{\rm QG}>7.43\times 10^{21}$ GeV \citep{Nemiroff:2012fk}, which corresponds to $\xi<1.6\times 10^{-3}$. Much tighter constraints can be obtained through the measurement of GRB polarization. The polarimetric observations show that photons in the prompt emission of some GRBs are highly linearly polarized. For example, The polarization degree of the first reported highly polarized burst, GRB 021206, is about $80\%\pm 20\%$ \citep{Coburn:2003}. However, a following re-analysis of the same data found no significant polarization signal \citep{Rutledge:2004}. The highest polarized burst reported so far is GRB 041219A, which has polarization degree $98\% \pm 33\%$ \citep{Kalemci:2007}. But again this result was criticized by a more detailed analysis \citep{McGlynn:2007}. In 2011, the gamma-ray burst polarimeter {\it GAP} \citep{Yonetoku:2011a} onboard the Japanese Interplanetary Kite-craft Accelerated by Radiation Of the Sun (IKAROS) detected two highly polarized bursts, GRB 110301A and 110721A. These two bursts have conform polarization degree of $70\%\pm 22\%(3.7\sigma)$ and $84_{-28}^{+16}\%(3.3\sigma)$, respectively \citep{Yonetoku:2012}. The temporal evolution of polarization has also been observed \citep{Greiner:2003,Gotz:2009,Yonetoku:2011b,Mundell:2013}. Although the uncertainty is still large and many controversies exist, the possibility that some GRBs are highly polarized can't be excluded. The high accuracy $\gamma$-ray polarimeter {\it POLAR} \citep{Xiao:2015} onboard the Chinese space laboratory Tiangong-II is fully designed to measure the GRB polarization in $50-500$ keV energy band. It is scheduled to launch in September, 2016. If {\it POLAR} is launched, the GRB polarimetric data will be significantly enlarged, and the statistical significance will be highly improved. In the theoretical aspect, several theoretical models have been proposed to explain the GRB polarization \citep{Sari:1999,Waxman:2003,Granot:2003dy,Lazzati:2004,Toma:2009,Mao:2013gha,Chang:2014,Chang:2014a,Chang:2014b,Lan:2016}. If GRBs are really highly polarized, it will give a strict constraint on LIV effect. The idea of using polarization to constrain LIV was first proposed by \citet{Gleiser:2001rm}. They analysed the polarimetric data in ultraviolet band for radio galaxy 3C 256 locating at a redshift of 1.82, and obtained $\xi<10^{-4}$. When Lorentz invariance is broken, the polarization vector (i.e., the electric component) of a photon will rotate during its propagation. Suppose a beam of photons emit from a GRB source and propagate to the observer on earth. Every photon will rotates its polarization vector by an angle $\Delta \theta(k)$, which depends on the photon energy. Let $\alpha$ to be the difference of rotation angles of polarization vectors between the highest energy photon and the lowest energy photon. If high polarization degree is observed, then $\alpha$ couldn't be too large, otherwise the net polarization will be severely suppressed. \citet{Toma:2012} set the upper limit of $\alpha$ to be $\pi/2$, and obtained a strict upper limit on the value $\xi$ in the order of $\mathcal{O}(10^{-15})$ from the polarimetric data of three GRBs. However, the GRBs used by \citet{Toma:2012} have no direct measurement of redshift, while redshift deduced from the empirical luminosity correlations has large uncertainty. \citet{Gotz:2013} used the polarization data of GRB 061122 locating at a redshift $z=1.33$, and obtained $\xi<3.4\times 10^{-16}$. Using the most distant polarized burst, GRB 140206A, which has a confirm redshift measurement of $z=2.739$, \citet{Gotz:2014vza} have obtained the strictest constraint to date, i,e., $\xi<1\times 10^{-16}$. All of these constraints are based on the assumption that the rotation angle $\alpha$ is smaller than $\pi/2$. Then a question arises: how much polarization degree can be conserved if the polarization vector changes an angle $\alpha$\,? The main aim of our paper is to address this question. We will give a detailed calculation on the evolution of GRB polarization as a function of $\alpha$, and show that a considerable amount of polarization (depending on the photon energy band) can be conserved even if $\alpha$ is larger than $\pi/2$. The rest parts of this paper are arranged as follows: In Section \ref{sec:general}, we present the general formulae for the evolution of polarization induced by LIV effect. In Section \ref{sec:special}, we employ the formulae to GRBs, which usually have power-law spectra in the energy band of $\sim$ keV. Three different cases are discussed: (1) photons are initially completely unpolarized, (2) photons are initially completely polarized, and (3) photons are initially partially polarized. Finally, discussions and conclusions are given in Section \ref{sec:conclusions}. | \label{sec:conclusions} In this paper, we have investigated the evolution of GRB polarization arising from LIV effect. The birefringence of light leads to the rotation of polarization vector duration propagation. We obtained the net polarization degree as a function of the rotation angle $\alpha$, where $\alpha$ represents the relative rotation angle of high-energy and low-energy photons. We showed that the net polarization degree decreases rapidly as $\alpha$ increases until $\alpha\approx \pi$. As $\alpha$ continuously increases, the polarization degree oscillates with a quasi-period $T\approx\pi$ and a gradually vanishing amplitude. More than $60\%$ of the intrinsic polarization degree can be conserved at $\alpha=\pi/2$. This is in conflict with the intuition that $\alpha$ couldn't be larger than $\pi/2$ when high polarization degree is observed. Hence, it is inappropriate to simply use $\pi/2$ as the upper limit to constrain LIV effect, especially when the photon energy band is wide and the spectrum is hard. Photons in a wider energy band have larger net polarization at the fixed $\alpha$. However, for a specific source, a wider energy band will also have a larger $\alpha$. The net effect is that photons in a wider energy band have a lower polarization degree. Therefore, the polarimetric observation in a wide energy band is favourable in constraining LIV. In addition, we found that GRBs with soft spectrum and high redshift are helpful to tightly constrain LIV. The compact space-borne Compton polarimeter {\it POLAR} onboard the Chinese space laboratory Tiangong-II is a high accuracy $\gamma$-ray polarimeter fully designed to measure the polarization of GRB in $50-500$ keV energy band. If a GRB at redshift $z\approx 1$ is observed by {\it POLAR} with $50\%$ polarization degree, and if the spectrum in the {\it POLAR} energy band follows the power-law distribution with index $p\approx 1$, then we can obtain the most conservative upper limit of LIV effect $\xi\lesssim 1\times 10^{-16}$. This is obtained by assuming that the GRB is intrinsically completely polarized. Otherwise, the constraint may be much tighter. We apply our formulae to some true GRB events. \citet{Yonetoku:2012} claimed to have detected a polarization degree of $84_{-28}^{+16}\%$ in GRB 110721A in the {\it IKAROS-GAP} energy band $[70,300]$ keV. The photon spectrum in this energy band can be well fitted by the simple power law, with the power-law index $p=0.94\pm 0.02$ \citep{Tierney:2011}. Unfortunately, the redshift of this burst has not been directly measured. The $2\sigma$ lower limit of redshift inferred from the Amati relation is 0.45 \citep{Toma:2012}. Using these observational values, and assuming that photons are initially completely polarized, we obtain the upper limit of LIV effect $\xi\lesssim 4\times 10^{-16}$. GRB 061122 is a highly polarized GRB with redshift measurement $z=1.33$ \citep{Gotz:2013}. The polarization degree measured by the IBIS on board {\it INTEGRAL} in the energy band $[250,800]$ keV is $>60\%$ at $1\sigma$ confidence level \citep{Gotz:2013}. The spectrum can be fitted by a power law with an exponential cut-off, i.e. $N(k)\propto k^{-\alpha}\exp(-k/k_c)$, where $\alpha=1.15\pm 0.04$ and $k_c=221\pm 20$ keV. Using these observational parameters, the most conservative upper limit of LIV effect constrained from this burst is $\xi\lesssim 5\times 10^{-17}$. Finally, it should be pointed out that in this paper, we only consider one type of LIV, i.e., doubly special relativity. The dispersion relation in equation (\ref{eq:dispersion}) breaks not only the Lorentz invariance, but also the CPT invariance. This case is not typically favored by theorists. It is widely discussed because it is one of the few theories of quantum gravity that can be tested. More importantly, the dispersion relation considered here is among the few which has the vacuum birefringence effect. If the vacuum birefringence effect does not exist, and the polarization vector does not rotate during the propagation, then the polarimetric observation couldn't be used to constrain LIV. | 16 | 9 | 1609.00193 |
1609 | 1609.02019_arXiv.txt | {The volume of radio-astronomical data is a considerable burden in the processing and storing of radio observations with high time and frequency resolutions and large bandwidths. For future telescopes such as the SKA, the data volume will be even larger.} {Lossy compression of interferometric radio-astronomical data is considered to reduce the volume of visibility data and to speed up processing.} {A new compression technique named ``Dysco'' is introduced that consists of two steps: a normalization step, in which grouped visibilities are normalized to have a similar distribution; and a quantization and encoding step, which rounds values to a given quantization scheme using a dithering scheme. Several non-linear quantization schemes are tested and combined with different methods for normalizing the data. Four data sets with observations from the LOFAR and MWA telescopes are processed with different processing strategies and different combinations of normalization and quantization. The effects of compression are measured in image plane.} {The noise added by the lossy compression technique acts like normal system noise. The accuracy of Dysco is depending on the signal-to-noise ratio of the data: noisy data can be compressed with a smaller loss of image quality. Data with typical correlator time and frequency resolutions can be compressed by a factor of 6.4 for LOFAR and 5.3 for MWA observations with less than 1\% added system noise. An implementation of the compression technique is released that provides a Casacore storage manager and allows transparent encoding and decoding. Encoding and decoding is faster than the read/write speed of typical disks.} {The technique can be used for LOFAR and MWA to reduce the archival space requirements for storing observed data. Data from SKA-low will likely be compressible by the same amount as LOFAR. The same technique can be used to compress data from other telescopes, but a different bit-rate might be required.} | \label{sec:conclusions} The Dysco technique for compressing visibilities is suitable for radio observations. The noise added by this compression technique acts like normal system noise. The accuracy of the compression is depending on the signal-to-noise ratio of the data: noisy data can be compressed with a smaller loss of image quality. Data with typical correlator time and frequency resolutions can be compressed by a factor of 6.4 for LOFAR and 5.3 for MWA observations with less than 1\% added system noise in image plane. After averaging observations in time and frequency to the typical resolutions used in processing, a compression factor of 3.2 to 4 can be reached with less than 1\% added system noise in image plane. The technique is in particular well suited to reduce the archival space requirements. So far, testing was performed only on low-frequency data from the MWA and LOFAR telescopes. The implementation is generic and can be applied to other telescopes. However, further experiments are required to determine acceptable bit-rates. | 16 | 9 | 1609.02019 |
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1609 | 1609.07120_arXiv.txt | We assess the effect of the local large scale structure on the estimation of two-point statistics of the observed radial peculiar velocities of galaxies. A large N-body simulation is used to examine these statistics from the perspective of random observers as well as ``Local Group (LG)-like" observers conditioned to reside in an environment resembling the observed universe within 20 Mpc. The local environment systematically distorts the shape and amplitude of velocity statistics with respect to ensemble-averaged measurements made by a Copernican (random) observer. The Virgo cluster has the most significant impact, introducing large systematic deviations in all the statistics. For a simple ''top-hat`` selection function, an idealized survey extending to $\sim 160\hmpc$ or deeper is needed to completely mitigate the effects of the local environment. Using shallower catalogues leads to systematic deviations of the order of $50$ to $200\%$ depending on the scale considered. For a flat redshift distribution similar to the one of the CosmicFlows-3 survey, the deviations are even more prominent in both the shape and amplitude at all separations considered $(\simlt 100\hmpc)$. Conclusions based on statistics calculated without taking into account the impact of the local environment should be revisited. | \label{sec:intro} A pillar of cosmology is the Cosmological Principle \citep{Milne1935} stating that the Universe approaches isotropy and homogeneity with increasing scales\footnote{A counter example to the Cosmological Principle is a distribution of particles in a random fractal encompassing empty volumes of the same size as the whole probed region \citep{1980Peebles,NusserLahav2000}}. This principle is incorporated in the modern hierarchical scenario for structure formation, where matter density fluctuations are well defined, with a correlation function approaching zero on large scales. In such a scenario, initial fluctuations are described by homogeneous Gaussian random fields, and thus measurements made by different random observers are equivalent. The difference in the statistical properties inferred by these observers is commonly denoted as ``cosmic variance''. Assuming that our position in the Universe is not privileged, which is expressed in terms of the Copernican Principle \citep[e.g.][]{Uzan2009}, deep large-scale galaxy redshift surveys \citep[\eg][]{2dfgrs,sdss,Scrimgeour2012,CMASS_corr2012,Nadathur2013,Guzzo2014,Alpaslan2014} as well as detailed analyses of the cosmic microwave background radiation \citep[][]{WMAP9,Planck2013,Planck2015} broadly support this picture. Recent years have witnessed the advent of high-quality and rich galaxy peculiar velocity data, \eg{} the SFI++ \citep{Springob2007}, 6dF \citep{Springob2014}, and CosmicFlows catalogues \citep{Courtois2011,Tully2013,Tully2016}. This re-kindled activity in the peculiar velocity field with the new data offering an unprecedented opportunity for cosmological measurements and theory testing. In late-time linear theory, peculiar velocities are proportional to the gravitational force field. Therefore, peculiar velocity catalogues are a direct probe of dark matter and can in principle provide valuable information on fundamental theories for structure formation \cite{StraussWillick}. Inference of cosmological information from local observations must take into account the uncertainties introduced by cosmic variance. This has been known for a long time, dating back to early studies of the density field of galaxies \citep[\eg][]{Sandage1978,Huchra1983,Soifer1984,Geller1989}. While cosmic variance in the statistical analysis of the galaxy distribution is well studied, its implications on peculiar velocity observations have received little attention \cite[but see][]{Tormen1993,Strauss1998,Bilicki2010,Hellwing2014} and remain poorly understood. Due to the long-range nature of gravity, local structures affect velocity correlations on much larger scales than those relevant to the density field \citep{Tormen1993,Borgani2000,Chodorowski2002}. With reliable velocity catalogues only available for galaxies out to distances of 100--200 $\hmpc$, the impact of nearby structures is likely very significant. A similar effect was already hinted for the case of a local velocity field dispersion measure \citep{Cooray2006,Marra2013,Wojtak2014}. Galaxy peculiar velocities are practically unbiased with respect to the underlying velocity field \citep[\eg][]{Vittorio1986,Gorski1988,Groth1989,StraussWillick,Nusser1998, streaming_vel_Omega,Juszkiewicz2000,Sarkar2007,Nusser2011,Hudson2012,Nusser2012,Feix2015}. This is in contrast to the galaxy distribution in redshift surveys which is a biased tracer of the mass density field. Thus, peculiar velocity catalogues are not merely complementary to redshift-space distortions, but provide an independent avenue towards testing fundamental physical theories of structure formation, dynamical dark energy and modified gravity \citep{infall_Zu2013,Li2013,Hellwing2014PhRvL,Berti2015,Bull2016}. Extracting cosmological information from the observed motions is, however, a highly non-trivial matter. Despite the recent increase in quality and number of distance indicator measurements, the corresponding peculiar velocity catalogues remain relatively sparse with significant observational and systematic errors especially at larger distances. There are several approaches for inferring cosmological information from the observations. One could make an attempt at reconstructing a 3D peculiar velocity field from which the underlying mass density can be derived. This would be very rewarding but the effort is hampered by the notorious inhomogeneous Malmquist bias \citep{Lynden-Bell1988a,Lynden-Bell1988b} leading to spurious enhancement of the derived density fluctuations. A more straightforward strategy which has provided important constraints on the standard paradigm is to compare between the measured velocities and the gravitational field associated with an independent redshift survey \citep[see \eg][]{Davis2011}. Although this analysis is free from cosmic variance uncertainties and is mainly free from Malmquist biases, it relies on redshift surveys and is therefore dependent on the biasing relation between mass and galaxies. Our main goal in this paper is to systematically assess the impact of cosmic variance and observer location on the peculiar velocity observables such as velocity correlation functions and mean streaming velocities (the first moment of galaxy pairwise velocity distribution). We neglect meagre redshift evolution which might be present in local ($z\approx 0$) peculiar velocity catalogues. Further, we make no attempt at incorporating observational errors on the measured velocities. These errors increase with distance and can obviously lead to large uncertainties. Subsequently, we do not model any inhomogeneous Malmquist bias related to these errors. This paper is organised as follows: in \S\ref{sec:simulation} we describe the numerical assets used in this work; section \S\ref{sec:vel_stats} introduces and describes velocity statistics we consider; in \S\ref{sec:biases} we discuss various theoretical biases, while in \S\ref{sec:LG_observers} we study the impact of observer location and galaxy radial selection on the velocity statistics. We conclude with a general discussion of our results and their implications in \S\ref{sec:conclusions}. | \label{sec:conclusions} In this paper we considered the estimation of two-point peculiar velocity statistics. We have refrained from assessing important effects related to observational errors such as Malmquist biases, and focused on the impact of cosmic variance and observer location. We have tested the ability of the radial velocity based estimators in Eqns.~(\ref{eqn:psi1}), (\ref{eqn:psi2}) and (\ref{eqn:v12_obs_estim}) at recovering the underlying correlations in the case of of complete coverage velocity catalogues. The $v_{12}$ estimator of \citet{Ferreira1999} performs very well by measuring the averaged infall velocity with a percent-level accuracy. The theoretical predictions for both correlations functions were off by a factor of $8-16\%$. Thus, even for perfect data the measured values of $\psi_1$ and $\psi_2$ should be compared with theoretical predictions of Eqn.~(\ref{eqn:psi_liner_comb}) with care. Further, since for realistic data these statistics depend strongly on the data completeness, a much better approach is to derive predictions for both \citet{Gorski1989} functions based on realistic mock catalogues, rather than a simplistic relation as the one expressed by Eqns.~(\ref{eqn:psi_per_lin}), (\ref{eqn:psi_par_lin}) and (\ref{eqn:psi_liner_comb}). Next we have checked if a sampling bias due to strong under-sampling would be an issue. This was a relevant test, as the currently available galaxy peculiar velocity catalogues contain a relatively small number ($\sim 10^4$) of objects. The tests show that all three velocity statistics are not sensitive to under-sampling. The ensemble averages of 10\% and $1\%$-sub-samples (with effective $\bar{n}=9\times10^{-4}$ and $9\times10^{-5}\,h^3 {\rm Mpc}^{-3}$ number densities) were statistically consistent with the full sample. Only in the case of a severe sub-sampling of the $0.1\%$-case (with $\bar{n}=9\times10^{-6}\,h^3{\rm Mpc}^{-3}$) the estimated mean showed some noticeable scatter around the true mean. In addition, we have found that the scatter around the mean is scale dependent, being a strong function of a pair separation for $v_{12}$. Albeit for both $\psi$'s, except the smallest scales of $R<30\hmpc$, the scatter shows only a very weak evolution with scale. All in all, we can report that all the three studied velocity statistics are performing well in the sparse sampling regime. Our most important result is related to the effect of the observed large scale environment on velocity statistics. We have performed a detailed analysis of cosmic variance in velocity statistics by considering differences in velocity observables as measured by a Copernican observer and LG-equivalents. We have considered four criteria compatible with LG properties and local environment. Velocity two-point statistics are found to be insensitive to the criteria related to the MW halo mass and the LG motion and its mean density (within $\sim 3\hmpc$). In contrast, the proximity of an observer to a Virgo-like cluster is highly significant, affecting the correlations up to scales of $\sim100\hmpc$. This has not been noticed by \citet{Tormen1993} since they only consider LG-analogue observers defined without imposing the presence of a nearby massive cluster. In the near future, peculiar velocity surveys are not likely to reach to much larger distances than currently, although the number densities will be growing. For instance, {\it CosmicFlows-4} is expected to contain of the order of $3\times10^4$ sources but still mostly within $R<150\hmpc$ as currently CosmicFlows-3 does\footnote{Tully, private communication.}. It is only the advent of all-sky HI radio surveys that can extend the reach of PV surveys to $\sim 2$ times larger distances, and the object number closer to $10^5$. Careful modelling of observer location, and survey selection strategy are necessary for obtaining reliable and unbiased velocity correlation estimates. Much more effort is required to extract cosmological information richly stored in galaxy velocity data. Towards this goal, constrained realization techniques \citep[][]{HoffmanRibak1991,vdWeygaert1996,Klypin2003,Courtois2012,Hess2013,Sorce2016}, aiming at incorporating prominent structures in the real Universe can be very rewarding. | 16 | 9 | 1609.07120 |
1609 | 1609.04409_arXiv.txt | We derive an analytic expression for the transitional column density value ($s_t$) between the lognormal and power-law form of the probability distribution function (PDF) in star-forming molecular clouds. Our expression for $s_t$ depends on the mean column density, the variance of the lognormal portion of the PDF, and the slope of the power-law portion of the PDF. We show that $s_t$ can be related to physical quantities such as the sonic Mach number of the flow and the power-law index for a self-gravitating isothermal sphere. This implies that the transition point between the lognormal and power-law density/column density PDF represents the critical density where turbulent and thermal pressure balance, the so-called ``post-shock density.'' We test our analytic prediction for the transition column density using dust PDF observations reported in the literature as well as numerical MHD simulations of self-gravitating supersonic turbulence with the Enzo code. We find excellent agreement between the analytic $s_t$ and the measured values from the numerical simulations and observations (to within 1.5 A$_V$). We discuss the utility of our expression for determining the properties of the PDF from unresolved low density material in dust observations, for estimating the post-shock density, and for determining the HI-H$_2$ transition in clouds. | \label{intro} Star formation in galaxies occurs in dense molecular environments and is governed by the complex interaction of gravity, magnetic fields, turbulence, and radiation pressure \citep{McKee07,Elmegreen11}. Despite decades of study, the fundamental conditions behind the transition of diffuse atomic gas to cold molecular gas are still relatively unconstrained \citep{Sternberg88,Krumholz09,MckeeKrum10,bialy15}. The initial conditions imprinted on the diffuse and molecular gas on parsec scales (i.e. the level of turbulence, the cloud density, the structure of the magnetic field) may determine the key properties of the initial mass function (IMF) and the star formation rates in galaxies \citep{Hennebelle11}. Therefore the properties of diffuse and molecular gas in and around star-forming clouds must be quantified in order to construct a theory of star formation that predicts the IMF. The density and column density probability distribution functions (PDFs) have been used extensively in understanding the properties of galactic gas dynamics, from the diffuse ionized medium to dense star-forming clouds. The application of the PDF in molecular clouds has included density tracers such as CO \citep{lee12,burkhart13a} and dust \citep{Kainulainen09,froebrich.rowles10, Schneider13b, Schneider14,schneider15,Lombardi15}. Tracing the PDF using dust emission and absorption provides the largest dynamic range of densities, in contrast to molecular line tracers such as CO, which suffer from depletion and opacity effects \citep{Goodman09,burkhart13a,burkhart13b}. Simulations of self-gravitating MHD turbulence have successfully reproduced the shape and properties of the observational PDFs \citep{Burkhart09,Federrath12,Federrath13,Collins12,burkhart15a} suggesting that the gas PDF stems from a combination of turbulence (which induces a lognormal PDF shape in density) and self-gravity (which is characterized by a power-law PDF in density). In more detail, observed and simulated PDFs of Giant Molecular Cloud (GMC) environments, which include supersonic turbulence and self-gravity, suggest that the highest column density regime of the PDF (i.e., above column densities of 1 A$_V$) has a power-law distribution \citep{Collins11, Schneider13b,Lombardi15,Burkhart09,Federrath12,Federrath13,Collins12,burkhart15a} while the lower column density material in the PDF is dominated by turbulent diffuse gas and takes on a lognormal form \citep{Vazquez-Semadeni94,Burkhart12,Padoan97a}. The implications for the shape of the gas density PDF in ISM clouds are profoundly linked to the kinematics, star formation rates and the chemistry of the gas \citep{Federrath12}. Kinematically, the PDF width of the lognormal density distribution can be related to the sonic Mach number of the gas in an isothermal cloud \citep{Federrath08,Burkhart09,Kainulainen13,Burkhart12,burkhart15a}. Star formation rates are linked to the gas density PDF in several analytic models which use the high density end of the PDF to provide the dense gas fraction to calculate star formation efficiencies \citep{Krumholz05,Hennebelle11,Padoan11}. More recently, the HI PDF in and around GMCs has been proposed as a tracer of the HI-H$_2$ transition \citep{burkhart15, imara16} as well as a more accurate tracer of the low density lognormal shape as opposed to dust emission/absorption data, which have difficulty tracing the lognormal form \citep{Lombardi15,Schneider13b}. \cite{burkhart15} and \cite{imara16} have shown that the lognormal portion of the column density PDF in a sample of Milky Way GMCs is comprised of mostly atomic HI gas while the power-law tail is built up by the molecular H$_2$. These studies suggest that the transition point in the column density PDF between the lognormal and power-law portions of the column density PDF traces important physical processes, such as the HI-H$_2$ transition and the density regime where self-gravity becomes dynamically important. In this work we derive an analytic formula for the transitional column density from the lognormal portion of the PDF to the power-law form (denoted $s_t$). We organize the paper as follows. In Section 2 we derive an expression for the transitional column density for a piecewise lognormal and power-law PDF distribution based on the assumption that the PDF is continuous and differentiable. We further demonstrate that $s_t$ is related to the physical parameters such as the sonic Mach number of the gas (i.e. kinematics), the post-shock density, and the power-law index for a self-gravitating isothermal sphere. In Section 3 we compare our analytic expression for the transitional column density to numerical simulations of self-gravitating MHD turbulence run using the Enzo code. In Section 4 we compare our analytic expression for the transitional column density to observations using data from the literature. In Section 5 we discuss our results, followed by our conclusions in Section 6. | \label{sec:con} The transition point between the turbulence-dominated (lognormal) portion of the PDF and the denser, self-gravitating (power-law) portion of the PDF is an important component of the star-formation process. In this paper we derived an analytic expression for the transitional point ($s_t$) of the column density PDF from a lognormal to a power-law. We find that: \begin{itemize} \item The expression for $s_t$ depends on the mean column density, width of the lognormal portion of the PDF (i.e. the sonic Mach number and driving parameter) and the slope of the power-law portion of the PDF (i.e. power-law index for a self-gravitating isothermal sphere) \item In the limit of strong collapse, $s_t$ represents the post-shock density given by the balance of turbulent and thermal pressure. \item The values predicted by the analytic expression for $s_t$ agree well with measurements from Herschel dust observations and Enzo AMR simulations. \item The analytic expression reported in Equation \ref{eqn.st} will be useful for determining the properties of the PDF from unresolved low density material in observations and for estimating the HI-H$_2$ transition in clouds. \end{itemize} | 16 | 9 | 1609.04409 |
1609 | 1609.08557_arXiv.txt | {} {We pose the question of how much information on the atmospheric parameters of late-type stars can be retrieved purely from color information using standard photometric systems.} {We carried out numerical experiments using stellar fluxes from model atmospheres, injecting random noise before analyzing them. We examined the presence of degeneracies among atmospheric parameters, and evaluated how well the parameters are extracted depending on the number and wavelength span of the photometric filters available, from the UV GALEX to the mid-IR WISE passbands. We also considered spectrophotometry from the Gaia mission.} {We find that stellar effective temperatures can be determined accurately ($\sigma \sim$ 0.01 dex or about 150 K) when reddening is negligible or known, based merely on optical photometry, and the accuracy can be improved twofold by including IR data. On the other hand, stellar metallicities and surface gravities are fairly unconstrained from optical or IR photometry: $\sim 1$ dex for both parameters at low metallicity, and $\sim 0.5$ dex for [Fe/H] and $\sim 1$ dex for $\log g$ at high metallicity. However, our ability to retrieve these parameters can improve significantly by adding UV photometry. When reddening is considered a free parameter, assuming it can be modeled perfectly, our experiments suggest that it can be disentangled from the rest of the parameters.} {This theoretical study indicates that combining broad-band photometry from the UV to the mid-IR allows atmospheric parameters and interstellar extinction to be determined with fair accuracy, and that the results are moderately robust to the presence of systematic imperfections in our models of stellar spectral energy distributions. The use of UV passbands helps substantially to derive metallicities (down to [Fe/H] $\sim -3$) and surface gravities, as well as to break the degeneracy between effective temperature and reddening. The Gaia BP/RP data can disentangle all the parameters, provided the stellar SEDs are modeled reasonably well.} | \label{Intro} The border between photometry and spectroscopy is poorly defined. As soon as two photometric measurements, a {\it color}, are available, photometry can be considered as very-low resolution spectroscopy. Multicolor photometric systems are now evolving from a few passbands to dozens of contiguous narrow windows (see, e.g., Aparicio Villegas et al. 2010). Obviously, narrower and more numerous passbands covering a wider spectral range are desirable, but the information content does not necessarily increase linearly with the number of filters. Bailer-Jones (2004) performed an optimization exercise varying the number, width, and location of the passbands to maximize the information content for stellar sources and minimize observing efforts. Introducing a new system is a luxury that only very large projects can afford, while in most cases data are available on one of a number of widely used photometric systems (Bessell 2005). In this paper we make an attempt to quantify the information that can be extracted from photometric or spectrophotometric observations of stars using several existing systems. We focus on some of the most useful ones with wide sky coverage, such as the SDSS $ugriz$ system (Fukugita et al. 1996), 2MASS $JHK_s$ (Skrutskie et al. 2006), WISE (Wright et al. 2010), and the GALEX passbands (Martin et al. 2005) as representatives of the optical, near-IR, mid-IR, and UV spectral windows, respectively (see Fig. \ref{flux}). We examine as well the potential of the Gaia BP/RP spectrophotometry (Jordi et al. 2010). While substantial work has already been devoted to studying the potential of many of these (spectro-)photometric systems in order to constrain stellar parameters and interstellar extinction (see, e.g., Lenz et al. 1998, Masana et al. 2006, Bailer-Jones et al. 2013), we provide here -- for the first time -- an attempt to make a fair comparison of all these systems and the incremental improvement that may be expected as different spectral windows are added. This is a purely theoretical study in which random and systematic errors are considered by means of simulations. Section \ref{model} describes the model spectra used in our analyses and the approximations adopted for including interstellar extinction. Section \ref{method} describes our analysis methodology and figure of merit. Section \ref{nored} is devoted to the case in which interstellar extinction is negligible, while Section \ref{red} takes that effect into account. Section \ref{xp} considers the important case of the spectrophotometric data that will soon become available from the ESA mission Gaia (de Bruijne 2012). Section \ref{sys} considers the impact of systematic errors in models, and Section \ref{summary} summarizes our results and closes the paper. \begin{figure*}[t!] \centering {\includegraphics[width=16cm]{f1.ps}} \caption{Three sample spectra with solar metallicity, surface gravity $\log g=4$, and effective temperatures (in decreasing order of flux level) of 6000, 5000, and 4000 K. For each model there are two curves, one corresponding to the Kurucz (1993) models (in black) and one to the more recent spectral energy distributions by Allende Prieto et al. (2016; red). Both are shown at a resolving power of $R\sim 200$ and smoothed (thick lines) to $R\sim 8$, corresponding to the effective resolution of a broad-band photometric filter. The position and shape of the responses for the different filters under consideration are shown with an arbitrary normalization. The GALEX FUV and the WISE W1 and W2 bandpasses fall outside of the spectral range represented.} \label{flux} \end{figure*} | \label{summary} We perform numerical experiments to explore the potential of broad-band photometry and spectrophotometry of stars to infer atmospheric parameters and the interstellar extinction along the line of sight. Our experiments are fairly simplistic in that we ignore systematic errors that result from an imperfect knowledge of the distortions introduced by instruments and, in the case of ground-based observations, the effect of Earth's atmosphere, but should hold as long as they are significantly smaller than the random uncertainties adopted (2-10\%, depending on the instrument and wavelength). As such, our results provide an upper limit to the performance achievable for real data. We find that photometry over a wide spectral range from the near-UV to the mid-IR can constrain the atmospheric parameters fairly well. While optical data are enough to constrain well the effective temperature of late-type stars, this parameter is retrieved with far better precision by including infrared observations. We show that the determination of surface gravity and metallicity can benefit substantially from the addition of near-UV photometry. For a similar signal-to-noise ratio, the performance expected for an instrument like the Gaia BP/RP spectrophotometer is better than standard photometric systems with a similar wavelength span. This is not surprising, since both resolution and sampling are much better for BP/RP than for the standard broad-band photometric systems. However, for the faintest stars observed by Gaia, photon noise will limit significantly the atmospheric parameters determined from the BP/RP observations. Photometric searches for ultra-metal-poor stars can benefit enormously from counting on UV passbands. In our tests without reddening, the success rate at [Fe/H]$\le -3$, defined as the fraction of stars correctly identified in that range, and the false-positive rate are 59\% and 42\%, respectively, using only SDSS photometry, but these figures change to 79\% and 22\%, respectively, when the GALEX filters are included in the observations. The effect of interstellar absorption on the observations hampers our ability to recover the stellar atmospheric parameters. Only with wide wavelength coverage -- in particular in the blue and into the near-UV -- and high signal-to-noise ratios can reddening be cleanly disentangled from variations in the atmospheric parameters. It will be very interesting to verify our expectations with real data, although it goes beyond the scope of the present paper and will likely require some sort of zero-point calibration of the synthetic photometry. In addition, in our numerical experiments we have sampled the spectral energy distributions at specific wavelengths corresponding to each of the bandpasses under study, but for practical applications it is necessary to convolve the model SEDs with the filter responses since a filter's effective wavelength depends on the stellar spectrum. Our main conclusion is that wide-area photometry at multiple wavelengths is a promising path for characterizing the stellar populations of the Milky Way and other nearby galaxies where stars can be resolved. Following up on pioneering studies (e.g., Ivezi\'c et al. 2008), projects that already do this in a homogeneous fashion such as ALHAMBRA (Moles et al. 2008), J-PLUS and JPAS (Ben\'{\i}tez et al. 2014), PAU (Castander et al. 2012), or the Gaia mission will demonstrate the potential and the limitations of this technique. | 16 | 9 | 1609.08557 |
1609 | 1609.05369_arXiv.txt | We consider scenarios in which the annihilation of self-conjugate spin-1 dark matter to a Standard Model fermion-antifermion final state is chirality suppressed, but where this suppression can be lifted by the emission of an additional photon via internal bremsstrahlung. We find that this scenario can only arise if the initial dark matter state is polarized, which can occur in the context of self-interacting dark matter. In particular, this is possible if the dark matter pair forms a bound state that decays to its ground state before the constituents annihilate. We show that the shape of the resulting photon spectrum is the same as for self-conjugate spin-0 and spin-1/2 dark matter, but the normalization is less heavily suppressed in the limit of heavy mediators. | It is well understood that dark matter (DM) annihilation or decay to a Standard Model (SM) fermion-antifermion pair $\bar f f$ can be chirality suppressed. Then, the dominant indirect detection process in the current epoch may instead involve the internal bremsstrahlung (IB) of an additional gauge boson (i.e., $\bar f f \gamma, \bar f f Z, \bar f f g, \bar f f' W^{\pm}$ final states)~\cite{Bergstrom:1989jr,Flores:1989ru,Bringmann:2007nk,Barger:2009xe,Kachelriess:2009zy, Ciafaloni2011,Bell:2011if,Barger2012,Garny2012,DeSimone:2013gj,Giacchino:2014moa,Bringmann:2015cpa}. These processes may not only dominate the annihilation/decay rate, but may also yield a hard boson spectrum which can be more easily distinguished from background. As a result, these internal bremsstrahlung processes have been well studied for the case of spin-0 or spin-1/2 dark matter. In this Letter, we discuss a case that has not been considered so far: chirality suppression of spin-1 dark matter annihilation lifted by internal bremsstrahlung. The annihilation of spin-1 dark matter particles $B$ to a fermion-antifermion pair has been studied in several specific models \cite{Cheng2002,Cheng:2003ju, Farzan2012, Yu2014}, and in these cases it is found that there is no chirality suppression, implying that the two-body final state is dominant. We point out that this unsuppressed contribution arises from the $J=2$ $s$-wave initial state. If the DM initial state is unpolarized, then there is indeed no way to avoid this unsuppressed contribution to the $2 \rightarrow 2$ annihilation cross section. But if the DM initial state is polarized, and the $J=2$ initial state is projected out, then the $s$-wave $BB \rightarrow \bar f f$ matrix element will be chirality suppressed, and the internal bremsstrahlung process will dominate. This scenario can be realized in a simple model in which the annihilation occurs through the formation of a $BB$ bound state, which decays to its ground state before the two constituents annihilate. If the ground state is not $J=2$, then the branching fraction for decay to the $\bar f f$ final state will be chirality suppressed, and the primary bound state decay channel will be to a three-body final state. We focus on the case in which internal bremsstrahlung involves the emission of a photon ($\bar f f \gamma$). There are general arguments which show that for self-conjugate spin-0 or spin-1/2 dark matter, the photon spectrum adopts a common universal form which depends only on $r$, the ratio of the mass of the mediating particle ($m_\Psi$) to the mass of the dark matter. We will show that this argument generalizes to the case of spin-1 dark matter with one key difference: for spin-0 or spin-1/2 dark matter, the annihilation matrix element necessarily scales as $m_\Psi^{-4}$ in the $r \gg 1$ limit, while for spin-1 dark matter, the matrix element only scales as $m_\Psi^{-2}$. The structure of this paper is as follows. In Section 2 we review the general arguments that underly the chirality suppression of dark matter annihilation to the $\bar f f$ final state and apply these arguments to the case of spin-1 dark matter, inferring that IB is only relevant for a polarized initial state. In section 3, we present the IB photon spectrum for the case of spin-1 dark matter, and demonstrate that its shape is necessarily the same as in the spin-0 and spin-1/2 cases. We also provide a physical realization of internal bremsstrahlung as the dominant annihilation channel in terms of the decay of a dark matter bound state. In section 4, we conclude with a discussion of our results. | We considered the annihilation of self-conjugate spin-1 dark matter to Standard Model fermions, assuming that flavor violation is minimal. We have shown that the $BB \rightarrow \bar f f$ annihilation cross section can be chirality suppressed, but only if the initial state is $J=0$; the $J=2$ state has an unsuppressed $2 \rightarrow 2$ annihilation cross section. For the $J=0$ state, the dominant $s$-wave annihilation process yields a three-body final state $\bar f f \gamma$ via internal bremsstrahlung, with a spectrum that is identical to the case of self-conjugate spin-0 or spin-1/2 dark matter. The typical suppression of the internal bremsstrahlung cross section by a factor of the fine structure constant implies that this process is unimportant for the case in which the initial DM state is unpolarized. But there are scenarios for which the DM state is polarized, as in the context of self-interacting dark matter. In particular, if dark matter forms a nonrelativistic $BB$ bound state, which decays to a $J=0$ ground state before the constituents annihilate, then internal bremsstrahlung could be the dominant annihilation process. We have shown that for spin-1 dark matter, the shape of the photon spectrum arising from internal bremsstrahlung is necessarily the same as for spin-0 and spin-1/2 dark matter, generalizing previous arguments regarding the universality of the photon spectrum. But unlike the spin-0 or spin-1/2 cases, in which the annihilation matrix element is suppressed by $m_\Psi^{-4}$ in the heavy mediator limit, in the spin-1 case the matrix element is only suppressed by $m_\Psi^{-2}$. This is particularly interesting because of its impact on complementary searches at the LHC. In order for internal bremsstrahlung to be possible, the mediator $\Psi$ must be charged, and there are tight constraints on new charged particles from collider experiments. If $m_\Psi$ is increased in order to evade those constraints, then the internal bremsstrahlung cross section becomes heavily suppressed. As a result, many studies of the IB photon spectrum have necessarily focused on the regime where the dark matter and the mediator are nearly degenerate; in this region of parameter space, the internal bremsstrahlung cross section is not heavily suppressed, and the mediator escapes collider searches because the jets/leptons produced by its decay are soft. For the case of spin-1 dark matter, however, $r$ can be made much larger without heavily suppressing the annihilation cross section. This opens a new window in parameter space in which one can search for dark matter annihilation via internal bremsstrahlung. | 16 | 9 | 1609.05369 |
1609 | 1609.00901.txt | {The present paper is a sequel to our previous work [Y. Ni et al., JCAP 1607, 049 (2016)] in which we studied the iron K$\alpha$ line expected in the reflection spectrum of Kerr black holes with scalar hair. These metrics are solutions of Einstein's gravity minimally coupled to a massive, complex scalar field. They form a continuous bridge between a subset of Kerr black holes and a family of rotating boson stars depending on one extra parameter, the dimensionless scalar hair parameter $q$, ranging from 0 (Kerr black holes) to 1 (boson stars). Here we study the limiting case $q=1$, corresponding to rotating boson stars. For comparison, spherical boson stars are also considered. We simulate observations with XIS/Suzaku. Using the fact that current observations are well fit by the Kerr solution and thus requiring that acceptable alternative compact objects must be compatible with a Kerr fit, we find that some boson star solutions are relatively easy to rule out as potential candidates to explain astrophysical black holes, while other solutions, which are neither too dilute nor too compact are more elusive and we argue that they cannot be distinguished from Kerr black holes by the analysis of the iron line with current X-ray facilities.} | The current paradigm for the nature of astrophysical black hole (BH) candidates is that they are well described by the Kerr metric~\cite{kerr}. Such a solution describes a rotating BH and is completely specified by only two parameters: the mass $M$ and the spin angular momentum $J$ of the compact object. On the theoretical side, this view is supported by the uniqueness theorems (see~\cite{Chrusciel:2012jk} for a review), stating that, in 4-dimensional general relativity, the only stationary, axisymmetric, asymptotically-flat, regular ($i.e.$ without geometric singularities or closed time-like curves on or outside the event horizon) solution of the vacuum Einstein equations is the Kerr metric. Dynamically, any initial deviations from the Kerr geometry occurring during gravitational collapse, should be radiated away through the emission of gravitational waves~\cite{price}. Moreover, the only other traditionally accepted degree of freedom -- electric charge -- can be, most likely, ignored, because a neutral equilibrium should be reached fairly quickly due to the highly ionized host environments of these objects, rendering any residual electric charge too small (for macroscopic bodies) to appreciably affect the spacetime metric~\cite{ec,neutralize,neutralize2}. Also, the possible presence of an accretion disk should have a negligible geometric backreaction, because the disk mass is typically many orders of magnitude smaller than the mass of the BH candidate~\cite{disk1,disk2}. In the end, deviations from the Kerr metric seem only to be possible in the presence of new physics, thus making their detection potentially very rewarding. Consequently, in the past 10~years, there has been a significant work to study how present and future facilities can test the Kerr BH hypothesis, both with electromagnetic observations~\cite{em1,em2,em3,Berti:2015itd} and gravitational waves~\cite{Berti:2015itd,gw1,gw2,gw3}. Among the electromagnetic approaches, the analysis of the iron K$\alpha$ line in the X-ray reflection spectrum of the accretion disk is a particularly powerful and promising tool to probe the strong gravity region around BH candidates~\cite{i1,i2,i3,i4,i5,i6,i7,i8}. The shape of the iron K$\alpha$ line is significantly affected by relativistic effects (Doppler boosting, gravitational redshift, light bending) occurring in the vicinity of the compact object. In the presence of high quality data and with the correct astrophysical model, this technique promises to provide stringent constraints on possible deviations from the Kerr metric because it can break the parameter degeneracy between the spin and possible non-Kerr features~\cite{ii1,ii2,ii3}. Other techniques meet, by contrast, serious difficulties to break such a parameter degeneracy~\cite{k1,k2,k3,k4}. Tests of the Kerr BH hypothesis fall, at least, into two general categories. The first one considers novel exact solutions of general relativity (or even modified gravity), by analysing more general matter contents beyond (electro-)vacuum, thus capable of producing non-Kerr metrics (see $e.g.$~\cite{ag2005, gp2009, m2015,hbh,Herdeiro:2015gia,Herdeiro:2015tia,Herdeiro:2016tmi}). The second one designs parametrized families of metric deformations from Kerr~\cite{p1,p2,p3,p4,p5,p6,p7,p8}, without worrying about which model they solve. Within the first approach, Kerr BHs with scalar hair (KBHsSH)~\cite{hbh} have recently gained a lot of attention due to their physically reasonable and astrophysically plausible matter sources. These solutions are regular on and outside an event horizon, and also obey all energy conditions. They are exact (albeit numerical) solutions of Einstein's gravity minimally coupled to a massive, complex scalar field, and interpolate between (a subset of) Kerr BHs -- when a normalized ``hair" parameter, $q$, vanishes -- and a family of gravitating solitons~\cite{Schunck:2003kk} -- when $q$ is maximal ($q=1$) --, the so-called boson stars; see, e.g., Ref.~\cite{Liebling:2012fv} for a review. They can bypass the various no-scalar-hair theorems that apply to this model by combining rotation with a harmonic time-dependence in the scalar field -- see~\cite{Herdeiro:2015waa} for a review of such theorems. Although these solutions are found numerically~\cite{hbh}, a proof of their existence in a small neighborhood of the Kerr family is available~\cite{proof}. In Ref.~\cite{yy}, we have computed the iron K$\alpha$ line expected in the reflection spectrum of a small sample of KBHsSH with $q<1$, to check whether present and future X-ray missions can constrain the scalar charge of astrophysical BH candidates. We found that some KBHsSH would have an iron line definitively different from those seen in the spectrum of BH candidates, and therefore such solutions are at tension with current data. Other considerably hairy BHs, however, are compatible with current data, but future X-ray mission will be able to put much stronger constraints. This work is the continuation of the study presented in Ref.~\cite{yy}. Here we consider the limiting case with $q=1$, in which there is no horizon and the solution describes an everywhere smooth, topological trivial configuration, $i.e.$ a boson star. As in Ref.~\cite{yy}, we study the iron line profile that should be expected in the reflection spectrum of the accretion disk around these objects. We simulate observations with XIS/Suzaku\footnote{http://heasarc.gsfc.nasa.gov/docs/suzaku/} to check whether it is possible to test the existence of these objects with the current X-ray facilities. We study a sample of 12~boson star solutions, whose lensing -- another property of phenomenological interest -- was considered in~\cite{Cunha:2015yba}. We find that some solutions can be surely ruled out, as Kerr models cannot provide a good fit of their iron line, while they do it with current X-ray data. Other solutions, assuming certain emissivity profiles, can have an iron line too similar to that of Kerr BHs to be tested by current X-ray missions. Last, there are a few solutions that may be tested by current X-ray facilities, but we cannot conclude they can be ruled out by current data without a more sophisticated analysis. We remark that boson stars have been suggested as BH mimickers, and studies in this context, including of the comparative phenomenology, have been report in, $e.g.$~\cite{Mielke:2000mh,Burt:2011pv,Guzman:2005bs,Berti:2006qt,Guzman:2009zz,Eilers:2013lla,Marunovic:2013eka,Vincent:2015xta,Meliani:2015zta,Meliani:2016rfe,v2-1,v2-2}. The content of the paper is as follows. In Section~\ref{s-2}, we briefly review the boson star solutions. We then describe the sample of 12~solutions that shall be studied in this paper. In Section~\ref{s-3}, we compute the shape of the iron line profile expected in the reflection spectrum of these 12~spacetimes. In Section~\ref{s-4}, we simulate observations with XIS/Suzaku and we study whether these boson starts solutions can be tested with present X-ray observatories. Summary and conclusions are presented in Section~\ref{s-5}. Throughout the paper, we employ natural units in which $c = G_{\rm N} =\hbar = 1$. | } In the present paper we have continued our study to observationally test the existence in the Universe of Kerr BHs with scalar hair, which are solutions of Einstein's gravity minimally coupled to a massive, complex scalar field found in Ref.~\cite{hbh}. We have considered the limiting case $q=1$, which yields rotating ($m=1$) boson stars. For comparison, here we have also considered non-rotating (spherically symmetric) boson stars ($m=0$). We have computed the profile of the iron K$\alpha$ line that should be expected from the reflection spectrum of the accretion disk of these objects. We have simulated observations with XIS/Suzaku, assuming that the source is a bright BH binary, and we have fitted the spectrum with a Kerr model. Since the available X-ray data of the reflection spectrum of BH binaries are regularly fitted with Kerr models, and there is no tension between data and theoretical predictions, we argue that the BH candidates in the Universe cannot be the boson star solutions with a bad fit, while those with an acceptable fit cannot be distinguished by Kerr BHs and are thus viable candidates. We remark, again, that should any future observations of BH candidates yield spectra incompatible with the Kerr model, the conclusions herein should be revisited. \begin{table}[h!] \begin{center} \begin{tabular}{|c|cc|cc|} \hline Solution & $1/r^3$, $t = 100$~ks & $1/r^3$, $t = 1$~Ms & LP, $t = 100$~ks & LP, $t = 1$~Ms \\ \hline \hline 1 & $\times$ & $\times$ & $\times$ & $\times$ \\ 2 & $\times$ & $\times$ & $\times$ & $\times$ \\ 3 & $\times$ & $\times$ & $\times$ & $\times$ \\ 4 & & $\times$ & & \\ 5 & & $\times$ & & \\ \hline 6 & $\times$ & $\times$ & $\times$ & $\times$ \\ 7 & & $\times$ & & $\times$ \\ 8 & & $\times$ & & $\times$ \\ 9 & & $\times$ & & $\times$ \\ 10 & & $\times$ & & $\times$ \\ 11 & $\times$ & $\times$ & $\times$ & $\times$ \\ 12 & $\times$ & $\times$ & $\times$ & $\times$ \\ \hline \end{tabular} \end{center} \caption{Summary of our results. $\times$ means that the Kerr model fails to give an acceptable fit. In the second and the third columns, the model assumes $I_{\rm e} \propto 1/r^3$. In the fourth and the fifth columns, we have $I_{\rm e} \propto h/(h^2 + r^2)^{3/2}$. The exposure time is 100~ks (second and fourth columns) and 1~Ms (third and fifth columns). See the text for more details. \label{tab2}} \end{table} Our results are summarized in Tab.~\ref{tab2}, where $\times$ is to indicate that the Kerr model fails to give an acceptable fit for the corresponding boson star solution, exposure time, and emissivity profile. The solutions 1-3, 6, 11, 12 can be ruled out. Even for an exposure time of 100~ks, we find that the Kerr model always provides a very bad fit. The iron line profiles of these boson stars are not compatible with those in the available X-ray data of BH binaries. For spherical boson stars, the qualitative conclusion is that acceptable solutions must be fairly compact. This qualitative behaviour also holds for the rotating solutions; in this case, however, both the most dilute and the most compact cases are ruled out for these ``short" exposure times. In a sense, our sample of solutions included more compact rotating than spherical boson stars, as the latter do not possess a light ring whereas the former do (beyond solution 10). Perhaps too compact spherical boson stars can also be ruled out. Unfortunately, in this case on has to go well into the spiral in order to get a light ring, which only appears in the third branch of solutions, i.e. after the second backbending in frequency, wherein the accuracy of solutions becomes poorer. The Kerr model provides an acceptable fit for the solutions 7-10 and for an exposure time of 100~ks, but fails to give an acceptable fit in the case of a longer observation of 1~Ms. In such a case, we should investigate these solutions better and analyze real data of specific sources to conclude whether the solutions 7-10 can be excluded or are consistent with the available data. Unfortunately, this is beyond the possibilities of the current version of our code. Last, we have the solutions 4 and 5. If we assume the lamppost emissivity profile $\sim h/(h^2 + r^2)^{3/2}$, even an observation of 1~Ms cannot distinguish these solutions from a Kerr BH, in the sense that the Kerr model provides a good fit. We thus argue that it impossible to test the existence of these objects with the current X-ray missions. %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% | 16 | 9 | 1609.00901 |
1609 | 1609.07193_arXiv.txt | {% The ARIANNA experiment is currently taking data in its pilot-phase on the Ross ice-shelf. Fully autonomous stations measure radio signals in the frequency range from 100 MHz to 1 GHz. The seven station HRA was completed in December 2014, and augmented by two special purpose stations with unique configurations. In its full extent ARIANNA is targeted at detecting interactions of cosmogenic neutrinos ($>10^{16}$eV) in the ice-shelf. Downward-pointing antennas installed at the surface will record the radio emission created by neutrino-induced showers in the ice and exploit the fact that the ice-water surface acts as a mirror for radio emission. ARIANNA stations are independent, low-powered, easy to install and equipped with real-time communication via satellite modems. We report on the current status of the HRA, as well as air shower detections that have been made over the past year. Furthermore, we will discuss the search for neutrino emission, future plans and the energy-dependent sensitivity of the experiment. } | The ARIANNA detector is aimed at detecting neutrinos at an energy above \unit[$10^{16}$]{eV} \cite{AstroP}. Whenever these neutrinos produce an electromagnetic cascade after interacting in the ice of Moore's Bay on the Ross Ice-shelf in Antarctica, the ensuing radio emission will be strong enough to be detectable in ARIANNA. The detector profits from the fact that the water-ice interface at the bottom of the shelf-ice acts as a mirror for the radio emission and the antennas can be installed on top of the ice (see Figure \ref{fig:ARIANNA}). Since the antenna geometry is not defined by deep holes in the ice, high-gain antennas can be installed that cover the frequency range from 50 to 1000 MHz. | The HRA pilot-array is performing according to expectations with a livetime fraction of more than 0.85. The radio environment is very quiet, the stations operate reliably on battery and solar power and first cosmic rays have been detected. The neutrino analysis only reveals cosmic rays as significant background, which can easily be suppressed using upward facing antennas. The upgrade of a station with a wind-generator and with four upward facing antennas to better be able to test reconstruction algorithms on measured air shower data is planned for the next season. ARIANNA with its 1296 stations will have a a projected sensitivity to cover many of the neutrino fluxes such as shown in \cite{KU}. | 16 | 9 | 1609.07193 |
1609 | 1609.07470_arXiv.txt | We present VLA and ALMA imaging of the deeply-embedded protostellar cluster \ngci\/ from 5~cm to 1.3~mm at angular resolutions as fine as 0\farcs17 (220~AU). The dominant hot core MM1 is resolved into seven components at 1.3~mm, clustered within a radius of 1000~AU. Four of the components have brightness temperatures $>200$~K, radii $\sim300$~AU, minimum luminosities $\sim10^4$~\lsun, and must be centrally heated. We term this new phenomenon a "hot multi-core". Two of these objects also exhibit compact free-free emission at longer wavelengths, consistent with a hypercompact HII region (MM1B) and a jet (MM1D). The spatial kinematics of the water maser emission centered on MM1D are consistent with it being the origin of the high-velocity bipolar molecular outflow seen in CO. The close proximity of MM1B and MM1D (440~AU) suggests a proto-binary or a transient bound system. Several components of MM1 exhibit steep millimeter SEDs indicative of either unusual dust spectral properties or time variability. In addition to resolving MM1 and the other hot core (MM2) into multiple components, we detect five new millimeter and two new centimeter sources. Water masers are detected for the first time toward MM4A, confirming its membership in the protocluster. With a 1.3~mm brightness temperature of 97~K coupled with a lack of thermal molecular line emission, MM4A appears to be a highly optically-thick 240~\lsun\/ dust core, possibly tracing a transient stage of massive protostellar evolution. The nature of the strongest water maser source CM2 remains unclear due to its combination of non-thermal radio continuum and lack of dust emission. | Massive star formation is a phenomenon of fundamental importance in astrophysics, but our understanding of this process is hampered by the heavy extinction and large distances to the nearest sites. Over the past decade, several examples of so-called ``massive protoclusters'', loosely defined as four or more (sub)millimeter continuum sources within $<10000$~AU, have been identified using interferometers \citep[e.g.][]{Hunter06,Rodon08,Palau13,Avison15}. The massive members of such protoclusters often span a wide diversity of evolutionary stages ranging from ultracompact HII~regions and hot molecular cores to cool dust sources \citep{Brogan07,Brogan11,Cyganowski07,Zinchenko12}. In these objects, the separations between the individual members are well-matched to those of the four principal members of the Trapezium cluster, suggesting that (sub)millimeter protoclusters trace the formation phase of the central massive stars of future OB clusters \citep{Hunter06}. Given that all four of the Trapezium stars themselves comprise compact multiple star systems, with separations of 15-400~AU \citep{Grellmann13,Schertl03}, it is critical to continue to study the members of these younger proto-Trapezia with higher angular resolution. Such observations are essential in order to probe for further multiplicity and to resolve the surrounding accretion structures, which are unlikely to be simple uniform disks, particularly in the context of binary and multiple systems. These observations must be undertaken at wavelengths long enough to penetrate the high column of obscuring dust. The \ngc\/ region, known as the Cat's Paw Nebula, contains multiple sites of high mass (M$_*>$8 \msun) star formation \citep{Straw89a,Persi08,Russeil10}. The IRAC/NEWFIRM survey of \citet{Willis13} identified 375 Class I YSOs and 1908 Class II YSOs, indicating a star formation rate several times that of Orion and just below the \citet{Motte03} criteria for a ``mini-starburst''. At the northeastern end of the region, the deeply-embedded source ``I'' was first identified in far-infrared images \citep{Emerson73,McBreen79,Gezari82}. \ngci\/ contains an embedded cluster of stars in near-infrared images \citep{Tapia96,Seifahrt08,Persi08} and a cometary ultracompact HII (UCHII) region \citep{dePree95,Carral02}. The UCHII region is detected in the mid-infrared along with a few other sources \citep{Kraemer99,Harvey83}. However, imaging at 10 and 18~$\mu$m with $0.4''$ resolution \citep{deBuizer02} yielded infrared luminosity estimates for these other sources that are several orders of magnitude less than the bolometric luminosity of the region, suggesting that additional more deeply embedded objects power the (sub)millimeter emission. Our previous 1.3~mm Submillimeter Array (SMA) observations of \ngci\ at $1\farcs6$ resolution revealed a cluster of four millimeter continuum sources in a Trapezium-like arrangement \citep{Hunter06}. Three of these objects are undetected in the infrared, and may represent a large fraction of the total luminosity of the protocluster. For clarity, these four sources have been renamed in this paper from SMA1..SMA4 to MM1..MM4; the UCHII region is known as MM3. The brightest two of these objects (MM1 and MM2) are the sources of hot core line emission seen in \ammonia\/ \citep{Beuther05,Beuther07}, HCN and CH$_3$CN \citep{Beuther08}, and many other organic molecules \citep{Thorwirth03,Kalinina10,Walsh10,Zernickel12}. Using CH$_3$CN from a {\it Herschel} line survey, \citet{Zernickel12} found an average gas temperature for \ngci\/ of 154~K, dominated by emission from the two hot cores MM1 and MM2. A high velocity bipolar outflow traced by various transitions of CO \citep{Qiu11,Leurini06,McCutcheon00,Bachiller90} and HCN 1--0 \citep{Beuther08} emanates from the region around MM1 and MM2, though it has been unclear which of the two objects is the powering source. In this paper, we present sub-arcsecond, comparable-resolution continuum imaging of \ngci\/ from 5~cm to 1.3~mm using the Karl G. Jansky Very Large Array (VLA) and the Atacama Large Millimeter/submillimeter Array (ALMA). These data have allowed us to detect and resolve a number of new likely protocluster members and to carry out detailed analysis of their spectral energy distributions. The spectral line information from these new data will be presented in a future paper. For the distance to \ngci\/, we adopt 1.3 kpc based on recent \water\/ and \methanol\/ maser parallax studies toward source~I(N), located $\sim 2\arcmin$ northeast of source~I: $1.34^{+0.15}_{-0.12}$~kpc \citep{Reid14,Wu14} and $1.26^{+0.33}_{-0.21}$~kpc \citep{Chibueze14}. In the past, the most commonly used value was 1.7~kpc from photometric estimates for the \ngc\/ region \citep{Neckel78,Pinheiro10,Russeil12}, implying a reduction by a factor of 1.7 for derived quantities based on the distance squared, such as mass and luminosity. The rescaled values for \ngci\/ are 700~\msun\/ \citep[c.f.][]{Russeil10} and $1.5\times10^5$~\lsun\/ \citep[c.f.][]{Sandell00}. | Our new ALMA and VLA radio continuum observations spanning 5~cm to 1.3~mm with comparable resolution--as fine as 220~AU--have revealed exciting new features and phenomena in the \ngci\/ massive protocluster. \begin{itemize} \item The dominant millimeter source, MM1, is resolved into seven compact (r$\sim$300~AU) components within a radius of 1000 AU, four of which have brightness temperatures exceeding 200~K, implying minimum luminosities of $10^4$~\lsun\/ and indicative of central heating. We interpret this structure as a ''hot multi-core'' in which at least four massive protostars are independently heating their surrounding cocoons of gas and dust. Several components of MM1 exhibit steep millimeter SEDs indicative of either unusual dust spectral properties or time variability. \item The two hottest dust sources, MM1B and MM1D, also exhibit free-free emission consistent with a HCHII region and a jet, respectively, as well as associated water maser structures. The close pairing of these two objects (440~AU separation) suggests that they form either a proto-binary star or a transient bound system. We suggest that MM1D is the driving source of the large-scale bipolar outflow from this region based on the water maser kinematics. \item The secondary hot core region, MM2, is less extended than MM1, yet also contains evidence for multiple components, MM2A and MM2B, with significantly different millimeter spectral indices consistent with optically-thin and optically-thick dust, respectively. A 6.7~GHz Class II \methanol\/ maser lies between the two objects. \item The enigmatic continuum source MM4A shows a 1.3~mm brightness temperature of 97~K, and a fit to its SED implies a dust optical depth of $\approx5$ at 1.3~mm. Our detection of water maser emission at the systemic velocity of the cluster proves that it is not a background object. The lack of thermal molecular lines remains difficult to understand, given the physical conditions, and we suggest that it may trace a rare evolutionary phase of a high mass protostar. \item In addition to the four previously known millimeter sources, we detect 5 new objects at 1.3~mm whose emission may originate from circumstellar disks around low to intermediate mass protostars. We also detect two new objects at centimeter wavelengths, one of which (CM1) has an infrared stellar counterpart and an SED consistent with an ionized stellar wind. The other (CM2) has a non-thermal spectral index (-0.5) and is associated with the strongest water maser emission in the cluster. While it shares these two characteristics with W3OH-H2O, CM2, in contrast, has no dust emission and may represent a unique object. \end{itemize} In summary, the picture of \ngci\/ emerging from these new observations is a diverse collection of young stellar objects in various states of formation and activity. In addition to the Trapezium-like arrangement of the four principle objects MM1-MM4, we now have evidence that two of these objects are themselves multiple systems, extending the analogy to the Orion Trapezium stars to a second level. Our results support the conclusion from $N$-body simulations that Trapezium-like systems appear frequently in the formation of a cluster \citep{Allison11}, while the close proximity of the seven components of MM1 highlights the inevitably dynamic nature of such proto-OB systems. Finally, the intriguing possibility that the dust emission from MM1 is time variable highlights the importance of future observations of both its continuum and maser emission. \bigskip | 16 | 9 | 1609.07470 |
1609 | 1609.00120_arXiv.txt | The origin of very low-mass hydrogen-burning stars, brown dwarfs, and planetary-mass objects at the low-mass end of the initial mass function is not yet fully understood. Gravitational fragmentation of circumstellar discs provides a possible mechanism for the formation of such low-mass objects. The kinematic and binary properties of very low-mass objects formed through disc fragmentation at early times ($<10$~Myr) were discussed in \cite{li2015} (hereafter L15). In this paper we extend the analysis by following the long-term evolution of disc-fragmented systems, up to an age of 10~Gyr, covering the ages of the stellar and substellar population in the Galactic field. We find that the systems continue to decay, although the rates at which companions escape or collide with each other are substantially lower than during the first 10~Myr, and that dynamical evolution is limited beyond 1~Gyr. By $t=10$~Gyr, about one third of the host stars is single, and more than half have only one companion left. Most of the other systems have two companions left that orbit their host star in widely separated orbits. A small fraction of companions have formed binaries that orbit the host star in a hierarchical triple configuration. The majority of such double companion systems have internal orbits that are retrograde with respect to their orbits around their host stars. Our simulations allow a comparison between the predicted outcomes of disc-fragmentation with the observed low-mass hydrogen-burning stars, brown dwarfs, and planetary-mass objects in the Solar neighbourhood. Imaging and radial velocity surveys for faint binary companions among nearby stars are necessary for verification or rejection for the formation mechanism proposed in this paper. | Low-mass stars and brown dwarfs are among the most common objects in the Galactic field \citep[e.g.,][]{kroupa2001, chabrier2005}. The majority of the neighbours of the Sun are brown dwarfs or are of spectral type~M. Notably, the closest star to our Sun, \object{Proxima Centauri}, is an M-dwarf that orbits the \object{$\alpha$~Cen A/B} system. Our closest neighbours beyond this system are primarily of very low mass -- including \object{Barnard's star} \citep{barnard1916}, the binary brown dwarf \object{Luhman~16} \citep{luhman2013}, which may even have a third companion \citep{boffin2014}, the brown dwarf \object{WISE~0855-0714} \citep{luhman2014}, and many others, such the M-stars \object{Wolf~359} and \object{Lalande~21185}, as well as many brown dwarfs \citep[e.g.,][and numerous others]{wolf1919, ross1926, luyten1979, strauss1999, burgasser2004, burningham2010, kirkpatrick2013, troup2016}. Given their faintness, it is likely that future surveys will reveal the presence of even more nearby low-mass stars and brown dwarfs. Finally, approximately one fourth of the nearby low-mass neighbours of the Sun are known to host one or more companions \citep[e.g.,][]{burgasser2007, luhman2012, duchene2013, ward2015}, while many others are companions to higher-mass stars \citep[see, e.g.,][and references therein]{kouwenhoven2006}. Despite their ubiquity, the formation mechanism for low-mass objects, particularly brown dwarfs, is still poorly understood. It may be possible that brown dwarfs form from core collapse, similar to higher-mass stars \citep[see, e.g.,][]{andre2014, riaz2014, lomax2015}. However, their masses are below or close to the Jeans mass in star forming regions \citep[e.g.,][]{palau2014, degregorio2016}. Another formation mechanism may be the gravitational fragmentation of circumstellar discs, and numerical simulations suggest that this is indeed possible \citep[e.g.,][]{stamatellos2009a, stamatellos2009b, tsukamoto2013, forgan2015, dong2016}. What fraction of circumstellar discs fragment, however, is still unknown. The decay of such disc-fragmented systems results in a population of low-mass stars, brown dwarfs and (free-floating) planetary-mass objects that contribute to shaping the low-mass end of the initial mass function \citep[see, e.g.,][]{thies2007, thies2008, thies2015, marks2015}. Disc fragmentation results in the formation of multiple secondaries around the central star with masses ranging from the planetary to the stellar regime. In this paper we follow the long-term (Gyr) evolution of such disc fragmented systems, using the results of \cite{li2015} (hereafter L15) as initial conditions and follow the dynamical evolution of these systems. Throughout this paper, we follow the classification of L15 by grouping the secondaries into three categories: (i) low-mass hydrogen-burning stars (LMSs) with masses over $80~\mjup$ ($\mjup$ is the mass of Jupiter), (ii) brown dwarfs (BDs) with masses in the range $13-80~\mjup$, and (iii) planetary-mass objects (PMOs) with masses below $13~\mjup$. We assume that all of these secondaries formed through the same mechanisms in our simulations, however each of these three categories may also form through other mechanisms \citep[see][]{whitworth2007, luhman2012}. L15 simulated the dynamical evolution of LMSs, BDs and PMOs formed through disc fragmentation based on the outcomes of the smoothed particle hydrodynamical (SPH) simulations of \cite{stamatellos2009a}. Their analysis covers the first 10~Myr of the dynamical evolution of these systems. They find that most systems attain a reasonably stable configuration at $t=10$~Myr, after a large number of (mostly the lowest mass) secondaries have escaped. A non-negligible fraction of secondaries have paired up into low-mass binaries, many of which escape and some of which remain in orbit around their host star. The simulations of L15 allow a comparison with observations of young stellar populations in or near star-forming regions and OB~associations. For a comparison with the much older field star population, however, a further analysis is necessary. In this paper we therefore carry out $N$-body simulations of disc-fragmented systems up to 10~Gyr, covering the age range of most stars in the Solar neighbourhood. This paper is organised as follows. We describe our methodology and initial conditions in Section~\ref{section:method}. We describe our results in Section~\ref{section:results}. Finally, we summarise and discuss our conclusions in Section~\ref{section:conclusions}. | \label{section:conclusions} We have studied the long-term evolution of disc-fragmented systems, in order to study the orbital configurations of LMSs, BDs and PMOs orbiting solar-type stars. This extends the simulations of L15 to cover 10~Gyr of dynamical evolution, which allows us to compare the products of the disc-fragmentation process with field population in the Solar neighbourhood. We refer to the time span studied by L15 as Stage~I ($0-10$~Myr) and the time span in this paper as Stage~II (10~Myr$-10$~Gyr). Stage~I roughly represents the period of time that the systems spend in or near their natal environment (star-forming regions and OB associations), while Stage~II allows us to compare our results with the stellar population in the Galactic field. Our main conclusions can be summarised as follows: \begin{enumerate} \item Systems continue to decay beyond 10~Myr due to the decay of higher-order systems, and also collisions between members of the system. Almost all of this dynamical evolution occurs in the first Gyr leaving a very stable population after this time. \item For approximately one third of the primaries (37\% in set~1 and 31\% in set~2), the host star ends up as a single star, despite the large ($3-11$) number of secondaries present during the phase of disc fragmentation. More than half of the host stars have one low-mass companion at $t=10$~Gyr, while 6\% (set~1) to 14\% (set~2) of the host stars have two companions. Only a small fraction ($\la 0.1\%$) of the host stars have three companions left, while no host star is able to retain four or more of its companions. \item The number of physical collisions with the host star is non-negligible during Stage~II. On average, each host star experiences 0.20 collisions in set~1, while the value is 0.25 for set~2. On the other hand, physical collisions between secondaries are very rare (less than 0.002 collisions per system). \item For all primaries that ultimately end up as a single star (37\% in set~1 and 31\% in set~2), the final dynamical process to occur is a scattering event involving two companions, which results in the dynamical ejection of one of the companions, and a merger between the host star and the other of the companions. All single host stars in our models are merger products. \item At 10~Gyr, most of the remaining single companions orbit their host star in wide orbits with periods between $100$ years and $1$~Myr (Figure~\ref{figure:period}). Very low-mass secondaries are potentially observable with imaging and radial velocity surveys. PMOs with separations less than 50~AU are absent, although higher-mass BDs and LMSs occur more frequent (ranging between, on average, 0.05 and 0.37 per system). The double companion systems contain two companions which have widely separated orbits, with period ratios mostly ranging between $P_{\rm out}/P_{\rm in}=10$ and $P_{\rm out}/P_{\rm in}=10^4$. \item Binary companion systems orbiting the host star at $t=10$~Gyr have external-to-internal period ratios mostly ranging between $P_{\rm ext}/P_{\rm int}=10$ and $1000$. The binary companion masses are usually comparable, and the mass ratio of the binary system with respect to the host star is typically between $q_{\rm ext}=0.15$ and 0.40. The binary companions show a more or less thermal internal eccentricity distribution ($e_{\rm int}$), while their external eccentricities ($e_{\rm ext}$) tend to be more circular (although high eccentricities also exist among these external orbits). About half of the external inclinations of the binary companions are below $i_{\rm ext}=20^\circ$ in set~1 and $i_{\rm ext}=10^\circ$ in set~2, while a large majority of the binary companions have internal inclinations ($i_{\rm int}$) beyond $90^\circ$, i.e., they have retrograde orbits. \end{enumerate} Most nearby stars are low-mass stars or brown dwarfs that formed billions of years ago. Our study allows us to speculate somewhat on how these low-mass neighbours may have formed and evolved over time. Disc fragmentation provides a possible (but certainly not the only) solution for the origin of many of these. The solar system itself is clearly a result of planet formation through core-accretion. Also our closest neighbour, the \object{$\alpha$ Centauri} triple system, has likely formed differently, with its very low-mass companion \object{Proxima Centauri} perhaps being either a result of a capture event \citep[e.g.,][]{kouwenhoven2010, moeckel2011, parker2014} or a result of a triple decay event \citep[e.g.,][]{reipurth2012}. The other known objects with a distance smaller than that of \object{Sirius} are the very low mass single objects \object{Barnard's Star} \citep{barnard1916}, \object{WISE~0855-0714} \citep{luhman2014}, \object{Wolf~359} and \object{Lalande~21185}, and the very low-mass binary system \object{Luhman~16} \citep{luhman2013}. The nearby population beyond \object{Sirius} is also dominated by very low-mass objects. Although disk fragmentation may not be the dominant scenario responsible for the origin of this low-mass population, it does explain many of its properties, including the origin of the very low-mass binaries and multiples in the proximity of the Sun, such as \object{Luhman~16}, \object{Luyten~726-8}, \object{EZ~Aquarii}, \object{Struve~2398}, \object{Gloobridge~34}, and \object{Epsilon~Indi}. Our study provides insight into a possible formation scenario for LMS, BD and PMO companions to solar-type stars in the more distant Galactic field, as well as their free-floating siblings, both in the form of singles and binaries. Many stars in the Galactic field are part of a binary or of a hierarchical multiple system \citep[e.g.,][]{tokovinin2002, tokovinin1997, tokovinin2008, tokovinin2014}. As dynamical capture is rare, these systems are almost certainly formed as such. Our theory provides an excellent explanation for the origin of several hierarchical multiples, for example, the hierarchical triple systems \object{HIP68532} and \object{HIP69113}, which both consist of a main sequence star, with a double companion made up of two low-mass stars \citep[see figure 9 in][]{kouwenhoven2005}. The scarcity of close-in brown dwarf companions predicted by our model is also reflected in observations \citep[e.g.,][]{kouwenhoven2007}. A deeper analysis and more detailed comparison with observations is necessary, as with our simulations we have only covered a subset of all possible initial conditions that may lead to disc fragmentation. In addition, we have not taken into account the effect of encounters with neighbouring stars and brown dwarfs, an affect that may be particularly important during first few million years, when a disc-fragmented system is still in its dense natal environment where the circumstellar disc is initially exposed to the violent interstellar medium \citep[e.g.,][]{bik2010}, and subsequently participates in rapid exchange of energy with its neighbours \citep[see, e.g.,][and numerous others]{allison2009}. Stellar encounters can disrupt existing stellar and planetary systems \citep[e.g.,][]{zheng2015, wang2015a, wang2015b, wang2016}, although in the Galactic field this mostly affects the widest companions. In the case of multi-companions systems, perturbations of outer companions can induce destabilization as a perturbation of outer component can propagate to the inner system \citep[e.g.,][]{hao2013, cai2015, cai2016}. The Galactic field itself also affects the evolution of wide binaries \citep[e.g.,][]{jiang2010, kaib2013}. Finally, other stars and brown dwarfs in the environment may be captured by the host system following a close encounter, and previously escaped secondaries may be captured by other stars as well \citep[e.g.,][]{kouwenhoven2010, perets2012}, which may provide an alternative solution to the origin of wide, low-mass companions. The stars and brown dwarfs in the solar neighbourhood likely represent a mixed population resulting from different formation mechanisms and from a different environmental interaction history. Despite the many unknowns, our model make clear predictions that can be statistically compared to nearby stars and brown dwarfs, to further constrain their origin. | 16 | 9 | 1609.00120 |
1609 | 1609.02914_arXiv.txt | The orbits, atmospheric parameters, chemical abundances, and ages of individual stars in the Milky Way provide the most comprehensive illustration of galaxy formation available. The Tycho-\gaia\ Astrometric Solution (\tgas) will deliver astrometric parameters for the largest ever sample of Milky Way stars, though its full potential cannot be realized without the addition of complementary spectroscopy. Among existing spectroscopic surveys, the RAdial Velocity Experiment (\rave) has the largest overlap with \tgas\ ($\gtrsim$200,000 stars). We present a data-driven re-analysis of \Nspectra\ \rave\ spectra using \thecannon. For red giants, we build our model using high-fidelity \apogee\ stellar parameters and abundances for stars that overlap with \rave. For main-sequence and sub-giant stars, our model uses stellar parameters from the \epic. We derive and validate effective temperature $\teff$, surface gravity $\logg$, and chemical abundances of up to seven elements (O, Mg, Al, Si, Ca, Fe, Ni). We report a total of \ReportedAbundances\ elemental abundances with a typical precision of 0.07~dex, a substantial improvement over previous \rave\ data releases. The synthesis of \raveon\ and \tgas\ is the most powerful data set for chemo-dynamic analyses of the Milky Way ever produced. | \label{sec:introduction} The Milky Way is considered to be our best laboratory for understanding galaxy formation and evolution. This premise hinges on the ability to precisely measure the astrometry and chemistry for (many) individual stars, and to use those data to infer the structure, kinematics, and chemical enrichment of the Galaxy \citep[e.g.,][]{Nordstrom_2004,Schlaufman_2009,Deason_2011,Casagrande_2011, Ness_2012,Ness_2013a,Ness_2013b,Casey_2012,Casey_2013,Casey_2014a,Casey_2014b, Boeche_2013,Kordopatis_2015,Bovy_2016}. However, these quantities are not known for even 1\% of stars in the Milky Way. Stellar distances are famously imprecise \citep[e.g.,][]{van_Leeuwen_2007,Jofre_2015,Madler_2016}, proper motions can be plagued by unquantified systematics from the first epoch observations \citep[e.g.,][]{Casey_Schlaufman_2015}, and stellar spectroscopists frequently report significantly different chemical abundance patterns from the same spectrum \citep{Smiljanic_2014}. The impact these issues have on scientific inferences cannot be understated. Imperfect astrometry or chemistry limits understanding in a number of sub-fields in astrophysics, including the properties of exoplanet host stars, the formation (and destruction) of star clusters, as well as studies of stellar populations and Galactic structure, to name a few. The \gaia\ mission represents a critical step forward in understanding the Galaxy. \gaia\ is primarily an astrometric mission, and will provide precise positions, parallaxes and proper motions for more than $10^9$ stars in its final data release in 2022. While this is a sample size about four orders of magnitude larger than its predecessor \hipparcos, both astrometry and chemistry are required to fully characterize the formation and evolution of the Milky Way. \gaia\ will also provide radial velocities, stellar parameters and chemical abundances for a subset of brighter stars, but these measurements will not be available in the first few data releases. Until those abundances are available, astronomers seeking to simultaneously use chemical and dynamical information are reliant on ground-based spectroscopic surveys to complement the available \gaia\ astrometry. The first \gaia\ data release will include the Tycho-Gaia Astrometric Solution \citep[hereafter \tgas;][]{Michalik_2015a,Michalik_2015b}: positions, proper motions, and parallaxes for approximately two million stars in the Tycho-2 \citep{Hog_2000} catalog. After cross-matching all major stellar spectroscopic surveys\footnote{Specifically we cross-matched the Tycho-2 catalog against the \apogee\ DR13 \citep{Zasowski_2013}, \ges\ internal DR4 \citep{Gilmore_2012,Randich_2013}, \galah\ internal DR1 \citep{DeSilva_2015}, \lamost\ DR1 \citep{Cui_2012}, and \rave\ DR4 \citep{Steinmetz_2006} catalogs.}, we found that the RAdial Velocity Experiment \citep[\rave;][]{Steinmetz_2006} survey is expected to have the largest overlap with the first \gaia\ data release: up to 264,276 stars. We used the \gaia\ universe model snapshot \citep{Robin_2012} to estimate the precision in parallax and proper motions that could be available in the first \gaia\ data release (DR1) for stars in those overlap samples. Comparing the expected precision to what is currently available, we further found that the \rave\ survey will benefit most from \gaia\ DR1: the distances of 63\% of stars in the \rave--\gaia\ DR1 overlap sample are expected to improve with the first \gaia\ data release, and 47\% of stars are likely to have better proper motions. Although the \gaia\ universe model assumes end-of-mission uncertainties --- and does not account for systematics in the first data release --- this calculation still provides intuition for the relative improvement that the first \gaia\ data release can make to ground-based surveys. The expected improvements for \rave\ motivated us to examine what chemical abundances were available from those data, and to evaluate whether we could enable new chemo-dynamic studies by contributing to the existing set of chemical abundances. We briefly describe the \rave\ data in Section~\ref{sec:data}, before explaining our methods in Section~\ref{sec:method}. In Section~\ref{sec:validation} we outline a number of validation experiments, including: internal sanity checks, comparisons with literature samples, and investigations to ensure our results are consistent with expectations from astrophysics. We discuss the implications of these comparisons in Section~\ref{sec:discussion}, and conclude with instructions on how to access our results electronically. | \label{sec:discussion} We have performed an independent re-analysis of \Nspectra\ \rave\ spectra, having derived atmospheric parameters ($\teff$, $\logg$, [Fe/H]) for all stars, as well as detailed chemical abundances for red giant branch stars. When combined with the \tgas\ sample, these results amount to a powerful compendium for chemo-dynamic studies of the Milky Way. However, our analysis has caveats. Inferences based on these results should recognize those caveats, and acknowledge that these results are subject to our explicit assumptions, some of which are provably incorrect. For practical purposes we adopted separate models: one for the giant branch and one for the main-sequence. A third model was used to derive relative weights for which results to use. The relative weighting we have used does not have any formal interpretation as a likelihood or belief (in any sense): it was introduced for practical reasons to identify systematic errors and combine results for multiple models. Because the relative weights have no formal interpretation, it is reasonable to consider this method is as \emph{ad hoc} as any other approach. The relative weighting has no warranty to be (formally) correct, and therefore may introduce inconsistencies or systematic errors rather than minimizing them. If we only consider the results from individual models, there are a number of cautionary remarks that stem from the construction of the training set. The labels for red giant branch stars primarily come from \apogee, where previous successes with \thecannon\ have demonstrated that \apogee\ labels based on high S/N data can be of high fidelity \citep{Ness_2015,Ness_2016,Ho_2016,Casey_2016b}. However, the lack of metal-poor stars in the \apogee/\rave\ overlap sample produced a tapering-off in the test set --- where \emph{no} stars had reliably reported metallicities below that of the training set --- which forced us to construct a heterogeneous training set. The metal-poor stars included in this sample are from high-resolution studies \citep{Fulbright_2010,Ruchti_2011}, but it is not known if the stellar parameters are of high fidelity because we have a limited number of quality statistics available. Moreover, there is no guarantee that the stellar parameters \emph{or} abundances are on the same scale as \apogee\ \citep[and good reasons to believe they will not be; see][]{Smiljanic_2014}. If the metallicities of metal-poor giant stars were on the same scale as the \apogee\ abundances, there is a larger issue in verifying that the main-sequence metallicities and giant branch metallicities are on the same scale. The training set for the main-sequence stars includes metallicities from a variety of sources, including \lamost, and the fourth \rave\ data release. Even on expectation value, there is no straightforward manner to ensure that the main-sequence model and the red giant branch model produce metallicities on the same scale. We see no systematic offset in metallicities of dwarf and giant stars that overlap between \rave\ and the \ges\ survey, suggesting that if there is a systematic offset, it must be small. Nevertheless, these are only verification checks based on $<1$\% of the data, and there is currently insufficient data for us to prove both models are on the same abundance scale. For some of the most metal-poor giant stars in \rave\, we \emph{know} the abundances are not on the same scale as \apogee, because we were forced to adopt abundances for specific elements when they were unavailable. Although we sought to adopt mean level of Galactic chemical enrichment at a given overall metallicity, this is not a representative abundance. Even if that is the \emph{mean} enrichment at that Galactic metallicity, there is no requirement for zero abundance spread. More fundamentally, we are incorrectly asserting that the element \emph{must} be detectable in the photosphere of the star. There may be no transition that is detectable in that star, even with zero noise, because it is too weak to have any effect on the spectrum. In the most optimistic case, this could be considered to be forcing the model to make use of correlated information between abundances. In a more representative (pessimistic) case, we are simply invoking what all abundances should be at low metallicity. This choice is reflected in the abundances of the test set. While we do recover trustworthy metallicities for ultra metal-poor ($[{\rm Fe/H}] \lesssim -4$) stars like CD-38~245, the individual abundances for all extremely metal-poor stars aggregate (in [X/Fe] space; Figure~\ref{fig:gce}) at the assumed abundances for the metal-poor stars in our sample. Thus, while the overall metallicities appear reliable, the individual abundances for extremely metal-poor stars in the test set cannot be considered trustworthy in any sense. For this reason we have updated the electronic catalog to discard these results as erroneous. In Section~\ref{sec:method} we assumed that any fibre- or time-dependent variations in the \rave\ spectra are negligible. This is provably incorrect. Indeed, \citet{Kordopatis_2013} note that the effective resolution of \rave\ spectra varies from $6{,}500 < \mathcal{R} < 8{,}500$, and that the effective resolution is a function of temperature variations, fibre-to-fibre variations, and thus position on the CCD \citep{Steinmetz_2006}. For this reason we ought to expect our derived stellar parameters or abundances to be correlated either with the fibre number, with the observation date, or both. If significant, the trend could produce systematically offset stellar abundances solely due to the fibre used. \citet{Kordopatis_2013} conclude that resolution-based effects on the \rave\ stellar parameters should be a second-order effect. We have not seen evidence of these resolution-based correlations in our results, however, we have only performed cursory (non-exhaustive) experiments to investigate this issue. We have shown some potential outcomes when the test set spectra differ significantly from the spectra in the training set. Test set spectra that is `unusual' from the training set can be projected as peculiar artefacts in label space. In other words, unusual spectra can appear as `clumps' in regions of parameter space that we could consider as being normal (e.g., an over-density of solar-type stars). We addressed this issue for the main-sequence and giant models by using a third model (Section \ref{sec:a-simple-model}) to calculate relative weights. However, spectra that are unusual from the training set used in the simple model could still project as systematic artefacts in label space. Indeed, there are two known artefacts in our data that are relevant to this discussion. The first is a small over-density at the base of the giant branch, which is likely a consequence of joining the 9-label and 3-label models. The second has an astrophysical origin: there are no hot stars ($\teff > 8000$~K) present in our training set, yet there are many in the \rave\ survey. However, the \rave\ pre-processing pipeline \citep[\texttt{SPARV};][]{Steinmetz_2006,Zwitter_2008} performs template matching against grids of cool \emph{and} hot stars, and therefore we can use that information to identify hot stars. In Figure \ref{fig:hot-stars} we show our derived effective temperatures $\teff$ and surface gravities $\logg$, where each hexagonal bin is colored by the \emph{maximum} temperature reported by \texttt{SPARV} for any star in that bin. We show the maximum temperature reported by \texttt{SPARV} to demonstrate that hot stars project into a single clump in our label space --- near the turn-off --- in a region where we may otherwise be deceived into thinking the observed over-density is consistent with expectations from astrophysics. This line of reasoning extends to spectra with other peculiar characteristics (e.g., chromospheric emission), and for these reasons we recommend the use of additional metadata to investigate possible artefacts. In our catalog we have included a column containing a boolean flag to indicate whether the labels pass very weak quality constraints. Specifically, we flag results as failing our quality constraints if \texttt{SPARV} indicates a $\teff > 8000$~K, or if $\chi_r^2 > 3$, or if $S/N < 10$~pixel$^{-1}$. These quality constraints represent the minimum acceptable conditions and should not be taken verbatim: judicious use of the morphological classifications \citep{Matijevic_2012} or additional metadata from the \rave\ pre-processing pipelines is strongly encouraged. We have presented a comprehensive collection of precise stellar abundances for stars in the first \gaia\ data release. In total we derive stellar atmospheric parameters for \ReportedStellarParameters\ stars, and report more than 1.69 million abundances. Despite the caveats and limitations discussed here, our validation experiments and comparisons with high-resolution spectroscopic studies suggests that our results have sufficient accuracy and precision to be useful for chemo-dynamic studies that become imminently feasible only in the era of \gaia\ data. We are optimistic that the \raveon\ catalog will advance understanding of the Milky Way's formation and evolution. \subsection*{Access the results electronically} \noindent{}Source code for this project is available at \texttt{\giturl}, and this document was compiled from revision hash \texttt{\githash} in that repository. Derived labels, associated errors, and relevant metadata are available electronically through the \rave\ database from 19 September 2016 onwards. Please note that it is a condition of using these results that the \rave\ data release by \citet{Kunder_2016} must also be cited, as the work presented here would not have been possible without the tireless efforts of the entire \rave\ collaboration, past and present. | 16 | 9 | 1609.02914 |
1609 | 1609.01126_arXiv.txt | An algorithm is used to generate new solutions of the scalar field equations in homogeneous and isotropic universes. Solutions can be found for pure scalar fields with various potentials in the absence and presence of spatial curvature and other perfect fluids. A series of generalisations of the Chaplygin gas and bulk viscous cosmological solutions for inflationary universes are found. Furthermore other closed-form solutions which provide inflationary universes are presented. We also show how the Hubble slow-roll parameters can be calculated using the solution algorithm and we compare these inflationary solutions with the observational data provided by the Planck 2015 collaboration in order to constraint and rule out some of these models. | Recent cosmological data indicate that the universe has undergone two acceleration phases: an early acceleration phase called 'inflation', prior to a radiation-dominated era, and a more recent era of accelerated expansion which appears to continue today \cite{Teg,Kowal,Komatsu,suzuki11,Ade15,planck2013,planck2015}. The gravitationally repulsive stress that is responsible for the current acceleration of the universe is called 'dark energy' and must possess sufficient negative pressure to exert gravitational repulsion. Its specific identity is still unknown and it may result from a modification of general relativity when gravity is very weak or the presence of a specific unknown matter field. Whilst a range of \textquotedblleft exotic\textquotedblright\ fluids and modifications of the gravitational action can provide cosmological acceleration, scalar fields are the simplest candidates to explain the acceleration phases of the universe. Moreover, scalar fields also have various applications in the inflationary phase of the universe, for instance in driving chaotic inflation \cite{lid02}. While the same scalar field might explain both of the periods of accelerated expansion, no convincing cosmological model has been found with this as a natural feature. In a scalar field cosmology the field equations are of second-order where the scalar field is introduced an extra degree of freedom, with a corresponding conservation equation. These equations display unexpected complexity. Simple power-law potentials for the scalar field can create finite-time singularities during inflation \cite{sing} and lead to chaotic dynamics \cite{page}, or singularity avoidance \cite{BMatz} if the universe is closed. Very few exact scalar-field solutions in a Friedmann-Lema\^{\i}tre-Robertson-Walker spacetime (FLRW) with spatial curvature are known \cite{hall,Ea}. In a spatially-flat FLRW spacetime closed-form solutions with or without sources for different scalar field potentials, or scalar fields which mimic other fluids, such as the Chaplygin gas, can be found in \cite{jdbchg, jdbmin,jdbnew,muslinov,ellis,barrow1,Mendez,barrowhyp,chimento,newref2,rubano,sahni,Kahya,basil,tscqg,palprd2015,pan} while some other classes of integrable scalar-field models are also given \cite{suzuki,fre,palia1,newref1}. Some solutions for three-dimensional FLRW spacetimes are given in \cite{bblanc,william,shaw}. However, scalar-field cosmology is conformally equivalent to other scalar-tensor theories, like Brans-Dicke or $f\left( R\right) $-gravity. \cite{bcot, maeda}. Hence, closed-form solutions of the conformally equivalent theories (see \cite{sol3,sol4} and references therein) can be used to construct closed-form solutions or to find new integrable systems for the non-minimally coupled scalar-field model. Recently, with the use of nonlocal conservation laws in \cite{dimakis}, the general analytical solution has been expressed for an arbitrary scalar field with an arbitrary number of independent perfect fluids possessing constant equation of state parameters in spatially flat or nonflat FLRW universes. These general results are applied in this paper to derive precise forms of the scalar field potential for various simple time-dependent forms for the expansion scale factor, or for particular equation of state parameters for the scalar field. Finally, the Hubble slow-roll parameters are studied for these closed-form solutions so that we can compare the inflationary parameters with the observable constraints provided by the Planck 2015 observations \cite{planck2015}. \ The plan of this paper is as follows. In section \ref{field}, the basic properties and definitions of scalar-field models are introduced. The cosmological metric we consider is the four-dimensional FLRW spacetime, while the gravitational action is that of general relativity with a minimally coupled scalar field. We review previous results in the literature and present the general analytical solution for the cosmological field equations for arbitrary scalar-field potential. Exact closed-form solutions obtained by using these general results are presented in sections \ref{exactS1} and \ref{exactS2}. Specific closed-form solutions are derived for spatially-flat FLRW universes when only one scalar field and a perfect fluid with constant equation of state parameter are present. For specific values of the barotropic parameter for the matter source, these results give closed-form solutions in the case of a nonflat FLRW universe. In section \ref{hsrpar}, we derive the Hubble slow-roll parameters for our models and compare them with that of the Planck 2015 data to isolate observationally allowed parameter ranges. Finally, in section \ref{conc}, we discuss our results and draw conclusions. | \label{conc} In this work we studied exact solutions in scalar field cosmology using a new mathematical approach, and with emphasis on inflationary models. We have found new closed-form solutions for spatially flat FLRW universes with or without an extra matter source. For the latter cosmological scenario, we determined exact solutions for the case in which the scalar field mimics the perfect fluid, the scalar field has a constant equation of state parameter different from that of the perfect fluid, and when the scalar field provides two perfect-fluid terms in the field equations. The first solution is the well known special solution of the exponential potential, while in the other two solutions the scalar field potentials are expressed in hyper trigonometric functions and the unified cold dark matter potential is recovered. \ Furthermore, these expressions can be applied in order to construct other solutions in a FLRW spacetime with spatial curvature. In the cosmological scenario in which the universe is dominated by the scalar field we determined the scalar field model in which the equivalent equation of state parameter is that of the Chaplygin gas, or some generalizations of the Chaplygin gas which have been proposed in the literature. \ We also considered solutions in which the Hubble function is expressed in terms of the Lambert function or by logarithmic function. These models provide exact inflationary universe solutions. We compared these solutions with the constraints on inflation from the Planck 2015 collaboration. In order to perform this analysis we expressed the Hubble slow-roll (HSR) parameters in terms of the expansion scale factor in the variables defined by our solution-generating functions. For every specific model and solution we calculated the HSR parameters and we derived the spectral indices in the first approximation. The diagrams for the density perturbations $\left( n_{s}\right) ~$with the scalar ratio $\left( r\right) $ and the variation $n_{s}^{\prime}$ have been derived and the subset of models which are compatible with the Planck 2015 data set are delineated. | 16 | 9 | 1609.01126 |
1609 | 1609.08836_arXiv.txt | Template-based extrapolations from only one photometric band can be a cost-effective method to estimate the total infrared (IR) luminosities ($L_{\mathrm{IR}}$) of galaxies. By utilizing multi-wavelength data that covers across 0.35--500\,$\mathrm{\mu m}$ in GOODS-North and GOODS-South fields, we investigate the accuracy of this monochromatic extrapolated $L_{\mathrm{IR}}$ based on three IR spectral energy distribution (SED) templates (\citealt[CE01]{Chary2001}; \citealt[DH02]{Dale2002}; \citealt[W08]{Wuyts2008a}) out to $z\sim 3.5$. We find that the CE01 template provides the best estimate of $L_{\mathrm{IR}}$ in {\it Herschel}/PACS bands, while the DH02 template performs best in {\it Herschel}/SPIRE bands. To estimate $L_{\mathrm{IR}}$, we suggest that extrapolations from the available longest wavelength PACS band based on the CE01 template can be a good estimator. Moreover, if PACS measurement is unavailable, extrapolations from SPIRE observations but based on the \cite{Dale2002} template can also provide a statistically unbiased estimate for galaxies at $z\lesssim 2$. The emission of rest-frame 10--100\,$\mathrm{\mu m}$ range of IR SED can be well described by all the three templates, but only the DH02 template shows nearly unbiased estimate of the emission of the rest-frame submillimeter part. | \label{sect:intro} The infrared (IR) sky is dominated by emission from star-forming galaxies (\citealt{Lagache2005,Viero2009,Viero2013}). Massive, young stars emit a large amount of ultraviolet (UV) radiation, which is absorbed by surrounding dust grains and then re-emitted at IR wavelength. Physical properties, such as dust components, dust temperature, star formation rate (SFR), can be decoded from the IR spectral energy distributions (SEDs) of galaxies. The total bolometric infrared luminosity $L_{\mathrm{IR}}$, which is simply an integration of the SED in a wavelength range that most often in 8--1000\,$\mathrm{\mu m}$ (e.g., \citealt{Kennicutt1998}), traces the total energy absorbed by dust. Hence $L_{\mathrm{IR}}$ can be used to estimate the obscured SFR (\citealt{Kennicutt1998,Kennicutt2012}), although non-star formation driven heating, like old stellar populations or active galactic nucleus (AGNs) heating, may have a non-negligible contribution in some galaxies. The most reliable method to estimate $L_{\mathrm{IR}}$ is deriving it from observations that sampling nearly the whole IR wavelength by directly integrating or multi-wavelength fitting of empirical templates (e.g., CE01, DH02) or physical dust models (e.g., \citealt{SiebenmorgenKruegel2007}). However, for the majority of galaxies we concerned, observations in IR bands are not enough to allow a direct integration even a reliable SED fitting. Using photometric data as less as possible to extract information about $L_{\mathrm{IR}}$ would be a cost-effective way to solve this problem. \cite{Sanders1996} provided an equation that using four {\it Infrared Astronomical Satellite} ({\it IRAS}) flux densities at 12, 25, 60 and 100\,$\mathrm{\mu m}$ to estimate $L_{\mathrm{IR}}$. DH02 derived two similar relations based on three {\it IRAS} bands (25, 60 and 100\,$\mathrm{\mu m}$) and three {\it Spitzer}/Multiband Imaging Photometer for {\it Spitzer} (MIPS) bands (24, 70 and 160\,$\mathrm{\mu m}$) but for the 3--1100\,$\mathrm{\mu m}$ bolometric luminosity. Recently, \cite{Dale2014} updated the previous results and added possible contribution from AGN, derived one {\it IRAS}-based and two {\it Spitzer}-based relations to estimate the total luminosity over 5--1100\,$\mathrm{\mu m}$ for a range of mid-infrared AGN fractions. Moreover, \cite{Boquien2010} provided detailed relations to estimate $L_{\mathrm{IR}}$ from just one or two {\it Spitzer} bands, especially from 8 and 24\,$\mathrm{\mu m}$ bands, calibrated by local galaxy samples. {\it Herschel}-based (from 70, 100, 160 and 250\,$\mathrm{\mu m}$ bands) relations were also developed by \cite{Galametz2013} using observations of local galaxies. We note that most of the empirical relations between $L_{\mathrm{IR}}$ and observed flux densities provided by above works require more than one photometric observations, while a few relations only need observation in one band but calibrated by local galaxies. One typical monochromatic method that easily generalizes to high-redshift galaxies is extrapolating $L_{\mathrm{IR}}$ based on IR SED templates. Following this method, \cite{Elbaz2010} have utilized observations from {\it Herschel} to check the self-consistency of the used IR SED templates. However, it is still necessary to carefully investigate the accuracy of this method by applying a more actual $L_{\mathrm{IR}}$ as reference. {\it Herschel Space Observatory}\footnote{{\it Herschel} is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA.} (\citealt{Pilbratt2010b}) provides observations with excellent angular resolution for many famous and well-studied extragalactic fileds from far-infrared (FIR) to submillimeter band. Photodetector Array Camera and Spectrometer (PACS; \citealt{Poglitsch2010}) onboard {\it Herschel} observed the FIR sky in 70, 100 and 160\,$\mathrm{\mu m}$ bands, while Spectral and Photometric Imaging Receiver (SPIRE; \citealt{Griffin2010}) present maps of the submillimeter sky in 250, 350 and 500\,$\mathrm{\mu m}$ bands. Combining with near-infrared (NIR) and mid-infrared (MIR) observations from {\it Spitzer Space Telescope} and optical data from other telescopes, such as {\it Hubble Space Telescope} ({\it HST}), we are able to perform a multi-wavelength SED fitting from UV to submillimeter. Because dust emission results from dust attenuation in optical band, we believe that SED fitting with constraints from optical observations could provide a more actual estimate of $L_{\mathrm{IR}}$ relative to that of fitting only included IR constraints as which is done by most of previous works (e.g., \citealt{Elbaz2010,Galametz2013}). In this paper, we use photometric data of the Great Observatories Origins Deep Survey northern (GOODS-North) and southern (GOODS-South) fields from UV to submillimeter to study the template-based monochromatic extrapolated total infrared luminosities $L_{\mathrm{IR}}$. Here, $L_{\mathrm{IR}}$ is defined as integration of the SED over the 8-1000\,$\mathrm{\mu m}$ wavelength range. We first compare the differences between the extrapolated $L_{\mathrm{IR}}$ from the bands that were frequently used (\citealt{Elbaz2010,Elbaz2011,Wuyts2011a}) based on different templates. A multi-wavelength SED fitting is taken in order to obtain a reference $L_{\mathrm{IR}}$ which is then used to compare with the extrapolated values from different templates. We also repeat this comparison in the rest-frame to study how well templates can describe the IR emission of galaxies. \par This paper is organized as follows. Section 2 presents a brief description of our multi-wavelength data. In Section 3, we introduce our method to calculate the monochromatic extrapolated $L_{\mathrm{IR}}$ and the code used for SED fitting. We then analyse our results and put forward some main conclusions in Section 4. We discuss how factors such as confusion noise of observations influence our conclusions in Section 5 and give a short summary in Section 6. Throughout this paper, we adopt a \cite{Chabrier2003} initial mass function (IMF) and a $\Lambda$CDM cosmology with $H_0=70$\,km\ s$^{-1}$\,Mpc$^{-1}$, $\Omega_{\Lambda}=0.7$, $\Omega_{\mathrm{m}}=0.3$. | To check the self-consistency of the {\sc MAGPHYS} results, we plot the distributions of stellar mass $M_*$, mass-weighted age $\mathrm{age}_M$, star formation timescale parameter $\gamma$, and the present star formation rate to initial star formation rate ratio $\psi/\psi_0$ derived from SED fitting for our clean sample in Figure \ref{fig5}. The median values of each distribution are labelled by vertical solid lines, and the dashed line in the $\psi/\psi_0$ distibution marks the value of $\mathrm{e}^{-1}\approx 0.37$. \placefigure{fig5} The stellar masses of our clean sources range from $1.68\times 10^9$ to $9.77\times 10^{11}\,M_{\odot}$, exhibit a media and $68\%$ dispersion of $\log(M_*/M_{\odot})=10.82_{-0.53}^{+0.41}$ which is the modest range of the main sequence galaxies at similar redshift \citep{Speagle2014}. The age of galaxy plotted in Figure \ref{fig5} is not the time since the onset of star formation which is widely used in simple stellar population (SSP) models but has no real physical meaning when the continuous SFH is applied (\citealt{daCunha2015}). In this work, we use the mass-weighted age, defined as \begin{equation} \mathrm{age}_M=\frac{\int_0^t\mathrm{d}t'\,t'\psi(t-t')}{\int_0^t\mathrm{d}t'\psi(t-t')} \end{equation} to describe the overall age of the galaxy model, where $\psi(t-t')$ is the star formation history of each model. For our clean sample, the distribution of this age has a median of 2.09\,Gyr, while the central 68th percentile range is 1.12--3.63\,Gyr. As mentioned above, {\sc MAGPHYS} assumes an exponentially declining model with a timescale paramter $\gamma$ to describe the continuous SFH. The resulting distribution of $\gamma$ in Figure \ref{fig5} shows that the median star formation timescale is 3.33\,Gyr ($\gamma=0.30\,\mathrm{Gyr}^{-1}$), larger than the median mass-weighted age. Combining the time when the star formation starts with $\gamma$, we are able to compute the ratio between the present instantaneous SFR $\psi$ and the initial SFR $\psi_0$. This ratio is in the mdeian larger than the e-folding value $\mathrm{e}^{-1}$ by 0.04, suggesting that more than half of our galaxies do not reach their e-folding time of star formation yet. In fact, nearly $90\%$ of our clean sample present a $\psi/\psi_0$ ratio of larger than $10\%$. Therefore, most of these galaxies still maintain a fairly active star formation, which lead them to product enough dust as well as become bright in infrared. We include {\it Herschel}/SPIRE observations in our SED fitting, but it is well known that these observations suffered serious confusion noise. To investigate how confusion noise affect the result of SED fitting, we perform a repeated fitting using data without SPIRE measurements (e.g., 250, 350 and 500\,$\mathrm{\mu m}$ bands) and compare some derived quantities, denoted as $Q^{\mathrm{ns}}$, with previous result in Figure \ref{fig6}. Statistically, for stellar mass $M_*$, dust luminosity $L_{\mathrm{dust}}$ and SFR\footnote{SFR from {\sc MAGPHYS} correspond to SFR averaged over the past 100 Myr.}, SED fitting without submillimeter data has nearly no effect on these quantities. That is because $M_*$ is almost determined by optical and NIR observations, and due to the energy balance technique of {\sc MAGPHYS} code, using data only from PACS is enough to constrain $L_{\mathrm{dust}}$ and SFR. However, owing to the fact that $M_{\mathrm{dust}}$ is dominated by large dust grains which is in thermal equilibrium at relatively low temperature, observations from SPIRE are necessary to constrain the cold component of dust emission. As a result, SED fitting without these observations could not give a reliable estimate of $M_{\mathrm{dust}}$. For all clean sources that have SPIRE observations, the $L_{\mathrm{dust}}^{\mathrm{ns}}/L_{\mathrm{dust}}$ ratio has a median value and the 16th-84th percentile range of $1.08_{-0.09}^{+0.33}$. Thus the confusion noise of SPIRE observations has no effect on our results about $L_{\mathrm{dust}}$. An additional SED fitting without using all {\it Herschel} observations is also performed for comparing. We find that only $M_*$ can hold a nearly unchanged result with a little larger dispersion, while $L_{\mathrm{dust}}$, $M_{\mathrm{dust}}$ and SFR can not trace the previous results anymore. The tendencies of the last three quantities are all similar to that of $M_{\mathrm{dust}}$ in Figure \ref{fig6}, present overestimate at low-value end and underestimate at high-value end. Therefore, constraint from FIR observations is necessary to obtain a reliable estimate of $L_{\mathrm{dust}}$ as well as SFR. In the meantime, we also check the differences of SED fitting results when apply the CB07 version of BC03 models to calculate stellar emission. We find that only stellar mass $M_*$ exhibit a significant systematic underestimate comparing with result of the previous BC03 models, while $L_{\mathrm{dust}}$, $M_{\mathrm{dust}}$ and SFR all present a nearly statistically unchanged result. Consequently, changing the choice of stellar emission models between BC03 and its CB07 version could not change our conclusions on dust luminosity. Furthermore, the main difference between BC03 and CB07 is the updated treatment of the thermally pulsing asymptotic giant branch (TP-AGB) phase of stellar evolution, which would lead to a lower stellar mass and a smaller age under the CB07 SSP models \citep{Bruzual2007}. As expected, the mass of stellar populaton $M_*[\mathrm{CB07}]$ is lower than $M_*[\mathrm{BC03}]$ by 25\% in the median which is smaller than the range of 50\%--80\% reported by \citet{Bruzual2007}. On the other hand, the $\mathrm{age}_M[\mathrm{CB07}]/\mathrm{age}_M[\mathrm{BC03}]$ ratio is close to unity with a median value of 0.94 for clean sources. These differences between our result and \citet{Bruzual2007} might be due to two reason: (1) our fitting applied composite stellar population (CSP) models which contain stars with different ages given by SFH introduced above, but not the SSP models used by \citet{Bruzual2007}; (2) the aforementioned mass-weighted age is larger than 2\,Gyr in the median when the differences between BC03 and CB07 models have gone through their maximum and become smaller. In definition, the mass-weighted age reflects the time when most of the stellar mass formed in one CSP model. Thus, this parameter of star-forming galaxies are highly depend on the assumed SFH employed in the SED fitting (\citealt{Conroy2013b}). However, the exponentially declining form of the SFH is unchanged when applied CB07 models and the resulting $\gamma_{\mathrm{CB07}}/\gamma_{\mathrm{BC03}}$ ration has a median of 1.0. Furthermore, for a given SFH form, \citet{Wuyts2011a} found that to strictly constrain $\mathrm{age}_M$ observations at the wavelength of $\lambda_{\mathrm{rest}}<0.2\,\mathrm{\mu m}$, where TP-AGB stars have almost no contribution (\citealt{Bruzual2007}), is required. Therefore, the mass-weighted ages of our clean sample show statistically little change when the CB07 models were used. In the case of stellar mass, the relative old stellar population contribute quite a bit mass of our clean sources, suggesting by their large median $\mathrm{age}_M$, and only a little in the difference between $M_*[\mathrm{CB07}]$ and $M_*[\mathrm{BC03}]$. In combination with younger population, especially those formed in recent 0.1--1\,Gyr \citep{Bruzual2007}, the median $M_*[\mathrm{CB07}]/M_*[\mathrm{BC03}]$ ratio becomes larger than the range given by \citet{Bruzual2007}. In this work, to investigate the accuracy of monochromatic extrapolated IR luminosity, we utilize multi-wavelength data of GOODS-North and GOODS-South fields to perform a UV-to-submillimeter SED fitting and use the output dust luminosity as a reference value. The main conclusions are as follows: \begin{enumerate} \renewcommand{\labelenumi}{(\theenumi)} \item We compare extrapolated $L_{\mathrm{IR}}^{24}$ and $L_{\mathrm{IR}}^{\mathrm{PACS}}$ based on different templates and conclude that CE01 and DH02 templates present nearly the same estimate in these two bands, while W08 template show large differences and should be used with caution for galaxies at $z\lesssim 0.5$. \item For CE01 template, extrapolations from PACS bands can estimate the actual IR luminosity out to $z\sim 3.5$ well. A few high-redshift galaxies hint that this consistency may be hold even at $z\sim 5$. However, extrapolations from MIPS 24\,$\mathrm{\mu m}$ for galaxies with $z>1.5$ result in a serious overestimate of the total IR luminosities. $L_{\mathrm{IR}}^{\lambda}$ from SPIRE bands also could not provide reliable estimate of $L_{\mathrm{dust}}$. For DH02 template, the $L_{\mathrm{IR}}^{\lambda}/L_{\mathrm{dust}}$ ratios from 24--160\,$\mathrm{\mu m}$ bands are similar to that of CE01 template, but show a little redshift evolution in 70\,\um\ band. For clean source at $z\lesssim 2$, extrapolations from SPIRE bands present a rough unbiased esitmate of $L_{\mathrm{dust}}$. In the case of W08 template, the absence of a significant ``mid-IR excess'' problem makes it useful to derive $L_{\mathrm{IR}}$ in MIPS 24\,\um\ band, while extrapolations from submillimeter bands also can provide unbiased estimate for galaxies at $z<2$ as DH02. \item Among the three templates we concerned, the CE01 template provides the best estimate of $L_{\mathrm{dust}}$ in PACS bands, while the DH02 and W08 templates perform better in SPIRE bands although the dispersion is still large. For extrapolations from MIPS 24\,\um\ band, only W08 template does not suffer a significant mid-IR excess problem. \item To obtain a reliable estimate of the actual IR luminosity using the monochromatic extrapolation method described in this work, our suggestions are as follow. To extrapolate from MIR bands (e.g., MIPS 24\,\um), CE01 template should be used for galaxies at $z<1.5$, and W08 template should be used for galaxies at $z>1.5$ although it will lead to an overestimate of nearly 90\%. Using FIR observations (e.g., PACS 70--160\,\um) to do this, CE01 template is the best choice out to $z\sim3.5$. Besides, if only submillimeter bands (e.g., SPIRE 250--500\,\um) are available, both DH02 and W08 template can be used for galaxies at $z\lesssim 2$, but the large uncertainty also should be kept in mind. Moreover, $L_{\mathrm{IR}}^{\mathrm{PACS}}$, which is derived from the available longest wavelength PACS band, based on CE01 template can be a good estimator. \item All the three templates exhibit different degrees of enhanced emission at $\lambda_{\mathrm{rest}}\lesssim 10\,\mathrm{\mu m}$, but well describe the emission of 10--100\,\um\ range of the IR SED. Only DH02 template show a nearly unbiased estimate of the emission of the rest-frame submillimeter part. \end{enumerate} | 16 | 9 | 1609.08836 |
1609 | 1609.03586_arXiv.txt | We present multi-telescope, ground-based, multiwavelength optical and near-infrared photometry of the variable L3.5 ultra-cool dwarf 2MASSW J0036159+182110. We present \INSERTNIGHTS \ nights of photometry of 2MASSW J0036159+182110, including \INSERTNIGHTSSIMULTANEOUS \ nights of simultaneous, multiwavelength photometry, spread over $\sim$\INSERTSPREADDAYS \ days allowing us to determine the rotation period of this ultra-cool dwarf to be \PeriodHoursTwoMassZeroZeroThirtySixAllbands \ $\pm$ \PeriodHoursErrorTwoMassZeroZeroThirtySixAllbands \ hr. Our many nights of multiwavelength photometry allow us to observe the evolution, or more specifically the lack thereof, of the light curve over a great many rotation periods. The lack of discernible phase shifts in our multiwavelength photometry, and that the amplitude of variability generally decreases as one moves to longer wavelengths for 2MASSW J0036159+182110, is generally consistent with starspots driving the variability on this ultra-cool dwarf, with starspots that are $\sim$100 degrees $K$ hotter or cooler than the $\sim$1700 $K$ photosphere. Also, reasonably thick clouds are required to fit the spectra of 2MASSW J0036159+182110, suggesting there likely exists some complex interplay between the starspots driving the variability of this ultra-cool dwarf and the clouds that appear to envelope this ultra-cool dwarf. | Detecting and characterizing the variability of ultra-cool dwarfs is an area that has attracted growing attention in recent years, with a wealth of variability detections from the late-M, L \& T spectral classes (e.g. \citealt{Gelino02,Irwin11,Radigan14,Buenzli14,Metchev15}). For L-dwarfs, ground and space-based optical, near-infrared and infrared observations have recently shown that low level variability of L-dwarfs is common \citep{BailerJonesMundt01,Gelino02,Lane07,Harding13,Koen13,Metchev15}; intriguingly, once viewing geometry is taken into account, all L-dwarfs might be variable \citep{Metchev15}. The questions this prompts are: what is the astrophysical cause of the observed variability, is it consistent across the L-spectral class, and if not where does the transition region lie between various astrophysical causes of the observed variability. On the stellar side of the hydrogen-fusing limit, starspots are the usual explanation for the observed variability. M-dwarfs are notoriously active, with detections of H$\alpha$, a common marker of activity, rising throughout the M-spectral class, including that nearly all very late M-dwarfs are active \citep{West04,Schmidt15}. % Rotation periods revealed by photometry \citep{Rockenfeller06,Irwin11}, and Doppler imaging techniques for M-dwarfs \citep{Barnes01,Barnes04}, have indicated that starspots are ubiquitous on M-dwarfs. Until recently there was reason to doubt that this starspot driven variability extended into the L-spectral class. The neutral atmospheres of L dwarfs were believed to be too electrically resistive, with too small of magnetic Reynolds numbers, for magnetic starspots to form \citep{Mohanty02,Gelino02}. This conclusion was previously supported by studies that found the frequency of H$\alpha$ detections fell sharply at the M/L transition, reaching negligible levels by the $\sim$L3 spectral class \citep{West04}. This did not stop speculation that the variability displayed by L-dwarfs might arise from magnetic starspots \citep{Clarke02,Lane07}. Such speculation might prove to be prescient, as a recent study has cast doubts on previous H$\alpha$ null-detections for L-dwarfs and therefore indicated that starspots might be present throughout the L spectral class: \citet{Schmidt15} analyzed higher signal-to-noise L-dwarf spectra and were able to detect H$\alpha$ for approximately 90\% of L0 dwarfs, and more than half of L-dwarfs as late as L5. Therefore, starspots might be driving the variability for cloudy early and even-late L dwarfs. Once ultra-cool dwarf effective temperatures drop and silicate clouds begin to clear at the L/T transition, cloud condensate variability has become the accepted explanation for the large amplitude variability that has been observed for these objects \citep{Artigau09,Radigan12,Gillon13,Radigan14,Buenzli14,Crossfield14Nature}. Multiwavelength photometry of these apparent cloudy ultra-cool dwarfs has returned light curves with different amplitudes at different wavelengths \citep{Radigan12}, and with significant temporal phase shifts in the observed variability at different wavelengths \citep{Buenzli12,Yang16}, including multiwavelength light curves that are anti-correlated in phase \citep{Biller13}. However, multiwavelength light curves of some of these brown dwarfs that display variability that is believed to be driven by heterogeneous cloud cover have not displayed significant phase offsets \citep{Apai13,Buenzli15}. The likely explanation is that multiwavelength phase shifts will not be caused by cloud variability if the clouds span multiple pressure layers, and specifically, if the clouds span the range of pressures that are probed by the various wavelengths of observation. Intriguingly, cloud condensate variability might not be constrained to the L/T transition. Spitzer/IRAC observations have indicated that nearly all L-dwarfs are likely variable \citep{Metchev15}, and anti-correlated light curves have been observed on a late-M dwarf \citep{Littlefair08} -- in both cases clouds are a likely explanation for the observed variability. Another possibility that has recently arose, is that the observed variability of ultra-cool dwarfs results from auroral activity, similar to the aurorae observed on planets in our own solar system (e.g. \citealt{Clarke80}) including the Earth. Such auroral activity has been observed at the end of the main sequence on an M8.5 dwarf \citep{Hallinan15}. Multiwavelength photometry of this dwarf displays light curves that are anti-correlated in phase, and \citet{Hallinan15} speculate that auroral activity might explain the anti-correlated light curves of another late M-dwarf \citep{Littlefair08}. Similar auroral activity may be indicated from radio detections of polarized, pulsed emissions from a T2.5 dwarf \citep{Kao16} and a T6.5 dwarf \citep{RouteWolszczan12,WilliamsBerger15}. Therefore, another possibility is that auroral activity may be responsible for some of the variability at the L/T transition \citep{Artigau09,Radigan12} that has previously been believed to be due to holes in condensate clouds. Finally, another explanation for the observed variability is atmospheric temperature variations, arising either deep in the atmosphere of these ultra-cool dwarfs, or at other pressure layers, that are communicated via radiative heating to the altitude regions that are probed by optical to near- and mid-infrared observations \citep{ShowmanKaspi13,RobinsonMarley14,Morley14}. It is also possible that ultra-cool dwarfs are variable due to more than one of the aforementioned astrophysical reasons. Time evolving clouds may periodically obscure magnetically driven cool or hot starspots \citep{Lane07,Heinze13,Metchev15}, or aurorae may play an occasional role on predominantly cloudy brown dwarfs \citep{Hallinan15}. Many variability studies of ultra-cool dwarfs to date have examined targets for only tens of minutes \citep{Buenzli14} to hours at a time \citep{Radigan14}. This method is likely sufficient for ultra-cool dwarfs that display constant variability \citep{Gizis13,Gizis15}, but will inadequately address the variability of objects with light curves that evolve rapidly, such as have been observed for several ultra-cool dwarfs \citep{Artigau09,Gillon13,Metchev15}. The rapid evolution of these brown dwarfs might be due to the properties of a single variability mechanism changing --- such as the size of starspots growing, or the thickness of clouds decreasing --- but could also be due to the growing or waning strength of a secondary variability mechanism compared to the primary mechanism. It is therefore imperative to monitor ultra-cool dwarfs for multiple rotation cycles spread over days, weeks, months and even years. Precise long-term monitoring of ultra-cool dwarfs is also a perfect data-set to search for transiting exoplanets that are as small or smaller than Earth-sized planets in the ``habitable zones'' of these dwarfs. Ultra-cool dwarfs typically have stellar radii similar to that of Jupiter \citep{Charbrier00}, meaning that even Earth-sized planets produce $\sim$1\% transit depths. For late M to early T dwarf spectral types the habitable zones stretch from periods of a day, up to a few days (depending on the influence of various atmospheric compositions and albedos, and the exact effects of tidal heating; \citealt{Bolmont11}; \citealt{BarnesHeller13}; \citealt{Zsom13}). The number of small planets appears to increase with decreasing stellar effective temperature \citep{Howard12,DressingCharbonneau15} -- trends that may continue into the substellar regime. There are already examples of M-dwarfs orbited by multiple rocky planets in short orbits: there are three rocky planets with periods less than two days orbiting an M4 dwarf \citep{Muirhead12}, and two rocky planets have been found to orbit an M8 dwarf with an additional rocky planet in a longer period orbit \citep{Gillon16}. Perhaps most importantly, a habitable rocky planet around an ultra-cool dwarf would be more favourable for follow-up than hosts with earlier spectral types; atmospheric spectral features of a planet around a brown dwarf should prove detectable with a feasible two weeks of total {\it James Webb Space Telescope} observing time \citep{Belu13}. Here we attempt to characterize and determine the astrophysical cause of the variability for 2MASSW J0036159+182110; in Section \ref{SecIntroTwoMassZeroZeroThreeSix} we provide an overview of this intriguing L3.5 ultra-cool dwarf. To determine the physical mechanism causing the variability of 2MASSW J0036159+182110 in Section \ref{SecObs} we present \INSERTNIGHTS \ nights of photometry of 2MASSW J0036159+182110 spread out over $\sim$\INSERTSPREADDAYS \ days, from optical to near-infrared wavelengths (the Ks, H, J, I, z' and R-bands), including \INSERTNIGHTSSIMULTANEOUS \ nights of simultaneous multiwavelength photometry in the optical and the near-infrared. In Section \ref{SecAnalysis} we analyze our multiwavelength photometry and determine that the rotation period of this ultra-cool dwarf is approximately \PeriodHoursTwoMassZeroZeroThirtySixAllbands \ $\pm$ \PeriodHoursErrorTwoMassZeroZeroThirtySixAllbands \ hr, and this is consistent from wavelength to wavelength for our optical and near-infrared photometry and the previously detected radio period, allowing us to place a limit on differential rotation for 2MASSW J0036159+182110. We also demonstrate a lack of significant phase offsets in our simultaneous multiwavelength variability, and therefore starspots rotating in and out of view, or gaps in clouds that span multiple pressure layers probed by our observations, are the most likely explanations for the observed variability on 2MASSW J0036159+182110. Our photometry does not reveal any significant flares, such as have been observed at radio wavelengths, and we are therefore able to place a limit on the frequency of flaring for this ultra-cool dwarf (Section \ref{SecFlares}). We also rule out the presence of transiting super-Earth sized planets in the habitable zone of this ultra-cool dwarf (Section \ref{SecPlanets}). In Section \ref{SecDiscussion} we demonstrate that the decreasing amplitudes of variability that we observe with increasing wavelength in our photometry are most consistent with starspots or gaps in clouds that are that are approximately $\sim$100 $K$ cooler or hotter than the photosphere of this ultra-cool dwarf. In Section \ref{SecConclusions} we conclude that the variability we observe from 2MASS 0036+18 is likely driven by starspots, but there probably exists some complex interplay between starspots and the clouds that appear to envelope this ultra-cool dwarf. | \label{SecConclusions} Our long-term, multiwavelength, ground-based light curves of the L3.5 dwarf 2MASS 0036+18 allow us to address a number of unique science cases. First, our many nights of photometry allow us to demonstrate that the light curves of 2MASS 0036+18 do not significantly evolve from rotation period to rotation period, or night-to-night. This lack of evolution allows us to precisely determine the rotation period of this L3.5 dwarf to be \PeriodHoursTwoMassZeroZeroThirtySixAllbands \ $\pm$ \PeriodHoursErrorTwoMassZeroZeroThirtySixAllbands \ hours. This rotation period is recovered from the R-band to the J-band, and is similar to the rotation period that has previously been recovered from radio \citep{Hallinan08} and infrared \citep{Metchev15} observations for this ultra-cool dwarf; therefore there is no strong evidence for differential rotation with latitude, or at different pressure layers as probed by shorter wavelength observations peering deeper into the atmosphere of this L3.5 dwarf. These light curves also allow us to constrain the rate of flares exhibited by 2MASS 0036+18; this ultra-cool dwarf must display significantly fewer optical/near-infrared flares than radio flares \citep{Berger02,Berger05}. Also, we are able to rule out transiting super-Earth, and even some Earth-sized planets in the habitable zone of this ultra-cool dwarf. Our \INSERTNIGHTSSIMULTANEOUS \ nights of simultaneous, multiwavelength photometry do not display discernible phase shifts, and therefore suggest that the variability of 2MASS 0036+18 is driven by starspots, or another mechanism (clouds, aurorae, etc.) that results in similarly modest phase shifts. The amplitude of variability generally decreases with increasing wavelength, a result consistent with starspots slightly warmer or cooler than 2MASS 0036+18 being responsible for the observed variability, and with the H$\alpha$ detection \citep{Pineda16} from this L3.5 dwarf. Our fit to the spectra of 2MASS 0036+18 suggests that considerable clouds envelope this ultra-cool dwarf. Clouds, or another mechanism, resulting in temperature differences could also be responsible for driving the observed variability of this ultra-cool dwarf, but this other mechanism would have to result in phase shifts as small or smaller than displayed in our multiwavelength photometry; therefore, if clouds are causing the variability of 2MASS 0036+18 then the gaps in clouds would likely have to span the multiple pressure layers probed by our observations and expose hotter layers of the atmosphere beneath the cloud layer. If a single mechanism is causing the variability of 2MASS 0036+18 then starspots or clouds that span multiple pressure layers are the leading explanations; however, other mechanism(s) that result in temperature differences, but a lack of significant multiwavelength phase shifts are also viable. The lack of significant evolution of the 2MASS 0036+18 light curve from rotation period to rotation period, or night-to-night, is in stark contrast to what has been observed for L/T transition brown dwarfs (e.g. \citealt{Artigau09, Radigan12, Gillon13,CrollUCDII}) and their variability that has generally been attributed to clouds. It is not clear if our detection of some, but not significant, evolution of the light curve of 2MASS 0036+18 favours an interpretation of clouds or spots driving the variability of this ultra-cool dwarf. Starspots on an active M-dwarf have been inferred to be quasi-stable for years (e.g. \citealt{Davenport15}). Naively, one might expect cloud structures to evolve rapidly, but multiwavelength photometric monitoring of an L1 dwarf \citep{Gizis13,Gizis15} have indicated the possibility of a cloud feature that is stable on that ultra-cool dwarf for up to two years; in addition, Jupiter's Great Red Spot has likely been present since, at least, soon after the invention of the telescope \citep{Hook1665,Cassini1666,Marcus93}, suggesting that one cannot rule out a priori the presence of long-lived cloud structures on 2MASS 0036+18. The most likely explanation for 2MASS 0036+18's variability is arguably a mixture of mechanisms. As our fit to the spectra of 2MASS 0036+18 suggests that significant clouds envelope this L3.5 dwarf, a mixture of clouds and starspots seems likely. The variability of 2MASS 0036+18 is likely driven predominantly by starspots, with some complex interplay between the hot/cool starspots and the clouds on this L3.5 dwarf that envelope this ultra-cool dwarf. We encourage further long-term, multiwavelength monitoring of this intriguing L-dwarf, as well as other L and ultra-cool dwarfs, to determine the long-term variability amplitudes, phases, and evolution, or lack thereof, of these objects. With the recent detection that nearly all L-dwarfs are variable \citep{Metchev15}, and that H$\alpha$ detections are common for many early L-dwarfs \citep{Schmidt15}, such monitoring may indicate whether the observed characteristics of 2MASS 0036+18 are common for other early L-dwarfs. | 16 | 9 | 1609.03586 |
1609 | 1609.03065_arXiv.txt | One of the outstanding challenges of cross-identification is multiplicity: detections in crowded regions of the sky are often linked to more than one candidate associations of similar likelihoods. We map the resulting maximum likelihood partitioning to the fundamental assignment problem of discrete mathematics and efficiently solve the two-way catalog-level matching in the realm of combinatorial optimization using the so-called Hungarian algorithm. We introduce the method, demonstrate its performance in a mock universe where the true associations are known, and discuss the applicability of the new procedure to large surveys. | Modern astronomy surveys tenaciously observe the sky every night and their automated software pipelines produce a huge number of exposures. The datasets from separate telescopes and instruments are then routinely combined to enable all multicolor and time-domain studies. Over the last decade there has been a tremendous progress in the field of catalog matching to provide reliable resources for astronomical and cosmological measurements. \citet{BS08-BayesCrossID} introduced a framework based on Bayesian hypothesis testing that properly incorporated astrometric uncertainties, which yielded superior results \citep{heinis}, and proved to be flexible to accommodate a variety of scenarios from associating stars with unknown proper motions \citep{kerekes} or even models for radio morphology \citep{fan}. These methods are also paired with efficient indexing ideas and search algorithms to make the process of matching fast on the largest catalogs, e.g., \citep{htm, healpix, igloo, zones, gpumatch} The methods before, however, were only concerned with assigning a reliable quality measure of the association for a given set of sources. The likelihood of the association would be compared to the likelihood of the detections belonging to separate objects using the Bayes Factor and the previously used approaches do not resolve situations where a single detection appears in multiple associations with similar likelihoods. Previous recommendations to deal with such situations resorted to including more data, e.g., photometric measurements to refine the likelihood measurements \citep{marquez}. In other words all candidate associations so far have been considered in isolation ignoring the fact that a star, for example, could only appear in a single match. Such exclusion rules can potentially affect a large fraction of the catalog especially in crowded fields. In this paper we study catalog-level cross-identification to provide strategies for finding the most likely matched catalogs in which every association is valid and no detections appear in multiple matches. In \csect{sec:catxid} we formulate the catalog-level matching problem and map the likelihood optimization onto discrete minimization that is solved efficiently by the so-called Hungarian algorithm. \csect{sec:app} details our implementation and its application to realistic mock catalogs. In \csect{sec:disc} we discuss the results and \csect{sec:sum} concludes the paper. \begin{figure*} \epsscale{1.2} \plotone{f1.pdf} \caption{\emph{Left:} A mock universe is constructed by randomly assigning coordinates (circles) and internal properties (color) to objects. \emph{Right:} Simulated surveys observe different subsets of the objects with overlap. The $+$ signs show the true directions of the objects but the \emph{red} and \emph{blue} dots illustrate the noisy measurements of the simulated sources. Objects that meet both selection functions are noted with small black dots on the true positions. } \label{fig:mock} \end{figure*} | \label{sec:disc} The results that correspond to the maximum likelihood solution of the entire catalog appear exceptionally reliable and go beyond the limitations of working with individual associations. With the marginal likelihoods for the winning associations we can go back and derive the probability of each match using the prior information on surface density of the sources in the catalogs \citet{BS08-BayesCrossID}. \subsection{Computational Cost} The proposed procedure is clearly more expensive than previous methods. % As mentioned earlier, the general assignment problem solved by the Hungarian algorithm requires $\O{(n_1\!+\!n_2)^3}$ % time, where $n_1$ and $n_2$ are the number of rows and columns of the input matrix $W$ (in our case \mbox{$|D|\!=\!n_1\!+\!n_2$}). For example, the 3600$\times$3600 weight matrices of our simulations were processed by the solver called \texttt{linear\_sum\_assignment()} in the \texttt{optimize} module of {SciPy} \citep{scipy} in approximately 2.2 seconds on a desktop computer. While this sounds prohibitively slow at first, when considering large astronomy surveys of hundreds of millions of sources, in practice the wall-clock time will be actually much shorter due to the sparsity of the weight matrix. In other words the list of candidate associations is actually only on the order of $\max(n_1,n_2)$ long instead of the number of pairs, which naively would be \O{n_1n_2}. In fact resolving the multiplicity issue can be reserved for the overlapping associations where friends-of-friends connected components emerge as problem cases. Luckily these can be dealt with separately, hence reducing the cardinality of the problem to much smaller numbers of sources. \citet{hsc} used a greedy approach to find the best the partitions using clustering algorithms and Bayesian hypothesis testing. In the 2-way matching case, we can actually afford to do a global search for the optimum with existing methods such as the Hungarian algorithm. \subsection{Likely Matched Catalogs} Finding the optimal matched catalog is unfortunately not the final answer. There can be other realizations with the same or similarly good likelihoods. In principle these alternatives should be reported for further statistical analysis and marginalized over in the study. While the approach applied in our study does not automatically provide the sub-optimal solutions, the optimum would a good starting point for some Markov chain Monte Carlo algorithm to explore alternative matched catalogs. The cost of storing or automatically generating these catalogs and providing the uncertainty estimates on the associations or their statistics needs to be studied further but is straightforward from the methodology's point of view. | 16 | 9 | 1609.03065 |
1609 | 1609.01851_arXiv.txt | We conducted an analysis of the distribution of elements from lithium to europium in 200 dwarfs in the solar neighbourhood ($\sim$ 20 pc) with temperatures in the range 4800--6200 K and metallicities [Fe/H] higher than --0.5 dex. Determinations of atmospheric parameters and the chemical composition of the dwarfs were taken from our previous studies. We found that the lithium abundances in the planet-hosting solar-analog stars of our sample were lower than those in the stars without planetary systems. Our results reveal no significant differences exceeding the determination errors for the abundances of investigated elements, except for aluminium and barium, which are more and less abundant in the planet-hosting stars, respectively. We did not find confident dependences of the lithium, aluminium and barium abundances on the ages of our target stars (which is probable because of the small number of stars). Furthermore, we found no correlation between the abundance differences in [El/Fe] and the condensation temperature ($T_{cond}$) for stars in the 16 Cyg binary system, unlike the case for 51 Peg (HD 217014), for which a slight excess of volatile elements and a deficit of refractories were obtained relative to those of solar twins. We found that one of the components of 16 Cyg exhibits a slightly higher average abundance than its counterpart ($<$[El/H](A--B)$>$ = 0.08 $\pm$0.02 dex); however, no significant abundance trend versus $T_{cond}$ was observed. Owing to the relatively large errors, we cannot provide further constraints for this system. | A large number of studies of planet-hosting stars \citep[][etc]{gonzalez:97, gonzalez:98, santos:00, fischer:05, udry:07, adibekyan:12a, adibekyan:12b} have been carried out over recent decades with the aim of increasing our understanding of the processes involved in the formation of planets. The earliest studies of the metallicity of planet-hosting stars \cite[e.g.][]{gonzalez:97}, as well as subsequent ones \citep[see][and reference therein]{udry:07} indicated that most of these stars are rich in metals. There are two possible explanations for the excess metallicity observed. The first is the infall of metal-rich (planetary) material into the stellar envelope \citep[e.g.][]{gonzalez:98} and the second is related to the prestellar enrichment of interstellar matter \citep[][and references therein]{fischer:05}. The issue has not yet been finally resolved; however, the second assumption is to be preferred, as the infall of matter cannot provide an appreciable increase in the metal abundance. The high metallicity of planet-hosting stars is well-established for stars with massive planets (Jupiter-like ones) \citep[][]{gonzalez:97, santos:00, santos:01, fischer:05, sousa:08}, while it is not for stars with less massive planets, namely those with planets the size of Neptune or Earth \citep[e.g.][]{udry:07, sousa:08, sousa:11, wang:15}. The results of \cite{buchhave:12} suggest that terrestrial planets have no special requirement for enhanced metallicity for their formation, and they support the hypothesis that stars hosting terrestrial planets have a metallicity similar to stars with no known transiting planets \citep{buchhave:15}. \cite{sousa:08} and \citep{adibekyan:12a, adibekyan:12b} drew attention to the possibility of planet formation for metallicities lower than solar. It was shown that there is an excess of $\alpha$-elements, especially magnesium, in the stars with low-mass planets, which is more pronounced in the thick-disc population than in the thin disc for metallicities below --0.3 dex \citep{adibekyan:12a, adibekyan:12b}. This implies that metals other than iron may noticeably contribute to planetary formation if the iron abundance is low. Lithium plays a unique role in the study of planet-hosting stars. Its abundance is apparently lower in the stars with planetary systems than in those without \citep[][]{gonzalez:00, gonzalez:08, gonzalez:10, israelian:04, israelian:09, delgado:14, figueira:14, delgado:15}. In particular, \cite{delgado:14} found that there is some evidence that lithium depletion in planet-hosting solar-type stars is higher when their planets are more massive than Jupiter. Hot stars that host Jupiter-like planets and have effective temperature \Teff\ in the range of 5900-6300 K show lithium abundances that are 0.14 dex lower than those in stars without detected planets \citep{delgado:15}. It should be noted, however, that lithium depletion is usually associated with stellar evolution \citep[][]{deniss:10, talon:05, anrassy:15}, a phenomenon that is confirmed by the dependence of lithium abundances on age \citep[e.g.][]{monroe:13, melendez:14, carlos:16}. \cite{carlos:16} found that the lithium abundances of solar twin stars are a function of stellar age, while there is no indication of any relationship between planet-hosting stars and enhanced lithium depletion. In order to identify either the presence or the absence of a possible relation between chemical abundances and mechanisms of planetary formation, various studies have been performed to examine the chemical peculiarities of planet-hosting stars \citep[e.g.][]{melendez:09, ramirez:09, adibekyan:15b}, the main properties of planets and their hosts (e.g. mass) \citep{kang:11, dorn:15, sousa:15}, their position in the Galaxy \citep{haywood:08, haywood:09, adibekyan:14}, and Galactic evolution \citep[e.g.][]{adibekyan:15b}. Studies of the chemical composition of the Sun, solar twins and solar analogues that have highly accurate abundance determinations ($\sim$ 0.01 dex), \cite{melendez:09} shown a decrease in the relative content of refractory elements in the Sun, namely one that is by 20$\%$ lower than those of solar-analogue stars and solar twins with giant planets. Because a correlation between the abundance differences and condensation temperature $T_{cond}$ was found, these authors speculated that the decrease in the relative content was associated with the presence of Earth-like planets. In the literature, there are several alternative explanations for the abundance trends with $T_{cond}$. \cite{adibekyan:14} found that chemical peculiarities (i.e. small refractory-to-volatile ratio) of planet-hosting stars are likely to reflect their older age and inner Galactic origin; hence, stellar age and, probably, Galactic origin are the key factors to establish the abundances of some specific elements. It has also been suggested that the $T_{cond}$ trend correlates strongly with stellar radius and mass \citep{mald:15, mald:16}. The trend may also depend on stellar environment \citep{onehag:14}, and internal processes, such as gas-dust segregation in the protostellar disc \citep{gaidos:15}. Recently, \cite{adibekyan:16} found that the $T_{cond}$ trend may depend strongly on the spectra of the stars used. In particular, these authors observed significant differences in the abundances of the same star as derived from different high-quality spectra. Previously, \cite{sousa:08} had suggested that the detectability of Neptune-class planets may increase in stars with a low metallicity. \cite{kang:11} also confirmed the presence of chemical abundance differences between stars with and without exoplanets, as well as the relationship between chemical abundances and planetary mass. \cite{sousa:15} studied the effect of stellar mass on the derived planetary mass and noted that the stellar mass estimates for giant stars should be employed with extreme caution when computing planetary masses. Binary systems of stellar twins are important objects in this investigation, as the effects of stellar age (chemical evolution) or birthplace in the Galaxy are similar for both components of the binary pair. \cite{nissen:15} showed that there are clear correlations between [El/Fe] ratios and stellar ages in solar twins. \cite{teske:16} found that both components in the binary system WASP 94 A and B are planet-hosting stars, and they differ in their chemical composition. The binary system HD 80606/HD80607 exhibits no remarkable differences in the abundances of its components, despite the fact that the star HD 80606 hosts a giant planet \citep{saffe:15, mack:16}. However, abundance variations were found in the XO-2 planet-hosting binary from independent observations \citep{biazzo:15, ramirez:15, teske:15}. As reported by \cite{spina:16} the abundance ratios [El/Fe] show signatures of both chemical evolution and planets. Our group has investigated nearly 600 stars, dwarfs and giants over several years. We determined their atmospheric parameters and chemical composition in order to study the chemical and dynamical evolution of the Galaxy and, primarily, the chemical enrichment of various Galactic substructures \citep[][]{mishenina:04, mishenina:06, mishenina:08, mishenina:12, mishenina:13, mishenina:15a, mishenina:15b}. We detected 14 planet-hosting stars among the target ones. It would be interesting to examine the behaviour of the abundances of various elements, including lithium, refractory and volatile elements, as well as neutron-capture elements for the stars with and without planets that are present in our data base. Such a study could enable an independent analysis of the correlation between the presence of planets and the chemical composition of stars, possibly shedding some light on the existing contradictions, as well as re-examining earlier determined constraints on the mechanisms of formation of planets and planetary systems. The paper is set out as follows. The observations and selection of stars, as well as determination of stellar parameters and elemental abundance are described in \S \ref{sec: stellar param}. The analysis of element abundances is presented in \S \ref{sec: el_analys}. In \S \ref{sec: Tcond_abun}, the correlations between elemental abundances and condensation temperatures are discussed. The connection with chemical enrichment of galactic disc is reported in \S \ref{sec: connect_struct}. \ref{sec: conclus} summarizes and concludes the paper. | \label{sec: conclus} Using homogeneous spectral data and techniques for determination of parameters and abundances of a number of elements, we compared results obtained for stars with and without planets. We have examined a total of about 200 stars, including 14 planet-hosting stars. Our main findings are as follows. \\ -- The lithium abundances in planet-hosting solar-analog stars (i.e. stars with parameters very similar to the solar) in our small sample were lower than those in the stars for which no planetary systems had been discovered. The lithium abundance does not exceed 1.7 according to the scale where hydrogen is 12.0, except one star, namely HD 9826. \\ -- For the binary star 16 Cyg, the star with a planet (16 Cyg B) has a lower lithium abundance than its companion without a detected planets.\\ -- No significant differences, exceeding determination errors, for the abundances of other elements were found between stars with and without planets, with possible exception for aluminum and barium abundances. \\ -- No statistically significant dependences of abundance differences in [El/Fe] and [El/H] with the condensation temperature $T_{cond}$ for the two stars in the 16 Cyg binary system were found; it was not possible to determine the slope in 16 Cyg because of the large error bars.\\ -- The abundance difference $<$[El/H](A--B)$>$ = 0.08 $\pm$0.02 for the two stars shows a slight surplus, the origin of which is not clear.\\ -- A slight excess of volatile elements and a deficit of refractories as compared with solar twins were obtained for 51 Peg (HD 217014).\\ The chemical composition of stars as well as atmospheric parameters are important diagnostic tools in studies of stellar evolution, nucleosynthesis computations, and studies of various physical processes both inside stars and on their surfaces, etc. The application of chemical composition analysis always requires careful consideration of the origin of the chemical composition and the mechanisms of its change. As noted above, much attention has been paid recently to the investigation of any correlation between the abundance of one or another element and the presence of planetary systems around stars. {\it Metallicity.} It is now clear that massive planets are observed preferentially around stars with solar or higher metallicities \citep[][]{fischer:05, sousa:08}, while Earth-like planets can be hosted by stars with different metallicities \citep{sousa:08, adibekyan:12a, adibekyan:12b, buchhave:12}. Our current sample contains several stars hosting Jupiter-like planets. Among them there is the only star (HD 154345) that hosts a single planet with mass M = 1M$_J$. The most metal-poor planet-hosting star in our sample has the metallicity of $\sim$ --0.3 dex and host a low-mass planet with a mass of $\sim$ 6.5 $M_{\oplus}$ (HD 97658). The requirement of high metallicity may be a significant factor for the presence of massive planets, but this is not the case for less massive planets. {\it Lithium.} This element is exposed to the effects of a number of factors to the larger extent than iron, which results in its abundance variations. In general, the lithium abundance is found to be lower in planet-hosting stars \citep[][etc.]{gonzalez:00, gonzalez:08, gonzalez:10, israelian:04, israelian:09, delgado:14, figueira:14, delgado:15}. The depletion of lithium owing to convection is observed in the stars with low temperatures (\Teff\ is lower that 5600 K) \citep[e.g.][]{randich:10}. With temperatures close to solar and higher, the change in the lithium abundance is caused by the chromospheric activity \citep[e.g.][]{tati:07} and (or) stellar rotation \cite{charb:92}. It should also be note that many authors \cite[e.g.][etc]{baumann:10, carlos:16} have found strong evidence for increasing lithium depletion with age. The low lithium abundance determinations obtained by us for stars with massive planetary systems sustain the hypothesis regarding lithium depletion during the formation of stars, but do not preclude the possibility of specific enrichment of the pre-stellar cloud from which the star was formed. The number of stars in our sample is not high enough to assert the presence of dependence of the lithium abundance on the age of stars. {\it Volatile and refractory elements.} The abundance determinations for elements with various condensation temperatures $T_{cond}$ in parent stars, which could be depleted during the formation of planets, means that we can suggest the possibility of the formation of rocky Earth-like planets or of rocky cores of massive planets \citep[e.g.][etc]{gonzalez:97, smith:01,melendez:09}. Our data showed the presence of a trend (a slope of the plotted dependence) for the relative elemental abundances (star -- mean elementary abundances in solar twins) with $T_{cond}$ only for one star among four. Thus, for the star HD 217014 (51 Peg) we found a slight excess of volatile elements and a deficit of refractories as compared with those obtained for the solar twins. It should be noted that 51 Peg has an age that is similar to the solar value \citep[e.g.][]{israelian:04}. The absence of such correlation for other three stars just sustains the assumption that the stellar chemical composition can be reckoned as better reflection of the chemical and dynamic evolution of the Galaxy, in particular the enrichment of the start's place of origin by supernovae of different types at the given time intervals \citep{adibekyan:14, nissen:15}. {\it Other elements}. \cite{kang:11} claimed different manganese abundance between the star with and without planets. Our studies have not confirmed any difference in the Mn abundances for our sample stars with and without planets \citep{kang:11}; however, we found that aluminium is more abundant while barium is less abundant in the stars with planetary systems. This fact requires further verification and confirmation. Using our data base, we can corroborate independently that the chemical composition is undoubtedly an important, albeit an auxiliary, factor to be taken into account when studying the presence and formation of planetary systems. On the one hand statistical approaches (based on relative studies) for large samples of stars are required; on the other hand it is necessary to investigate each single star and each chemical element individually. | 16 | 9 | 1609.01851 |
1609 | 1609.06195_arXiv.txt | A programme of worldwide, multi-wavelength electromagnetic follow-up of sources detected by gravitational wave detectors is in place. Following the discovery of GW150914 and GW151226, wide field imaging of their sky localisations identified a number of candidate optical counterparts which were then spectrally classified. The majority of candidates were found to be supernovae at redshift ranges similar to the GW events and were thereby ruled out as a genuine counterpart. Other candidates ruled out include AGN and Solar System objects. Given the GW sources were black hole binary mergers, the lack of an identified electromagnetic counterpart is not surprising. However the observations show that is it is possible to organise and execute a campaign that can eliminate the majority of potential counterparts. Finally we note the existence of a ``classification gap'' with a significant fraction of candidates going unclassified. | As part of the effort to discover and characterise astronomical gravitational wave (GW) sources, a worldwide programme of electromagnetic (EM) follow-up has been established (\cite[Abbot et al. 2016a]{abbott-a}). The programme is organised under the auspices of a series of Memoranda of Understanding (MOU) with over 70 groups who have access to observing resources that can participate in the follow-up. The MOU maintains confidentiality until the discoveries are announced. In the first aLIGO/Virgo observing run (O1) a number of candidate optical counterparts were identified by various wide field optical imaging facilities. In this paper we mainly discuss the work carried out by the 2.0 metre robotic Liverpool Telescope (LT) to spectroscopically follow up a number of those candidates. We also draw some broad conclusions about a potential spectroscopic ``classification gap''. More details of the work presented here can be found in \cite{copper-gw}. | Given that the O1 GW sources were both BH+BH systems (where we do not expect an EM signature) the overall results of the follow-up and classification programme are encouraging. We have shown that a 2-metre class telescope with a high-throughput, low resolution spectrograph can eliminate many candidate counterparts at redshifts similar to the GW sources. In future aLIGO/Virgo observing runs it is proposed that a distance estimate (``3d sky localisation'' - \cite[Singer et al. 2016]{singer}) and {\sc ``EM-BRIGHT''} flag (indicating the possible presence of a neutron star) will be distributed with the alerts. Localisations will also improve as more GW stations come on-line. All of this will help with targeted follow-up of candidates associated with potential host galaxies in the correct redshift range. For GW151226 a total of 77 candidate optical counterparts were announced via GCN. We note that (over all of the facilities involved): \begin{itemize} \item 37 of these received a firm spectral classification, \item a further 18 had a more tentative classification based on photometric light curves, \item there were 3 cases where the transient had faded into the host galaxy before spectroscopy could be attempted. \end{itemize} It follows that 19 candidates were not followed up. It is therefore clear that there is a significant danger of a ``classification gap'' opening up, where potential counterparts will be discovered at a significantly faster rate than can be spectroscopically followed up. While localisation error boxes are anticipated to reduce in size (from the current hundreds of square degrees to tens of square degrees) as more GW stations come on line, it is also anticipated that event numbers will increase by an order of magnitude over the next few years as detector sensitivities improve. The gap is therefore likely to remain a problem. In addition we note that new sources of transients such as LSST (predicted to discover $\sim10^6$ transients/night -- \cite[LSST Science collaboration et al. 2009]{lsst}) will also be in competition for spectroscopic follow-up time. To reduce the classification gap we propose the community must (a) start moving to more automated and efficient methods of triggering spectroscopic follow-up of candidates using technologies such as RTML (\cite[Hessman 2006]{rtml}) and VOEvent (\cite[Williams \& Seaman 2006]{vovenet}) and (b) consider the construction of optimized spectroscopic follow-up facilities with large apertures and fast slew speeds (e.g. Liverpool Telescope 2 - \cite[Copperwheat et al. 2015]{lt2}). | 16 | 9 | 1609.06195 |
1609 | 1609.01060_arXiv.txt | {We report the detection of OH$^+$ and H$_2$O$^+$ in the $z=0.89$ absorber toward the lensed quasar \PKS1830. The abundance ratio of OH$^+$ and H$_2$O$^+$ is used to quantify the molecular hydrogen fraction (\fH2) and the cosmic-ray ionization rate of atomic hydrogen (\zzH) along two lines of sight, located at $\sim 2$~kpc and $\sim 4$~kpc to either side of the absorber's center. The molecular fraction decreases outwards, from $\sim 0.04$ to $\sim 0.02$, comparable to values measured in the Milky Way at similar galactocentric radii. For \zzH, we find values of $\sim 2 \times 10^{-14}$~s$^{-1}$ and $\sim 3 \times 10 ^{-15}$~s$^{-1}$, respectively, which are slightly higher than in the Milky Way at comparable galactocentric radii, possibly due to a higher average star formation activity in the $z=0.89$ absorber. The ALMA observations of OH$^+$, H$_2$O$^+$, and other hydrides toward \PKS1830 reveal the multi-phase composition of the absorbing gas. Taking the column density ratios along the southwest and northeast lines of sight as a proxy of molecular fraction, we classify the species ArH$^+$, OH$^+$, H$_2$Cl$^+$, H$_2$O$^+$, CH, and HF as tracing gases increasingly more molecular. Incidentally, our data allow us to improve the accuracy of H$_2$O$^+$ rest frequencies and thus refine the spectroscopic parameters.} | Hydrides, that is, molecules or molecular ions containing a single heavy element with one or more hydrogen atoms, are formed in the interstellar medium (ISM) from the initial chemical reactions starting from atomic gas, making them fundamental to interstellar chemistry. Their rotational spectrum starting at submm/FIR wavelengths (a few exceptions, such as H$_2$S and H$_2$Cl$^+$, have ground-state transitions in the millimeter band) makes them difficult or even impossible to study from the ground. For this reason, many hydrides have only recently been detected in the ISM (see the review by \citealt{ger16}). The two species of interest here, hydroxylium (OH$^+$) and oxidaniumyl (H$_2$O$^+$), were first detected in the Galactic interstellar medium by \cite{wyr10} with the Atacama Pathfinder EXperiment (APEX) telescope, and by \cite{oss10} with the Herschel Space Observatory, respectively. These two species, and the third oxygen-bearing ion in the family, hydronium (H$_3$O$^+$), are formed from O$^+$ by successive hydrogen abstraction reactions with H$_2$. There can also be a significant contribution from the H$_3^+$+O reaction, in any place where the molecular fraction is relatively high. OH$^+$ and H$_2$O$^+$ are destroyed by dissociative recombination with electrons and by reactions with H$_2$. H$_3$O$^+$ does not react with H$_2$ and is mostly destroyed by dissociative recombination. The relative abundances of these species are thus controlled by the molecular hydrogen fraction and cosmic-ray ionization rate \citep{hol12}, providing a powerful diagnostic of these two quantities. \cite{ind15} analyzed Herschel observations of OH$^+$, H$_2$O$^+$, and H$_3$O$^+$ in multiple lines of sight toward bright submillimeter continuum sources in the Galaxy. Their results confirm that OH$^+$ and H$_2$O$^+$ are primarily tracing gases with a low molecular hydogen fraction ($\fH2 = 2 \times N({\rm H}_2) / [N({\rm H})+2 \times N({\rm H}_2)]$) of a few percent. They also determined the cosmic-ray ionization rate of atomic hydrogen (\zzH) across the Galactic disk and found an average $\zzH \sim 2 \times 10^{-16}$~s$^{-1}$ that shows little variation over the disk outside of Galactocentric radii of 5~kpc. Closer to the Galactic center, \zzH\ is found to increase by up to two orders of magnitude in the central region itself. To date, there are only a few observations of OH$^+$ and/or H$_2$O$^+$ in extragalactic sources (e.g., \citealt{vdwer10, yan13, gon13}), including detections at very high redshift \citep{wei13, rie13}. \cite{gon13} report ionization rates toward Arp\,220 and NGC\,4418 comparable to, or possibly higher than, values in the Galactic center. Although the Herschel Space Observatory is no longer in operation, it is now possible to extend the studies of these hydrides to other galaxies, e.g., with the Atacama Large Millimeter/submillimeter Array (ALMA). The molecular-line rich absorber at $z=0.89$ located in front of the quasar \PKS1830\ is certainly a prime target in this prospect. \PKS1830\ is lensed by the absorber, a nearly face-on typical spiral galaxy (e.g., \citealt{wik98,win02,koo05}). Molecular absorption is seen along the two independent lines of sight toward the southwest and northeast images of the quasar, at galactocentric radii of $\sim 2$~kpc and $\sim 4$~kpc in opposite directions from the center of the absorbing galaxy, respectively. The apparent size of the lensed images of the quasar is a fraction of a milliarcsecond at mm wavelengths (\citealt{jin03}). The volume of gas seen in absorption in each line of sight is thus roughly enclosed in a cylinder with a base of $< 1$~pc in diameter and a depth of a few tens to hundreds of pc, depending on the inclination ($i=17^\circ-32^\circ$, \citealt{koo05}) and the thickness of the absorber's disk. More than forty molecular species have been detected toward the southwest image, where the H$_2$ column density is $\sim 2 \times 10^{22}$~\cm-2\ (e.g., \citealt{mul11,mul14a}). In contrast, only about a dozen species have been detected toward the northeast image, where the molecular gas has a lower column and is more diffuse. Several submm lines of hydrides have already been observed toward \PKS1830\ with ALMA, such as those for CH, H$_2$O, NH$_2$, NH$_3$ \citep{mul14a}, H$_2$Cl$^+$ \citep{mul14b}, ArH$^+$ \citep{mul15}, as well as HF \citep{kaw16}. | \label{sec:conclusions} We present ALMA observations of OH$^+$ and H$_2$O$^+$ along two independent sight lines in the $z=0.89$ absorber toward the quasar \PKS1830. We find: \begin{itemize} \item OH$^+$ and H$_2$O$^+$ are both detected along the two lines of sight and, by comparison with other hydrides, show evidence for a multi-phase gas composition in the absorbing gas. \item The column density ratios between the SW and NE lines of sight increase in a sequence ArH$^+$, OH$^+$, H$_2$Cl$^+$, H$_2$O$^+$, CH, and HF, consistent with our understanding that these hydrides trace gas with an increasing molecular fraction. \item The abundance ratios between OH$^+$ and H$_2$O$^+$ are used to infer the fraction of molecular hydrogen and cosmic-ray ionization rate in the disk of the absorber. We find values ($\fH2=0.04$ and 0.02; $\zzH \sim 2 \times 10^{-14}$~s$^{-1}$ and $\sim 3 \times 10^{-15}$~s$^{-1}$ for the southwest and northeast lines of sight, respectively) which are comparable to those in the Milky Way at similar galactocentric radii ($\sim 2$~kpc and $\sim 4$~kpc, respectively). The slightly higher ionization rates found in the $z=0.89$ absorber could be related to a higher average star formation activity than in the Milky Way. \item We revise the rest frequencies of H$_2$O$^+$ lines, although precise laboratory microwave measurements are required to further increase the accuracy. \end{itemize} The number of sources for which observations of OH$^+$ and H$_2$O$^+$ have been reported and where \fH2 and \zzH\ have been estimated is still limited. The ALMA observations of \PKS1830\ presented in this paper allow us to investigate \fH2 and \zzH\ in the disk of a distant $z=0.89$ galaxy, and compare them to those in the Milky Way. With the discovery of new high-redshift molecular absorbers, such studies could be drastically extended with ALMA. \begin{acknowledgement} We thank the anonymous referee for her/his useful comments which helped us to improve the clarity of the manuscript. This paper makes use of the following ALMA data: ADS/JAO.ALMA\#2012.1.00056.S, \#2013.1.00020.S, and \#2015.1.00075.S. ALMA is a partnership of ESO (representing its member states), NSF (USA) and NINS (Japan), together with NRC (Canada) and NSC and ASIAA (Taiwan) and KASI (Republic of Korea), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, AUI/NRAO and NAOJ. \end{acknowledgement} | 16 | 9 | 1609.01060 |
1609 | 1609.01256_arXiv.txt | We study in this paper the capture of a massless particle into an isolated, first order Corotation Eccentric Resonance (CER), in the framework of the Planar, Eccentric and Restricted Three-Body problem near a $m+1:m$ mean motion commensurability ($m$ integer). While capture into Lindblad Eccentric Resonances (where the perturber's orbit is circular) has been investigated years ago, capture into CER (where the perturber's orbit is elliptic) has not yet been investigated in detail. Here, we derive the generic equations of motion near a CER in the general case where both the perturber and the test particle migrate. We derive the probability of capture in that context, and we examine more closely two particular cases: $(i)$ if only the perturber is migrating, capture is possible only if the migration is outward from the primary. Notably, the probability of capture is independent of the way the perturber migrates outward; $(ii)$ if only the test particle is migrating, then capture is possible only if the algebraic value of its migration rate is a decreasing function of orbital radius. In this case, the probability of capture is proportional to the radial gradient of migration. These results differ from the capture into Lindblad Eccentric Resonance (LER), where it is necessary that the orbits of the perturber and the test particle converge for capture to be possible. Possible applications for planetary satellites are discussed. | Orbital captures into Mean Motion Resonance (MMR) is the key to understanding the orbital evolution of satellites, rings and planets. The special case of the Lindblad Eccentric Resonance (LER) has been investigated by many authors in the context of the Planar, \textit{Circular} and Restricted Three-Body Problem. In this case, a secondary object orbiting a massive primary body perturbs a massless particle, so that the critical angle $\phi_L=(m+1)\lambda_s - m\lambda_p - \varpi_p$ librates, where $m$ is an integer, $\lambda_s$ and $\lambda_p$ are the longitudes of the secondary and the test particle, respectively, and $\varpi_p$ is the longitude of pericenter of the test particle. In a general context, \cite{henrard82} estimated the probability of capture into a first order LER, while \cite{borderies84} extended this work and derived capture probabilities into $m+1:m$ and $m+2:m$ (second order) LERs. If the orbit of the secondary is \textit{eccentric}, a Corotation Eccentric Resonance (CER) appears close to and associated with each LER. It is dynamically described by the critical angle $\phi_c=(m+1)\lambda_s - m\lambda_p - \varpi_s$, where $\varpi_s$ is the longitude of the secondary pericenter. The physical effects of LERs and CERs are not the same: the CER mainly affects the semi-major axis of the test particle and keeps its orbital eccentricity almost constant, forcing the particle to librate inside so-called corotation sites like a simple pendulum. In contrast, the LER acts on its eccentricity but keeps the semi-major axis almost constant \citep{elmoutamid14}. In this work, we derive the probability of capturing a test particle into an isolated CER, as both the secondary and the particle suffer orbital migration that secularly change their semi-major axes. This is a novel calculation that complements what has been done before in the case of LERs. The term ``isolated" means here that the CER and LER are sufficiently pulled apart so that their coupling is negligible. Eventually, the goal is to extend the study of captures into MMR to more realistic cases where both the CER and LER act in concert on the particle, but this will not be considered here. This work can be applied to many situations, in the context of planetary rings and satellites. In the Saturnian system for example, the satellites Aegaeon, Anthe and Methone are respectively captured into 7:6, 10:11 and 14:15 CERs with Mimas \citep{cooper08,hedman09,hedman10,elmoutamid14}. Atlas is in a 54:53 CER with Prometheus \citep{renner16}. In the case of the Neptunian system, Adams ring arcs may be dynamically confined by the satellite Galatea via corotation resonances \citep{renner14,nicholson95,foryta96,sicardy99}. \if{ Moreover, thanks to the Kepler mission, it has been observed that in the context of exoplanets pairs, commensurabilities configurations are very rare. Here we show that the probability of capture is very small \red [I do not see anywhere that we show that!] \bla which is in agreement with the exoplanet pairs results of \cite{fabrycky14}. \red [I do not agree, see the discussion] \bla Future work in this context will be done to study those cases. \red [This whole paragraph belongs to the discussion, it is too early here]. \bla }\fi We note here that the capture of a particle into a CER bears some resemblance with the capture of a rotating body into spin-orbit resonance \citep{gold66}. In both cases, a slow effect (orbital migration in our case, tidal friction for the spin-orbit resonances) drives a pendulum-like system into a librating state, possibly capturing it permanently into that state. This will be commented later. The paper is structured as follows: In section 2, we describe the dynamical structure of our problem based on the so-called CorALin model, taking into account the dissipation parameters. In Section 3, we study the particle and the secondary migration terms involved in the derivation of the probability of capture. The latter is derived in section 4. In section 5, we discuss the similarity of our problem with the capture in spin-orbit resonance and finally, a summary and conclusions are given in Section 6. | We have studied the mechanism of capture into an isolated first order Corotation Eccentric Resonance (CER) in the context of the Elliptical Restricted Three-body Problem, in which a test particle orbits around a central mass $M_c$, near a first order mean motion resonance $m+1:m$ with a perturbing secondary of mass $m_s$. We derive a formula for the probability of capture given by Eq.~\ref{eq_proba_gen_approx} in a general context. Then we apply this result to two particular cases, first where we consider a migration of the particle only (Eq.~\ref{eq_proba_eg}), and secondly, in the case of migration of the secondary (Eq.~\ref{eq_proba_es}). We point out the noteworthy fact that in this case, the capture probability does not depend on the way the secondary migrates. Under realistic assumptions, we rewrite our formula for capture in the second case as Eq.~(\ref{eq_proba_SCER}). This equation has the interesting consequence that if the corotation radius $a_0$ sweeps a total area $S$ in a region where particles are uniformly distributed with surface density $N$, then the corotation sites will eventually be populated (due to captures) with a particle surface density of $NS/4S_0$. Conversely, if we observe today a certain number of particles trapped in corotation sites, we may estimate the number of particles originally present in that region. This work can be applied in many cases, specially in the context of planetary rings, satellites and in the context of exoplanets. As an example, in the context of the Saturn system, \cite{elmoutamid14} study the case of Aegaeon which is trapped in a 7:6 CER caused by Mimas, with $a_0 \sim 167,500$~km. Since $W_{\rm CER} \sim 30$~km, we find a probability of capture of $\sim5\times10^{-5}$. This very low probability suggests that originally, there were many more such objects in the Saturn system, but that only a few of them were captured into CERs. Note finally that this study assumes that the CER is isolated, i.e. that the condition~(\ref{eq_condition}) is fulfilled. In the cases where $D_{\rm CL} \sim W_{\rm CER}$, the motion in phase space is chaotic, as the LER and CER are strongly coupled, see Fig.~\ref{fig_chaos}. This occurs when the effects of the oblateness of the central body, the presence of a massive disk or another companion in the system, are not strong enough to split and thus isolate the resonances from each other. The chaotic nature of motion then prevents an easy analytical derivation of the probability of capture into the $m+1:m$ mean motion resonance. In that case, numerical integrations might be useful to see whether the probability derived here (Eq.~\ref{eq_proba_gen}) remains valid, at least in order of magnitude. | 16 | 9 | 1609.01256 |
1609 | 1609.06689_arXiv.txt | Using numerical models for star clusters spanning a wide range in ages and metallicities ($\metal$) we study the masses of binary black holes (BBHs) produced dynamically and merging in the local universe ($\redshift\lesssim0.2$). After taking into account cosmological constraints on star-formation rate and metallicity evolution, which realistically relate merger delay times obtained from models with merger redshifts, we show here for the first time that while old, metal-poor globular clusters can naturally produce merging BBHs with heavier components, as observed in GW150914, lower-mass BBHs like GW151226 are easily formed dynamically in younger, higher-metallicity clusters. More specifically, we show that the mass of GW151226 is well within $1\sigma$ of the mass distribution obtained from our models for clusters with $\metal/\Zsun\gtrsim0.5$. Indeed dynamical formation of a system like GW151226 likely requires a cluster that is younger and has a higher metallicity than typical Galactic globular clusters. The LVT151012 system, if real, could have been created in any cluster with $\metal/\Zsun\lesssim0.25$. On the other hand, GW150914 is more massive (beyond $1\sigma$) than typical BBHs from even the lowest-metallicity ($\metal/\Zsun=0.005$) clusters we consider, but is within $2\sigma$ of the intrinsic mass distribution from our cluster models with $\metal/\Zsun\lesssim0.05$; of course detection biases also push the observed distributions towards higher masses. | \label{S:intro} Detection of gravitational waves (GWs) from merging black hole (BH) binaries has reignited widespread interest in understanding the astrophysical implications and the origins of BBHs \citep[][]{PhysRevLett.116.061102,PhysRevLett.116.241103,2041-8205-818-2-L22}. Current theoretical estimates indicate that detectable BBH merger events may be rather frequent, few--$500\,\mergerrate$, with large uncertainties depending on production channels and model assumptions (e.g., compare \citealt{2016PhRvD..93h4029R,2017MNRAS.464L..36A} with \citealt{2016MNRAS.460.3545D}). In the first observing run itself, the advanced LIGO observatories (aLIGO) have detected GW signals from two BBH mergers and a lower significance `trigger' event \citep{PhysRevLett.116.061102,PhysRevLett.116.241103,2016PhRvX...6d1015A}. These detections already show a large diversity in the masses of BBHs merging in the local universe: the chirp masses ($\mchirp$) at source for GW150914, LVT151012, and GW151226 are $28_{-2}^{+2}$, $15^{+1}_{-1}$, and $8.9^{+0.3}_{-0.3}$, respectively. Broadly speaking, two major channels have been proposed for BBH formation and subsequent merger. High-mass stellar binaries may evolve in isolation to create merging BBHs, for example, by going through a specific sequence of events involving low-kick supernovae (SNe) and common envelope (CE) evolution \citep[e.g.,][]{2012ApJ...759...52D,2013ApJ...779...72D,2015ApJ...806..263D,2014ApJ...789..120B,2016ApJ...819..108B,2015A&A...574A..58K,2016Natur.534..512B}, and via chemically homogeneous evolution of tidally distorted binaries \citep[e.g.,][]{2016MNRAS.458.2634M,2016A&A...588A..50M}. Alternatively, merging BBHs could be produced dynamically at the centers of dense star clusters \citep[e.g.,][]{2010MNRAS.402..371B,2014MNRAS.441.3703Z,2015PhRvL.115e1101R,2016PhRvD..93h4029R}. The process involved here is fundamentally different from BBH formation in isolation. A negligible fraction of BBHs formed in dense massive star clusters are primordial, and none with $\tdelay\geqslant1\,\gyr$ are composed of BHs formed from stars that were born in that binary \citep[e.g.,][]{2017ApJ...834...68C}. As massive stellar binaries evolve in dense star clusters, even if they were initially hard, mass loss from stellar winds and compact object formation can make these binaries soft. Consequently, stellar encounters and natal kicks during BH formation disrupt these primordial binaries. Later, single BHs dynamically acquire other BH companions via three-body binary formation and binary-mediated exchange interactions which preferentially insert the relatively more massive BHs into a binary ejecting a less massive non-BH member from it \citep[e.g.,][]{2003gmbp.book.....H,2012ApJ...754..152C}. Because of this, the BBH properties and their merger times, unlike those formed in isolation, do not depend on the assumptions of initial binarity or binary orbital properties. While the rate of mergers are affected by the assumptions for the IMF and natal kicks, the BBH masses and merger delay times ($\tdelay$) are insensitive to those as well \citep{2017ApJ...834...68C}. While several studies have modeled BBH formation in isolation for a wide range in metallicities taking into account the cosmological evolutions of the star formation rate ($\sfr$) and metallicity \citep[e.g.,][]{2016Natur.534..512B,2016MNRAS.461.3877D,2016MNRAS.460.3545D}, due to primarily the computational cost, numerical studies of dynamically formed BBHs have so far restricted themselves to either a narrow range in metallicities and ages typical of the Galactic globular clusters (GGCs) \citep[e.g.,][]{2015PhRvL.115e1101R,2017ApJ...834...68C} or to low-mass (initial $N\sim5\times10^3$), young ($\sim100\,\myr$) star clusters \citep[e.g.,][]{2014MNRAS.441.3703Z}. These studies also assumed that all model clusters formed roughly at the same epoch independent of the metallicity to evaluate the redshifts of BBH mergers from $\tdelay$ found in the models. This of course is a simplification. Stars form with wide ranges in metallicities at any redshift \citep[e.g.,][]{2014ARA&A..52..415M}. Massive star clusters are also observed today with a large range in ages and metallicities, for example, in M51, M101, and the LMC \citep[e.g.,][]{2005A&A...431..905B,2006AJ....132..883B,2007A&A...469..925S}. Even for the GGCs, the metallicity distribution has a long tail extending to $\Zsun$ \citep{1996AJ....112.1487H}. We relax past assumptions and consider BBH formation and merger in clusters spanning a wide range in metallicities and metallicity-dependent distributions for cluster-formation redshifts ($\zform$). Our goal is to investigate whether all hitherto detected GW sources could have been formed dynamically in star clusters. Furthermore, we study effects of star cluster metallicity and age on the detectable properties (mass and eccentricity) of BBH mergers. In \S\ref{S:numerical} we describe our numerical setup. In \S\ref{S:result} we show the key results. We conclude in \S\ref{S:conclude}. | \label{S:conclude} We have studied the effects of the parent cluster's metallicity (and metallicity-dependent age) on the BBH masses merging in the local universe. Assuming cluster origin, we have found likely cluster properties of detected GW sources by comparing detected masses with mass distributions from models (\S\ref{S:BBHprop}, Fig.\ \ref{fig:zvsm}). We find that $\mchirp$ of GW150914 is not within $1\sigma$ of the intrinsic $\mchirp$ distributions for BBHs merging in $\redshift\leq0.2$ for any metallicities we consider, but is within $2\sigma$ for mergers from clusters with $\metal/\Zsun\leq0.05$. Since below $\metal/\Zsun=0.05$ the $\mchirp$-distribution is insensitive to metallicity and dynamically created BBHs are generally heavier than those produced in isolation for any given metallicity, mergers of BBHs as massive as GW150914 in $z\leq0.2$ are likely intrinsically rare. Of course, detection biases push the observed distributions towards higher masses. Since the lower the metallicity, the larger the range in merging BBH masses, detection biases would affect the mass distributions from lower-metallicity clusters more. Thus, the {\em detection} of mergers like GW150914 would be less rare \citep{2016ApJ...824L...8R}. $\mchirp$ of LVT151012 is near the peak of the distribution from clusters modeled with $\metal/\Zsun=0.25$, and is within $1\sigma$ of the distributions from all clusters modeled with $\metal/\Zsun\leq0.25$. $\mchirp$ of GW151226 is closest to the peak of $\mchirp$-distribution from clusters with $\metal/\Zsun=0.75$ and is within $1\sigma$ from clusters with $\metal/\Zsun\geq0.5$ Thus, assuming cluster origin, GW151226 likely formed in a higher-metallicity, younger cluster than typical GGCs. We find several additional notable trends. Less massive BBHs have longer $\tdelay$ for any metallicities (Fig.\ \ref{fig:tdelay}) since clusters dynamically form, process, and eject heavier BBHs earlier due to mass segregation \citep[e.g.,][]{2015ApJ...800....9M}. Lower-metallicity clusters typically have higher $\zform$ (Fig.\ \ref{fig:zdist}). Hence, BBHs from lower-metallicity clusters require longer $\tdelay$ to merge in $\redshift\lesssim0.2$. Lower metallicity leads to the formation of heavier BBHs \citep[e.g.,][]{2012ApJ...749...91F}, but longer $\tdelay$ decreases merging BBH masses (Fig.\ \ref{fig:tdelay}). Hence, while the expected trend is an increase of BBH masses merging in $\redshift\lesssim0.2$ as metallicity decreases, the mass distributions for local BBH mergers become insensitive to metallicity for $\metal/\Zsun\lesssim0.05$. Furthermore, we find that the masses of local BBH mergers are not very sensitive to the initial $N$ or $r_v$ even when $N$ is varied over an order of magnitude and $r_v$ by a factor of $2$ (Table\ \ref{T:props}). This indicates that the metallicity and metallicity-dependent age of the parent cluster are likely the most important properties to determine the peaks and distributions of BBH masses merging in the local universe. | 16 | 9 | 1609.06689 |
1609 | 1609.01768_arXiv.txt | The High Altitude Water Cherenkov (HAWC) $\gamma$-ray observatory is a wide field of view (1.8~Sr) and high duty cycle ($>95\%$ up-time) detector of unique capabilities for the study of TeV gamma-ray sources. Installed at an altitude of 4100m in the Northern slope of Volc\'an Sierra Negra, Puebla, by a collaboration of about thirty institutions of Mexico and the United States, HAWC has been in full operations since March 2015, surveying 2/3 of the sky every sidereal day, monitoring active galaxies and mapping sources in the Galactic Plane to a detection level of 1~Crab per day. This contribution summarizes the main results of the first year of observations of the HAWC $\gamma$-ray observatory. | The advent of new spectral windows for the study of the Universe during the second half of the XX$^{\rm th}$ Century brought a remarkable expansion of our knowledge of the cosmos. Today, astronomy is performed across more than twenty decades of the electromagnetic spectrum, from tens of neV, the photon energy of 100 MHz radio-waves, to at least 100 TeV. Ranging from the MeV to the PeV, gamma rays represent the most extreme form of electromagnetic radiation. Its upper boundary may be even further in energy, having been sought up to the $10^{19}\,\rm eV$ range with the Pierre Auger observatory~\cite{auger-limit}. From the observational point of view, $\gamma$ rays are highly non-thermal radiation associated to astrophysical particle accelerators that produce cosmic rays. And, in the same way as high-energy $\gamma$-ray astronomy is complemented at other wavelengths, it complements searches in the dawn of the multi-messenger era. Astronomical observations of $\gamma$ rays are performed with space and ground-based instruments. There are three types of detectors with their own particular characteristics, as illustrated in the diagram~\ref{gamma-tel}. Their differences in energy coverage, field of view and sensitivity makes them highly complementary: in the same manner as space-borne pair production telescopes, like {\em Fermi}-LAT or {\em AGILE}, perform continuous observations of large portions of the sky in the GeV range, HAWC permanently surveys the sky transiting above it in the TeV range; both types of instruments provide useful information to atmospheric Cherenkov telescopes, which are highly sensitive but require to be pointed at specific parts of the sky, which are inaccessible part of the year. \noindent \setlength{\unitlength}{1cm} \begin{picture}(16,12) \put(1.5,6.5){\framebox(7,4.5)[t]{}} \put(1.5,10.4){\makebox(7,0.5)[t]{\underline{\bf Pair production telescopes}}} \put(1.5,9.8){\makebox(7,0.5)[t]{0.1 - 100 GeV}} \put(1.5,9.3){\makebox(7,0.5)[t]{Small effective area}} \put(1.5,8.8){\makebox(7,0.5)[t]{Background free}} \put(1.5,8.3){\makebox(7,0.5)[t]{Large FoV and high duty cycle}} \put(1.5,7.8){\makebox(7,0.5)[t]{\dotfill}} \put(1.5,7.5){\makebox(7,0.5)[t]{All sky surveys \& monitoring}} \put(1.5,7.0){\makebox(7,0.5)[t]{Transients (AGN, GRB)}} \put(1.5,6.5){\makebox(7,0.5)[t]{Extended diffuse emission}} \put(8.5,6.5){\framebox(7,4.5)[t]{}} \put(8.5,10.4){\makebox(7,0.5)[t]{\underline{\bf Air shower arrays}}} \put(8.5,9.8){\makebox(7,0.5)[t]{0.1 - 100 TeV}} \put(8.5,9.3){\makebox(7,0.5)[t]{Large effective area}} \put(8.5,8.8){\makebox(7,0.5)[t]{Good noise rejection}} \put(8.5,8.3){\makebox(7,0.5)[t]{Large FoV and high duty cycle}} \put(8.5,7.8){\makebox(7,0.5)[t]{\dotfill}} \put(8.5,7.5){\makebox(7,0.5)[t]{Partial sky surveys \& monitoring}} \put(8.5,7.0){\makebox(7,0.5)[t]{Transients (AGN, GRB)}} \put(8.5,6.5){\makebox(7,0.5)[t]{Extended diffuse emission}} \put(8.5,2.0){\framebox(7,4.5)[t]{}} \put(8.5,5.9){\makebox(7,0.5)[t]{\underline{\bf Atmospheric Cherenkov telescopes}}} \put(8.5,5.3){\makebox(7,0.5)[t]{30 GeV - 30 TeV}} \put(8.5,4.8){\makebox(7,0.5)[t]{Large effective area}} \put(8.5,4.3){\makebox(7,0.5)[t]{Excellent noise rejection}} \put(8.5,3.8){\makebox(7,0.5)[t]{Small FoV and low duty cycle}} \put(8.5,3.3){\makebox(7,0.5)[t]{\dotfill}} \put(8.5,3.0){\makebox(7,0.5)[t]{Detailed studies of known sources}} \put(8.5,2.5){\makebox(7,0.5)[t]{Deep surveys of limited regions}} \put(8.5,2.0){\makebox(7,0.5)[t]{High resolution spectra}} \put(1.5,2.0){\framebox(7,4.5){}} \thicklines \put(2.5,2.5){\vector(1,0){5}} \put(7.7,2.4){\mbox{\bf E}} \put(2.5,2.5){\vector(0,1){3}} \put(2.2,5.7){\mbox{\bf FoV}} \put(0.7,3.0){\mbox{Deg}} \put(0.7,8.5){\mbox{Sr}} \put(5.0,1.6){\mbox{GeV}} \put(12.0,1.6){\mbox{TeV}} \put(1.5,0.5){\makebox(14,1){Diagram 1: three types of astronomical $\gamma$ ray telescopes separated in terms}} \put(1.5,0.0){\makebox(14,1){ of their energy coverage ({\em horizontally}) and field of view (FoV) ({\em vertically}).\label{gamma-tel}}} \end{picture} The Large Aperture Telescope (LAT) on board of the {\em Fermi $\gamma$-Ray Space Telescope} has been one of the main driving forces in high energy astrophysics in the last years. Since its 2008 launch, {\em Fermi}-LAT has obtained a deeper and deeper exposure of the whole sky, leading to the finding of about 3000 sources of $\gamma$ rays listed in various catalogs. In particular the 3FGL catalog is the deepest all-sky survey in the 0.1-300 GeV range performed to date, while the 1FHL catalog contains sources detected using photons of energies $>10~\rm GeV$~\cite{3fgl,1fhl}. Of particular relevance for high energy studies is the 2FHL catalog of sources detected in the $50~\rm GeV - 2~TeV$ range, which offers the possibility of complementary studies with atmospheric Cherenkov telescopes and air shower arrays, in particular HAWC~\cite{2fhl}. Atmospheric Cherenkov telescopes (ACTs) provide unique detailed information about very high-energy $\gamma$-ray sources~\cite{tev-astro2008}. They employ the atmosphere as the upper part of the detector, taking advantage of particle cascades triggered by the interaction of cosmic and $\gamma$ rays with atmospheric nuclei. ACTs sample the Cherenkov light produced by the secondary particles of the cascade as they travel down the atmosphere faster than the speed of light in the air. As the emission cone of these particles is rather narrow, about $1.4^\circ$, these telescopes have to be pointed to individual objects for their study. Furthermore, the emitted Cherenkov light is faint and can only be detected during night-time, on clear dark nights. To their advantage, ACTs are much more sensitive than air shower arrays: they can detect in a few minutes point sources that require more than one transit with HAWC. In the last couple of decades ACTs have found tens of TeV sources through pointed observations, mostly active galactic nuclei, in particular BL Lacs \& flat spectrum radio quasars, and also starburst galaxies and Galactic sources, as listed in the TeVCat~\cite{tevcat}. An important input to TeVCat is the Galactic Plane survey performed by the HESS collaboration. This is a very deep survey in the TeV band covering Galactic latitudes $|b|<3.5^{\circ}$, and longitudes $250^{\circ} < \ell < 360^{\circ}$ and $\ell < 65^\circ$~\cite{gplane-hess}. As a consequence of the large number of Galactic sources found in the HESS Galactic Plane survey, the distribution of source types in TeVCat departs from the one derived from the 2FHL, highlighting the bias of (deep) pointed observations and highlighting the relevance of TeV wide field of view and high duty cycle observatories. Extensive air shower (EAS) arrays can perform continuous survey and monitoring observations complementary to the deeper pointed ACT exposures. EAS arrays have been employed as cosmic-ray detectors for decades; the cosmic-ray experiment installed at Haverah Park in the 1960s is one of the earliest examples of water Cherenkov detectors~\cite{haverah-park}. EAS with the capability of separating photons and hadrons have been developed in United States (MILAGRO), Mexico (HAWC) and China (Tibet AS-$\gamma$ and ARGO). In particular MILAGRO proved the feasibility of astronomical observations with the water Cherenkov technique by performing a successful multi-steradian survey of the $\gamma$-ray TeV sky. It detected photons of median energy 40 TeV from Galactic objects, in particular extended diffuse sources associated to GeV pulsars. At the Galactic anti-center MILAGRO detected the Crab Nebula and an extended source at the location of Geminga; the first Galactic quadrant showed prominent emission at the Cygnus region and the previously unreported source MGRO~J1908+06. The BL~Lac object Mrk~421 was also detected and sampled in timescales of months during the lifetime of the observatory~\cite{milagro-survey}. \begin{figure}[t] \begin{center} \vspace{-8mm}\includegraphics[width=0.9\hsize]{figure1} \end{center} \vspace{-6mm} \caption{\label{skymap}HAWC skymap of statistical significances above background, in celestial coordinates. The declination coverage is approximately $-30^{\circ}<\delta<+70^{\circ}$. The emission band at the first Galactic quadrant is clearly visible on the left. The three outstanding point sources from left to right, Mrk~501, Mrk~421 and the Crab Nebula. Next to the Crab, with lower significance, is Geminga.} \end{figure} | 16 | 9 | 1609.01768 |
|
1609 | 1609.08680_arXiv.txt | For 32 central stars of PNe we present their parameters interpolated among the new evolutionary sequences. The derived stellar final masses are confined between 0.53 and 0.58 $M_\odot$ in good agreement with the peak in the white dwarf mass distribution. Consequently, the inferred star formation history of the Galactic bulge is well restricted between 3 and 11\,Gyr and is compatible with other published studies. The new evolutionary tracks proved a very good as a tool for analysis of late stages of stars life. The result provide a compelling confirmation of the accelerated post-AGB evolution. | The main stellar evolutionary tracks for post-AGB evolution are now over 20 years old. Recent results show that they predict the wrong masses and ages for central stars of planetary nebulae (PNe). \citet{GZHS2014} based on a detailed analysis of 31 Galactic bulge PNe proposed that the post-AGB evolution should be accelerated by a factor of 3. Coincidentally, during the last White Dwarf Workshop, \citet{M3B2015} presented new models for the evolution of central stars of PNe obtaining similarly shorter post-AGB timescales on purely theoretical grounds. Here we re-analyse the data sample (recently partially upgraded) however this time we interpolate among the new evolutionary sequences aiming to verify their timescales. | The most important outcome of our work are the much lower final masses, as compared with the old models, which is a robust result. The obtained final masses are in good correspondence with the most common white dwarf masses, $\sim0.55\,M_\odot$. Our results agree well with the previous suggestion \citep{GZHS2014} that the \citet{B1995} tracks were far too slow. The Galactic bulge SFH, derived from the total ages, points to a burst-like old population, plus some younger objects which is in agreement with the recent Galactic disk SFH and with the idea that the bulge originated from disk instabilities. These numerous accordances constitute a strong support for the short timescales of the new models of \citet{M3B2016}. The kinematical/photoionization reconstruction method elaborated recently by \citet{GZM2016} proved to be efficient. | 16 | 9 | 1609.08680 |
1609 | 1609.06148_arXiv.txt | Owing to technological advances, the number of exoplanets discovered has risen dramatically in the last few years. However, when trying to observe Earth analogs, it is often difficult to test the veracity of detection. We have developed a new approach to the analysis of exoplanetary spectral observations based on temporal multifractality, which identifies time scales that characterize planetary orbital motion around the host star, and those that arise from stellar features such as spots. Without fitting stellar models to spectral data, we show how the planetary signal can be robustly detected from noisy data using noise amplitude as a source of information. For observation of transiting planets, combining this method with simple geometry allows us to relate the time scales obtained to primary and secondary eclipse of the exoplanets. Making use of data obtained with ground-based and space-based observations we have tested our approach on HD 189733b. Moreover, we have investigated the use of this technique in measuring planetary orbital motion via Doppler shift detection. Finally, we have analyzed synthetic spectra obtained using the SOAP 2.0 tool, which simulates a stellar spectrum and the influence of the presence of a planet or a spot on that spectrum over one orbital period. We have demonstrated that, so long as the signal-to-noise-ratio $\ge$ 75, our approach reconstructs the planetary orbital period, as well as the rotation period of a spot on the stellar surface. | \label{sec:intro} The last three decades have seen the birth of exo-planetary science. With the advent of various techniques, which include, but are not limited to, pulsar timing \citep{Wolszczan:1992}, Doppler measurements \citep{Mayor:1995}, transit photometry \citep{Charbonneau:2000}, micro-lensing \citep{Beaulieu:2006} and direct imaging \citep{Chauvin:2004}, thousands of planets have been detected orbiting distant stars. A central focus is the detection of so-called \emph{Exo-Earths}, Earth-like planets in terms of mass and radius, orbiting around a star at a distance that, given sufficient atmospheric pressure, would allow for the existence of liquid water on its surface \citep[e.g][and references therein]{Kopparapu:2013}. The techniques that are most commonly used in discovering other planets are transit photometry \citep[e.g.][and references therein]{Lissauer:2014aa} or Doppler measurements \citep[e.g.][and references therein]{Mayor:2014aa}. Recently, using the transit method, detection of nine candidates for habitable planets was announced that may fall within the habitable zones of their host stars \citep{Anglada:2016,Morton:2016aa}. Whilst the combination of these approaches have provided an impressive range of observations, the detection of Earth analogs is a challenging problem. Indeed, the presence of instrumental and astrophysical noise are sources of uncertainty for such discoveries \citep[e.g.][]{Fischer:2016aa}. The fingerprints of an exo-Earth could easily be hidden in stellar noise, or stellar signals might mimic the presence of an exoplanet. Moreover, when such noise is modeled, there is a risk of introducing spurious signals in the analysis of data \citep[e.g.][]{Dumusque:2015aa, Rajpaul:2016aa}. Contemporary studies aim to understand and correctly evaluate the imprint of stellar activity on exoplanet detections \citep[e.g.][]{Lanza:2007,Lanza:2011,Aigrain:2012,Korhonen:2015}. Nonetheless, radial velocity and transit photometry methods can still produce false detections, especially when dealing with Earth-like planets, or planets characterized by a signal $\leq$1m s$^{-1}$ \citep{Dumusque:2016,Morton:2016aa}, which is, in part, due to the models upon which these methods are based. Such exoplanet evolution and stellar models over-fit parameters to the data, which turns out to be crucial when one aims to detect terrestrial-like planets \citep{Dumusque:2012aa,Rajpaul:2016aa}. These parameters include, but are not limited to, planet and stellar radius, their masses, eccentricity, impact factor, transit duration, transit ingress/egress duration, transit depth, orbital inclination, distance of the planet from the star, limb-darkening parameters (which themselves may vary with the law used), and shape of the transit curve. In some cases, instrumental signals such as the wobble of the instrument cluster aboard the \emph{Spitzer} satellite or the latent charge build-up in the pixels (ramp) are also modeled \citep{Grillmair:2007aa,Todorov:2014aa}. Additionally, noise sources such as granulation over the stellar surface employ model fitting to estimate the effect of that noise on the data \citep{Dumusque:2012aa}, stellar activity concurrent with the stellar rotation are modeled by fitting sine waves through the radial velocity data, and long term stellar activity, light contamination from near by stars, among others are parametrically modeled. This fitting of models has led to the introduction of spurious signals in the observed data and hence to false detections. In order to fit these stellar models to data, one must begin by considering a particular system, for example an unblended eclipsing binary or a transiting planet. Hence, the list of models that the data can represent needs to be complete, which constitutes a weakness for these types of fitting schemes \citep{Morton:2016aa}. Moreover, these methods cannot always distinguish whether the signal is instrumental or from stellar activity \citep{Morton:2016aa}. Finally, because it affects the details of the fitting and hence the results, the observations must have a high signal-to-noise ratio (S/N). A key aspect of the analysis of exo-planetary systems is the identification of periodic timescales. This is also central to the study of multi-planetary systems \citep[e.g.][]{Millholland:2016aa}, and often requires the combined use of the two most successful approaches in exoplanet searches; the transit and radial velocity methods. The most widely used method for identifying periodic timescales is the Lomb--Scarge periodogram, which is based on the assumption that the data can be interpreted as a sum of periodic signals. However, this technique is used on data that have been ``filtered'' using models for stellar noise and terrestrial atmospheric contamination. Hence, it is again possible to accumulate artifacts in the data through the use of such models. Many approaches have been developed by the exoplanetary science community to address these matters. For example, \cite{Dumusque:2016} have compared state-of-the-art methods on simulated radial velocity data and concluded that the detection of planets below the 1 m s$^{-1}$ threshold is still controversial. Moreover, the extraction of radial velocities from spectra relies on cross-correlation techniques that are based on the use of spectral templates based on stellar models. Therefore, the determination of the Doppler shift remains a method-dependent challenge. For example, in the recent radial velocity study of Proxima Centauri \citep{Anglada:2016}, the use of the TERRA algorithm \citep{Anglada:2012} rather than the HARPS pipeline, has been crucial for the quantification of the signal. Here, to extract the timescales that characterize a planetary system and stellar features without {\em a priori} assumptions about the data itself, we introduce a new approach to spectral analysis. Namely, in the spirit of the Langevin theory of Brownian motion, we quantify a signal coming from a star as the combination of a deterministic dynamics and stochastic noise, but make no {\em a priori} assumptions about the nature of these processes, and thereby examine an unfiltered time series $X_i$ of a stellar spectral signal. If the star hosts a planetary system, the time scales associated with stellar rotation and activity, as well as with planetary motions must be present. We need not (a) make assumptions regarding the combination of periodic signals, or (b) use stellar models. The goal is to identify the dominant timescales of the observed system as agnostically as possible, and then use elementary geometry to reveal the underlying dynamics of the system. The flexibility of this approach, based on temporal multifractality, allows one to identify stellar signals that would otherwise be missed by fitting sine waves to the data. We begin with a description of the method in \S \ref{sec:method}. The proposed approach can be used both for transiting planets and for radial velocity measurements. We show two examples in \S \ref{sec:data}; one for a transiting planet, and the other for a simulated observation of a planet detectable only via radial velocity measurements. Finally, we discuss the results and their robustness in \S \ref{sec:discuss}, and conclude in \S \ref{sec:conc}. | \label{sec:conc} We have presented and tested a new multi-fractal approach for the analysis of exo-planetary spectral observations. The goal is to use a fit-free procedure to identify robust time scales associated with the exo-planetary orbital motion around a host star, as well as to detect time scales associated with stellar features. With these timescales in hand, one can compute key system parameters such as the ratio of the size of the planet to that of the star and the latitude of transit, {\em without} use of stellar evolution models, data fitting, noise filtering, and the additional wide variety of other assumptions about the system that are typically made. The concept of the approach is to take an agnostic (or model free) view of the observed spectral structure. The method makes no {\em a priori} assumptions about the temporal structure in any observed spectra and makes use of only one number, the generalized Hurst exponent, the value of which underlies the identification of the key time scales. We have reconstructed the primary and secondary transit times of the exoplanet HD 189733b, using data from both ground-based (HARPS spectrograph) and space-based (\emph{Spitzer}) observations. Using the SOAP 2.0 tool, which simulates a stellar spectrum and how it can be affected by the presence of a planet across one orbital period, our method is further tested in the context of measuring planetary orbital motion via Doppler shift detection. Because the SOAP tool can also simulate the presence of a spot on the stellar surface across one rotational period of the star, we have demonstrated that our approach reconstructs the planetary orbital period, as well as the rotation period of a spot covering 5\% of the stellar surface. Moreover, we have tested the analysis with a wide range of S/Ns. We can reconstruct: (a) the period of a planet producing a 40 m s$^{-1}$ Doppler shift of the stellar spectrum, and (b) the rotational period of the star, based on the presence of a spot on its surface, provided that the S/Ns are $\ge 75$ and $\ge 30$ respectively. Importantly, to avoid introduction of errors arising from intrinsic irregularities in the analyzed time series, this method has the highest fidelity when observations are carried out at sufficiently high frequency that the time-difference between each measurement is less than the shortest relevant timescale that may be present in the system. In conclusion this method based on Multi-fractal Temporally Weighted Detrended Fluctuation Analysis of time series is a robust way to measure planetary orbital motion. It provides a fertile framework to examine other data sets and to explore trying to systematically distinguish stellar noise from planetary motion. | 16 | 9 | 1609.06148 |
1609 | 1609.03592_arXiv.txt | We study the atomic physics and the astrophysical implications of a model in which the dark matter is the analog of hydrogen in a secluded sector. The self-interactions between dark matter particles include both elastic scatterings as well as inelastic processes due to a hyperfine transition. The self-interaction cross sections are computed by numerically solving the coupled Schr\"{o}dinger equations for this system. We show that these self-interactions exhibit the right velocity dependence to explain the low dark matter density cores seen in small galaxies while being consistent with all constraints from observations of clusters of galaxies. For a viable solution, the dark hydrogen mass has to be in 10--100 GeV range and the dark fine-structure constant has to be larger than 0.01. This range of model parameters requires the existence of a dark matter--anti-matter asymmetry in the early universe to set the relic abundance of dark matter. For this range of parameters, we show that significant cooling losses may occur due to inelastic excitations to the hyperfine state and subsequent decays, with implications for the evolution of low-mass halos and the early growth of supermassive black holes. Cooling from excitations to higher $n$ levels of dark hydrogen and subsequent decays is possible at the cluster scale, with a strong dependence on halo mass. Finally, we show that the minimum halo mass is in the range of $10^{3.5}$ to $10^7 M_\odot$ for the viable regions of parameter space, significantly larger than the typical predictions for weakly interacting dark matter models. This pattern of observables in cosmological structure formation is unique to this model, making it possible to rule in or rule out hidden sector hydrogen as a viable dark matter model. | While the standard $\Lambda$CDM cosmological model with collisionless, cold dark matter (CDM) is successful in explaining the observed large-scale structure in the Universe, there are many puzzles on galactic scales yet to be explained convincingly by a CDM-based scenario. $N$-body simulations of CDM structure growth predict cuspy radial density profiles ($\rho\sim r^{-1}$) with high central densities~\cite{Navarro:1995iw,Bullock:1999he,Diemand:2007qr,Stadel:2008pn}; however, observations of rotation curves reveal cored ($\rho\sim$ constant) or otherwise low-density inner regions in dark matter dominated galaxies, from dwarfs to low surface brightness (LSB) galaxies ~\cite{Moore:1994yx,Persic:1995ru,deBlok:1996jib,deBlok:2001hbg,Swaters:2002rx,Simon:2004sr,Spekkens:2005ik,KuziodeNaray:2007qi,deBlok:2008wp,Donato:2009ab,Oh:2010ea,Adams:2014bda}. Galaxy clusters also show evidence for a deficit of dark matter within the effective stellar radius of the central galaxy, with the mass profile outside being consistent with CDM predictions~\cite{Newman:2009qm, Newman:2011ip, Newman:2012nw}. The dark matter density profiles in dwarf spheroidal galaxies (dSphs) are a subject of current debate, with various studies finding that the stellar data for various dSphs is most consistent with a core~\cite{Walker:2011zu,Amorisco:2011hb,Jardel:2011yh,Amorisco:2012rd,Amorisco:2013uwa}, or a cusp~\cite{Jardel:2012am,Richardson:2014mra}, or is unconstrained~\cite{Breddels:2013qqh}. However, it seems clear that these galaxies are less dense than expected in pure CDM models~\cite{BoylanKolchin:2011de,BoylanKolchin:2011dk,Papastergis:2014aba}. The inclusion of supernovae and/or black hole feedback processes in cosmological simulations may ameliorate these anomalies in dwarf galaxies by significantly altering the central gravitational potential~\cite{Governato:2012fa,Brooks:2012vi,Maccio:2006nu,Stadel:2008pn,Oh:2010mc,Onorbe:2015ija,Katz:2016hyb,Read:2015sta,Wetzel:2016wro}. However, it is unclear whether such effects are simultaneously able to affect the halo structure to the extent observed in low-mass ($M_\ast \sim10^6$--$10^7 \msolar$) isolated dwarf galaxies~\cite{Ferrero:2011au,Weinberg:2013aya,Papastergis:2015tg} and low surface brightness galaxies~\cite{deNaray:2011hy}. It is also not clear if the diversity of cores observed inferred from rotation curves can be explained in the context of these models \cite{deNaray:2009xj,Oman:2015xda,Pace:2016oim}. Self-interacting dark matter (SIDM) is an attractive solution~\cite{Spergel:1999mh} to these anomalies that works across the range of mass scales under consideration, from dwarf galaxies to galaxy clusters~\cite{Rocha:2012jg}. In such a scenario, scatterings between dark matter particles allow for energy to be transferred from the hotter outer regions of the halo into the colder innermost regions. SIDM halos thus have hotter cores with higher velocity dispersions than the cold, cuspy interiors of collisionless dark matter halos, which lack a mechanism to heat the inner cusp into a core. $N$-body simulations~\cite{Vogelsberger:2012ku,Rocha:2012jg} and analytic models based on these simulations find that the aforementioned issues in small-scale structure (on the dSph and LSB scales) may be alleviated if the dark matter is strongly self-interacting with a hard-sphere scattering cross section of $0.5~\cm^2/\g \lesssim \sigma /m \lesssim 5.0~\cm^2/\g$~\cite{Zavala:2012us,Elbert:2014bma,Kaplinghat:2015aga,Vogelsberger:2015gpr}, where $m$ is the mass of the dark matter. In order to produce cores of radius $10$--$50$ kpc in cluster-sized halos in which the relative dark matter particle velocity is $v\sim 1000~\km/\s$, the required cross section is $\sigma\sim 0.1~\cm^2/\g$~\cite{Kaplinghat:2015aga}. We are thus motivated to consider SIDM models with \textit{velocity-dependent} cross sections that are suppressed at cluster-scale velocities. Upper limits on the SIDM cross section may also be derived through the observed ellipticities of cluster-scale halos ($\sigma/m \lesssim 1~\cm^2/\g$)~\cite{Peter:2012jh} and the measured center-of-mass offsets in merging cluster systems ($\sigma/m<0.47~\cm^2/\g$)~\cite{Harvey:2015hha}, but we find both these constraints to be weaker than those obtained from measurements of the inner density profiles of galaxy clusters~\cite{Kaplinghat:2015aga}. Atomic dark matter~\cite{Goldberg:1986nk, Mohapatra:2001sx, Feng:2009mn, Kaplan:2009de, Behbahani:2010xa, CyrRacine:2012fz, Cline:2013pca,Foot:2014uba}, in which the features of Standard Model (SM) hydrogen are copied to a dark sector, has all the features required of an SIDM model to solve the small-scale puzzles. We consider a dark proton and dark electron, which are charged under an unbroken U(1) gauge group and may combine to form dark hydrogen. If the formation of dark hydrogen bound states is efficient, these dark atoms constitute approximately all of the dark matter. Since the dark hydrogen is a composite particle with an extended, finite size, its self-interaction cross section can be naturally large, as required for SIDM. In this work, we consider atomic dark matter that exists today exclusively in bound states---dark recombination was fully complete. This model has uniquely testable phenomenology due to the ability of atomic dark matter to dissipate energy. We calculate and explore the cosmological consequences of both collisional scattering (which transfers energy between between dark atoms) and inelastic hyperfine upscattering (which results in energy loss through excitations and subsequent decays). These cross sections are velocity dependent, allowing the self-interactions to modify the halo profile to varying degrees in different astrophysical systems. The general trend is that the cross sections of particles in dwarf halos with characteristic velocities of $\sim 40~\km/\s$ will be larger than those in cluster halos with characteristic velocities of $\sim 1000~\km/\s$, which allows for regions of parameter space in which this model may resolve the aforementioned issues in small-scale structure. The heating rate from scatterings as well as the cooling losses from inelastic collisions vary widely depending on both the model parameters and the radial position in a halo of a given mass. The combined effects of both types of scatterings may thus lead to nontrivial effects on dark matter halo structure and evolution. At higher particle energies, additional atomic interactions, such as collisional excitations to the $n=2$ state and ionization, may begin to affect the structure of cluster-sized halos. For this interesting range of parameter space, where we see competing effects from collisional heating and cooling processes on the evolution of halos, we find additional features in the small-scale halo mass function that allow us to distinguish atomic dark matter from CDM cosmologically. Coupling between the dark matter and dark radiation produces dark acoustic oscillations, which are weakly constrained by measurements of the matter power spectrum and the cosmic microwave background (CMB)~\cite{CyrRacine:2013fsa}. The most interesting effect of acoustic oscillations in the dark plasma would be a cutoff in the matter power spectrum set by the size of the fluctuations entering the horizon before the time of matter-radiation decoupling, resulting in a minimum dark matter halo mass that is significantly larger than in the typical weak-scale models without hidden sectors. Altering small-scale structure with SIDM neither assumes nor requires any interactions between dark matter and SM particles beyond gravitational interactions; thus, we take a minimal approach and seclude the dark sector from the visible sector. Atomic dark matter has been presented in other contexts, such as a mirror universe~\cite{Mohapatra:2001sx,Foot:2014mia} and supersymmetry~\cite{Behbahani:2010xa}. The effect of hidden sector dissipative dark matter on small-scale structure has been previously studied in Refs.~\cite{Foot:2014uba, Foot:2014mia}. We note that our approach differs from prior works~\cite{Foot:2014uba, Foot:2014mia}: the cored profiles in this work result from the collisional scatterings of the neutral dark atoms, whereas the density profiles in Ref.~\cite{Foot:2014uba} are shaped by a combination of bremsstrahlung cooling processes in the dark sector as well as energy injection from visible supernovae [which is made possible through the inclusion of a kinetic mixing interaction between the dark U(1) and the SM hypercharge]. Kinetic mixing has also been used to explain DAMA and CoGeNT~\cite{Kaplan:2009de,Kaplan:2011yj,Cline:2012ei,Cline:2012is,Foot:2014uba} and the $3.5~\keV$ line~\cite{Cline:2014eaa}. The paper is structured as follows. In Sec.~\ref{sec:model} we describe the atomic dark matter model and the scattering properties of dark hydrogen. In Sec.~\ref{sec:apps} we consider dark hydrogen as an SIDM candidate and determine the parameter space allowed to accommodate SIDM and satisfy cosmological constraints. In Sec.~\ref{sec:inelastic} we investigate how inelastic scattering processes can affect halo formation. In Sec.~\ref{sec:clusters} we discuss the possibility of upscatterings to the $n=2$ excited state as well as collisional ionization in cluster-scale halos. We conclude in Sec.~\ref{sec:conclusions}. | \label{sec:conclusions} In this paper, we have investigated a model of self-interacting dark matter that mimics the properties of atomic hydrogen. Dark matter in the late universe takes the form of dark hydrogen, which is neutral under a new U(1) gauge force. We do not assume a specific interaction between this new U(1) and the SM for the predictions in this paper. The key features of our work are the inclusion of a hyperfine interaction, which induces an energy splitting in the ground state of dark hydrogen, and the calculation of the basic heat transport properties in halos, which allows us to identify the viable regions of parameter space where the small-scale puzzles can be solved. Collisions of dark atoms in halos may induce hyperfine excitations, which then decay by emitting dark photons. Halo cooling from this upscattering and subsequent energy loss works against halo heating that occurs from the scattering processes. To study these effects on halo structure, we calculated the cross sections for dark hydrogen scattering over a wide range of parameter space, using techniques from standard hydrogen to aid in numerically solving the Schr\"{o}dinger equation. The velocity dependence of the cross sections allows the heating and cooling mechanisms to operate differently on scales of dwarf spheroidal galaxies ($\vrms = 40~\km/\s$) compared to scales of galaxy clusters ($\vrms = 1000~\km/\s$). We argue that the viscosity cross section where both the forward and backward scattering are suppressed is the better quantity, compared to the momentum-transfer cross section, to use when comparing to SIDM simulation results and observational constraints. The velocity dependence of the viscosity cross section shows a sharp drop for kinetic energies larger than about $0.1\; E_0 \simeq 0.1\; \alpha^2 m_H/R$ as contributions from higher partial waves become important. This allows the model to be consistent with cluster constraints. The typical cross section at $E=0.1\; E_0$ is roughly $10\; a_0^2$ and scales approximately as $E^{-1.3}$ above these energies. For kinetic energies below $0.1\; E_0$, we see a steady increase in the viscosity cross section with decreasing relative velocity, which implies that the scattering processes are very important in small halos. The viscosity cross section in this regime scales roughly as $E^{-0.4}$. We have found regions of parameter space for the atomic dark matter model in which dark matter self-interactions can explain the measured core sizes in both dwarfs and clusters, while being consistent with all other observations including cluster halo shapes. The solutions are not fine-tuned; for a hyperfine splitting that is about $10^{-4}E_0$, we find that much of the parameter space with $\chi_e < 0.01$ and dark hydrogen mass in the 10--100 GeV range is viable. In this part of parameter space, the dark matter is in atomic form and we find that cooling mechanisms are generically important for the structure of low-mass halos (masses below $10^{10} M_\odot$) but not important enough to completely disrupt these halos. An immediate consequence of this observation is that the collapse of small halos at early times will be affected by the cooling and, therefore, it is likely that the growth of the seeds of supermassive black holes will also be altered. We leave this discussion for another paper. The kinetic energy of dark matter particles in galaxy clusters is large enough to allow for additional atomic physics. We find that collisional excitations to $n=2$ and ionizations could be significant processes in galaxy clusters for $\Ehf = 10^{-5}$. For $\Ehf = 10^{-4}$, we show that the cooling rate due to these processes is subdominant to the heating rate and our predictions, which assume negligible scattering to $n=2$ and fully atomic dark hydrogen, are robust. Thus, galaxy clusters are important astrophysical laboratories for testing atomic dark matter models. The interactions between the dark matter and the light mediator in the early Universe modifies the kinetic decoupling of the dark matter. The kinetic decoupling temperature may be used to estimate the minimum halo mass in the universe. Assuming that the ratio of the hidden sector temperature to the visible photon temperature at late times is 0.6 (close to the maximum allowed by BBN constraints), we find that the range of halo minimum masses in the viable regions of parameter space are between $10^{3.5}$ and $10^7 M_\odot$. These minimum masses are smaller than the host masses of the currently observed dwarf galaxies, but much larger than the minimum masses predicted for dark matter in weak-scale theories. If the ratio of the temperatures is smaller (due to the fact that the two sectors were reheated to different temperatures and remained decoupled), then the minimum halo masses will be lower by a factor of $(\xi/0.6)^{9/5}$. In summary, we have shown that an analog of hydrogen in the hidden sector is a viable self-interacting dark matter candidate that can alleviate the small-scale structure formation puzzles, and the dissipative nature of atomic dark matter provides a phenomenologically rich foundation to make observational predictions. | 16 | 9 | 1609.03592 |
1609 | 1609.06979_arXiv.txt | {The determination of the properties of neutron stars from the underlying theory, QCD, is still an unsolved problem. This is mainly due to the difficulty to obtain reliable results for the equation of state for cold, dense QCD. As an alternative route to obtain qualitative insights, we determine the structure of a neutron star for a modified version of QCD: By replacing the gauge group SU(3) with the exceptional Lie group G$_2$, it is possible to perform lattice simulations at finite density, while still retaining neutrons. Here, results of these lattice simulations are used to determine the mass-radius relation of a neutron star for this theory. The results show that phase changes express themselves in this relation. Also, the radius of the most massive neutron stars is found to vary very little, which would make radius determinations much simpler if this would also be true in QCD.} \FullConference{34th annual International Symposium on Lattice Field Theory\\ 24-30 July 2016\\ University of Southampton, UK} \begin{document} | The properties of neutron stars \cite{Glendenning:1997wn,Steiner:2010fz,Kapusta:2006pm} have been the subject of astronomical studies and observations since decades. Nevertheless, a satisfying explanation of their properties is still out of reach. The main reason for this is that it was not yet possible to derive the equation of state governing neutron stars from the fundamental theory, i.\ e.\ QCD, at finite density and small or zero temperature \cite{Friman:2011zz}. The reason is that lattice gauge theory, the mainstay of non-perturbative QCD calculations, suffers from the sign problem \cite{Gattringer:2010zz,deForcrand:2010ys}. Alternative methods are either still not sufficiently far developed or make heavy use of modeling, which reduces their predictivity \cite{Leupold:2011zz,Buballa:2003qv,Pawlowski:2010ht,Braun:2011pp}. Turning this around, data of neutron stars could provide us with a laboratory for strong interactions in this cold and dense regime of the phase diagram \cite {Lattimer:2006xb,Glendenning:1997wn,Steiner:2010fz}. Hence, they make it possible to observe large-scale effects of non-Abelian gauge theories. Also the recent discovery of gravitational waves \cite{Abbott:2016blz} will provide data on the bulk properties, giving further constraints \cite{DelPozzo:2013ala}. Pending a solution for QCD, it would be highly interesting to know whether any generic signatures of a realistic equation of state of a non-Abelian gauge theory can manifest in the features of neutron stars. To make progress, the avenue exploited here will be to study a QCD-like theory without sign problem. To have any chance at qualitative similarity requires this theory to have fermionic baryons, especially neutrons, and have rather similar features as QCD itself. A theory with these feature is G$_2$-QCD, i.\ e.\ QCD where the gauge group SU(3) is replaced by the exceptional Lie group G$_2$ \cite{Holland:2003jy,Maas:2012wr}, see \cite{Maas:2012ts} for a review of the properties of this theory. A rough outline of the full phase diagram \cite{Maas:2012wr}, detailed spectroscopical data and properties of the structure at zero temperature and finite baryon density \cite{Wellegehausen:2013cya} are available for this theory from lattice simulations. Based on these results, we determine the mass-radius relation of a G$_2$-QCD neutron star. Because the data are for one flavor and comparatively heavy pions, it will probably not be a semi-quantitatively good guideline, but some qualitative features can be expected, and seem to be, rather robust. This will be discussed in more detail elsewhere \cite{Hajizadeh:unpublished}. There, we will also consider the impact of varying parameters like the pion mass. In total, the features obtained are not too different from the expectations from models \cite{Glendenning:1997wn,Steiner:2010fz,Kapusta:2006pm}, which is by itself promising. Most interesting is that we find evidence for different phases inside this model neutron star to surface in astronomically observable quantities. | We have provided, for the first time, a mass-radius relation for a neutron star based on the equation of state of a QCD-like theory with fermionic baryons, using lattice data as an input. The results show a number of expected features. In particular, the curves become steeper and the maximum mass increases, without exceeding the maximum value possible from general relativity. Of course, given the fact that the theory is different from QCD and that only one flavor with a rather heavy Goldstone is included, the increase to roughly 1.2 solar masses instead of the more than two possible in QCD is not entirely unexpected. However, even this could already hint that either the masses of the quarks or the number of degrees of freedom play a significant role for the maximum neutron star mass. Less speculative and more visible in the data is that the phase change seen in the baryon density impresses itself into the mass-radius relation of the neutron star. It also survives in the pressure gradient. This implies that already observing the mass-radius relation in sufficient detail could resolve the old question of whether a neutron star is a, more or less, monolithic object, or whether it contains a multitude of different phases. These kind of qualitative insights show how important the study of QCD-like theories is for the understanding of possible features of the real case. Of course, the results here could be improved in various ways, some of which we will investigate in \cite{Hajizadeh:unpublished}, while others will have to await improved results for the input equation of state. Still, this gives eventually the hope that we could learn how to interpret qualitative features of neutron stars from such model systems. | 16 | 9 | 1609.06979 |
1609 | 1609.01304_arXiv.txt | By means of the fossil record method implemented through Pipe3D, we reconstruct the global and radial stellar mass growth histories (MGHs) of an unprecedentedly large sample of galaxies, ranging from dwarf to giant objects, from the ``Mapping Nearby Galaxies at the Apache Point Observatory" survey. We confirm that the main driver of the global MGHs is mass, with more massive galaxies assembling their masses earlier (downsizing), though for a given mass, the global MGHs segregate by color, specific star formation rate (sSFR), and morphological type. From the inferred radial mean MGHs, we find that at the late evolutionary stages (or for fractions of assembled mass larger than $\sim 80\%$), the innermost regions formed stars on average earlier than the outermost ones (inside-out). At earlier epochs, when the age resolution of the method becomes poor, the mass assembly seems to be spatially homogeneous or even in the outside-in mode, specially for the red/quiescent/early-type galaxies. The innermost MGHs are in general more regular (less scatter around the mean) than the outermost ones. For dwarf and low-mass galaxies, we do not find evidence of an outside-in formation mode; instead their radial MGHs are very diverse most of the time, with periods of outside- in and inside-out modes (or strong radial migration), suggesting this an episodic SF history. Blue/star-forming/late-type galaxies present on average a significantly more pronounced inside-out formation mode than red/quiescent/early-type galaxies, independently of mass. We discuss our results in the light of the processes of galaxy formation, quenching, and radial migration. We discus also on the uncertainties and biases of the fossil record method and how they could affect our results. | The study of how galaxies did assemble spatially their stellar masses is of paramount relevance for understanding the overall picture of galaxy evolution. {A first inference on this can be provided by the mean mass- or luminosity-weighted age gradients obtained from line-strength indices and photometric studies of local galaxies. A more complete analysis is provided by the study of the color-magnitude diagram from resolved stars in nearby galaxies of by the fossil record method using Integral Field Spectroscopy observations (IFS; see below). Several studies have shown that a significant fraction of the local galaxies present negative mean age gradients \citep[for recent works see e.g.,][]{Wang:2011aa,Lin:2013aa,Li:2015aa,Dale+2016}, which can be interpreted as earlier star formation (SF) of the inner regions with respect to the outer ones; this could imply that these galaxies assembled their stellar masses from inside to out.} Other kind of studies based on look back observations, suggest rather an uniform radial mass growth at early/intermediate redshifts \citep[][]{vanDokkum+2013,Patel+2013}. More recently, \citealp[][]{Perez+2013} have reported that the differences in the inner-to-outer stellar mass assembly histories of local galaxies studied with IFS depends on the total galaxy stellar mass; for their less massive galaxies, the trend even inverts suggesting an outside-in stellar formation mode. Some studies of dwarf galaxies report indeed positive age gradients for them \citep{Bernard+2007,Gallart:2008aa,Zhang:2012aa}. The radial way in which galaxies assemble their mass (or quench their SF) seems to depend also on their morphological type. For instance, for a small sample of early-type galaxies, \citet[][]{Sanchez-Blazquez:2007aa} have concluded that a solely inside-out or outside-in scenario is ruled out for these galaxies \citep[see also ][]{Mehlert:2003aa,Kuntschner:2010aa}. Summarizing, it is not yet well established how did galaxies on average assemble radially their stellar masses or quench their SF. The aim of this work is to provide a significant contribution on this question. According to the current cosmological paradigm, galaxies form from gas that cools and falls into the center of dark matter halos. Therefore, the galaxy gas accretion and the consequent SFR rate are expected to be associated to the cosmological dark matter accretion rate \citep[see e.g.,][]{vandenBosch2002,Faucher-Giguere+2011,Rodriguez-Puebla+2015}. On the other hand, the disks formed inside growing dark matter halos are predicted to assemble from the inside out \citep[][and more references therein]{Gunn1982,Silk1987,Mo+1998,Avila-Reese+2000,Roskar+2008,Mo+2010}. The formation of spheroids is believed to occur mostly from the morphological transformation of disks, either by mergers or by internal dynamical processes \citep[see for reviews e.g.,][]{Mo+2010,Brooks+2016}. During early gas-rich mergers, strong bursts of SF {that consume most of the gas are expected to happen, resembling this the so-called ``monolithic collapse'' model \citep{Eggen+1962}; late major mergers, responsible also of spheroid formation, add an spatially extended population of \textit{ex situ} (likely old) stars \citep[e.g.,][]{Hopkins+2009,Rodriguez-Gomez+2016}, and produce a significant stellar radial mixing in the primary galaxy}. Therefore, the spatial mass assembly of spheroid-dominated galaxies is expected to depart from the inside-out mode of disk galaxies. Moreover, galaxies along their evolution may suffer several phases of disk destruction and rebuilding, leaving this a complex imprint in their { present-day stellar population spatial distributions}. The other aspects that may play a relevant role in assessing the observed { stellar population spatial distribution} of galaxies are when, where, and why the SF is shut off in galaxies. This shut off process is commonly called in the literature as "quenching". According to the literature, there is not a general mechanism that explains the quenching of galaxies; rather exist multiple hypotheses that try to explain the shut off as a function of mass and environment \citep[e.g.,][]{Lin:1983aa,Dekel:1986aa,Efstathiou:1992aa,Birnboim:2003aa,Di-Matteo:2005aa,Slater:2014aa,Tal:2014aa}. Ram-pressure, strangulation, and harassment are some of the mechanisms related to environmental conditions that can explain the shutoff of SF in satellite galaxies \citep[e.g.,][]{Gunn:1972aa,Larson:1980aa,Farouki:1981aa,Hidalgo+2003,Tal:2014aa}. Mass dependent processes related to the heating up of the intrahalo gas can help to stop the SF of central galaxies \citep[e.g.,][]{Guo:2014aa,Tal:2014aa}; for example, the virial shock heating in massive halos \citep[the halo quenching model,][]{Birnboim:2003aa,Dekel:2006aa} and the AGN feedback \citep[e.g.,][]{Kauffmann:2004aa,Bower:2006aa,Guo:2014aa}, in particular the powerful outflows in quasars that deplete the gas content in massive galaxies \citep{Sanders:1988aa,Di-Matteo:2005aa}. On the other hand, since the SFR of galaxies is expected to primarily depend on the halo dark matter accretion rate (see above), then if the latter significantly drops, then the former shoould drop too \citep[cosmological quenching; e.g.,][]{vandenBosch2002,Feldmann:2015aa}. Whatever the dominant process is, the form of how the SF can be deactivated correlates with the galaxy mass \citep[e.g.,][]{Woo:2013aa}. In order to constrain the scenarios and processes of galaxy evolution above described, it is crucial to obtain information about how the stellar mass grows during the evolution of the galaxy as a function of radius. Considering that the light of a galaxy is an assembly of multiple star contributions, it is possible to infer how these stars assembled the observed galaxy light through its life. This method uses stellar evolutionary models to get the energy spectra distribution (SED) of single stellar populations (SSP) at different ages and metallicities. Thus, the problem is reduced to look for the best combination of SSPs that mimics the observed galaxy SED. Then, it is possible to obtain information of the stellar mass per SSP at different time steps. As a result, it is feasible to reconstruct the SF or stellar mass growth history (SFH or MGH) of an observed galaxy. This approach based in stellar population synthesis (SPS) is known in the literature as the \textit{fossil record method} \citep[e.g.,][]{Tinsley:1980aa,Buzzoni:1989aa,Bruzual:2003aa,Walcher:2011aa}. The fossil record method has been widely used to recover the global SFH and stellar masses by using spectral and photometric data of galaxies in large surveys like the Sloan Digital Sky Survey (SDSS) one \citep[e.g.,][]{Kauffmann:2003aa,Kauffmann:2003ab,Cid-Fernandes:2005aa,Gallazzi:2005aa,Tojeiro+2007}. With multi-band and spectroscopic data of good spatial resolution, it is possible to resolve the SFH (or MGH) in space \citep[e.g.,][]{Brinchmann:2000aa,Kong:2000aa,Perez-Gonzalez:2008aa,Lin:2013aa}. These studies used to be limited by the size of the slit or fiber (for the spectroscopy data) or by the number of photometric filters that are used to adjust the SED. These limitations can retrieve important bias in the interpretation of the data. A fixed fiber (or slit) size cannot locally resolve the galaxy properties because the observed galaxy spectrum integrates all or a substantial fraction of the light within the aperture. For the case of photometric studies, it is required a large enough number of photometric bands to reproduce the galaxy spectrum \citep[e.g.,][]{Benitez:2009aa}. It is important to correct the observed photometric fluxes by emission lines to accurate fit the observed galaxy SED. However, with the advent of the IFS, % it is possible to access to spectral data resolved through different galaxy regions at the same time and obtain the spectral information of the local galaxy properties \citep[e.g.,][]{Bacon:2001aa,Cappellari+2011,Croom:2012aa,Sanchez:2012aa,Cid-Fernandes:2013aa,Bundy:2015aa}. With IFS data and SSP decomposition, it is feasible to obtain information about how galaxies assembled globally and locally their stellar masses. % \citet{Perez+2013} used the advantages of the IFS to resolve the radial gradient of the stellar mass assembly. They analyzed $105$ galaxies from the Calar Alto Legacy Integral Field Area survey \citep[CALIFA,][]{Sanchez:2012aa} and explored the mass growth gradient from $10^{9.8}$ to $10^{11.26}\ \solarm$. With the MaNGA survey \citep[Mapping Nearby Galaxies at the Apache Point Observatory,][]{Bundy:2015aa}, it is possible to perform a spatially resolved study of the MGHs for an unprecedentedly large galaxy sample both in number and in mass range. The MaNGA survey is one of the three core programs of the fourth generation of the SDSS, and plans to get IFS observations for $10000$ galaxies (selected by stellar masses) within a redshift range of $0.01<z<0.18$ \citep{Bundy:2015aa,Li:2015aa,Law:2015aa}. MaNGA performs dithered observations with integral field units (IFU) that cover a spectroscopic range of $3600$ to $10300\ \AA$, with a resolution of $R\sim2000$. MaNGA provides $17$ fiber-bundle IFUs with a field of view that covers from $12''$ (19-fiber IFU) to $32''$ (137-fiber IFU). For a more detailed description of the MaNGA survey and its instrumentation, see \citet{Gunn:2006aa}, \citet{Bundy:2015aa}, \citet{Drory:2015aa}, and \citet{Yan:2016aa}. Some Early results with MaNGA are the study of \citet{Li:2015aa} that maps the gradients of stellar population and star formation indicators D4000 and Halpha covering a large wavelength range.\citet{Wilkinson:2015aa} maps the gradients of stellar population parameters and dust attenuation, discussion of impact of different fitting techniques and IFU size. Finally, \citet{Belfiore:2015aa} performs an spatially resolved study of emission line diagnostic diagrams, maps of gas metallicities and element ratios. {In this paper we present the global and spatially-resolved normalized MGHs of the so far observed galaxies in the MaNGA survey, namely those to be reported in the SDSS Data Release 13. We focus our analysis on galaxies from the Primary sample, covering a mass range from dwarf to giant galaxies, and with numbers that overcome those of the few previous related works. Our goal is to explore how galaxies did assemble globally and locally their stellar masses as a function of mass. The large number of galaxies allows us to study this question also as a function of color, specific SFR (sSFR), and morphology, by separating the galaxies in each mass bin at least in two groups according to these properties. It should be said also that our analysis differs in several aspects to those presented in previous works, for instance, those of \citet{Perez+2013} and \citep[][]{Gonzalez-Delgado+2016}.} The outline of the paper is as follows. In Section \ref{analysys}, we explain the sample selection, its analysis and the methodology. In Section \ref{results-meanMGHs}, we discuss both the global and radially-resolved mean stellar MGHs as a function of the final galaxy mass, color, SF activity, and morphology. In Section \ref{inMGH}, we recover the individual radial mass assembly gradient for the final sample of galaxies at the times when the $90\%$, $70\%$ and $50\%$ of their total masses were assembled. In Section \ref{discussion}, we compare our results with previous works, we discuss the possible implications of our inferences for the galaxy evolution paradigm, and comment on the caveats of these inferences. Finally, in Section \ref{conclusions}, we present the summary of the paper and our main conclusions. | \label{conclusions} We have inferred the global and spatially-resolved stellar MGHs of the first set of galaxies observed in the MaNGA/SDSS-IV survey (to appear in the SDSS Data Release 13) by means of the fossil record method using Pipe3D. We used only the Primary sample and excluded from our analysis the strongly interacting/merging galaxies as well as those from the Color-Enhanced MaNGA sample. Then, a sub-sample of 454 galaxies at redshifts lower than 0.037 was studied in particular in order to obtain normalized MGHs with the final mass $\mathcal{M}_0$ defined at the \textit{same} redshift for all galaxies ($z_{\rm lim}=0.037$ corresponding to an initial LBT of $\approx 500$ Myr). The sample was divided into four mass bins and, for each mass bin, the mean normalized MGHs at different radial regions (up to 1.5$R_{50}$) and their scatter were calculated. We have also calculated the mean normalized MGHs for galaxies separated into blue/red colors, star forming/quiescent, and late/early types. Moreover, the differences between the assembly times of a given mass fraction of the innermost and outermost regions, $\Delta T_{\rm i-o}$, was presented for each individual galaxy and for different values of the attained mass fraction. The main goal of the present work was to probe the general way galaxies assembled their stellar masses, both globally and at different radial regions, as a function of mass and other galaxy properties. Our study is aimed to explore only general trends; a more detailed separation of galaxies according to both their masses and properties, as well as their environment, will be presented in forthcoming papers. The main results from our analysis are as follow. \begin{itemize} \item The larger the final galaxy mass, the earlier on average was assembled most of this mass. For example, galaxies in the $\log$($\mathcal{M}_0/ \solarm$) = 8.5--9.3, 9.3--10.0, 10--10.7 and 10.71--11.2 bins assembled 70\% of their final masses at LBTs of 3.7, 10, 11, and 11 Gyr on average. Dwarf and low-mass galaxies present actually a large diversity of global stellar MGH shapes, showing that their stellar mass assembly is not only delayed with respect to more massive galaxies but also more episodic. On the other hand, in a given mass bin, red/quiescent/early-type galaxies assemble on average their masses earlier than blue/star forming/late-type ones. However, these differences are not so large as those among the different mass bins. \item {At late evolutionary stages (or high fractions of assembled mass), most of galaxies in all the mass bins show that the innermost regions formed stars earlier than the outermost ones (inside-out formation mode).} For instance, in the $\log$($\mathcal{M}_0/\solarm$) = 8.5--9.3, 9.3--10.0, 10--10.7, and 10.7--11.2 bins, the regions inside 0.5 $R_{50}$ attain 90\% of their final masses on average 0.5, 0.5, 1.5, and 1.0 Gyr before than regions outside 1 $R_{50}$ do it, respectively. For galaxies more massive than $\sim 10^{10} \solarm$, the inside-out trend in the mean radial MGHs remains until LBTs of $9-10$ Gyr, when 70--80\% of the inner and outer regions were in place. At earlier stages (lower assembled mass fractions), this trend tends to disappear; the inner and outer MGH differences become actually diverse {\bf and within the large age uncertainty of the method}. In the case of dwarf galaxies, the radial MGHs are very diverse most of the time, showing periods of outside-in and inside-out growth modes (or strong radial migration). % \item {The outermost regions of galaxies present on average less regular MGHs (more scattered around the mean) than the innermost ones, being this trend more pronounced as more massive are the galaxies. } \item The way galaxies assemble their mass radially depends more on the galaxy color/sSFR/type than on its mass: blue/star-forming/late-type galaxies follow on average a significantly more pronounced inside-out formation mode than red/quiescent/early-type galaxies. {This is true specially for the late evolutionary stages; for the old stellar populations that are more common in the later galaxies, the divergent solutions of the fossil record methods at the earliest epochs of mass assembly, does not allow to establish a clear difference between the inner and outer MGHs of these galaxies}. For these galaxies, the outermost MGHs present also a large diversity of shapes, in many cases with signs of {older assembly than the innermost regions, specially for the small galaxies.} \item {For galaxies more massive than $\sim 5\times 10^{9} \solarm$, the age differences $\Delta T_{\rm i-o}$ at epochs when $50\%$ of the innermost/outermost masses were formed are mostly below the age uncertainty. The situation improves at later evolutionary stages, when larger mass fractions form. Our results suggest a common trend of transiting from a flat/outside-in (or undetermined) gradient mode to an inside-out mode; this trend is more frequent for the more massive galaxies. Galaxies less massive than $\sim 5\times 10^{9} \solarm$ present large differences between the formation epochs of the innermost and outermost regions at early epochs, both in the direction of the inside-out or outside-in mode; these differences decrease at later epochs. } \end{itemize} From the general trends obtained in our analysis, we can conclude that mass is the main driver of the \textit{global} stellar MGHs, while the way mass is \textit{radially} assembled is more related to the galaxy color/sSFR/type. The trend is that more massive galaxies form earlier than less massive galaxies (downsizing), and galaxies with radial MGHs more regular (less scattered around the mean) and with a stronger inside-out mode are more related to bluer/star-forming/later-type objects. The stellar mass assembly of galaxies less massive than $10^{10} \solarm$, specially the least massive ones, reveals itself as episodic and stochastic, both at the global and local level, with periods of outside-in and inside-out formation modes. However, as discussed in \S\S \ref{migration} this diversity of radial MGHs could be also due to the expected strong radial migration processes in these galaxies, where the older stellar populations suffer more net outward migration. {After taking into account such an outward migration, the intrinsic radial MGHs could follow instead an inside-out growth mode. } For $\mathcal{M}_0>10^{10} \solarm$, the marked inside-out formation mode seen for most of the blue/star-forming/late-type galaxies, at least since $\sim 10$ Gyr ago, is in agreement with predictions of inside-out disk formation in the context of the hierarchical $\Lambda$ cold dark matter cosmology. {If takes into account a possible net outward flow of stars due to migration mechanisms, a more pronounced intrinsic inside-out mode would be inferred.} % In the case of red/quiescent/early-type galaxies, {due to the poor age resolution when the ages are old, we can not say too much about the formation mode of these early-forming galaxies. However, at later epochs, when more than $70-80\%$ of the mass was assembled, they seem to follow on average a weak inside-out formation mode,} which could be produced by quenching of a powerful AGN, which shuts-off the SF first in the central regions and then in the outer ones; given the relatively small differences between the innermost and outermost MGHs (or relatively small values of $\Delta T_{\rm i-o}$), this effect is not expected to long not too much ($<<1$ Gyr). On the other hand, the weak hint of early outside-in formation in most of the red/quiescent/early-type galaxies may be incorrect, and the fact that very old stellar populations are present in the outermost regions of these galaxies could be rather due to the accretion of satellites with old populations (late dry mergers). {It should be said that the main conclusions from our study remain the same when including galaxies from the MaNGA Secondary sub-sample, i.e., those with a radial coverage that reaches 2.5 effective radii. We have not included these galaxies in order to avoid possible aperture biases, and because the mass dependence on redshift is not trivial for this sub-sample and it may introduce selection effects.} The fossil record method applied to IFS observations is a unique technique to infer the local and global evolution of \textit{individual} galaxies. However, we should be borne in mind that these inferences are subject to several uncertainties and degeneracies. For example, the age uncertainty in the spectrum inversion is large, specially for the contributions coming from old stellar populations. In the present work the inferred global MGHs are biased to high mass fractions assembled at the earliest epochs with respect to inferences based on look back empirical and semi-empirical studies or even with respect to some previous archaeological inferences (see \S\S \ref{comparison}). As discussed in \S\S \ref{caveats}, for intermediate/large LBTs, because of the low age resolution the stellar MGHs can be systematically modified, specially toward older ages. The choice of the SPS models also affects significantly the cumulative SFHs or MGHs; for instance, according how the thermally pulsating asymptotic giant branch stars are treated, shifts of some Gyrs in the MGHs are obtained. More work is necessary to evaluate and control the effects of uncertainties and degeneracies in the fossil record inferences of galaxy stellar mass assembly. In defense of the general results presented here, we should say that, once the same method and analysis is applied to all galaxies and radial regions, the results related to comparisons of the means among different masses and among different radial regions should be reliable at least in a qualitative level. | 16 | 9 | 1609.01304 |
1609 | 1609.03906_arXiv.txt | {We present FORS2 (attached to ESO's Very Large Telescope) observations of the exoplanet WASP-17b during its primary transit, for the purpose of differential spectrophotometry analysis. We use the instrument in its Mask eXchange Unit (MXU) mode to simultaneously obtain low resolution spectra of the planet hosting star, as well as several reference stars in the field of view. The integration of these spectra within broadband and smaller 100\AA~bins provides us with ‘white' and spectrophotometric light curves, from 5700 to 8000\AA. Through modelling the white light curve, we obtain refined bulk and transit parameters of the planet, as well as wavelength-dependent variations of the planetary radius from smaller spectral bins through which the transmission spectrum is obtained. The inference of transit parameters, as well as the noise statistics, is performed using a Gaussian Process model. We achieve a typical precision in the transit depth of a few hundred parts per million from various transit light curves. From the transmission spectra we rule out a flat spectrum at >3$\sigma$ and detect marginal presence of the pressure--broadened sodium wings. Furthermore, we detect the wing of the potassium absorption line in the upper atmosphere of the planet with 3-$\sigma$ confidence, both facts pointing to a relatively shallow temperature gradient of the atmosphere. These conclusions are mostly consistent with previous studies of this exo--atmosphere, although previous potassium measurements have been inconclusive.} | In the short number of years in which extrasolar planetary systems have been detected and studied, there has been great progress made in understanding these alien worlds and characterising their intrinsic properties. Methods such as radial velocity and transit monitoring, have facilitated the measurement of bulk densities, as well as some initial approximations of the structure and atmospheric properties of extrasolar planets. It is the latter with which we are concerned here, whereby spectroscopic observations of exoplanetary primary transits lead to setting constraints on the chemical composition and physical processes within exo-atmospheres. Through transmission spectroscopy one analyses the wavelength dependent variations of the exoplanetary radius. Such minute variations are caused by the presence of an optically thick atmosphere at specific wavelengths, dictated by the absorption and scattering characteristics of the gas and aerosols present near the planet's terminator. These wavelength dependent opacity variations are scanned across the spectrum, and used to probe chemical compositions of the planets' atmospheres \citep{Seager2000,Brown2001}. The ideal candidates for such studies are those planets with extended atmospheres, i.e. with large atmospheric scale heights, such as the so-called hot Jupiters \citep{Seager1998}. The search for signatures of exo-atmospheres began as early as the time when predictions were being made. For instance \cite{Rauer2000} and \cite{Harris2012}, among others, looked for spectral signatures due to absorption by extended atmospheres surrounding 51 Peg b and $\tau$ Bo\"otis Ab. Traditionally, transmission spectroscopy has been dominated by space-based facilities, due to their obvious advantage of not being affected by telluric extinction \citep{Charbonneau2002,Ehrenreich2007,Pont2008,Gibson2012a,Deming2013,Knutson2014,Swain2014}. However, large ground-based telescopes, with instruments meeting the requirements of transmission spectroscopy science goals, VLT/FORS \citep{Appenzeller1998}, Gemini/GMOS \citep{Hook2004} or Magellan/IMACS \citep{Bigelow1998} being the most effective of those, have also been able to contribute to such studies. \cite{Gibson2013a,Gibson2013b} among others have shown the capabilities of GMOS in producing transmission spectra for a number of transiting hot Jupiters. With its MOS (Multi Object Spectroscopy) and MXU (Mask eXchange Unit) modes, the FORS2 (FOcal Reducer and low dispersion Spectrograph) instrument on the Unit Telescope 1 (UT1 -- Antu) of ESO's Very Large Telescope (VLT) has also been a key player in producing such spectra. \cite{Bean2010,Bean2011} were able to obtain pioneering results for the GJ1214b transmission spectrum, an exoplanet in the mini Neptune regime, from the visible to NIR bands. However, subsequent attempts at transmission spectroscopy analysis with this instrument were mostly unsuccessful due to systematics present in the data associated with prisms of the Longitudinal Atmospheric Dispersion Corrector (LADC). Therefore, a project was started at ESO Paranal to address such issues \citep{Boffin2015}, by removing the inhomogeniously degrading coating from the prisms of the decommissioned FORS1's LADC, replacing them with their FORS2 counterpart. The improvements were subsequently highlighted through transit observations of WASP-19b, the results of which were presented in \citet{Sedaghati2015}. An essential issue plaguing all the aforementioned instruments, has been the role of instrumental systematics and their manifestation in the final transit light curves. As transmission spectroscopy heavily relies on extremely precise measurement of the transit depth (a few hundreds of ppm\footnote{Parts per million.} for hot Jupiters, down to a few ppm for terrestrial planets), and subsequently the relative planetary radius, consideration of such factors is crucial in determining the correct scaled radius values, as well as pragmatic estimation of their error budgets. This fact is of paramount importance when a transmission spectrum is utilised together with atmospheric models in determining the physical properties of an exo-atmosphere. The time-dependent correlated noise\footnote{Sometimes also refered to as `red' or systematic noise.} due to systematics is somewhat reduced through differential spectroscopy techniques employed in this work, however some effects are still expected to remain which originate from poorly understood sources. In this paper we utilise a method with this apparent correlated noise component in mind, to model and analyse transit light curves. The method is based on the Gaussian process (GP) of \citet{Gibson2012b}, adapted for modelling time-correlated noise, which has been shown to provide conservative and realistic error estimations \citep{Gibson2013b} as compared to other parametric methods such as the Wavelet decomposition techniques of \citet{CarterWinn2009}. This is an essential point to consider when detecting exoplanetary atmospheric features with transmission spectroscopy, since underestimating the precision of radius measurements can lead to false or inaccurate characterisation of the atmosphere. Here we report FORS2 observations of WASP-17b in MXU mode. WASP-17b \citep{Anderson2010} is an ultra low density Jupiter size planet on a $\sim$3.74d, possibly retrograde orbit around an 11.6V magnitude F6V star. It has a mass and radius of 0.486$\pm$0.032 M$_\text{Jupiter}$ and 1.991$\pm$0.081 R$_\text{Jupiter}$ respectively \citep{Anderson2011}, making it an extremely bloated hot Jupiter at only 0.06 $\rho_\text{Jupiter}$. This together with an equilibrium temperature of 1771$\pm$35 K \citep{Anderson2011}, mean a relatively large atmospheric scale height, making it an ideal candidate for transmission spectroscopy studies. This paper is structured as follows. Section 2 highlights our observations and data reduction, in Section 3 we present detailed transit data analysis steps, in Section 4 we show the results of the analysis and discuss the atmospheric characteristics of WASP-17b , in Section 5 we briefly discuss our atmospheric results for this exoplanet and finally in Section 6 a brief summary of our conclusions is presented. | In this work, we have presented the transmission spectrum of the hot-Jupiter WASP-17b using the FORS2 instrument at ESO's VLT, in its multi-object spectroscopy mode. Using a combination of light from multiple comparison stars, we obtained the broadband differential transit light curve of this planet, whereby the bulk orbital and physical parameters were derived. From the broadband light curve, we obtained refined non-wavelength dependent transit parameters, consistent with previous analyses \citep{Anderson2011,Southworth2012,Bento2014}. Through detailed study of parameter inference and correlations, we modelled all the obtained light curves using the quadratic limb darkening law for the host star. Spectrophotometric light curves are analysed as independent GP noise models together with an analytical transit function, where strict, delta--function, priors for non-wavelength dependent parameters were assumed based on the broadband solution. We explored the posterior distribution for the remaining free parameters, those being the scaled radius, the two coefficients of the limb darkening law, noise model parameters and three coefficients of a second order polynomial describing the colour--dependent out of transit flux variations, to quote the best fit solutions as a function of channel wavelength. We take a non-parametric approach to modelling the time--correlated noise in the data, with time taken as the only input of our GP model \citep{Gibson2012b} in calculation of the covariance matrix. This procedure was performed on two sets of light curves, where for the second set (CMC), we applied what is known as the common mode correction by removing the systematic noise common to all the wavelength channels. Through comparison of transmission spectra (produced from both sets of light curves) with synthetic atmospheric models, we rule out a cloudy makeup of WASP-17b's upper atmosphere with high significance (>3$\sigma$). From fitting a Rayleigh scattering slope we estimate an atmosphere with a mean molecular weight consistent with prevalence of H$_2$. Further observations with the 600B grism of FORS2, extending the spectrum towards the ultraviolet, will be required to confirm and quantify this aspect of the exo-atmosphere with higher precision. Additionally we looked closer at possibility of enhanced absorption towards the two main optical absorbers, sodium and potassium. We do not detect a significant variation of the planetary radius at the sodium core with a 50\AA~bin, consistent with previous conclusions of \cite{Wood2011} and \cite{Zhou2012}. Due to the low significance levels we not able to confirm nor rule out the presence of the pressure--broadened wings of the sodium absorption line. Further, higher precision and resolution observations will be required to confirm this feature. Similar conclusions are made for potassium, although we were not able to probe the absorption in the core of the line due to the telluric $O_2$ feature. However, we do confirm the presence of the pressure--broadened wing of the potassium line with 3-$\sigma$ significance, which amounts to a significant detection. Ultimately, our observations and analysis highlight the importance and capability of ground-based facilities in detecting and characterising exoplanetary atmospheres. FORS2 will play an important role in those efforts, providing the wider community an essential outlet for followup of fascinating current and future targets. | 16 | 9 | 1609.03906 |
1609 | 1609.00134_arXiv.txt | {} {The magnetic field of coronal mass ejections (CMEs) determines their structure, evolution, and energetics, as well as their geoeffectiveness. However, we currently lack routine diagnostics of the near-Sun CME magnetic field, which is crucial for determining the subsequent evolution of CMEs.} {We recently presented a method to infer the near-Sun magnetic field magnitude of CMEs and then extrapolate it to 1 AU. This method uses relatively easy to deduce observational estimates of the magnetic helicity in CME-source regions along with geometrical CME fits enabled by coronagraph observations. We hereby perform a parametric study of this method aiming to assess its robustness. We use statistics of active region (AR) helicities and CME geometrical parameters to determine a matrix of plausible near-Sun CME magnetic field magnitudes. In addition, we extrapolate this matrix to 1 AU and determine the anticipated range of CME magnetic fields at 1 AU representing the radial falloff of the magnetic field in the CME out to interplanetary (IP) space by a power law with index ${\alpha}_{B}$.} {The resulting distribution of the near-Sun (at 10 \rs) CME magnetic fields varies in the range [0.004, 0.02] G, comparable to, or higher than, a few existing observational inferences of the magnetic field in the quiescent corona at the same distance. We also find that a theoretically and observationally motivated range exists around ${\alpha}_{B}$ = -1.6 $\pm 0.2$, thereby leading to a ballpark agreement between our estimates and observationally inferred field magnitudes of magnetic clouds (MCs) at L1.} {In a statistical sense, our method provides results that are consistent with observations.} \date{Received ........ / Accepted .......} | \label{sec:intro} Knowledge of the magnetic field entrained in coronal mass ejections (CMEs) is a crucial parameter for their energetics, dynamics, structuring, and eventually of their geoeffectiveness. For instance, the overall CME energy budget is dominated by the energy stored in non-potential magnetic fields \citep[e.g.,][]{forbes2000,avour2000}. In addition, given that CMEs and interplanetary (IP) counterparts (interplanetary CMEs (ICMEs)) are magnetic configurations with a low-$\beta$ plasma parameter, their structural evolution as they propagate and expand into the IP space is dictated by the balance and interactions between their magnetic field and the ambient solar wind \citep[e.g.,][]{dem2009}. Moreover, upon arrival at 1 AU, the magnitude of the southward magnetic field of earth-directed interplanetary CMEs (ICMEs) is the most important parameter determining their geoeffectiveness \citep[e.g.,][]{wu2005}. Therefore, the near-Sun magnetic field magnitude is a key parameter for both space weather studies and applications, for example, by constraining the properties of coronal flux ropes ejected into the IP medium \citep[e.g.,][]{shiota2016}, and in anticipation of the observations of upcoming solar and heliospheric missions. Unfortunately, very few direct observational inferences of near-Sun ($\sim$ 1-7 \rs) CME magnetic fields exist currently \citep[e.g.,][]{bast2001,jensen2008,tun2013}. These are based on relatively rare radio emission configurations, such as gyrosynchrotron emission from CME cores and Faraday rotation, and require detailed physical modeling of relevant radio emission processes to infer the magnetic field that is sought after. We recently proposed a new method to deduce the near-Sun magnetic field magnitude (hereafter, magnetic field) of CMEs \citep[][]{case16}. This method relies on the conservation of magnetic helicity in cylindrical flux ropes and uses as inputs the magnetic helicity budget of the source region and geometrical parameters (length and radius) of the associated CME. It supplies an estimation of the near-Sun CME magnetic field which is then extrapolated to 1 AU using a power-law fall-off dictated by the radial (heliocentric) distance. We have successfully applied this method to a major geoeffective CME launched from the Sun on 7 March 2012, which triggered one of the most intense geomagnetic storms of solar cycle 24. Recently, two other methods to infer the CME-ICME magnetic field vectors were proposed \citep[][]{kunkel2010,savani2015}. \citet{kunkel2010} use a flux-rope CME model, driven by poloidal magnetic flux injection, which is constrained by the height-time profile of the associated CME. The \citet{savani2015} method is based on the heliospheric magnetic helicity rule, the tilt of the source active region (AR), and the magnetic field strength of the compression region around the CME. In this work we perform a parametric study of the method to assess its robustness before applying it to observed cases. We essentially use distributions of input parameters derived from observations to determine the near-Sun and 1 AU magnetic fields for a set of synthetic CMEs. This study offers statistics sufficient to determine the range of the anticipated CME magnetic fields both near-Sun and at 1 AU. The latter distribution is compared to actual magnetic-cloud (MC) observations at 1 AU. In the following, Section \ref{sec:method} describes how we infer the near-Sun CME magnetic field, while Section \ref{sec:b1au} describes how this value is extrapolated to 1 AU. Section \ref{sec:parametric} describes our parametric study, Section 5 includes some further tests and uncertainty estimations, while Section 6 summarizes our results, their limitations, and an outlook for future revisions. | \label{sec:discussion} Developing methods for the practical estimation of the magnetic field of CMEs, both near the Sun and at 1 AU, is a timely and important task for assessing the near-Sun energetics and dynamics of CMEs and for providing clues of the possible geoeffectiveness of their ICME counterparts. We recently developed one such method and we hereby perform a parametric study of it. Our study only applies to the CME magnetic-field magnitude and not its orientation, hence, reaching results pertinent to the CME geoeffectiveness requires an extension of this work. Our major conclusions are the following: \begin{enumerate} \item The predicted near-Sun CME (at 10 \rs) magnetic fields (Figure \ref{fig:near_sun}) exhibit a FWHM range of [0.004, 0.03] G and their distribution shows values that are comparable to, or higher than, magnetic fields measured in the quiescent corona by a handful of observations. For solar AR sources prone to eruptions, the FWHM of CME magnetic fields is [0.02, 0.07] G, which is clearly higher than the quiescent-corona magnetic field at 10 \rs$\;$(Figure \ref{fig:near_sun}). \item The extrapolated CME-ICME magnetic field at 1 AU depends more sensitively on the power-law index ${\alpha}_{B}$ of its radial dependence than on the near-Sun CME magnetic fields (Figure \ref{fig:b1aucolor}). \item Considering the full range of literature-suggested ${\alpha}_{B}$ values ([-2.7,-1.0]), we find that the extrapolated near-Sun magnetic fields at 1 AU do not match MC magnetic-field measurements (Figure \ref{fig:b1auall}). \item For $\alpha_{B}$ varying in the range [-1.9,-1.4], we obtain a considerable ballpark agreement with MC magnetic-field measurements at 1 AU in terms of both the similarity of the corresponding distributions and the high fraction of $B_{1AU}$ values falling within the $B_{MC}$ value range. A best-fit $\alpha_{B}$ attains a value of -1.6 (Figures \ref{fig:b1au0}, \ref{fig:bcc}). \end{enumerate} Statistically, therefore, our method is able to reproduce the ballpark of the ICME magnetic field magnitudes at 1 AU reasonably well. This result encourages us to seek further opportunities to apply the method to observed CME cases in the future. Interestingly, the best-fit $\alpha_{B}=$-1.6 stems independently from the analytical model of \citet{dem2009}, which treats CMEs as expanding force-free magnetic flux ropes in equilibrium with the total pressure of the ambient solar wind. In the following, we summarize our assumptions and simplifications that could represent areas of future method improvements. We used AR helicity values taken from \citet{tzio2012}, which were calculated via the \citet{geor2012} method. Several methods exist to calculate $H_{m}$ (Section \ref{sec:theory}). Application of these methods to the same data set, i.e., a sequence of HMI vector magnetograms, spanning over a two-day period (6 - 7) of March 2012 for the supereruptive NOAA AR 11429, showed that the $H_{m}$ determinations, even thought they stem from very different methods, show an overall agreement within a factor $\sim 2.5$ (Tziotziou et al., 2016, in preparation). In addition, while the employed $H_{m}$ values refer to entire ARs, it is known that no AR sheds its entire helicity budget in a single eruption; it is more appropriate to attribute a fraction of the helicity budget to eruptions. This fraction seems to be relatively small, typically one order of magnitude less, but can be up to 40$\%$ of the total helicity budget in some models \citep[e.g.,][]{kliem2011,morait2014}. Nevertheless, as a first-order approximation, the used AR-wide $H_{m}$ value should not dramatically overestimate the CME magnetic field in our analysis, given the large dispersion of helicity values ($\sim$3 orders of magnitude) and the large number of synthetic CMEs ($10^4$). In future works, nonetheless, it is meaningful to search for eruption-related $H_{m}$ changes that should then be attributed to the ensuing CME. As observed in coronagraph field of view, CMEs have curved fronts. In Section \ref{sec:theory} we nonetheless assumed a straight, cylindrically shaped CME front. This is because the employed flux-rope model is cylindrical. However, CMEs may flatten during IP propagation \citep[e.g.,][]{savani2010}. At any rate, adopting a curved CME-front shape most likely introduces a (small) scaling factor in the derived $B_{*}$ distributions. In the parametric model of Section \ref{sec:parametric}, we assumed that the distributions of the magnetic ($H_{m}$) and geometrical parameters ($\alpha$ and $\kappa$) are statistically independent, that is, they do not exhibit statistical correlations. This may be not entirely true. In spite of the agreement found, the FWHM range of ${\alpha}_{B}$ values still allows $\sim$30 \%$\;$ of the projected $B_{1AU}$ values to lie outside the observed $B_{MC}$ value range (Figure \ref{fig:bcc}). For example, there is a high-$B_{1AU}$ tail that is not present in the observations (Figure \ref{fig:b1au0}). This suggests that our single (and simple) power-law description of the radial evolution of CME-ICME magnetic field of Equation \ref{eq:scaleb} may require future improvement. In particular, it is possible that the CME magnetic field experiences a stronger radial decay closer to the Sun, hence a single-power law description may not be entirely realistic. In addition, CMEs-ICMEs could experience magnetic erosion during their IP travel \citep[e.g.,][]{dasso2006,ruffe2015}, and they may thus end up with reduced magnetic fields at 1 AU. Finally, while excessively high MC magnetic fields are only very rarely reported \citep[e.g.,][]{liu2014}, the lowest $B_{1AU}$ values, below the lower limit of the $B_{MC}$ distribution, may call for complex processes in the IP medium that could enforce the magnetic field of an ICME. This could be a CME-CME interaction, for example, or an interaction between a CME and trailing fast solar wind streams \citep[e.g.,][]{lugaz2008,shen2011,harrison2012,temmer2012,liu2014,lugaz2014}. Clearly, more detailed analysis is required to tackle the above issues. Such a major task would be the application of the method to a set of carefully selected, well-observed CME-ICME cases, where both photospheric coverage would be sufficient and detailed GCS modeling would exist, along with satisfactory MC measurements at L1. This exercise would also aim to attribute eruption-related helicity changes to CMEs, hence tackling issues (1), (2) above. In addition, it would enable one to determine whether magnetic and geometrical parameters are correlated, therefore addressing issue (3) above. Further theoretical and modeling work is required to understand the radial evolution of CME-ICME magnetic fields, hence tackling issue (4) above. One such avenue would be to to analyze simulations of CME propagation in the IP medium and monitor the evolution of their magnetic fields with heliocentric distance, and at the same time investigating whether and how ${\alpha}_{B}$ depends on CME properties (e.g., speed, width), background solar wind (e.g., speed, density), and IP magnetic field. Here we treated ${\alpha}_{B}$ in a rather ad hoc manner; however, ${\alpha}_{B}$ appears as the single most important parameter for describing the ICME magnetic field at 1 AU, which apparently enables one to encapsulate most of the relevant physics into a simple form of self-similar IP expansion. That said, one should not dismiss the role of the near-Sun CME magnetic field in the determination of the ICME magnetic field at 1 AU. We also need observational inferences of this important parameter at sufficient numbers and our proposed method is one such promising avenue. Our framework may be generalized to non-force-free states \citep[e.g.,][]{hidalgo2002,chen2012,berdi2013,subra2014,case16,nieves2016} and cylindrical geometries (i.e., curved flux-ropes) \citep[e.g.,][]{janv2013,vand2015} provided that the existence of explicit relationships connect its geometrical ($R$ and $L$) and magnetic ($H_m$) parameters. Ultimately, we will perform the most meaningful tests of the two central parameters, $B_{*}$ and ${\alpha}_{B}$, and key assumptions of our model, when pristine observations by the two forthcoming, flagship heliophysics missions, Solar Orbiter and Solar Probe Plus, become available. | 16 | 9 | 1609.00134 |
1609 | 1609.02307_arXiv.txt | We propose a novel particle physics model in which vector dark matter (VDM) and dark radiation (DR) originate from the same non-Abelian dark sector. We show an illustrating example where dark $SU(3)$ is spontaneously broken into $SU(2)$ subgroup by the nonzero vacuum expectation value (VEV) of a complex scalar in fundamental representation of $SU(3)$. The massless gauge bosons associated with the residual unbroken $SU(2)$ constitute DR and help to relieve the tension in Hubble constant measurements between {\tt Planck} and Hubble Space Telescope. In the meantime, massive dark gauge bosons associated with the broken generators are VDM candidates. Intrinsically, this non-Abelian VDM can interact with non-Abelian DR in the cosmic background, which results in a suppressed matter power spectrum and leads to a smaller $\sigma_8$ for structure formation. | \label{sec:intro} It has been well established that about $25\%$ of energy density in our Universe is made of non-baryonic dark matter (DM). From the perspective of particle physics, weakly-interacting massive particle (WIMP) is one of the nicely motivated candidates. In WIMP scenario, DM is in thermal equilibrium with standard model (SM) particles at high temperature and freezes out at later time. Such an optimistic framework has triggered enthusiastic DM searches in direct, indirect and collider detection experiments. However, we have to admit that so far all confirmed evidence for DM is only from gravitational interaction, which leaves wide possibilities for DM's particle identities. Recently, there are renewed interests in interacting DM--DR models~\cite{Bringmann:2013vra, Boehm:2014vja, Ko:2014bka, Chu:2014lja, Buen-Abad:2015ova, Lesgourgues:2015wza, Tang:2016mot, Bringmann:2016ilk, Ko:2016uft} which could have distinguishing effects on large scale structure. Depending on the DM--DR interactions, these effects can be similar to baryonic acoustic oscillation or dramatically different. Motivations for such models are at least twofold. One is that the DR component could help to resolve the conflict between \planck ~\cite{Planck:2015xua} and Hubble Space Telescope (HST) data~\cite{Riess:2016jrr}. The other is that interaction between DM and DR can give a smaller $\sigma_8$ for structure growth, suggested by low redshift measurements, such as weak lensing survey CFHTLenS~\cite{Heymans:2012gg}. These tensions have stimulated various investigations on cosmological models~\cite{Pourtsidou:2016ico, DiValentino:2016hlg, Qing-Guo:2016ykt, Archidiacono:2016kkh, Wyman:2013lza, Zhang:2014lfa, Lesgourgues:2015wza, Ko:2016uft, DiValentino:2016ucb}. In this paper, we propose a new scenario where DM and DR have the same origin from a single Yang-Mills dark sector, unlike early attempts where DM and DR have different identities~\cite{Bringmann:2013vra, Boehm:2014vja, Ko:2014bka, Chu:2014lja, Buen-Abad:2015ova, Lesgourgues:2015wza, Tang:2016mot, Bringmann:2016ilk, Ko:2016uft}. In our framework presented below, a non-Abelian gauge group is spontaneously broken into its non-Abelian subgroup. The massless gauge boson associated with the residual subgroup constitutes non-Abelian DR, while other massive gauge bosons make non-Abelian VDM candidates. Naturally, VDM can interact with DR through the original Yang-Mills gauge interactions, inducing some observable effects on cosmology and astrophysics. This paper is organized as follows. In Sec.~\ref{sec:su3}, we start with an explicit example where dark $SU(3)$ is broken to its subgroup $SU(2)$ by nonzero VEV of a complex scalar belonging to the fundamental representation of $SU(3)$. Then we generalize to dark $SU(N)$ that is broken into $SU(N-1)$, and give a brief proof why the massive gauge bosons are stable and therefore make good DM candidates. Next in Sec.~\ref{sec:constraint}, we discuss some phenomenologies and constraints on such a class of models, especially on DM--DR scattering, DM self-interaction and DR's contributions to $N_\eff$. Then in Sec.~\ref{sec:relic}, we estimate how DM's relic density can be satisfied with freeze-in process. Later in Sec.~\ref{sec:numeric} we illustrate the effects on matter power spectra in the interacting DM--DR scenario. Finally, we give our conclusion. | \label{sec:conclusion} In this paper, we have proposed a particle physics model in which vector dark matter (VDM) and dark radiation (DR) have a common origin, namely a Yang-Mills dark sector. We have explicitly shown an illustrating case where dark $SU(3)$ gauge group is spontaneously broken to its $SU(2)$ subgroup. The residual massless gauge bosons constitute DR while other massive ones make up the VDM. Interestingly, VDM naturally can scatter with DR through the original Yang-Mills self-interaction, which can lead to a suppressed matter power spectrum and give rise to a smaller $\sigma_8$ for structure growth. On the other hand, DR also helps to resolve the tension in Hubble constant measurements between {\tt Planck}~\cite{Planck:2015xua} and Hubble Space Telescope~\cite{Riess:2016jrr}. In the minimal model with one fundamental scalar that breaks the dark gauge symmetry, the new massive Higgs boson only couples to residual massless gauge boson at loop level, so the coupling would be too small to thermalize DR. However, production of DR could be due to the decay of this new Higgs. The viable way to produce DM with correct relic abundance can be realized in freeze-in mechanism, provided the reheating temperature is about one order-of-magnitude smaller than DM mass. | 16 | 9 | 1609.02307 |
1609 | 1609.09479_arXiv.txt | {In large SEP events, ions can be accelerated at CME-driven shocks to very high energies. Spectra of heavy ions in many large SEP events show features such as roll-overs or spectral breaks. In some events when the spectra are plotted in energy/nucleon they can be shifted relative to each other to make the spectral breaks align. The amount of shift is charge-to-mass ratio (Q/A) dependent and varies from event to event. This can be understood if the spectra of heavy ions are organized by the diffusion coefficients \citep{Cohen2005}. In the work of \citet{Li2009}, the Q/A dependences of the scaling is related to shock geometry when the CME-driven shock is close to the Sun. For events where multiple in-situ spacecraft observations exist, one may expect that different spacecraft are connected to different portions of the CME-driven shock that have different shock geometries, therefore yielding different Q/A dependence. In this work, we examine one SEP event which occurred on 2013 November 4. We study the Q/A dependence of the energy scaling for heavy ion spectra using Helium, oxygen and iron ions. Observations from STEREO-A, STEREO-B and ACE are examined. We find that the scalings are different for different spacecraft. We suggest that this is because ACE, STEREO-A and STEREO-B are connected to different parts of the shock that have different shock geometries. Our analysis indicates that studying the Q/A scaling of in-situ particle spectra can serve as a powerful tool to remotely examine the shock geometry for large SEP events.} \authorrunning{Zhao, Li, \& Mason et al. } % \titlerunning{Shock geometry from multiple spacecraft observation} % | % \label{sec.intro} Understanding Solar Energetic Particles (SEPs) is a central topic of Space Plasma research. Studying SEPs provides a unique opportunity to examine the underlying particle acceleration process which exists at a variety of astrophysical sites. Furthermore, understanding SEPs is of practical importance since SEPs are a major concern of space weather. It is now widely accepted that these high energy particles are accelerated mostly at solar flares and shocks driven by coronal mass ejections (CMEs). Events where particles are accelerated mainly at flares are termed ``impulsive'' \citep{Cane1986} as the time intensity profile shows a rapid rise and fast decay. In contrast, events where particles are accelerated mainly at CME-driven shocks are termed ``gradual'' \citep{Cane1986,Reames1999} where the time intensity profiles vary gradually compared to impulsive events. For large SEP events, recent studies \citep[e.g.][]{Reames2009, Cliver2006, Gopalswamy2012, Mewaldt2012} suggested that energetic particles that are observed near Earth in these events are mostly accelerated at the shocks driven by CMEs rather than in flare active regions. In many large SEP events, particle fluence spectra exhibit exponential rollover or double power law features \citep[e.g.][]{Mewaldt2005, Mewaldt2012}. The break energy or the roll-over energy, $E_0$, is between a few to a few 10's of MeV/nucleon \citep{Mazur1992, Cohen2005, Mewaldt2005, Tylka2005, Desai2016}. Simulations \citep{Li2005} show that spectral breaks can occur naturally for particle acceleration at a CME-driven shock. In examining these features, \citet{Cohen2003, Cohen2005} and \citet{Mewaldt2005} noted that the break energies are nicely ordered by $(Q/A)^{\sigma}$. They suggested that this ordering can be understood if the energy breaks or roll-overs for different heavy ions occur at the same values of the diffusion coefficient $\kappa$. Later, \citet{Li2009} attempted to relate $\sigma$ to shock geometry. They showed that the value of $\sigma$ is usually in the range of $1$ to $2$ for parallel shocks, but can become as small as $\sim 1/5$ for perpendicular shocks. For the most general case of an oblique shock, the total diffusion coefficient is given by, \beq \kappa = \kappa_{||} \cos^2(\theta_{BN}) + \kappa_{\perp} \sin^2(\theta_{BN}). \label{eq:kappa_total} \eeq In the above, $\kappa_{||}$ and $\kappa_{\perp}$ are the parallel and perpendicular diffusion coefficients and $\theta_{BN}$ is the angle between the upstream magnetic field and the shock normal. Since in general $\kappa_{||}$ and $\kappa_{\perp}$ have different $Q/A$ dependence \citep{Li2009}, equation~(\ref{eq:kappa_total}) yields a complicated $Q/A$ dependence for the break energy at an oblique shock. Recently, \citet{Desai2016} have surveyed $0.1$-$100$ MeV/nucleon H-Fe fluence spectra for $46$ isolated large gradual SEP events observed at ACE during solar cycles 23 and 24. They found that the range of $\sigma$ for heavy ion spectra in these events is mostly between $0.2$ to $2$, although some events have a $\sigma$ value larger than $2$. In the work of \citet{Li2009}, it is assumed that the spectral break or roll-over from in-situ observations reflects the same feature of the escaped particle spectra at the shock. We note that some recent calculations have suggested that spectral breaks can emerge as a transport effect \citep{Li2015, Zhao2016}. However, in these calculations, the size of the spectral index change, i.e. $\delta \gamma = \gamma_{a} - \gamma_{b}$, where $\gamma_{a}$ and $\gamma_{b}$ are the spectral indices above and below the break energy $E_0$, is very small. This is in contrast to the observations where $\delta \gamma$ can be large and varies noticeably from one event to another. Furthermore, the transport effect shown in \citep{Li2015, Zhao2016} predicts a $Q/A^{\sigma}$ dependence of the spectra break energy $E_0$ with an upper limit of $\sigma$ to be $1.3$. In a recent statistical suryey however, \citet{Desai2016b} found that $\sigma$ in $33$ SEP events ranged between $\sim 0.2-3$, which clearly exceeds the upper limit of $\sim 1.3$ predicted by scatter-dominated transport models \citep{Li2015, Zhao2016}. Here we follow \citet{Li2009} and assume that the break is a feature of the escaped particle spectrum at the shock. Note that the Diffusive Shock Acceleration (DSA) does not predict a spectral break for the shock-accelerated particle spectrum. Nevertheless, there could be a variety of reasons for such a break. For example, the break may represent the maximum energy given a finite acceleration time. In this case we expect the break energies to be high, $>$ several $10$'s of MeVs for protons. It could also represent the cut-off energy for escape, i.e., particles with energy lower than the break energy are trapped more within the shock complex. In this case, the break energies may be low, $\sim <$ several MeV for protons. In both cases, however, the break energy is decided by the diffusion coefficient $\kappa$, so that the same $Q/A$ analysis discussed in the work of \citet{Li2009} applies. Here we do not discuss the underlying mechanism that leads to the spectral break, but use the $Q/A$ scaling of heavy ion spectra to remotely infer the shock geometry. We note that since particles are continuously accelerated at the CME-driven shock, this shock geometry reflects only an ensemble average of the shock geometry over a period. If however, the energetic particles near the break energy are mostly accelerated at early times (i.e. in the case of the break energy representing the maximum energy), then we expect that the spectral break reflects the shock geometry when the shock is still close to the Sun. Since for the same CME the shock geometry differs at different longitudes, then with simultaneous in-situ observations from multiple spacecraft, one may obtain different Q/A-scaling. This is illustrated in the cartoon shown in Figure~\ref{fig.Cartoon}. In the cartoon, the two field lines colored in blue (assumed here to be unperturbed Parker field) intersect with the shock at a quasi-perpendicular configuration and the two field lines colored in green intersect with the shock at a quasi-parallel configuration. \begin{figure}[htb] \centering \noindent \includegraphics[width=0.6\textwidth]{ms0149fig1.eps} \caption{A cartoon showing that two spacecraft can be magnetically connected to different portions of a CME-driven shock which have different shock geometries. } \label{fig.Cartoon} \end{figure} The above discussion indicates that studying the Q/A scaling of heavy ion spectra simultaneously at multiple spacecraft may be used to infer shock geometry of the CME-driven shock for SEP events where heavy ion spectra are well organized by Q/A. In this work, we examine one such SEP event that occurred on 2013 November 4. In the following, we describe the observation in section~\ref{sec.obs} and present the fitting results in section~\ref{sec.res}. We conclude in section~\ref{sec.disc}. | \label{sec.disc} In this paper, we examine the spectra of heavy ions of the 2013 November 4 SEP event from STA, STB, and ACE. We select this event because the time-integrated heavy ion spectra from all three spacecraft show spectral break features above $\sim 1$ Mev/nucleon. Although the pre-event background is elevated due to the presence of SEPs from another event that had occurred two days earlier, time intensity profiles from all three spacecraft show clear increases from the background. Our analyses show that: 1) for all three spacecraft the He, O, and Fe spectra can be well organized by Q/A and when scaled by $(Q/A)^{\sigma}$, spectra for different heavy ions overlap nicely, 2) the scaling parameter $\sigma$ is sensitive to the charge to mass ratio, and for the sets of heavy ion charge states we choose in section~\ref{sec.res}, we obtain $\sigma$ for STA, ACE, and STB to be $1.12$, $1.01$, and $0.66$, respectively; 3) Under the framework of \citet{Li2009}, these values of $\sigma$ suggest that STA (and ACE) are connected to the quasi-parallel part of the shock and STB is connected to the quasi-perpendicular part of the shock. This is qualitatively in agreement to the configuration shown in Cartoon~\ref{fig.Cartoon}. For ACE and STEREO-A, we integrated reasonably long periods (see Figure 2) to obtain the fluence spectra. As the CME-driven shock propagates out from the Sun to 1 AU, it continues to accelerate particles, but the maximum energy decreases with distance. As a result, the event-integrated spectrum represents an ensemble average of the shock spectra at different times. Since the spectral break feature is around $10$ MeV, a reasonably high energy for this event, we expect the dominant contributing particles for ACE and STA spectra are accelerated close to the Sun, e.g, within $0.3$ AU. For STEREO-B, the period of integration is shorter and close to the shock arrival. Particles of all energies show significant increases from the background at around the same time, indicating that this is due to magnetic connection. This is consistent with the fact that the event is an eastern event as seen from STEREO-B. For STEREO-B observations, one may speculate that the spectrum is more local, i.e. the spectrum represents the shock spectrum when the shock is close to 1 AU. However, particles accelerated at earlier times but are trapped by the CME-driven shock may also contribute. Consequently for the STEREO-B observervation, the range of the radial distance of the shock it samples should be larger than those by ACE and STEREO-A. Therefore, when we interpret the result of the Q/A scaling, one needs to be careful in that the shock geometry for the STEREO-B observation may suffer a larger variation than those from STEREO-A and ACE. We remark that, as revealed by Figure~\ref{fig.Qplot}, our analysis may be used not only to obtain shock geometry estimates for multiple spacecraft, but also to examine charge state variability of heavy ions in an event. In principle, one can compare our charge state results with that of ion charge state measurements. However, ACE/SEPICA, which measured energetic particle charge states, does not have data after 2005; and both ACE/SWICS and STEREO/PLASTIC make only charge state measurements at solar wind energies. | 16 | 9 | 1609.09479 |
1609 | 1609.05182_arXiv.txt | The number of nonrelativistic axions can be changed by inelastic reactions that produce photons or relativistic axions. Any odd number of axions can annihilate into two photons. Any even number of nonrelativistic axions can scatter into two relativistic axions. We calculate the rate at which axions are lost from axion stars from these inelastic reactions. In dilute systems of axions, the dominant inelastic reaction is axion decay into two photons. In sufficiently dense systems of axions, the dominant inelastic reaction is the scattering of four nonrelativistic axions into two relativistic axions. The scattering of odd numbers of axions into two photons produces monochromatic radio-frequency signals at odd-integer harmonics of the fundamental frequency set by the axion mass. This provides a unique signature for dense systems of axions, such as a dense axion star or a collapsing dilute axion star. | 16 | 9 | 1609.05182 |
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1609 | 1609.04247_arXiv.txt | We consider thick domain walls in a de Sitter universe following paper by Basu and Vilenkin~\cite{BV}. However, we are interested not only in stationary solutions found in~\cite{BV}, but also investigate the general case of domain wall evolution with time. When the wall thickness parameter, $\delta_0$, is smaller than $H^{-1}/\sqrt{2}$, where $H $ is the Hubble parameter in de Sitter space-time, then the stationary solutions exist, and initial field configurations tend with time to the stationary ones. However, there are no stationary solutions for $\delta_0 \geq H^{-1}/\sqrt{2}$. We have calculated numerically the rate of the wall expansion in this case and have found that the width of the wall grows exponentially fast for $\delta_0 \gg H^{-1}$. An explanation for the critical value $\delta_{0c} = H^{-1}/\sqrt{2}$ is also proposed. | As is well known, domain walls could be created in the universe if/when a discrete symmetry is spontaneously broken. An interesting example of this kind was suggested in ref.~\cite{lee-p-spont} where the idea of spontaneous $CP$ violation was put forward. However, in the simplest version this mechanism encounters serious cosmological problems because even a single domain wall inside the present day cosmological horizon would strongly distort the observed isotropy of CMB~\cite{ZKO}. To cure this cosmological disaster a few mechanisms of wall destruction were proposed~\cite{kuzmin1, kuzmin2, kuzmin3, kuzmin4}. In our recent paper~\cite{dgrt} we explored the idea of spontaneous $CP$ violation to construct a (nearly) baryo-symmetric cosmology which might be compatible with observations. According to this scenario domains with opposite signs of $CP$ violation appeared during inflation and survived at the stage of reheating when the baryogenesis operated. Later the walls between these domains dissolved and therefore the domain wall problem did not arise. As a result, this model could lead to baryo-symmetric universe with cosmologically large regions of matter and antimatter. Some general features of such dynamical $CP$ violation are described in refs~\cite{ad-dyn-CP-1, ad-dyn-CP-2, AD-Varenna}. For successful implementation of such cosmological model it is imperative that the distance between the matter-antimatter domains is also cosmologically large. It could be realized if the width of the domain wall which existed during baryogenesis was cosmologically large (we understand by the domain wall the piece of space between two regions in the universe, where the field has not yet relaxed to its equilibrium value, though the double minimum in the potential has already disappeared). The matter-antimatter domains should be separated by at least several megaparsec in terms of the present day scale to avoid excessive matter-antimatter annihilation. On the other hand, the distance should not be too large, otherwise the scenario would lead to too large angular fluctuations of CMB~\cite{CdRG}. It means in particular, that domains with opposite $CP$ symmetry breaking must be created during inflationary stage, otherwise both the size of matter-antimatter domains and the transition regions between them would be too small. In contrast, baryogenesis must proceed after inflation was over to avoid strong dilution of the baryon asymmetry. The evolution of the domain walls in de Sitter space-time was considered by Basu and Vilenkin~\cite{BV}. The authors argued that the width of the domain wall is determined by the ratio $C \equiv \lambda\eta^2/H^2>0$, where $H$ is the Hubble parameter, which was assumed to be constant, $\eta$ is the vacuum expectation value of the Higgs-like field which induced the spontaneous symmetry breaking, and $\lambda$ is the coupling constant in the double-well potential, see \eqref{eq:lagrangian}. If $C \gg 2$, the width of the domain wall would be close to its flat space-time value, $\delta_0=1/(\sqrt{\lambda}\eta)$, which is microscopically small, because $\sqrt{\lambda}\eta$ is essentially the mass of the Higgs-like boson. To create astronomically wide domain wall this boson must be practically massless and thus it would generate long range forces most probably excluded or strongly restricted by experiment. On the other hand, if $C < 2$, it is not excluded that the width of the domain wall may be astronomically large. In presented paper we show that this is indeed the case. In ref.~\cite{BV} only the stationary problem was considered, when the shape of the domain wall was a function of a single variable, $l=zH\exp(Ht)$, which is the length interval in de Sitter space. In this case the equation of motion is reduced to an ordinary differential equation which makes the problem much simpler technically. In the paper~\cite{BV} it was found numerically that the stationary solution exists only if $C>2$. In what follows we lift the assumption of the stationarity and consider the general plane solution being a function of both variables: the distance from the wall and time. It allows us to see how the solution approaches the stationary one and, in particular, what happens with initial configurations if $C\leq 2$, when the stationary solution does not exist. Our paper is organized as follows. In Section~\ref{sec:stationary} we reproduce the calculation of Basu and Vilenkin~\cite{BV} and provide a simple explanation why the critical point is at $C=2$. In Section~\ref{sec:evolution} we consider the field evolution with respect to both time and coordinate. Finally, in Section~\ref{sec:conclusions} we conclude. | \label{sec:conclusions} Time evolution of thick domain walls in a de Sitter universe is considered. We have shown that for large values of parameter $C>2$ the initial kink configuration in a de Sitter background tends to the stationary solution obtained by Basu and Vilenkin~\cite{BV}. We also confirmed the BV result that the width of the stationary wall rises with decreasing value of $C$. For $C<2$ the stationary solution does not exist and the width of the wall infinitely grows with time. For $C \lesssim 0.1$ the rise is close to the exponential one within the precision of our numerical calculations. This result confirms the assertion made in ref.~\cite{dgrt} that the transition region between matter and antimatter domain might be cosmologically large. This result is essential for application of spontaneous breaking of symmetry between particles and antiparticles to realistic cosmology. In our version of this scenario~\cite{dgrt} the walls between matter and antimatter domains dissolved and hence the huge energy density of domain walls, which was a stumbling block of the traditional approach, did not destroy isotropy and homogeneity of the universe. Baryogenesis in our model proceeded after the double well potential returned to the ''normal'' potential with the single minimum at $\chi =0$ (here $\chi$ is field which made the wall). Correspondingly the field $\chi$ started to move from the initial values $\langle \chi \rangle = \pm \eta $ to zero, so the domain walls disappeared. However, if the classical field $\chi$ did not completely relax down to zero prior to baryogenesis, it would induce $CP$-violation of different signs at different domains which in turn led to an excess of matter or antimatter in this remnants of the regions with non-zero $\chi$. A necessary condition to make this model realistic is a sufficiently large distance between matter and antimatter domains. Otherwise the annihilation on the boundaries would create too high gamma ray background, according to ref.~\cite{CdRG}. The calculations presented above demonstrate that there is some range of the parameter values, namely $C = \lambda \eta^2 /H^2 \lesssim 0.1$, for which the width of the domain wall exponentially rises and so the distance between matter and antimatter could be large enough. On the other hand, too wide domain walls lead to large universe regions devoid of baryons, so it might distort the observed quasi-isotropy of CMB at the angular distance above the diffusion (Silk) damping scale. However, this statement can be questioned, since the temperature contrast between baryon/antibaryon and empty regions might push the domains apart diminishing the baryon-antibaryon diffusion towards each other and this would allow for smaller separation between the domains below the diffusion damping scale. The latter problem is under investigation now. \vspace{5mm} {\bf Note added.} After this paper was submitted to journal, we were informed about the paper \cite{Voronov}, where the related problems were considered. We thank A.E. Kudryavtsev for this reference. \appendix | 16 | 9 | 1609.04247 |
1609 | 1609.06242_arXiv.txt | { We obtained Arecibo \HI\ line follow-up observations of 154 of the 2600 galaxies in the \nan\ Interstellar Baryons Legacy Extragalactic Survey (NIBLES) sample. These observations are on average four times more sensitive than the original observations at the \nan\ Radio Telescope. The main goal of this survey is to characterize the underlying \HI\ properties of the NIBLES galaxies which were undetected or marginally detected at \nan. Of the \nan\ non-detections, 85\% were either clearly or marginally detected at Arecibo, while 89\% of the \nan\ marginal detections were clearly detected. Based on the statistics of the detections relative to $g$-$i$ color and $r$-band luminosity (\Lr) distribution among our Arecibo observations, we anticipate $\sim$60\% of our 867 \nan\ non-detections and marginal detections could be detected at the sensitivity of our Arecibo observations. Follow-up observations of our low luminosity (\Lr\ < 10$^{8.5}$ \Lsun) blue sources indicate that they have, on average, more concentrated stellar mass distributions than the \nan\ detections in the same luminosity range, suggesting we may be probing galaxies with intrinsically different properties. These follow-up observations enable us to probe \HI\ mass fractions, log(\MHI/\Mstar) 0.5 dex and 1 dex lower, on average, than the NIBLES and ALFALFA surveys respectively.} | \label{introduction} The optical luminosity function (LF) and the \HI\ mass function (HIMF) are two of the most important and fundamental tracers of the volume density distribution of galaxies in the universe. They yield clues to both the baryonic and dark matter content of galaxies, as well as their evolutionary histories. Consequently, there are many applications for which the LF and HIMF can be used, for example, as constraints in galaxy formation models \cite[see, e.g.,][]{benson2003, lu2014}. Many studies have attempted to constrain both of these functions over the years. Since the LF was first fitted to an analytic form by \cite{schechter1976}, many subsequent studies have attempted to analyze its various properties and constrain its parameters (see, e.g., \citealt{felten1985, efstathiou1988, loveday1992, loveday2015, blanton2001, blanton2003, dorta09, mcnaught2014}). The HIMF, having the same functional form as the LF, has also been analyzed in detail, although to a somewhat lesser extent \cite[see, e.g.,][]{zwaan97, zwaan03, kilborn99, kovac2005, springob2005a, martin2010, hoppmann2015}. To date, both these functions have been treated separately in their analyses. One of the main goals of the \nan\ Interstellar Baryons Legacy Extragalactic Survey (NIBLES) is to study the inter-relation between these two fundamental population tracers. More specifically, we want to analyze the HIMF and other galaxy properties as a function of optical luminosity. To achieve this goal, we carried out a 21cm \HI\ line survey at the 100m class \nan\ Radio Telescope (NRT). The final observed sample consists of 2600 galaxies selected from the Sloan Digital Sky Survey (SDSS; \citealp[see e.g.,][]{york00}) with radial velocities 900$<$cz$<$12,000 \kms. The galaxies were selected to be distributed evenly over their entire range of absolute $z$-band magnitudes ($\sim-13.5$ to $-24$), which was used as a proxy for total stellar mass --- see \citet[Paper I]{vandriel2016} for further details. The NIBLES galaxy selection criteria are: \begin{enumerate} \item{Must have both SDSS magnitudes and optical spectrum;} \item {Must lie within the local volume (900$<$cz$<$12,000 \kms);} \item {Uniform sampling of each 0.5 magnitude wide bin in absolute $z$-band magnitude, \Mz}; \item {Preferentially observe nearby objects;} \item {No a priori selection on color.} \end{enumerate} NIBLES, with its relatively uniform selection of galaxies that are based on total stellar mass, is aimed to complement other recent and/or ongoing large \HI\ surveys in the local volume, in particular, blind surveys such as ALFALFA \citep[e.g.,][]{haynes2011}. One main advantage of NIBLES over blind \HI\ surveys is our increased on-source integration time, which enables us to reduce the $rms$ noise of the observations. Each NIBLES source was initially observed at \nan\ for about 40 minutes of telescope time, resulting in a mean $rms$ noise of $\sim$3 mJy at $18$ \kms\ resolution. In the case of weak or non-detections, observations were repeated (as time allowed) resulting in a target $rms$ noise between 1.5 and 1.8 mJy for the majority of our undetected sample. However, there were a number of sources where follow-up time was unavailable to achieve the desired $rms$, which resulted in a mean of 2.3 mJy for the remainder of the undetected sample, yielding a bimodal $rms$ noise distribution (see Paper I). Of the 2600 NIBLES galaxies, 1733 (67\%) were clearly detected, 174 (7\%) marginally detected, and 693 (27\%) were not detected. To adequately quantify our \HI\ distribution across the optical LF, we need to gain a statistical understanding of the underlying \HI\ distribution of sources which were undetected at \nan. We therefore carried out pointed observations of 90 undetected or marginally detected galaxies at the 305m Arecibo radio telescope, which gives us a noise level reduction by about a factor of four. Additionally, we had a number of sources suffering from observational problems at \nan\ which we re-observed at Arecibo, and during periods of time when primary target sources were unavailable, we observed detected NIBLES sources to compare flux calibrations at the two observatories. In total, we observed 154 galaxies from the NIBLES sample (see Sect. \ref{sample selection} and Paper I for details). Here we present the results from these follow-up observations along with a brief synopsis of the differences in the data between the \nan\ and Arecibo samples. The main purpose of this paper is data presentation. Further analysis will be carried out in subsequent papers. In Sect. \ref{sample selection} we describe the selection of the observed sample of galaxies and in Sect. \ref{observations}, the observations and data reduction. The results are presented in Sect. \ref{results} and discussed in Sect. \ref{discussion}. An analysis of this data regarding the impact on our \HI\ distribution as a function of optical luminosity will be presented in Paper III (Butcher et al., in prep.). All source numbers presented in this paper refer to the NIBLES source number, which can be cross-referenced with other common source names in the tables presented here and in Paper I. | We obtained about four times more sensitive follow-up \HI\ observations at Arecibo of 90 NIBLES galaxies that were either not detected or marginally detected at \nan. We detected 80\% of these sources, which has enabled us to probe their underlying \HI\ distribution. The Arecibo detections have on average five times lower \HI\ masses than the \nan\ upper limits estimated in Paper I. Contributing to this factor of five lower mass is not only the lower peak flux densities we are able to detect with Arecibo, but also the $\sim$37\% narrower line widths in our follow-up sample compared to the \nan\ detections of sources with the same optical luminosity. This average difference in line width is primarily driven by the low luminosity (\Lr\ $<$ 10$^{8.5}$ \Lsun) sources which correspondingly show a higher central concentration of light. This may be an indication that these relatively gas-poor galaxies have, on average, a more centrally confined \HI\ mass distribution compared to the \nan\ detected sample in the same luminosity range. If we assume the $g$-$i$ color and \Lr\ distribution of Arecibo detection fractions are representative of the entire \nan\ undetected and marginally detected samples, we estimate $\sim$60\% (520) could be detected with the four times better sensitivity of our Arecibo observations. This would put the over-all NIBLES detection rate at about 86\%. Lastly, our Arecibo follow-up observations enabled us to sample our \nan\ undetected sample to \HI\ mass fractions 0.5 dex lower, on average, than our \nan\ detections. Some of these galaxies with low \MHI/\Mstar\ fractions lie in virtually unexplored parameter space (e.g., around log(\Lr) = 11.5) and could potentially be used to shed further light on galaxy evolution processes studied by modelers, e.g., \cite{Kannappan2013}. | 16 | 9 | 1609.06242 |
1609 | 1609.01654_arXiv.txt | {The Hamamatsu R11410-21 photomultiplier tube is the photodetector of choice for the XENON1T dual-phase time projection chamber. The device has been optimized for a very low intrinsic radioactivity, a high quantum efficiency and a high sensitivity to single photon detection. A total of 248 tubes are currently operated in XENON1T, selected out of 321 tested units. In this article the procedures implemented to evaluate the large number of tubes prior to their installation in XENON1T are described. The parameter distributions for all tested tubes are shown, with an emphasis on those selected for XENON1T, of which the impact on the detector performance is discussed. All photomultipliers have been tested in a nitrogen atmosphere at cryogenic temperatures, with a subset of the tubes being tested in gaseous and liquid xenon, simulating their operating conditions in the dark matter detector. The performance and evaluation of the tubes in the different environments is reported and the criteria for rejection of PMTs are outlined and quantified.} | The direct detection of dark matter particles scattering off target nuclei is one of the most sought-after measurements in modern physics. In the last decade, dual-phase time projection chambers (TPC), operated with liquefied xenon (LXe), have reached the highest sensitivity for Weakly Interacting Massive Particle (WIMP) interactions above masses of $5\,\rm{GeV}/c^2$~\cite{Undagoitia:2015gya}\cite{Baudis:2016qwx}. The XENON1T experiment~\cite{Aprile:2015uzo} is designed to further improve this sensitivity, searching for WIMPs through their scattering off xenon nuclei. At the expected recoil energies of a few keV, the recoiling nucleus can create excimer molecules in xenon which, upon dissociation, produce scintillation of vacuum ultraviolet (VUV) photons with a wavelength of 178\,nm~\cite{Cheshnovsky1972}. In addition, particle recoils produce free electrons that are drifted and extracted from the liquid by an applied electric field to be subsequently amplified in the gaseous xenon volume. % The amplified electron showers % produce a secondary scintillation signal through the process of electroluminescence\,\cite{Lansiart1976}. It becomes evident that a successful detection of particle interactions in liquid xenon can only be achieved by an efficient detection of VUV photons. To reach the highest possible sensitivities to dark matter interactions, the photomultiplier tubes (PMTs) are required to have a high photon detection efficiency (quantum efficiency) at a wavelength of 178\,nm, ultra low radioactivity levels~\cite{Aprile:2015lha} and stable long-term performance at operating temperatures of $-100$\,$^{\circ}$C. In this article, the evaluation tests of 321 Hamamatsu R11410-21 PMTs are described. The selection criteria for the final 248 tubes for XENON1T are discussed, based on characteristic parameters such as the dark count rate, quantum efficiency, afterpulse rate, light emission, long-term stability, peak-to-valley ratio and single photoelectron (SPE) resolution. The general testing procedure and experimental facilities are introduced in section~\ref{sec:exp}. The distributions of measured and derived PMT parameters are shown in section~\ref{sec:res}. Section~\ref{sec:lightemission} describes the emission of light by some of the tested PMTs, being one of the major rejection criteria. Dedicated measurements of the long-term stability in cryogenic xenon environments are presented in section~\ref{sec:xetests}. Here the evolution of the dark count rates and gains is reported. Section~\ref{sec:ap} describes a detailed study of the afterpulse spectra and the methods used to determine the presence of leaks in faulty PMTs by identification of residual gas molecules. The conclusions of this article are summarized in section~\ref{sec:sum}. Complementary studies of earlier versions of the R11410 phototube have been reported in~\cite{Lung:2012pi}\cite{Baudis:2013xva}. This PMT is also operated in experiments using liquid and gaseous xenon, such as PandaX\,\cite{Cao:2014jsa}, NEXT\,\cite{Alvarez:2012flf} and RED\,\cite{Akimov:2012aya}\cite{Akimov:2015aoa}, while a version for operation in liquid argon is used in DarkSide\,\cite{Agnes:2014nla} and GERDA\,\cite{Agostini:2015boa}. The long term performance of 37 R11410-MOD PMTs within the PandaX dark matter experiment is presented in\,\cite{Li:2015qhq}. | \label{sec:sum} In this article we report on the extensive testing campaign of 321 R11410-21 photomultiplier tubes, out of which 248 have been selected for long-term operation in the XENON1T experiment. The evaluation of the tubes is based on the dark count rate, SPE response, transit time spread, quantum efficiency, and level of light emission. Detailed measurements have been performed on a subset of PMTs in gaseous and liquid xenon. The evolution of the dark count rate and gain have been studied, showing the viability to operate the tubes long-term in a dark matter experiment. The identification of residual gas molecules within the PMTs through the analysis of afterpulses has been studied in detail, matching the experimental measurements with analytical calculations and simulations. The results have been used for diagnosis of the PMT vacuum quality and identification of leaks. A total of 73 tubes, 22\,\% of the amount tested, have been rejected and thus excluded for use in XENON1T. Out of these, 12 tubes have been rejected due to an elevated or unstable dark count rate, either at cryogenic or room temperature. A larger amount, 53 in total, presented light emission. Out of the 44 PMTs tested in LXe, 8 were identified to show afterpulses from xenon ions, indicating a leak in the tube and thus being unsuited for long-term operation. It must be noted that the leak statistics is not representative of the whole set, since most PMTs tested in LXe were selected for their unstable performance during previous tests or suspicion of a leak after operation in nitrogen gas at $-100^{\circ}$C. The selected 248 PMTs have an average quantum efficiency of (34\,$\pm$\,3)\,\% and show a low average dark count rate of (40\,$\pm$\,13)\,Hz at $-100\,^{\circ}$C. At a gain between $(2-3)\times 10^{6}$ the peak-to-valley ratio ranges from 2.5 to 4.5, proving the large signal to noise separation of this tube. During their operation in XENON1T, the evolution of the gain and dark count rates will be monitored for all PMTs. The methods presented here will also allow to identify the appearance of light emission or leaks in the tubes throughout the course of the experiment. | 16 | 9 | 1609.01654 |
1609 | 1609.04592_arXiv.txt | We analyze 15,000 spectra of 29 stellar-mass black hole candidates collected over the 16-year mission lifetime of \rxte\ using a simple phenomenological model. As these black holes vary widely in luminosity and progress through a sequence of spectral states, which we broadly refer to as hard and soft, we focus on two spectral components: The Compton power law and the reflection spectrum it generates by illuminating the accretion disk. Our proxy for the strength of reflection is the equivalent width of the Fe-K line as measured with respect to the power law. % A key distinction of our work is that for {\em all} states we estimate the continuum under the line by excluding the thermal disk component and using only the component that is responsible for fluorescing the Fe-K line, namely the Compton power law. We find that reflection is several times more pronounced ($\sim 3$) in soft compared to hard spectral states. This is most readily caused by the dilution of the Fe line amplitude from Compton scattering in the corona, which has a higher optical depth in hard states. Alternatively, this could be explained by a more compact corona in soft (compared to hard) states, which would result in a higher reflection fraction. | \label{section:intro} During the course of its 16-year mission, the {\it Rossi X-ray Timing Explorer} (\rxte) detected far more photons (30 billion in PCU-2 alone) from accreting black holes than any other X-ray observatory. The sample of black holes (BHs) targeted by {\it RXTE} is chiefly comprised of nearby stellar-mass systems. While the total Galactic population of stellar BHs is believed to be many millions, only a tiny subset of approximately 50 are known to us, namely those located in X-ray binaries. A wondrous property of BHs, their utter simplicity, is the essence of the famous no-hair theorem: Each BH in nature is fully described by just its mass and spin. Roughly half of the known stellar BHs have a dynamically-determined mass. The measured masses range from $\sim5-20~\msun$ \citep{Ozel_2010,Reid_2014,Laycock_2015,Wu_2016}. Meanwhile, estimates of spin have been obtained for many of them during the past decade, principally by modeling either the thermal continuum emission of the accretion disk \citep[e.g.;][]{Zhang,MNS14}, or the relativistically-broadened reflection spectrum \citep[e.g.;][]{Fabian_1989,Reynolds_2014}. Our focus is primarily on transient BH systems that cycle between a minuscule fraction of the Eddington limit upward to the limit itself. During an outburst, a transient BH progresses through a sequence of spectral-timing states, which are broadly termed ``hard'' or ``soft,'' based on a measure of X-ray hardness \citep{Fender_2004}. As a source evolves over the course of months and its hardness varies, sweeping changes occur in many properties of the system including the composition of its spectrum, the intensity of Fourier flicker noise, and the presence or absence of quasi-periodic oscillations and jets \citep[e.g.;][]{Homan_2005,RM06,Heil_2015}. Stellar BHs emit a complex multicomponent X-ray spectrum. A {\it thermal} blackbody-like component is produced in the very inner accretion disk. The disk is truncated at a radius $\rin$ before reaching the event horizon. A hard {\it power-law} component results from Compton scattering of the thermal disk photons in hot coronal gas that veils the disk. The third principal component is a {\it reflection} spectrum generated by illumination of the cold disk ($kT\sim0.1-1$\,keV) by the power-law component. The reflection component is a rich mix of radiative recombination continua, absorption edges and fluorescent lines \citep{Ross_1993,Garcia_Kallman_2010}. An analysis of these three interacting spectral components provides constraints on the source properties including geometry (e.g., on $\rin$ and the scale of the corona). The relationships between these components across the full range of behavior displayed by accreting stellar BHs is the focus of this paper. Our results are based on an analysis for 29 stellar BHs (10 dynamically-confirmed BHs and 19 BH candidates) of all the data collected using \rxte's prime detector unit (PCU-2), some 15,000 spectra in all, with a total net exposure time of 30\,Ms. Importantly, we recalibrate the data using our tool {\sc pcacorr}, which greatly reduces the level of systematic error \citep{pcacorr}. Given the scope of our study, relativistic reflection models are too complex and computationally slow for our purposes \citep[e.g.; {\tt reflionx, xillver, relxill;}][]{reflionx,relxill2}. We therefore employ a simplistic, phenomenological model and estimate the strength of the reflection spectrum by determining the equivalent width with respect to the Compton continuum of its most prominent reflection feature, namely the $6.4-7.0$~keV Fe-K line. The paper is organized as follows: In Section~\ref{section:data} we describe the data sample and our approach to modeling the data. Our results are presented in Section~\ref{section:results}, followed by a discussion in Section~\ref{section:disc} and our conclusions in Section~\ref{section:conc}. \begin{figure}[] \begin{center} \includegraphics[width=1\columnwidth]{fig_tmp_testplot_f1master1} \caption{({\it top}): Hardness-intensity diagrams for all data and ({\it bottom:}) for six well-known BHs with abundant data (where for reference the gray background shows all data). For reference, the count rate of the Crab Nebula is $\approx 2600$~s$^{-1}$. Note that a HID does not allow one to compare the luminosities of sources because the intensity is in detector units.}\label{fig:qdiag} \end{center} \end{figure} | \label{section:conc} We have examined the strength of reflection in a global study of stellar BHs using a simplistic, phenomenological spectral model. We directly validate the reflection paradigm, wherein power-law flux induces reflection emission. In separating possible contribution from disk self-irradiation, we demonstrate that the power law's contribution is dominant. Most importantly, we show that the corona produces reflection features up to an order of magnitude more pronounced in soft rather than hard states. The data suggest an ordered transition in which the line-to-continuum strength declines gradually with spectral hardness. This is the first time the ``$R-\Gamma$'' correlation has been shown to extend through (and increase in) BH soft states. One possible explanation is that a more compact disk-coronal geometry in soft states would produce the observed trend. However, the most natural explanation for this trend is suggested by \citet{Petrucci_2001}, who describe the dilution of line features emitted by the disk due to Compton-scattering in the corona. In our case, because hard states have corona with higher optical depth than soft states, their line features are correspondingly weakened resulting in the observed anti-correlation between $\HR$ and reflection strength. | 16 | 9 | 1609.04592 |
1609 | 1609.08543_arXiv.txt | We show for the first time that sustained turbulence is possible at low magnetic Prandtl number for Keplerian flows with no mean magnetic flux. Our results indicate that increasing the vertical domain size is equivalent to increasing the dynamical range between the energy injection scale and the dissipative scale. This has important implications for a large variety of differentially rotating systems with low magnetic Prandtl number such as protostellar disks and laboratory experiments. | Differentially rotating flows are ubiquitous in nature, from accretion flows around young stars and compact objects to convection zones inside stars. Magnetic fields are important in many of these flows. Understanding the stability properties of these flows can shed light into the processes driving the transport of angular momentum and energy in these environments. The Magnetorotational Instability (MRI) is a linear instability that has emerged as a potential explanation for the origin of turbulence in weakly magnetized differentially rotating flows (e.g. magnetized Taylor-Couette flows, accretion disks \cite{1959velikhov, 1960PNAS...46..253C, 1991ApJ...376..214B}). MRI requires a weak magnetic flux to work but given the paucity of direct observations of magnetic field amplitudes and geometries, it is of interest to explore a range of cases for the magnetic flux threading an accretion flow, including the case of zero mean magnetic flux. The earliest 3D simulation of ideal magnetohydrodynamic (MHD) Keplerian shear flow without a mean magnetic flux by \cite{hawley1996} showed that the saturated level of turbulent stresses decreases with resolution (see also \cite{pessah2007}). Later work by \cite{fromang2007} used physical diffusion coefficients and showed that turbulence does not persist for magnetic Prandtl number, $Pm = Rm/Re \lesssim 2$. Recent work on linearly stable hydrodynamic shear flows (e.g., \cite{hof2006}) suggests that turbulence lifetime is finite and increases exponentially with the Reynolds number, $Re$. \cite{rempel} explored the question of turbulence lifetime in Keplerian shear flows with zero net magnetic flux with a small vertical domain size ($L_z=1$) and found that the turbulence lifetime increases exponentially with $Rm$ between $9,000$ and $11,000$ ($Re=3,125$). Examples of $Pm \ll 1$ systems include astrophysical systems such as protostellar disks (for more astrophysical examples, see table 1 of \cite{brandenburg2005}) and laboratory plasmas (e.g., \cite{sisan2004}, \cite{exp2014}). The $Pm \ll 1$ limit has been studied extensively in non-helically forced isotropic MHD turbulence simulations and theory, with the consensus that a larger $Rm_{\text{crit}}$ is required in this limit compared to the $Pm \gg 1$ \cite{sch2007}. Earlier work by \cite{sch2004} had claimed that the small scale dynamo (growth of the magnetic field on dissipation scales) does not exist in the low $Pm$ limit but later work \citep{iskakov2007} demonstrated growth and sustenance of magnetic fields (at a higher $Rm_{\text{crit}}$). In their study, \cite{sch2004} conjectured that a low $Pm$ dynamo might only be possible in the presence of a mean field. Indeed for MRI simulations with net magnetic flux, turbulence can be sustained in the low $Pm$ regime (\cite{lesur2007}, \cite{meheut2015}). The question of whether a large domain size can have a significant effect on flow stability has been long under discussion (e.g., \cite{pomeau1986}, \cite{philip2011}). The shearing box is a local approximation \cite{1965MNRAS.130...97G} for differentially rotating flows such as accretion disks (or Taylor-Couette flows) and it is unclear whether it can exhibit similar behavior as the spatiotemporal chaotic patterns that might emerge in a realistic accretion flow with its very large spatial extent and very large $Re$ (\cite{cross1993}). Nevertheless, within the shearing box framework, it is of interest to explore how a small domain size (`minimal flow unit' \cite{jimenez1991}, \cite{rincon2007}) is different from a system with larger degrees of freedom as a result of larger domain size, resolution, $Re$ and $Rm$. Using an asymptotic analysis with net toroidal and vertical flux, \cite{julien2007} showed that scale separation between the most unstable mode and the vertical domain size leads to a saturated state of MRI where the energy is dominated by the box scale in the vertical direction. In this Letter, we explore the question of whether sustained turbulence can be found in unstratified\footnote{ \cite{davis2010} suggest that density stratification due to gravity lowers the critical $Pm$ slightly but they did not explore the $Pm<1$ regime.} incompressible MHD Keplerian shear flows with zero mean magnetic flux with a variety of box sizes and dissipation coefficients. We find that the turbulence lifetime is very sensitive to the vertical box size \footnote{We do not explore elongated azimuthal domains as \cite{riols2015} already explored $L_y=20$ and did not find evidence for sustained turbulence at $Pm < 1$.} and that large scale structures develop both in the velocity and the magnetic fields. | Recent numerical simulations and experiments on hydrodynamic shear flows suggest that transition to turbulence mimics a second order phase transition and can be described by the directed percolation universality class (\cite{shih2016NatPh, lemoult2016NatPh, sano2016NatPh}). The basic idea is that linearly stable flows develop two kinds of domains: laminar and turbulent corresponding to dead and active states. It is through the interaction of these two types of domains that turbulence starts to spread and eventually fills the whole domain as the Reynolds number is increased. This might explain why the box size plays a special role: a larger box size allows for more turbulent and laminar domains to fit in the box and thus allows for complex pattern formation (\cite{cross1993}). A potential analogy for our simulations is the generation of $B_x$ (active state) through stochastic processes on small scales and the subsequent generation of $B_y$ (dead state) through the $S B_x$ term (see also the recent study by \cite{riols2016} who describe the dynamo using the concept of `active' and `slave' perturbations). We have demonstrated for the first time that Keplerian shear MHD turbulence (without using any external forcing or net magnetic flux) can sustain for several thousand shear times at $Pm \leq 1$ (for $Re=10,000$), if one uses domain sizes with $L_z \geq 4$. We find that with the $Re$ and $Rm$ up to $10,000$ there is no turbulence at $Pm \leq 1$ when $L_z = 1$, consistent with previous work. It seems that by increasing $L_z$, one increases the separation between the outer scale ($\sim L_z$) and the dissipation scale, which is qualitatively similar to the effect of increasing $Rm$. But the increase in $L_z$ at fixed $Re$ and $Rm$ has a more dramatic effect on turbulence lifetime than increasing $Re$ and $Rm$ at $L_z=1$. Our work clearly illustrates that flow stability studies should not be confined to very small domains and has especially important implications for laboratory plasmas, which can only explore the $Pm\ll 1$ regime {\cite{sisan2004}, \cite{exp2014}}. | 16 | 9 | 1609.08543 |
1609 | 1609.08069_arXiv.txt | The current state of the link problem between radio and optical celestial reference frames is considered. The main objectives of the investigations in this direction during the next few years are the preparation of a comparison and the mutual orientation and rotation between the optical {\it Gaia} Celestial Reference Frame (GCRF) and the 3rd generation radio International Celestial Reference Frame (ICRF3), obtained from VLBI observations. Both systems, ideally, should be a realization of the ICRS (International Celestial Reference System) at micro-arcsecond level accuracy. Therefore, the link accuracy between the ICRF and GCRF should be obtained with similar error level, which is not a trivial task due to relatively large systematic and random errors in source positions at different frequency bands. In this paper, a brief overview of recent work on the GCRF--ICRF link is presented. Additional possibilities to improve the GCRF--ICRF link accuracy are discussed. The suggestion is made to use astrometric radio sources with optical magnitude to 20$^m$ rather than to 18$^m$ as currently planned for the GCRF--ICRF link. In addition, the use of radio stars is also a prospective method to obtain independent and accurate orientation between the Gaia frame and the ICRF. | The ESA's {\it Gaia} space astrometry mission \citep{Perryman2001, Lindegren2008} commenced successfully in December 2013 and its main scientific program in July 2014. One of the most important results of the {\it Gaia} mission will be a new highly-accurate optical celestial reference frame; {\it Gaia} Celestial Reference Frame (GCRF). Although the final GCRF version is expected to be available in the early 2020s, intermediate releases are planned, the first of them (DR1) is expected to be released in 2016. A new release of the VLBI-based celestial reference frame of similar accuracy, the 3rd realization of the International Celestial Reference Frame (ICRF3) is planned for 2018 \citep{Jacobs2014}. Both radio (ICRF) and optical (GCRF) frames must be realizations of the same concept of the International Celestial Reference System (ICRS), \citet{Arias1995}) with an expected accuracy at the level of a few tens of $\mu$as. The link between the ICRF and GCRF should be realized at a similar level of accuracy, which is not a trivial task. This problem is similar to that of the link between the {\it Hipparcos} Celestial Reference Frame (HCRF) and the ICRF \citep{Kovalevsky1997}. Generally speaking, both GCRF and ICRF object positions are time-dependent. Therefore, analogously to HCRF, both the orientation and rotation of the GCRF with respect to ICRF are to be defined. In this paper, the general term 'orientation' is used to avoid a non-principal discussion related to spin. Interested readers can find more theoretical and practical details in \citet{Lindegren1995,Kovalevsky1997}. On the one hand, the link task is more straightforward for the GCRF than for the HCRF, as most of the ICRF objects will be directly observed by {\it Gaia}. On the other hand, the task is much more complicated due to the requirement that a much higher level of accuracy for the GCRF--ICRF link is needed so as to not compromise the high level of precision of the two frames. The basic method to tie the {\it Gaia} catalog to the ICRF, and hence to the ICRS, is using {\it Gaia} observations of compact extragalactic ICRF objects that have accurate radio astrometric positions. With the help of these common objects, the orientation angles between the ICRF and GCRF will be determined. Finally, the GCRF catalog will be aligned to the ICRS by applying these orientation angles. The accuracy of this link depends on many factors, such as random and systematic errors of both radio and optical catalogs. The objective of this paper is to briefly overview recent work on the ICRF--GCRF link and to discuss new possibilities to improve the link accuracy. It should be noted that the link between the GCRF and ICRF is not a task currently planned for completion before the end of this decade \citep{Jacobs2014}. Based on the {\it Hipparcos} experience, it can be envisioned that the work on improving such a link will be continued for a prolonged period after completion of the {\it Gaia} mission. The improvements will be primarily based on a new VLBI-based celestial reference frame (CRF) realization of which the accuracy can improve over time. New after-mission {\it Gaia} data reductions are also possible. Therefore, research and development of the methods for the linking of radio and optical reference frames will remain one of the primary tasks of fundamental astrometry throughout the next decade. The paper is structured as follows. In section~\ref{sect:basic}, basic equations used to link two CRF realizations are described. Section~\ref{sect:overview} contains a brief overview of recent investigations regarding aspects concerning the ICRF--GCRF link. The following three sections are devoted to a discussion of new possibilities and possible improvements in both theoretical analysis and the final ICRF--GCRF link accuracy, such as the choice of the ICRF realization used for modelling and simulation (Section~\ref{sect:radio_frame}), using more link sources (Section~\ref{sect:more_sources}), using radio stars (Section~\ref{sect:radio_stars}), and proper accounting for the galactic aberration in proper motions (Section~\ref{sect:ga}). | A great milestone in the construction of the celestial reference frame is the {\it Gaia} mission, which will result in the GCRF with expected accuracy of a few tens of micro-arcsecond for final release in the early 2020s. A new ICRF release based on VLBI observations of extragalactic radio sources of similar accuracy can be also expected by that time. Constructing of a single multi-frequency celestial reference frame based on the ICRF and GCRF is the primary task of fundamental astronomy for the next decade \citep{Gaume2015}. The first step toward this goal is an accurate alignment of the {\it Gaia} catalog to the ICRF. The desired goal is to achieve mutual orientation between the two frames with an accuracy of 0.1~mas or better, which is a challenging task. Currently, the International VLBI Service for Geodesy and Astrometry (IVS) is conducting a special program on observation of prospective link sources not having a sufficient number of observations in the framework of regular observing programs \citep{LeBail2016}. This forms part of the plan to prepare to achieve the GCRF--ICRF link using 195 selected sources selected on basis of several criteria, such as, inclusion in the ICRF2, optical magnitude $\le 18^m$, symmetric compact structure, sufficient radio flux density. The two first criteria seem to be outdated. Firstly, the ICRF2 is currently not the most appropriate radio source position catalog that can serve as an ICRF3 prototype, which is planned to be used for initial alignment of the {\it Gaia} catalog to the ICRS as discussed in Section~\ref{sect:radio_frame}. Secondly, as shown in Section~\ref{sect:more_sources}, using optically fainter sources up to 20$^m$ provides more precise determination of the orientation angles. It is even more important that the use of more sources with reliable VLBI positions is necessary to mitigate the impact of most of the negative factors mentioned below, see Fig.~\ref{fig:err_mag_add} and related text. As discussed, there are possibilities to multiply the number of link sources taking into account the substantial increase of the number of radio sources with reliably determined positions and the number of sources having photometry measurements. The following problems were identified in the literature that dilute the accuracy of the link between radio and optical frames: \begin{itemize} \item Structure effects (discussed below). \item Systematic errors of radio position catalogs. \item Multiple radio sources related to a single object in optics, e.g., binary black holes, and vice versa. \item Gravitational lenses. \item Errors in ICRF--GCRF cross-identification. \end{itemize} Most probably, the main factor that will deteriorate the accuracy of the link between the radio and optical frames, is the source structure. It can manifest at both radio and optical wavelengths as complex, asymmetric distribution of the brightness over the source map, spatial bias between the optical and radio brightness centroids, core-shift effects, spatial bias between the core/AGN brightness centroid and with respect to the optical centroid of the host galaxy. Moreover, the structure effects are often variable. Although many studies are devoted to this factor, see, \citet{Fey1997,SilvaNeto2002,Moor2011,Bouffet2013,Zacharias2014,Berghea2016} and papers cited therein, there are insufficient data to quantify the impact of source structure on the accuracy of the orientation angles between optical and radio frames. Evidently, the most complete study is provided by \citet{Zacharias2014} used in the current work. It should be noted that though the structure effect can reach several mas for an individual source, it can hardly be systematic and thus will be averaged over the sky during the computation of the orientation angles between the GCRF and ICRF. However, supplementary observations and theoretical considerations are needed to quantify this effect more accurately. Two main methods to obtain the link between the GCRF and ICRF were considered in this paper. The first method is direct {\it Gaia} observations of the sufficiently optically bright ICRF sources. This method allows for a straightforward solution of the task. However, there are serious constraints on the accuracy of this method caused by the previously mentioned factors. These factors can limit the real accuracy of the GCRF--ICRF link to 0.1 mas or worse. Observations of radio stars can serve as an alternative equally accurate method. It was successfully used to link the HCRF to the ICRF. However this method is also affected by some accuracy-limiting factors \citep{Malkin2016e}. Many radio stars comprise double or multiple systems, and thus their orbital motions must be accounted for. The accurate {\it Gaia}-derived orbits can be used for this purpose. Moreover, radio stars may have complex and variable structures, which might cause a time-dependent bias between the radio and optical positions. \citet{Lestrade1995} estimated the impact of radio star structure and possible variations in the radio star emitting centre to be within the error budget of approximately 0.5 mas. \citet{Lestrade1999} found that the structure-induced systematic errors in the VLBI positions of 12 stars ranged from 0.07~mas to~0.54 mas, with a median value of 0.18~mas. Provided that several tens of radio stars have been observed, this factor should not significantly impact on the errors in the orientation angles. Combination of two methods of linking the GCRF to the ICRF should facilitate improvement of the systematic accuracy of the link between radio and optical frames. Finally, it should be noted that the preparation for the aligning of the {\it Gaia} catalog to ICRF3 planned for 2018--2019 is only an intermediate stage in construction of the multi-frequency celestial reference frame. Certainly, the main work on the link between radio and optical frames is to be done in the early 2020s, after the final {\it Gaia} catalog is prepared. It is desirable to plan preparation of a new ICRF realization, will it be called ICRF4 or ICRF3 extension, at the same time, i.e. immediately before the final GCRF link to ICRF. Comparison of gsf2015b with ICRF2 has clearly shown that the latest radio catalog should be used for this work due to the accuracy of the VLBI-derived CRF solutions which rapidly improves with time. In the framework of this activity, it appears to be very important to fast-track the plans on the ICRF improvements in the Southern Hemisphere \citep{Jacobs2014,Plank2015}, and to start a program of observations of radio stars \citep{Malkin2016e}. As the {\it Hipparcos} experience has shown, it can be expected that the work on improving the link between radio and optical frames will be continued during a long period after completing the {\it Gaia} mission. Continuous quality improvement of a VLBI-based ICRF promises corresponding improvement of this link with time. | 16 | 9 | 1609.08069 |
1609 | 1609.08633_arXiv.txt | The nearly circular (mean eccentricity $\bar{e}\sim 0.06$) and coplanar (mean mutual inclination $\bar{i}\sim 3^{\circ}$) orbits of the Solar System planets motivated Kant and Laplace to put forth the hypothesis that planets are formed in disks, which has developed into the widely accepted theory of planet formation. Surprisingly, the first several hundred extrasolar planets (mostly Jovian) discovered using the Radial Velocity (RV) technique are commonly on eccentric orbits ($\bar{e}\sim 0.3$). This raises a fundamental question: Are the Solar System and its formation special? The {\it Kepler} mission has found thousands of transiting planets dominated by sub-Neptunes, but most of their orbital eccentricities remain unknown. By using the precise spectroscopic host star parameters from the LAMOST observations, we measure the eccentricity distributions for a large (698) and homogeneous Kepler planet sample with transit duration statistics. Nearly half of the planets are in systems with single transiting planets (singles), while the other half are multiple-transiting planets (multiples). We find an eccentricity dichotomy: on average, {\it Kepler} singles are on eccentric orbits with $\bar{e}\approx$ 0.3, while the multiples are on nearly circular $(\bar{e} = 0.04^{+0.03}_{-0.04})$ and coplanar $(\bar{i} = {1.4}^{+0.8}_{-1.1}$ degree) orbits similar to the Solar System planets. {\xie Our results are consistent with previous studies {\dong of smaller samples} and individual systems.} We also show that {\it Kepler} multiples and Solar System objects follow a common relation ($\bar{e}\sim$(1-2)$\times\bar{i}$) between mean eccentricities and mutual inclinations. {\xie The prevalence of circular orbits and the common relation may imply that the Solar system is not so atypical in the Galaxy after all.} | LAMOST has been performing large-scale Galactic surveys with the spectral resolution $R\sim 1800$ \cite{Zha12, Liu15}, and the LAMOST observations of the {\it Kepler} field started in 2011 \cite{Dec15}. We use the stellar parameters extracted from the LAMOST Stellar Parameter Pipeline (LASP) (see Sec 4.4 of \cite{Luo15} and also \cite{Wu14}) in the ``AFGK high quality stellar parameter catalog'' of LAMOST DR1, DR2 and DR3-alpha data releases. There are 29553 unique {\it Kepler} targets that have LAMOST/LASP stellar parameters. We perform several internal and external examinations on the accuracy of LAMOST/LASP stellar parameters for the dwarfs. There are 5924 {\it Kepler} targets that have LAMOST/LASP stellar parameters from more than one epochs of LAMOST observations. We assess the internal errors by making comparisons of the multi-epoch observations for the same objects. We use the unbiased estimator \cite{Guo15} $\Delta Q_i= \sqrt{n/(n-1)} (Q_i-\bar{Q})$ , with $i=1,2,...,n$, where $i$ denotes each of the individual measurement of $n$ repeated measurements for the stellar parameter $Q$ for each star. Fig.~\ref{fig_internal} shows $\Delta{T_{\rm eff}}$, $\Delta{log(g)}$ and $\Delta{\rm [Fe/H]}$ using the unbiased estimator (black dots). We calculate the $68.3\%$ confidence interval in various bins of $g$-band Signal-to-Noise-Ratio per pixel (SNR$_g$) from LAMOST/LASP and find that they are well described by the second-order polynomials as a function of SNR$_g$ shown in Fig.~\ref{fig_internal}. For measurements with high SNR (${\rm SNR}_g > 50$), the internal errors in $T_{\rm eff}$, $log(g)$ and [Fe/H] are less than $35$K, $0.05\,$dex and $0.03\,$dex, respectively. Following a previous approach of external examination of LAMOST/LASP stellar parameters \cite{Don14}, we compare the LAMOST/LASP parameters with those obtained from the high-resolution spectroscopy with the SPC method \cite{Buc12}. There are 87 stars in common between the SPC and LAMOST/LASP samples. The results of the comparison are shown in the top three panels of Fig.~\ref{fig_external}. The mean differences are small: ${\Delta{T_{\rm eff}}}= 27 K$, ${\Delta{log(g)}}= -0.04 {\,\rm dex}$, ${\Delta{{\rm[Fe/H]}}}=0.015{\,\rm dex}$, respectively. For those with ${\rm SNR}_g > 50$, the standard deviations in $\Delta{T_{\rm eff}}$, $\Delta{log(g)}$, $\Delta{\rm [Fe/H]}$ are $101 K$, $0.15{\,\rm dex}$, $0.074{\,\rm dex}$, respectively. Because there are a small number of common stars with low SNRs, it is difficult to calibrate the errors directly from external calibrators for low-SNR measurements. In order to estimate the error bars for both high- and low-SNR measurements, the standard deviations derived from high-SNR measurements are added in quadrature with the internal error bars in the form of the second-order polynomials shown in Fig.~\ref{fig_internal}. Since the standard deviations for high-SNR measurements from the external calibrations are much larger than the internal errors at similar SNRs, this approach keeps external error calibrations at high SNRs while takes the internal calibrations into account for low-SNR measurements. We have also made comparisons in $log(g)$ with the {\it Kepler} asteroseismology sample for solar-type stars \cite{Cha14}. There are 260 common stars between the LAMOST/LASP and the seismology samples. We find that the $log(g)$ determinations from LAMOST/LASP are in excellent agreement with asteroseismic values, with $\Delta{log(g)}= 0.03 \pm 0.09$ for ${\rm SNR}_{g} > 50$ (see the red points in the bottom panel of Fig.~\ref{fig_external}). The dispersion (0.09 dex) is smaller than that from the comparison with the spectroscopic sample (0.15 dex). The larger dispersion for the latter likely reflects the systematic uncertainties in the SPC method, as demonstrated by comparison with the {\it Kepler} seismology sample \cite{Hub13}. We apply the same quadrature corrections taking into account for the internal errors and the resulting values as a function of SNR are shown in Fig.~\ref{fig_external} (dashed line in the bottom panel). Note that similar comparisons have been made before for LAMOST $log(g)$ but mostly for giant stars \cite{Ren15} or a mixture of giant and dwarf stars \cite{Luo15}. For giant stars, LAMOST $log(g)$ appears to have larger uncertainties compared to that for the results for the dwarfs studied here. Fig.~\ref{fig_teff_logg} shows the $T_{\rm eff}$ and $log(g)$ distributions of the LAMOST (black), high-resolution spectroscopy (blue) and asteroseismology (red) samples discussed above. For the planet hosts studied in our main work, we only include dwarfs with $log(g)>4$ ($log(g) = 4$ is shown as dashed line). Even though asteroseismology provides higher precision in $log(g)$ than high-resolution spectroscopy, the available seismology stars cover poorly for $log(g) > 4.4$ thus possibly limiting the parameter space for its applicability. Given the limitation for both calibrators, we adopt two sets of $log(g)$ with uncertainties determined from high-resolution spectroscopy and seismology respectively, and we derive the eccentricity distributions using both sets of $log(g)$ uncertainties separately. We have also make similar comparisons with the SPC sample published in 2014 \cite{Buc14}, and there are twice as many common stars available as compared to the 2012 sample \cite{Buc12} used above. The mean differences and standard deviations of $\Delta{T_{\rm eff}}, \Delta{log(g)}, \Delta{\rm [Fe/H]}$ are $15 K \pm 111 K, -0.04 \pm 0.15, -0.05 \pm 0.14$ for the 2014 sample. The standard deviations in $\Delta{T_{\rm eff}}$ and $\Delta{log(g)}$ are similar to the 2012 sample while twice larger in [Fe/H]. This likely due to the new prior in $log(g)$ introduced to the 2014 study \cite{Buc14} and the covariance between $log(g)$ and [Fe/H]. In this work, we adopt the uncertainties derived from earlier SPC sample \cite{Buc12}. Fig.~\ref{fig_kic} shows the comparison in $log(g)$ between the {\it Kepler} input catalog (KIC) \cite{Bro11} and LAMOST/LASP. KIC stellar parameters are widely used for studies of {\it Kepler} planets, including previous studies of {\it Kepler} planet distributions using transit duration statistics \cite{Moo11}. From the comparison, for stars with $log(g)_{\rm LAMOST}>3.5$, the standard deviation of $\Delta{log(g)} = log(g)_{\rm KIC} - log(g)_{\rm LAMOST}$ is 0.3 dex, translating to 0.45 dex in uncertainties for $\rho_*$. In addition, there are serious trends of $\Delta{log(g)}$ as a function of stellar parameters, in particular $log(g)$. The average $\Delta{log(g)}$ is close to zero for stars with close to solar gravity $log(g) \sim 4.4$, but for the stars bigger than the Sun, the KIC $log(g)$ values tend to be under-estimated while for the stars smaller than the Sun, the KIC $log(g)$ tend to be over-estimated. The dynamical range of KIC $log(g)$ is smaller than the spectroscopic $log(g)$. The large dispersion and severe systematic render any statistical studies based on KIC $log(g)$ likely untrustworthy. We determine the stellar mass, radius and density with the LAMOST/LASP $T_{\rm eff}$, $log(g)$, ${\rm [Fe/H]}$ using isochrone fitting on a dense grid of isochrones. We use the 2012 version of the interpolated isochrones from ``The Dartmouth Stellar Evolution Database'' \cite{Dot08} with a range of [Fe/H] from -1.5 dex to 0.5 dex (grid size of 0.02 dex) and stellar age from 1 to 13 Gyrs (grid size of 0.5 dex). We have also applied a separate method \cite{Ser13} with the Dartmouth isochrones and found good consistency between the two. | 16 | 9 | 1609.08633 |
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1609 | 1609.08319_arXiv.txt | We present the results of the analysis of the impact of bulges on the radial distributions of the different types of supernovae (SNe) in the stellar discs of host galaxies with various morphologies. We find that in Sa--Sm galaxies, all core-collapse (CC) and vast majority of SNe Ia belong to the disc, rather than the bulge component. The radial distribution of SNe Ia in S0--S0/a galaxies is inconsistent with their distribution in Sa--Sm hosts, which is probably due to the contribution of the outer bulge SNe Ia in S0--S0/a galaxies. The radial distributions of both types of SNe are similar in all the subsamples of Sa--Sbc and Sc--Sm galaxies. These results confirm that the old bulges of Sa--Sm galaxies are not significant producers of Type Ia SNe, while the bulge populations are significant for SNe Ia only in S0--S0/a galaxies. | 16 | 9 | 1609.08319 |
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1609 | 1609.06304_arXiv.txt | We present a hierarchical Bayesian method for estimating the total mass and mass profile of the Milky Way Galaxy. The new hierarchical Bayesian approach further improves the framework presented by \cite{2015EHW, eadie2016} and builds upon the preliminary reports by \cite{eadieJSM, eadieIAU}. The method uses a distribution function $f(\mathcal{E},L)$ to model the galaxy and kinematic data from satellite objects such as globular clusters (GCs) to trace the Galaxy's gravitational potential. A major advantage of the method is that it not only includes complete and incomplete data simultaneously in the analysis, but also incorporates measurement uncertainties in a coherent and meaningful way. We first test the hierarchical Bayesian framework, which includes measurement uncertainties, using the same data and power-law model assumed in \cite{eadie2016}, and find the results are similar but more strongly constrained. Next, we take advantage of the new statistical framework and incorporate all possible GC data, finding a cumulative mass profile with Bayesian credible regions. This profile implies a mass within $125$kpc of $4.8\times10^{11}\msun$ with a 95\% Bayesian credible region of $(4.0-5.8)\times10^{11}\msun$. Our results also provide estimates of the true specific energies of all the GCs. By comparing these estimated energies to the measured energies of GCs with complete velocity measurements, we observe that (the few) remote tracers with complete measurements may play a large role in determining a total mass estimate of the Galaxy. Thus, our study stresses the need for more remote tracers with complete velocity measurements. | In our two previous papers, \citet[][hereafter Paper I]{2015EHW} and \citet[][hereafter Paper II]{eadie2016}, we estimated the Galaxy's mass and mass profile using a new Bayesian method and the kinematic data of Milky Way globular clusters (GCs) and dwarf galaxies (DGs). Paper I laid the groundwork: we tested the method on simulated data and then applied the method to Milky Way satellite data in a preliminary analysis. A main advantage of the new Bayesian method ws that both complete and incomplete velocity vectors were included in the analysis simultaneously. Furthermore, the tests on simulated data showed that our Galactic mass estimates were insensitive to incorrect velocity anisotropy assumptions. Paper I incorporated an analytic Hernquist model (for simplicity and testing of the method), and used GCs and DGs as tracers of the Milky Way's potential. The satellites were assumed to follow the same spatial distribution as the dark matter. Despite the simplicity of the model, the results were in agreement with many other studies \citep[see][for a comparison figure]{wang2015}. The promising results of Paper I led us to implement an arguably more realistic model for the Milky Way in Paper II, in which the distributions of the dark matter and the Galactic tracers are allowed to differ. The Paper II model uses power-law profiles with different parameters for the dark matter and tracers, and also includes velocity anisotropy as a parameter. This model is explained in detail by \cite{evans1997}, and previous applications of the model to the Milky Way and other galaxies were completed by \citet[][note that the notations vary between Evans' and Deason's papers]{deason2011, deason2012,Deason2012ApJ}. Because the model includes a spatial profile for only a single population of tracers, we used GC kinematic data alone instead of a mixture of DGs and GCs. The results in Paper II suggested a mass estimate for the Milky Way that was significantly lower than the mass found in Paper I under the Hernquist model, but closer in agreement to recent studies which suggest a ``light'' Milky Way \citep[e.g.][]{gibbons2014}. An issue that is not fully addressed in Paper I or II is the inclusion of measurement uncertainty. Measurement uncertainties can differ substantially from object to object, with some tracers having very precise radial velocities or proper motions and others having very imprecise ones. Using a sensitivity analysis, we found in Paper I that measurement uncertainties can play a significant role in the mass estimate of the Galaxy, contributing up to $50\%$ of the uncertainty in the estimate. In addition, we found that certain individual objects had very high leverage. For example, when the single GC Palomar 3 was removed from the analysis, the mass estimate of the Galaxy decreased by more than 12\%. Thus, it seems prudent to include measurement uncertainties in a rigorous and consistent way when estimating the mass and mass profile of the Galaxy. Here, we substantially improve upon Paper II by introducing a \emph{hierarchical} Bayesian method that includes the measurement uncertainties of proper motions and line-of-sight velocities in a measurement model. Preliminary tests of this method have been reported by \citet*{eadieIAU} and \citet*{eadieJSM} using the Hernquist model and data from GCs and DGs in Paper I, but here we apply the arguably more realistic tracer model from Paper II, and also use all of the available GC data. | We have described a coherent, hierarchical Bayesian method for estimating the mass profile of the Milky Way Galaxy, and applied this method to the Galaxy using GC data. This statistical framework allows us to take full advantage of all of the available GC kinematic data, and also provides a meaningful and coherent probabilistic way to incorporate measurement uncertainties. Using the assumptions of the power-law model (Section~\ref{sec:df}), the hierarchical framework for including uncertainties (Section~\ref{sec:methods}), and the prior distributions (Section~\ref{sec:priors}), and confronting this model with data from 143 GCs around the Milky Way, we arrive at a cumulative mass profile for the Galaxy with uncertainties (Figure~\ref{fig:143gcs}) and a mass estimate within $125\kpc$ of $4.8\times10^{11}\msun$ (the 95\% Bayesian credible regions are $(4.0-5.8)\times10^{11}\msun$). When we extrapolate the mass profile to the virial radius $(\approx 179\kpc)$, we find $M_{vir}=6.2\times10^{11}\msun$ with a 95\% Bayesian credible region of $(4.7-7.8)\times10^{11}\msun$. This mass estimate is notably lower than those in other studies. The statistical framework presented here will be highly useful and appropriate for other tracer objects around the Milky Way, such as halo stars and DGs. Using our approach with data sets from large programs, such as Gaia, could yield a well-constrained mass estimate for the Galaxy. Incorporating large data sets in this analysis will present some computational challenges, but given the effectiveness of our MCMC sampler we are confident that this will be a tractable problem through parallelization. The first order of business, however, is to better understand what tracer populations will provide the most information about the Milky Way's gravitational potential. Thus, in our next paper Eadie, Keller, et al. (in preparation) we perform a series of blind tests of simulated data of Milky Way-type galaxies that were created through hydrodynamical simulations \citep{keller2015,keller2016}, and investigate which tracer information is necessary for constraining the mass of the Milky Way. | 16 | 9 | 1609.06304 |
1609 | 1609.06903_arXiv.txt | We use the concept of the spiral rotation curves universality to investigate the luminous and dark matter properties of the dwarf disc galaxies in the local volume~(size~$\sim11$~Mpc). Our sample includes 36 objects with rotation curves carefully selected from the literature. We find that, despite the large variations of our sample in luminosities ($\sim$ 2 of dex), the rotation curves in specifically normalized units, look all alike and lead to the lower-mass version of the universal rotation curve of spiral galaxies found in Persic et al. We mass model the double normalized universal rotation curve $V(R/R_{opt})/V_{opt}$ of dwarf disc galaxies: the results show that these systems are totally dominated by dark matter whose density shows a core size between 2 and 3 stellar disc scale lengths. Similar to galaxies of different Hubble types and luminosities, the core radius $r_0$ and the central density $\rho_0$ of the dark matter halo of these objects are related by $ \rho_0 r_0 \sim 100\hspace{0.1cm} \mathrm{M_\odot pc^{-2}}$. The structural properties of the dark and luminous matter emerge very well correlated. In addition, to describe these relations, we need to introduce a new parameter, measuring the compactness of light distribution of a (dwarf) disc galaxy. These structural properties also indicate that there is no evidence of abrupt decline at the faint end of the baryonic to halo mass relation. Finally, we find that the distributions of the stellar disc and its dark matter halo are closely related. | It is widely believed that only 15 per cent of the total matter in the Universe is in the form of ordinary baryonic matter. Instead the other 85 per cent is provided by dark matter (DM), which is detectable, up to now, only through its gravitational influence on luminous matter. The paradigm is that DM is made by massive gravitationally interacting elementary particles with extremely weak, if not null interaction via other forces \citep[e.g.,][]{white1982, jungman}. In this framework the well known ($\Lambda$)CDM scenario, successfully describing the large structure of the Universe, has emerged \citep{kolb90}: accurate N-body simulations have found that the DM density profile of the virialized structures such as galactic halos is universal and well described by the Navarro-Frenk-White profile \citep[hereafter NFW;][]{nfw}. However, at the galactic scales, this scenario has significant challenges. First, the apparent mismatch between the number of the detected satellites around the Milky Way and the predictions of the corresponding simulations, known as the "missing satellite problem" \citep{klypin99,moore99}, which also occurs in the field galaxies \citep{zavala09,papastergis11,klypin15}. This discrepancy widens up when the masses of the detected satellites are compared to those of the predicted subhalos (i.e. "too big to fail problem") \citep[see][]{boylan-kolchin12,ferrero12,garrisonkimmel,papastergis15}. Furthermore, there is the "core-cusp" controversy: the inner DM density profiles of galaxies generally appear to be cored, and not cuspy as predicted in the simplest ($\Lambda$)CDM scenario \citep[e.g.,][to name few]{salucci01,deblok02,gentile05,weinberg13, bosma04,simon05,gentile04,gentile07,donato09,delpopolo09,oh11}. These apparent discrepancies between the observations and the predictions of the DM-only simulations suggest to either abandon the ($\Lambda$)CDM scenario in favour of the others \citetext{e.g., selfinteracting DM \citealp{vogelsberger14,elbert15} or warm DM \citealp{devega13,devega133,lovell14,devega14}} or upgrade the role of baryonic physics in the galaxy formation process. The latter can be done including strong gas outflows, triggered by stellar and/or AGN feedback that are thought to strongly modify the original ($\Lambda$)CDM halo profiles out to a distance as large as the size of the stellar disc \citep[e.g.,][]{navarro96,read05, mashchenko06,pontzen12,pontzen14,dicintio14}. Although these issues are present in galaxies of any luminosity, however in low luminosity systems they emerge more clearly and appear much more difficult to be resolved within the ($\Lambda$)CDM scenario. Thus, galaxies with I-band absolute magnitude $M_I \gtrsim -17$ play a pivotal role in that, observationally, these objects are dark matter dominated at all radii. Moreover in the ($\Lambda$)CDM scenario they may be related to the building blocks of more massive galaxies. In light of this the importance of dwarf spheroidal galaxies in various DM issues is well known \citep[see, e.g.,][]{gilmore07}. However, down to $M_{I} \sim -11$ there is no shortage of rotationally supported late-type systems, although a systematic investigation is lacking. These rotationally supported systems have a rather simple kinematics suitable for investigating the properties of their dark matter content. In normal spirals, one efficient way to represent and model their rotation curves (RCs) comes from the concept of a universal rotation curve (URC). Let us stress that the concept of universality in RCs does not mean that all of them have a unique profile, but that all the RCs of $10^9$ local spirals (within $z \simeq 0.1$) can be described by a same function of radius, modulated by few free parameters. They depend on the galaxy's global properties, namely magnitude (or luminosity/mass) and a characteristic radius of the luminous matter\footnote{i.e. optical radius $R_{opt}$ defined as the radius encompassing 83 per cent of the total luminosity.} so that: $V(R)= V(R,L,R_{opt})$. This concept, implicit in \citet{rubin85}, pioneered by \citet{persic91}, set by \citet{pss96} (PSS, Paper I) and extended to large galactocentric radii by \citet{salucci2007} has provided us the mass distribution of (normal) disc galaxies in the magnitude range $-23.5 \lesssim M_I \lesssim -17$.\footnote{Extensions of the URC to other Hubble types are investigated in \citet{salucci97,noordermeer07}.} This curve $V(R,L,R_{opt})$, therefore, encodes all the main structural properties of the dark and luminous matter of every spiral \citep[PSS,][]{yegorova07}. In this paper, we work out to extend the RCs universality down to low-mass systems and then, to use it to investigate the DM distribution in dwarf disc galaxies. Noticeably, for this population of galaxies the approach of stacking the available kinematics is very useful. In fact, presently, for disc systems with the optical velocity $V_{opt}\lesssim 61\hspace{0.1cm} km/s$, some kinematical data have become available (galaxies of higher velocities are investigated in the PSS sample). However, there are not enough individual {\it high quality high resolution extended} RCs to provide us with a solid knowledge of their internal distribution of mass. Instead, we will prove that the 36 selected in literature {\it good quality good resolution reasonably extended} RCs (see below for these definitions), once coadded, provide us with a reliable kinematics yielding to their mass distribution. In this work, we construct a sample of dwarf discs from the local volume catalog (LVC) \citep[][hearafter K13]{karachentsev13}, which is $\sim 70$ per cent complete down to $M_{B}\approx-14$ and out to 11 Mpc, with the distances of galaxies obtained by means of primary distance indicators. Using LVC, we go more than 3 magnitudes fainter with respect to the sample of spirals of PSS. Moreover the characteristics of the LVC guarantee us against several luminosity biases that may affect such faint objects. The total number of objects in this catalog is $\sim$900 of which $\sim180$ are dwarf spheroidal galaxies, $\sim500$ are dwarf disc galaxies and the rest are ellipticals and spirals. All our galaxies are low mass bulgeless systems in which rotation, corrected for the pressure support, totally balances the gravitational force. Morphologically, they can be divided into two main types: gas-rich dwarfs that are forming stars at a relatively-low rate, named irregulars (Irrs) and starbursting dwarfs that are forming stars at an unusually high rate, named blue compact dwarfs (BCD). The dwarf Irr galaxies are named "irregulars" due to the fact that they usually do not have a defined disc shape and the star formation is not organized in spiral arms. However, some gas-rich dwarfs can have diffuse, broken spiral arms and be classified as late-type spirals (Sd) or as Magellanic spirals (Sm). The starbursting dwarfs are classified as BCD due to their blue colours, high surface brightness and low luminosities. Notice that it is not always easy to distinguish among these types since the galaxies we are considering often share the same parameters space for many structural properties \citep[e.g.,][]{kormendy85,binggeli94,tolstoy09}. In this paper, we neglect the morphology of the baryonic components as long as their stellar disc component follows a radially exponential surface density profile; the identifiers of a galaxy are $V_{opt}$, its disc length scale $R_D$ and its $K$-band magnitude $M_K$ that can be substituted by its disc mass. We refer to these systems of any morphologies and $M_K \gtrsim -18 $ as dwarf discs (dd). In order to compare galaxy luminosities in different bands, we write down the dd relations between the magnitudes in different bands $<B-K>\simeq2.35$ \citep{jarrett03} and $<B-I>\simeq1.35$ \citep{fukugita95}. The plan of this paper as follows: in Section 2 we describe the sample that we are going to use; in Section 3 we introduce the analysis used to build the synthetic RC; in Section 4 we do the mass modelling of the synthetic RC; in Section 5 we denormalize the results of the mass modelling in order to describe individually our sample of galaxies and then we define their scaling relations; in Section 6 we discuss our main results. | We have compiled literature data for a sample of dd galaxies in the local volume ($\lesssim 11 Mpc$) with HI and $H_{\alpha}$ RCs. Then for these galaxies we establish the corresponding URC in normalized and physical units and investigate the related dark and luminous matter properties, not yet studied statistically in these objects. Our sample spans $\sim 2$ decades ($\sim 10^6-3 \times 10^8 L_{\odot}$) in luminosity, which coincides with the faint end of the luminosity function of disc galaxies. In magnitude extension is as large as the whole range of normal spirals usually investigated in terms of URC. For example, the galaxies in the sample are up to $\sim4$ magnitudes fainter than the lowest limit in the PSS sample. We find that, the large variations of our sample in luminosity and morphologies require double normalization. Notably, after this noralization we have that all RCs in double normalized units are alike. This implies that the structural parameters of the dark and luminous matter of these galaxies do not have any explicit dependence on luminosity except those coming from the normalizing process. Additionally, the good agreement of our coadded RC with that of the first PSS's luminosity bin indicates that in such small galaxies the mass structure is already dominated by a dark halo with a density core as big as a stellar disc. Then by applying to the double normalized rotation curve the standard $\chi^2$ mass modelling, we tested three DM density profiles. Wherein the NFW profile fails to reproduce the coadded curve, while the Burkert and DC14 profiles show excellent quality fits with $\chi_{red}^2<1.$ This result points towards the cored DM distribution in dwarf disc galaxies. The same conclusion was drawn in the papers on Things and Little Things samples \citep[see, e.g.,][]{oh11,oh15}, where the authors found for their dwarfs much shallower inner logarithmic DM density slopes than those predicted by DM-only ($\Lambda$)CDM simulations. The present analysis has the advantages of bigger statistics, but above all, is immune from systematics that can affect the mass modelling of individual galaxies. We also defined, galaxy by galaxy, the values of the dark and luminous matter structural parameters. Surprisingly, a new actor enters the scene of the distribution of matter in galaxies, the compactness of the stellar component, which allows us to establish the URC in these low-mass galaxies. However, in order to understand better the role of this compactness it is required to investigate galaxies in the transition regime which appears at about $V(R_{opt})\simeq60 km/s $. As a consequence of the derived mass distributions, there is no evidence for the sharp decline in the baryonic to halo mass relation. Similar result, for dwarf galaxies in the field, was found by \citet{ferrero12}. Nevertheless, notice that in DC14 case the estimated baryonic mass is slightly lower than that of the URC mass model, which brings it closer to the abundance matching relation inferred from e.g. \citet{papastergis12}. Furthermore, since the fit resulting from the baryon-influenced DC14 profile has a lower value of the disc mass, it agrees, within the errors, with the extrapolation of the $M_D-M_{vir}$ relation derived from the abundance matching by \citet{moster13}. Let us also recall, that the \citet{dicintio14} model has been already tested against observations in works by \citet{katz16} and \citet{pace16}. Although both groups use similar methods, the drawn conclusions are different \citep[see also][]{read16}. Therefore, the consistency level between observations and the ($\Lambda$)CDM model of galaxy formation, specifically the abundance matching technique deserves further investigation. At the same time, the S-shape of $M_{vir}-M_{bar}$ relation may be interpreted as different physical mechanism occurring along the mass sequence of disc galaxies. Theoretically, it has been shown that the energetics of star formation differs among different galaxies with a characteristic dependence on the halo-to-stellar mass ratio \citep{dicintio14,chan15} and possibly also on star formation history \citep{onorbe15}. We remark that we found that the DM and the stellar distribution follow each other very closely out to the level for which, in log-log frame, the compactness of the stellar disc is proportional to that of the DM halo. We believe that here we are touching a crucial aspect in the DM issue, whose investigation, however, much exceed the scope of this paper. Finally we would like to stress that the results of this work \citep[and of the previous works, see, e.g., ][]{donato09,gentile09} indicate that the DM around galaxies should be considered, rather than the final product of the cosmological evolution of the massive components of the Universe, galaxies, today, but as the direct manifestation of one of its most extraordinary mysteries. | 16 | 9 | 1609.06903 |
1609 | 1609.08018.txt | Near-Earth asteroids (NEAs) in the 1--100 meter size range are estimated to be $\sim$1,000 times more numerous than the $\sim$15,000 currently-catalogued NEAs, most of which are in the 0.5--10 kilometer size range. Impacts from 10--100 meter size NEAs are not statistically life-threatening but may cause significant regional damage, while 1--10 meter size NEAs with low velocities relative to Earth are compelling targets for space missions. We describe the implementation and initial results of a real-time NEA-discovery system specialized for the detection of small, high angular rate (visually-streaked) NEAs in Palomar Transient Factory (PTF) images. PTF is a 1.2-m aperture, 7.3-deg$^2$ field-of-view optical survey designed primarily for the discovery of extragalactic transients ({\it e.g.}, supernovae) in 60-second exposures reaching $\sim$20.5 visual magnitude. Our real-time NEA discovery pipeline uses a machine-learned classifier to filter a large number of false-positive streak detections, permitting a human scanner to efficiently and remotely identify real asteroid streaks during the night. Upon recognition of a streaked NEA detection (typically within an hour of the discovery exposure), the scanner triggers follow-up with the same telescope and posts the observations to the Minor Planet Center for worldwide confirmation. We describe our ten initial confirmed discoveries, all small NEAs that passed 0.3--15 lunar distances from Earth. Lastly, we derive useful scaling laws for comparing streaked-NEA-detection capabilities of different surveys as a function of their hardware and survey-pattern characteristics. This work most directly informs estimates of the streak-detection capabilities of the Zwicky Transient Facility (ZTF, planned to succeed PTF in 2017), which will apply PTF's current resolution and sensitivity over a 47-deg$^2$ field-of-view. | A near-Earth asteroid (NEA) is by definition any asteroid with perihelion $q<1.3$ AU and aphelion $Q>0.983$ AU. From the largest NEA (of diameter $D\approx30$ km) down to $D\approx0.5$ km in size---for which the known population is largely complete---the cumulative size-frequency distribution (Figure~1) goes roughly as $N(D)\propto D^{-2}$, where $N(\text{0.5 km})\approx 10^4$. Harris (\citeyear{har08}, \citeyear{har13}) presents these statistics, and describes how the original `Spaceguard' goal to catalog 90\% of all $D>1$ km NEAs was achieved by the mid-2000s, while the current congressional mandate is to find 90\% of all $D>140$ m NEAs by 2020. The incrementally-decreasing target size in the NEA census has been mostly motivated by risk mitigation. Over the quarter-century that began with our realization of an asteroid's role in the dinosaurs' extinction ({\it e.g.}, \citealp{alv80}) through to our fulfillment of the 1-km Spaceguard goal, the estimated risk of an individual's death from asteroid impact---initially believed comparable to that of a commercial airplane accident---dropped by an order of magnitude. Surveying to the currently recommended $D>140$ m can decrease this risk by yet another order of magnitude \citep{har08}. Hence, discovery of $D<100$ m NEAs will contribute only minimally to any further significant reduction in the risk of death to any individual. However, events like the Tunguska and Chelyabinsk airbursts \citep{bro13} could have caused caused a significant numbers of deaths if the impact parameters had been different and did cause environmental or other damage. This suggests that impacts from 10--100 m objects qualify as `natural disasters' that merit advance warning, and possibly prevention via space-based manipulation of hazardous NEAs. However, the size-frequency distribution informing these estimates is uncertain across orders of magnitude in impactor size, and constrained on the small end ($D\lesssim10$ m) by infrasound detections of bolide fluxes (\citealp{sil09}). Besides impact mitigation ({\it e.g.}, \citealp{ahr92}; \citealp{lul05}), other space-based activities benefiting from small NEA discoveries include in-situ compositional studies \citep{mue11} and resource utilization \citep{elv14}. NEAs have also been declared a major component of NASA's manned spaceflight program \citep{oba10}. NEA rendezvous feasibility depends critically on mission duration and fuel requirements, these in turn are functions of the NEA's orbit and relative velocity ($\Delta v$) with respect to Earth (\citealp{sho78}; \citealp{elv11}). Robotic missions that have been proposed may facilitate or complement the manned program. One example was the proposed Asteroid Retrieval Mission (ARM; \citealp{bro10}), which evolved into the NASA Asteroid Redirect Mission. At the time of the writing of this paper, NASA has chosen to retrieve a boulder from a larger asteroid, rather than retrieve an entire small asteroid. Natural temporary capture of meter-scale NEAs into Earth-centric orbits, if confirmed via the discovery of `mini-moons' (\citealp{gra12}; \citealp{bol14}), would present another appealing class of targets. Cleary, discovery of 1--100 m sized NEAs is motivated by different (and more diverse) applications than those which have driven the census of larger NEAs. The discovery \emph{method} often likewise differs. Most large NEAs were found via the `tracklet' method of linking several serendipitously-observed positions within a night or across several nights. This is the basis of `MOPS'-like detection software ({\it e.g.,} \citealp{den13}), which in its present state is most efficient at detecting NEAs moving slower than $~\sim$5 deg/day \citep{jed13}. Below this rate, an NEA's individual detections are nearly point-like for typical survey exposure times ({\it e.g.}, 30--60s), and sufficiently localized on the sky given typical intra-night pointing cadences ({\it e.g.}, 15--45 minutes). Hazardous NEAs occupy a range of orbits with moderate eccentricities, and so they spend most of their time far from the Earth and Sun, where their sufficiently slow apparent motions allow them to be easily detected with this technique. Searches for NEAs with DECam (Allen et al. 2013 \& 2014) using conventional MOPS techniques on a large telescope appear to be particularly promising for finding small NEAs. Recent searches using NEOWISE (Mainzer et al. 2014) have also yielded interesting new detection of small NEAs. In contrast, the method of \emph{streak detection} enables discovery of much smaller and closer ({\it i.e.}, brighter and faster-moving) NEAs. Whereas slower-moving NEAs can be mistaken for main-belt asteroids, streaked asteroids having an angular rate of larger than 1 deg/day are likely to be NEAs. Unlike the tracklet method, discovery via streak detection is possible on the basis of a \emph{single} exposure via recognition of the streak morphology, meaning repeat visits to the same patch of sky are unnecessary and more area can be searched. Lastly, NEAs that are detected as streaks are typically closer and therefore are 2 to 3 times brighter than those found by the tracklet method when they are tracked non-sidereally, making them more convenient for follow-up from dedicated (including amateur-class) facilities once an approximate orbit is determined. Survey-scale application of the streak-detection method for NEA discovery was pioneered by \cite{hel79} using photographic plates on the Palomar 18-inch Schmidt telescope in the 1970s. \cite{rab91} was the first to apply this method with CCD detectors in near real-time with the Spacewatch survey. Combining Spacewatch's streaked NEA detections ({\it e.g.}, \citealp{sco91}) with its tracklet-detected NEAs \citep{jed95} produced a debiased NEA number-size distribution \citep{rab00} spanning four orders of magnitude in size (10 km $>D>$ 1 m). In 2005, Spacewatch initiated a public Fast Moving Object (FMO) program (\cite{mcmillan+05}) that allowed public access to the survey's imagery on the internet so that members of the public could scan images and report detections of streaks. \begin{figure}[t] \centering \includegraphics[scale=1]{fig1_nea_paper.pdf} \caption{Cumulative NEA population distribution models compared to discovered objects. Plot adapted from a figure in \cite{ru14}.} \end{figure} \begin{figure}[t] \centering \includegraphics[scale=0.85]{fig2_nea_paper.pdf} \caption{Number of streaked NEA discoveries as a function of time (bins include 1990--2014) and survey, where `streaked' is here defined as any discovered streak greater than 10 seeing-widths in length (see text for details).} \end{figure} Figure~2 breaks down the number of streaked NEA discoveries as a function of time and survey, from 1991 through 2014-Oct. Here `streaked' is taken to mean any detection wherein the length of the imaged streak is greater than 10 seeing widths. The counts in Figure~2 were compiled by first retrieving all NEA discovery observations from the Minor Planet Center (MPC) database and then using JPL's HORIZONS service \citep{gio96} to compute the on-sky motion at the discovery epoch. These rates were then converted into streak lengths in units of seeing widths, where the continental surveys all have assumed $2''$ seeing and Pan-STARRS has assumed $1''$ seeing. The assumed exposure times come mostly from a table in \cite{lar07}, except for PTF and Pan-STARRS, which have assumed exposure times of 60s and 45s, respectively. \begin{figure*}[t] \centering \includegraphics[scale=0.53]{fig3_nea_paper.pdf} \caption{Some known small NEAs serendipitously detected by PTF. These observations were retrieved solely by computing these known objects' positions at the epochs of archival PTF images and visually verifying the streak's presence. All images are 200$''$ $\times$ 200$''$ with linear contrast scaling from $-0.5$$\sigma$ to 7$\sigma$.} \vspace{10pt} \end{figure*} Before 2005, Spacewatch was the only contributor of significantly streaked NEA discoveries, and it is also the most prolific streaked-NEA discoverer overall. There are two likely reasons for this: (1) Spacewatch's relatively long 120s exposure time, and (2) the active role of a human screener (`observer') during data collection, as documented by \cite{rab91}. The Catalina Sky Survey also has a dedicated human operator to scan candidates and conduct same-night follow-up \citep{lar07}, which explains its similarly consistent contribution of streaked discoveries. Some major NEA surveys of the past two decades \emph{not} contributing to the streaked discoveries in Figure~2 include LINEAR \citep{sto00}---likely because of its short 8s exposures, as well as NEAT \citep{pra99} and LONEOS \citep{sto02}---which to our knowledge lacked real-time human interaction with their respective data flows. The years 2013 and 2014 marked a clear upturn in the discovery of streaked NEAs. The purpose of this paper is to document a new streak-discovery pipeline which has contributed in part to this increased discovery rate. | 16 | 9 | 1609.08018 |
|
1609 | 1609.00187_arXiv.txt | {We report on {\igr}, {\swift} and {\xmm} observations of {\src} performed during the outburst that occurred between March 23 and April 25, 2015. The source reached a peak flux of $0.7(2)\times10^{-9}$~{\fluxcgs} and decayed to quiescence in approximately a month. The X-ray spectrum was dominated by a power-law with photon index between 1.6 and 1.8, which we interpreted as thermal Comptonization in an electron cloud with temperature $> 20$ keV . A broad ($\sigma\simeq1$~keV) emission line was detected at an energy ($E=6.9^{+0.2}_{-0.3}$ keV) compatible with the K-$\alpha$ transition of ionized Fe, suggesting an origin in the inner regions of the accretion disk. The outburst flux and spectral properties shown during this outburst were remarkably similar to those observed during the previous accretion event detected from the source in 2009. Coherent pulsations at the pulsar spin period were detected in the {\xmm} and {\igr} data, at a frequency compatible with the value observed in 2009. Assuming that the source spun up during the 2015 outburst at the same rate observed during the previous outburst, we derive a conservative upper limit on the spin down rate during quiescence of $3.5\times10^{-15}$~Hz~s$^{-1}$. Interpreting this value in terms of electromagnetic spin down yields an upper limit of $3.6\times10^{26}$~G~cm$^{3}$ to the pulsar magnetic dipole (assuming a magnetic inclination angle of $30^{\circ}$). We also report on the detection of five type-I X-ray bursts (three in the {\xmm} data, two in the {\igr} data), none of which indicated photospheric radius expansion.} | Accreting millisecond pulsars (AMSPs hereafter) are neutron stars (NS) that accrete matter transferred from a low mass ($M_2\la M_{\odot}$) companion star \citep{wijnands1998}; their magnetospheres are able to truncate the disk in-flow and channel the in-falling matter to the regions of the surface close to the magnetic poles, producing coherent pulsations in the X-ray light curve. The extremely quick rotation of AMSPs is attained during a previous Gyr-long phase of sustained mass accretion, and these sources are considered the most immediate progenitors of millisecond radio pulsars \citep{bisnovatyikogan1974,alpar1982,radhakrishnan1982,archibald2009,papitto2013nat}. So far, accretion driven coherent pulsations have been detected from 17 transient low-mass X-ray binaries \citep[see, e.g.][for a review]{patruno2012c}. Pulsations from 15 of these objects have been observed during relatively bright ($L_X\approx \mbox{few}\times 10^{36}$ erg s$^{-1}$) X-ray outbursts lasting few-weeks to few-months. In two cases, PSR J1023+0038 and XSS J12270-4859 \citep{archibald2015,pap2015}, the coherent signal was detected when the source was at a much lower luminosity level ($L_X\approx \mbox{few}\times 10^{33}$ erg s$^{-1}$); these two sources were also detected as radio pulsars during X-ray quiescence \citep{archibald2009,roy2015}, similar to the AMSP IGR J18245--2452 \citep{papitto2013nat}. {\src} was discovered by {\igr} during an outburst in September 2009 \citep{baldovin2009,bozzo2010}. The detection of 4.1 ms coherent pulsations in the {\it Rossi X-ray Timing Explorer} light curve allowed the identification of the source as an AMSP in a binary system with a 3.47 hr orbital period \citep{markwardt2009}. The measured pulsar mass function indicated a main-sequence companion star with a mass between 0.15 and 0.44 $M_{\odot}$ \citep{papitto2010}. At discovery, 18 type-I X-ray bursts were observed from the pulsar \citep{altamirano2010,bozzo2010,papitto2010,falanga2011}. Evidence of photospheric radius expansion was not displayed by any of these bursts, and an upper limit to the source distance (6.9 kpc) was provided by \citet{altamirano2010}. A new outburst of {\src} was detected by {\igr} on March 23, 2015 \citep{bozzo2015,bozzo2015b}. Here, we report on a series of {\igr} and {\swift} observations performed throughout the event, as well as on a {\xmm} Target of Opportunity observation performed three days after the onset of the outburst. | We have presented an analysis of {\xmm}, {\inte} and {\swift} observations performed during the outburst detected from {\src} during early 2015, the second observed from the pulsar after the discovery outburst in 2009. The outburst profile, spectral and burst properties were remarkably similar to those observed during the last accretion event detected in 2009, suggesting that the properties of the accretion flow did not change much in the two episodes. The frequency of the coherent signal detected by {\xmm} three days after the beginning of the outburst was compatible with the value measured by the same observatory during the 2009 outburst. Therefore, a firm assessment of the spin evolution of the pulsar during the time elapsed between the two outbursts was not possible. However, taking into account the accretion driven spin up observed during the 2009 outburst, and assumed for the 2015 outburst, we derived an upper limit of $3.5\times10^{-15}$~Hz~s$^{-1}$ on the spin down rate during quiescence. Electromagnetic spin-down of a NS with a magnetic field weaker than $3.5\times10^{8}$ G (at the equator of the NS and assuming a magnetic inclination of $30^{\circ}$) can account for this inferred spin-down. A magnetic field of the same order has been inferred also for other AMSPs. Observations of future outbursts will allow to derive more stringent constraints on the long-term spin evolution of the pulsar, as well as enabling a first estimate of the orbital period derivative. | 16 | 9 | 1609.00187 |
1609 | 1609.09290_arXiv.txt | The F3 main sequence star KIC 8462852 (Boyajian's Star) showed deep (up to 20\%) day-long brightness dips of unknown cause during the 4 years of the \textit{Kepler} mission. A 0.164~mag (16\%) dimming between 1890 and 1990 was claimed, based on the analysis of photographic plates from the Harvard Observatory. We have gathered an independent set of historic plates from Sonneberg Observatory, Germany, covering the years 1934 -- 1995. With 861 magnitudes in B, and 397 magnitudes in V, we find the star to be of constant brightness within 0.03~mag per century (3\%). Consistent outcomes are found using by-eye estimates of the best 119 plates. Results are supported by data from Sternberg Observatory, Moscow, which show the star as constant between 1895 and 1995. The previously claimed century-long dimming is inconsistent with our results at the $5\sigma$-level, however the recently reported modest dimming of 3\% in the Kepler data is not inconsistent with our data. We find no periodicities or shorter trends within our limits of 5\% per 5-year bin, but note a possible dimming event on 24 Oct 1978. | The \textit{Kepler} Space Telescope's exquisite photo\-metry has allowed for the detection of more than a thousand exoplanets \citep{2010Sci...327..977B}. The data have also been used for a wide range of stellar studies \citep[e.g.,][]{2011ApJ...743..143H,2015ApJ...798...42H}. One very peculiar object is Boyajian's Star (KIC 8462852, TYC 3162-665-1). This is an F3 main-sequence star, which was observed by the NASA Kepler mission from April 2009 to May 2013. An analysis by \citet{2016MNRAS.457.3988B} shows inexplicable series of day-long brightness dips up to 20\%. The behavior has been theorized to originate from a family of large comets \citep{2016ApJ...819L..34B}, or signs of a Dyson sphere \citep{2016ApJ...816...17W}. Subsequent analysis found no narrow-band radio signals \citep{2016ApJ...825..155H} and no periodic pulsed optical signals \citep{2016ApJ...825L...5S, 2016ApJ...818L..33A}. The infrared flux is equally unremarkable \citep{2015ApJ...815L..27L, 2015ApJ...814L..15M, 2016MNRAS.458L..39T}. Recently, the star has been claimed to dim by $0.164$~mag ($\sim16\%$) between 1890 and 1990 \citep{2016ApJ...822L..34S}, and lost $\sim3\%$ of brightness during the 4.25yrs of Kepler mission \citep{2016ApJ...830L..39M}. The century-long dimming has been challenged by \citet{2016ApJ...825...73H} and \citet{2016arXiv160502760L}, who find systematic noise to be the most likely cause. Further effort has been put in regular cadence photometry of the star from ground-based facilities. In parallel to observations from AAVSO amateurs, a community-financed initiative has collected $100,000$ USD, and currently spends the money purchasing telescope time\footnote{\url{https://www.kickstarter.com/projects/608159144/the-most-mysterious-star-in-the-galaxy}}. The distance estimate from Gaia's DR1 parallax ($391.4\substack{+122.1 \\ -75.2}$~pc, \citet{2016arXiv160905492H}) can not constrain the absolute magnitude of the star. To resolve the controversy whether this star has dimmed considerably over 130 years (and perhaps more so earlier), it is crucial to gather more data. | We have analyzed photo\-metry from the Sonneberg observatory (1934 -- 1995), and other sources (1895 -- 1995), and found strong evidence that Boyajian's Star has not dimmed, but kept a constant flux within a few percent. Variations on 5-year scale are found to be of order 5\% or less. We found a possible $>=8\%$ dip on 24 Oct 1978, resulting in a period of the putative occulter of 738 days. | 16 | 9 | 1609.09290 |
1609 | 1609.00378_arXiv.txt | We present spatially--resolved H{\sc i} kinematics of 32 spiral galaxies which have Cepheid or/and Tip of the Red Giant Branch distances, and define a calibrator sample for the Tully-Fisher relation. The interferometric H{\sc i} data for this sample were collected from available archives and supplemented with new GMRT observations. This paper describes an uniform analysis of the H{\sc i} kinematics of this inhomogeneous data set. Our main result is an atlas for our calibrator sample that presents global H{\sc i} profiles, integrated H{\sc i} column--density maps, H{\sc i} surface density profiles and, most importantly, detailed kinematic information in the form of high--quality rotation curves derived from highly--resolved, two--dimensional velocity fields and position--velocity diagrams. | The Tully-Fisher relation (TFr) is one of the main scaling relations for rotationally supported galaxies, describing an empirical correlation between the luminosity or visible mass of a spiral galaxy and its rotational velocity or dynamical mass \citep{tf77}. Numerous observational studies of the statistical properties of the TFr have been undertaken in the past. Generally, their aim was to find observables that reduce the scatter in the relation, and thus improve the distance measure to spiral galaxies. An accurate distance measure is necessary to address some of the main cosmological questions pertaining to the local Universe, such as the value of the Hubble constant, the local large--scale structure and the local cosmic flow field, e.g. \citealt{TC12}, hereafter \citetalias{TC12}. Use of the implied scaling law and understanding its origin is one of the main challenges for theories of galaxy formation and evolution. In particular the detailed statistical properties of the TFr provide important constraints to semi--analytical models and numerical simulations of galaxy formation and evolution. It is an important test for theoretical studies to reproduce the slope, scatter and the zero point of the TFr in different photometric bands simultaneously. Recent cosmological simulations of galaxy formation have reached sufficient maturity to construct sufficiently realistic galaxies that follow the observed TFr \citep{mar14, vog14, schaye15}. An important outstanding issue, however, is how to connect the results from simulations to the multitude of observational studies, which often do not agree with each other \citep{mcg12, zar14}. Moreover, while performing these comparison tests, it is important to ensure that the galaxy parameters derived from simulations and observations have the same physical meaning. For example, for simulated galaxies the circular velocity of a galaxy is usually derived from the potential of the dark matter halo, while the observed velocity, based on the width of the global H{\sc i} profile obtained with single--dish telescopes, is a good representation for the dark matter halo potential only in rather limited cases. Note that the rotational velocity of a galaxy, inferred from the width of its global H{\sc i} profile, can be directly related to the dark matter potential only when the gas is in co--planar, circular orbits and the rotation curve has reached the turnover to a constant velocity at radii properly sampled by the gas disk. However, this is often not the case as galaxies display a variety of rotation curves shapes, which can not be characterised with a single parameter such as the corrected width of the global H{\sc i} line profile. Furthermore, the number of observational studies that take into account the two--dimensional distribution of H{\sc i} and the detailed geometry and kinematics of a galaxy's disk is rather limited. With the developments of new methods and instruments, the observed scatter in the Tully-Fisher relation has decreased significantly, mainly due to more accurate photometric measurements. Photographic magnitudes \citep{tf77} were improved upon by CCD photometry. The advent of infra--red arrays shifted photometry to the JHK bands, minimizing extinction corrections, and then to space--based infra--red photometry with the Spitzer Space Telescope \citep{wer04}, which better samples the older stellar population and maximizes photometric calibration stability. As was shown by \citet{sorce12}, the measurement errors on the total $3.6 \mu m$ luminosity are reduced to a point where they no longer contribute significantly to the observed scatter in the Tully--Fisher relation. Hence, other measurement errors, such as uncertainties on a galaxy's inclination and/or rotational velocity, combined with a certain intrinsic scatter are responsible for the total observed scatter. \begin{figure} \begin{center} \includegraphics[scale=0.75]{histMT.pdf} \caption{ The distribution of morphological types of our sample galaxies. \label{fig_ttype}} \end{center} \end{figure} It was shown by \citet{v01}, with a study of nearly equidistant spiral galaxies in the Ursa Major cluster with deep K'--band photometry, that the observed scatter in the Tully--Fisher relation can be reduced significantly when the measured velocity of the outer, flat part of the H{\sc i} rotation curve (V$_{\rm flat}$) is used instead of the rotational velocity as estimated from the width of the global H{\sc i} profile. Indeed, the spatially-resolved, two--dimensional velocity fields of gas disks, obtained with radio interferometric arrays, often reveal the presence of warps, streaming motions and kinematic lopsidedness of the gas disks, while the global H{\sc i} profiles do not allow the identification of such features. Therefore, the width of a global H{\sc i} line profile, albeit easily obtained with a single--dish telescope, may not be an accurate representation of V$_{\rm flat}$. Moreover, the analysis of spatially resolved H{\sc i} kinematics shows that the rotation curves of spiral galaxies are not always flat and featureless. Rotation curves may still be rising at the outermost observable edge of the gas disk, in which case the width of the global H{\sc i} profile provides a lower limit on V$_{\rm flat}$. Rotation curves may also be declining beyond an initial turnover radius, in which case V$_{\rm flat}$ derived from the rotation curve is systematically lower than the circular velocity derived from the corrected width of the global H{\sc i} profile. In this paper, we investigate the differences between three measures of the rotational velocity of spiral galaxies: The rotational velocity from the corrected width of the global H{\sc i} profile, the maximal rotation velocity of the rotation curve (V$_{\rm max}$) and the rotational velocity of the outer flat part of the rotation curve (V$_{\rm flat}$). The effects of these different velocity measures on the observed scatter and the intrinsic properties of the TFr will be discussed in a forthcoming paper. To derive V$_{\rm max}$ and V$_{\rm flat}$ we analyse in detail the H{\sc i} synthesis imaging data for a Tully-Fisher calibrator sample of 32 nearby, well--resolved galaxies and present these data in the form of an Atlas and tables. We will discuss the identification and corrections made for warps and streaming motions in the gas disks when deriving the H{\sc i} rotation curves. Forthcoming papers will use these data to investigate the luminosity--based TFr using panchromatic photometric data, as well as the baryonic TFr for the same calibrator sample. This paper is organised as follows: Section 2 describes the sample of selected galaxies. Section 3 presents the data collection and analysis, including subsections on the comparison between the rotational velocities obtained from the rotation curves and the corrected width of the global H{\sc i} profiles. Section 4 presents the properties of the gaseous disks of the sample galaxies. A summary and concluding remarks are presented in Section 5. The Atlas is described in the Appendix, together with notes on individual galaxies. \begin{table*} \begin{tabular}{llrlllllrl} \hline Name & Hubble type & P.A. & Incl. & Log$\,d_{25}$ & $m_{T}^{b,i} [3.6]$ & $\mu_{0}^{i} [3.6]$ & Log$\,h_{r} [3.6]$ & Distance \\ & & deg & deg & arcsec & mag &mag arcsec$^{-2}$& arcsec & Mpc \\ \hline NGC 0055 & SBm & 101 & 86 & 2.47 & 5.64 & 17.54 & 2.32 & 2.09 \\ NGC 0224 & SAb & 35 & 78 & 3.25 & 0.48 & 17.74 & 2.46 & 0.76 \\ NGC 0247 & SABd & 166 & 76 & 2.29 & 6.27 & 18.52 & 2.26 & 3.45 \\ NGC 0253 & SABc & 52 & 81 & 2.42 & 3.37 & 14.36 & 2.10 & 3.45 \\ NGC 0300 & SAd & 114 & 46 & 2.29 & 5.56 & 25.30 & 2.18 & 1.93 \\ NGC 0925 & SABd & 104 & 57 & 2.03 & 7.24 & 20.73 & 1.95 & 9.16 \\ NGC 1365 & SBb & 20 & 54 & 2.08 & 5.91 & 27.71 & 1.85 & 17.94 \\ NGC 2366 & IBm & 32 & 74 & 1.64 & 9.76 & 21.20 & 2.07 & 3.31 \\ NGC 2403 & SABc & 126 & 60 & 2.30 & 5.52 & 19.51 & 2.47 & 3.19 \\ NGC 2541 & SAc & 169 & 63 & 1.48 & 9.16 & 20.22 & 1.60 & 11.22 \\ NGC 2841 & SAb & 148 & 66 & 1.84 & 5.70 & 16.90 & 1.79 & 14.01 \\ NGC 2976 & SAc & 143 & 60 & 1.86 & 6.95 & 19.15 & 1.78 & 3.56 \\ NGC 3031 & SAb & 157 & 59 & 2.33 & 3.40 & 18.68 & 2.15 & 3.59 \\ NGC 3109 & SBm & 93 & 90 & 2.20 & 7.76 & 18.25 & 2.08 & 1.30 \\ NGC 3198 & SBc & 32 & 70 & 1.81 & 7.44 & 17.37 & 1.61 & 13.80 \\ IC 2574 & SABm & 50 & 69 & 2.11 & 8.60 & 22.18 & 2.04 & 3.81 \\ NGC 3319 & SBc & 34 & 59 & 1.56 & 9.05 & 22.80 & 1.74 & 13.30 \\ NGC 3351 & SBb & 10 & 47 & 1.86 & 6.45 & 24.77 & 1.68 & 10.91 \\ NGC 3370 & SAc & 143 & 58 & 1.38 & 8.70 & 20.43 & 1.23 & 0.23 \\ NGC 3621 & SAd & 161 & 66 & 1.99 & 6.37 & 17.57 & 1.71 & 7.01 \\ NGC 3627 & SABb & 173 & 60 & 2.01 & 5.39 & 19.01 & 1.78 & 10.04 \\ NGC 4244 & SAc & 42 & 90 & 2.21 & 7.39 & 18.34 & 2.47 & 4.20 \\ NGC 4258 & SABb & 150 & 69 & 2.26 & 4.95 & 18.89 & 2.24 & 7.31 \\ NGC 4414 & SAc & 166 & 55 & 1.29 & 6.56 & 23.03 & 1.69 & 17.70 \\ NGC 4535 & SABc & 180 & 45 & 1.91 & 6.80 & 25.28 & 1.62 & 15.77 \\ NGC 4536 & SABc & 118 & 71 & 1.85 & 7.17 & 18.03 & 1.70 & 15.06 \\ NGC 4605 & SBc & 125 & 69 & 1.77 & 7.22 & 18.29 & 1.66 & 5.32 \\ NGC 4639 & SABb & 130 & 55 & 1.46 & 8.28 & 26.64 & 1.27 & 21.87\\ NGC 4725 & SABa & 36 & 58 & 1.99 & 5.93 & 21.91 & 1.79 & 12.76 \\ NGC 5584 & SAB & 157 & 44 & 1.50 & 8.87 & 25.27 & 1.39 & 22.69 \\ NGC 7331 & SAb & 169 & 66 & 1.96 & 5.45 & 19.12 & 1.79 & 14.72 \\ NGC 7793 & SAd & 83 & 53 & 2.02 & 6.43 & 22.30 & 1.85 & 3.94 \\ \hline \end{tabular} \caption{The Tully-Fisher Calibrator Sample. Column (1): galaxy name (as shown in \href{https://ned.ipac.caltech.edu/forms/byname.html}{NED}); Column (2): Hubble type (as shown in \href{https://ned.ipac.caltech.edu/forms/byname.html}{NED}); Column (3): optical position angle of the disk from \citetalias{TC12}; Column (4): optical inclination of the disk from \citetalias{TC12}; Column (5): optical diameter (as shown in \href{https://ned.ipac.caltech.edu/forms/byname.html}{NED}); Column (6): total apparent magnitude, derived using 3.6 $\mu$m Spitzer band (corrected for Galactic and internal reddening); Column (7): disk central surface brightness, derived by fitting the exponential disk to the surface brightness profile in 3.6 $\mu$m Spitzer band (corrected for inclination); Column (8): disk scale length, calculated from the 3.6 $\mu$m Spitzer band; Column (9): Distance in Mpc, measured using TRGB or Cepheid distance estimation method. An averaged value is taken, when distance is measured by both methods (provided by The Extragalactic Distance Database (EDD) \href{http://edd.ifa.hawaii.edu/}{http://edd.ifa.hawaii.edu/}). } \label{tbl_samp} \end{table*} \begin{table*} \begin{tabular}{lllllr@{.}lr@{.}lr@{$\times$}ll} \hline Name & Data source & Array/ & Obs. date & FREQ & \multicolumn{2}{l}{B-width} & \multicolumn{2}{l}{Ch-width} & \multicolumn{2}{l}{Beam size} & RMS noise\\ & & configuration & dd/mm/yy & MHz & \multicolumn{2}{l}{MHz} & \multicolumn{2}{l}{km$s^{-1}$} & \multicolumn{2}{l}{arcsec$^{2}$}& mJy/beam\\ \hline NGC 0224 &R. Braun$^{7}$& WSRT & 08/10/01 & 1421.82 & 5 & 00 & 2& 06 & 60&60 &2.70\\ NGC 0247 & ANGST$^{2}$ & VLA(BnA/CnB) & 06/12/08 & 1419.05 & 1 & 56 & 2& 60 & 9&6.2 &0.73\\ NGC 0253 & LVHIS & ATCA & 08/03/94 & 1419.25 & 8 & 00 & 3& 30 & 14&5 &1.47\\ NGC 0300 & LVHIS & ATCA(EW352,367)& 03/02/08 & 1419.31 & 8 & 00 & 8& 00 &180&87 &7.18\\ NGC 0925 & THINGS$^{3}$ & VLA (BCD/2AD) & 08/01/04 & 1417.79 & 1 & 56 & 2& 60 & 5.9&5.7 &0.58\\ NGC 1365 & S. J\"{o}rs\"{a}ter$^{8}$ &VLA& 26/06/86 & 1412.70 & 1 & 56 &20& 80 & 11.5&6.3 &0.06\\ NGC 2366 & THINGS & VLA (BCD/2AD) & 03/12/03 & 1420.02 & 1 & 56 & 2& 60 & 13.1&11.8&0.56\\ NGC 2403 & THINGS & VLA (BCD/4) & 10/12/03 & 1419.83 & 1 & 56 & 5& 20 & 8.7&7.6 &0.39\\ NGC 2541 & WHISP$^{4}$ & WSRT & 16/03/98 & 1417.86 & 2 & 48 & 4& 14 & 13&10&0.59\\ NGC 2841 & THINGS & VLA (BCD/4) & 30/12/03 & 1417.38 & 1 & 56 & 5& 20 & 11&9.3 &0.35\\ NGC 2976 & THINGS & VLA (BCD/4AC) & 23/08/03 & 1420.39 & 1 & 56 & 5& 20 & 7.4&16.4 &0.36\\ NGC 3031 & THINGS & VLA (BCD/2AD) & 23/08/03 & 1420.56 & 1 & 56 & 2& 60 & 12.9&12.4 &0.99\\ NGC 3109 & ANGST & VLA(2AD/2AC) & 07/12/08 & 1418.52 & 1 & 56 & 1& 30 & 10.3&8.8 &1.31\\ NGC 3198 & THINGS & VLA (BCD/4) & 26/04/05 & 1417.28 & 1 & 56 & 5& 20 & 13&11.5 &0.34\\ IC 2574 & THINGS & VLA(BCD) & 18/01/92 & 1420.13 & 1 & 56 & 2& 60 & 12.8&11.9 &1.28\\ NGC 3319 & WHISP & WSRT & 20/11/96 & 1416.87 & 2 & 48 & 4& 14 & 18.4&11.9 &0.91\\ NGC 3351 & THINGS & VLA (BCD/4) & 06/01/04 & 1416.72 & 1 & 56 & 5& 20 & 9.9&7.1 &0.35\\ NGC 3370 & this work & GMRT & 06/03/14 & 1414.32 & 4 & 16 &16& 00 & 30&30 &2.39 \\ NGC 3621 & THINGS & VLA (BnAC/4) & 03/10/03 & 1416.95 & 1 & 56 & 5& 20 & 15.9&10.2&0.70\\ NGC 3627 & THINGS & VLA (BCD/2AD) & 28/05/05 & 1416.96 & 1 & 56 & 5& 20 & 10.6&8.8&0.41\\ NGC 4244 & HALOGAS$^{5}$& WSRT & 23/07/06 & 1419.24 & 10 & 00 & 2& 60 & 21&13.5 &0.18\\ NGC 4258 & HALOGAS & WSRT & 29/03/10 & 1418.28 & 10 & 00 & 2& 60 & 30&30 &0.25\\ NGC 4414 & HALOGAS & WSRT & 25/03/10 & 1417.02 & 10 & 00 & 2& 60 & 30&30 &0.19\\ NGC 4535 & VIVA$^{6}$ & VLA (CD) & 20/01/91 & 1411.16 & 3 & 12 &10& 00 & 25&24 &0.61\\ NGC 4536 & VIVA & VLA (CS) & 22/03/04 & 1411.89 & 3 & 12 &10& 00 & 18&16 &0.34\\ NGC 4605 & WHISP & WSRT & 11/11/98 & 1419.47 & 2 & 48 & 4& 14 & 13&8.6 &0.93\\ NGC 4639 & this work & GMRT & 06/03/14 & 1415.62 & 4 & 16 &16& 00 & 30&30 &2.00\\ NGC 4725 & WHISP & WSRT & 18/08/98 & 1414.49 & 4 & 92 &16& 50 & 30&30 &0.60\\ NGC 5584 & this work & GMRT & 10/03/14 & 1412.78 & 4 & 16 &16& 00 & 30&30 &2.83\\ NGC 7331 & THINGS & VLA (BCD/4) & 20/10/03 & 1416.55 & 1 & 56 & 5& 20 & 6.1&5.6 &0.44\\ NGC 7793 & THINGS & VLA (BnAC/2AD) & 23/09/03 & 1419.31 & 1 & 56 & 2& 60 & 15.6&10.8 &0.91\\ \hline \end{tabular} \caption{Observational parameters of the sample. Column (1): galaxy name; Column (2): name of the survey or of individual author, where from the data were taken (see the reference below); Column (3): instrument and the set up; Column (4): date of observations dd/mm/yy; Column (5): observed frequency, MHz; Column (6): band--width, MHz; Column (7): channel-- width, MHz; Column (8): synthesised beam, arcsec; Column (9): RMS noise, mJy/beam. \textbf{References}: 1. The Local Volume H{\sc i} Survey \citep{lvhis} 2. ACS Nearby Galaxy Survey Treasury \citep{angst} 3. The H{\sc i} Nearby Galaxy Survey \citep{things} 4. Westerbork observations of neutral Hydrogen in Irregular and SPiral galaxies \citep{whisp} 5. Hydrogen Accretion in LOcal GAlaxieS survey \citep{halogas} 6. VLA Imaging of Virgo in Atomic gas \citep{viva} 7. \citet{andromeda} 8.\citet{1365} } \label{tbl_obs} \end{table*} \begin{table} \begin{tabular}{lcccl} \hline Name & Obs. Date&& Calibrators& $T_{obs}$\\ &dd/mm/yy& & &hrs\\ \hline NGC 3370& 06/03/14&& 3C147; 0842+185&4\\ & 07/03/14&&3C147; 0842+185&4.5\\ &11/03/14&&3C147; 0842+185&2.5\\ NGC 4639&06/03/14&&3C286; 1254+116&1.1 \\ &07/03/14&&3C286; 1254+116&2\\ &10/03/14&&3C147;3C286;&4\\ & &&1254+116 & \\ &11/03/14&&3C286; 1254+116&3.9\\ NGC 5584&10/03/14&&3C286; 3C468.1;& 6.75\\ & && 1445+099 & \\ &11/03/14&&3C286; 3C468.1;& 4.25\\ & && 1445+099 & \\ \hline \end{tabular} \caption{The GMRT observations. Column (1): galaxy name; Column (2): observational dates; Column (3): flux and phase calibrators; Column (4): integration time.} \label{tbl_calib} \end{table} | We have presented the analysis of 21-cm spectral--line aperture synthesis observations of 32 spiral galaxies, which are representing a calibrator sample for studying the statistical properties of the Tully--Fisher relation. The data were collected mostly from the literature and obtained with various observational facilities (VLA, ATCA, WSRT). We observed for the first time three galaxies in our sample ourselves with the GMRT. The most important aspect of this work is that, despite the broad range in data quality, we analysed the entire sample in the same manner. Although previously many of these galaxies were studied individually, we present for the first time a set of H{\sc i} synthesis imaging data products for all these galaxies together, analysed in a homogeneous way. The data products consist of total H{\sc i} maps, velocity fields in which asymmetric velocity profiles are identified and removed, H{\sc i} global profiles, radial H{\sc i} surface--density profiles, position--velocity diagrams and, most importantly, high--quality extended H{\sc i} rotation curves, all presented in an accompanying Atlas. The rotation curves and the derived measurements of V$_{\rm max}$ and V$_{\rm flat}$ will be used in forthcoming papers investigating the statistical properties of the Tully--Fisher relation. The radial H{\sc i} surface--density profiles will be used for rotation curves decompositions, aimed at studying the mass distributions within spiral galaxies. Overall, our kinematical study of the gas disks of our sample galaxies shows excellent agreement with previously reported values for the position and inclination angles, and systemic velocities. We find a good agreement with literature values for both H{\sc i} fluxes and the widths of the global H{\sc i} profiles at the 50\% level, obtained from single--dish observations. However, in some cases we find somewhat larger fluxes than previously reported in the literature (see Section 3.3.1 for details). This can be easily explained given the large angular size of our galaxies, and the fact that single--dish observations can miss some flux due to the relatively small beam sizes. All galaxies in our sample have extended H{\sc i} disks, which serve our purpose to probe the gravitational potential of their dark matter halos. Moreover, our galaxies follow the well--known correlation between their H{\sc i} mass and the diameter of their H{\sc i} disks, but with a slightly shallower slope then shown in previous studies. Studies of the H{\sc i} mass and its correlation with the absolute $K_{s}$ magnitude show that while more luminous galaxies tend to have more H{\sc i} gas, its fraction decreases with luminosity. We find hints for similar trends with Hubble--type: late--type spiral galaxies tend to have a larger fraction of H{\sc i} gas than more early--type spirals, but the total mass of the H{\sc i} gas decreases. A qualitative comparison with the volume--limited sample of Ursa Major galaxies shows that our calibrator sample is respresenatative for a population of field galaxies. The radial H{\sc i} surface--density profiles were scaled radially with $R_{HI}$ and divided into five groups according to the morphological type of the galaxies. The gas disks of early--type spirals tend to have central holes while the has disks of late--type spirals and irregulars tend to be centrally concentrated. Velocity fields were constructed by fitting Gaussian and Gauss--Hermite polynomial functions to the velocity profile at each position in the data cube. We identified the pixels with skewed velocity profiles, using the difference between these two velocity fields, and censored them in the Gaussian velocity field, thereby creating a third type of velocity field. Rotation curves were derived from these censored velocity fields using tilted--ring modelling. The procedure was done in four steps, which allowed to follow the geometry of a possible warp in the outer region. Rotation curves were identified into 3 categories: rising, flat and declining. The obtained values of V$_{\rm max}$ and V$_{\rm flat}$ were compared with the velocities derived from the corrected H{\sc i} global profiles at the 50\% and 20\% level. The comparison tests show that while the width of the H{\sc i} profile can be a good representation for the maximal rotation velocity of the galaxy, it may significantly overestimate the velocity at the outer, flat part of the rotation curve, especially for high--mass galaxies which tend to have declining rotation curves. For the rising rotation curves, where V$_{\rm max}$ is lower than V$_{\rm flat}$, the width of the profile tends to underestimate V$_{\rm flat}$. The statistical properties of the Tully-Fisher relation based on different velocity measures ($W$, V$_{\rm max}$ and V$_{\rm flat}$) will be studied in forthcoming papers. | 16 | 9 | 1609.00378 |
1609 | 1609.08945_arXiv.txt | { The {\it Kepler\/} space mission led to a large number of high-precision time series of solar-like oscillators. Using a Bayesian analysis that combines asteroseismic techniques and additional ground-based observations, the mass, radius, luminosity, and distance of these stars can be estimated with good precision. This has given a new impetus to the research field of galactic archeology. } { The first data release of the Gaia space mission contains the TGAS (Tycho-Gaia Astrometric Solution) catalogue with parallax estimates for more than 2 million stars, including many of the {\it Kepler\/} targets. Our goal is to make a first proper comparison of asteroseismic and astrometric parallaxes of a selection of dwarfs, subgiants, and red giants observed by {\it Kepler\/} for which asteroseismic distances were published. } { We compare asteroseismic and astrometric distances of solar-like pulsators using an appropriate statistical errors-in-variables model on a linear and on a logarithmic scale. } { For a sample of 22 dwarf and subgiant solar-like oscillators, the TGAS parallaxes considerably improved on the Hipparcos data, yet the excellent agreement between asteroseismic and astrometric distances still holds. For a sample of 938 {\it Kepler\/} pulsating red giants, the TGAS parallaxes are much more uncertain than the asteroseismic ones, making it worthwhile to validate the former with the latter. From errors-in-variables modelling we find a significant discrepancy between the TGAS parallaxes and the asteroseismic values. } { For the sample of dwarfs and subgiants, the comparison between astrometric and asteroseismic parallaxes does not require a revision of the stellar models on the basis of TGAS. For the sample of red giants, we identify possible causes of the discrepancy, which we will likely be able to resolve with the more precise Gaia parallaxes in the upcoming releases. } | The seismic study of stars has undergone a revolution during the past decade, thanks to the space missions CoRoT \citep[launched in 2006;][]{Auvergne2009} and {\it Kepler\/} \citep[launched in 2009;][]{Borucki2010}. Not only did these space data confirm the method of asteroseismology \citep[for an extensive monograph; see][]{Aerts2010}, they also allowed powerful applications to thousands of stars across stellar evolution for a wide variety of stellar birth masses. Major breakthroughs of relevance to the current study of stellar distances were the discovery of acoustic non-radial pulsation modes in red giants \citep{DeRidder2009} and the excitation of dipole mixed modes probing both the deep interior and the structure of the outer envelope of such stars \citep[e.g.][]{Beck2011}. Thanks to their mixed gravity and acoustic character, mixed modes allow the core properties of a star to be tuned and therefore can be used to pinpoint the evolutionary status \citep{Bedding2011}. Asteroseismology of red giants offers the unique opportunity of providing stellar ages for studies of the Milky Way, termed {galactic archeology\/} \citep[e.g.][]{Miglio2013}. Indeed, the measurement of the frequency at maximum oscillation power and of the large frequency separation, along with a spectroscopic estimate of the effective temperature, can be transformed into high-precision estimates of the stellar mass and radius by assuming that the input physics of solar models is also applicable to solar-like stars. Under this reasonable assumption, stellar masses and radii can be derived with relative precisions of merely a few per cent, while further comparison with stellar models provides a seismic age estimate with a precision below 20\% when systematic uncertainties due to modest variations in the input physics are taken into account \citep{Chaplin2014,Metcalfe2014}. Proper computation of the apparent CoRoT or {\it Kepler\/} magnitude according to the passbands of these satellites then allows the luminosity of the stars to be transformed into an ``asteroseismic'' distance \citep{SilvaAguirre2012,Rodrigues2014,Anders2016}. So far, the asteroseismic distances of stars in the solar neighbourhood have been compared {a posteriori\/} with Hipparcos values for the distances whenever available, with good agreement \citep[e.g.][]{SilvaAguirre2012}. With the Gaia mission in full swing, we foresee a quantum leap forward in this research, both in the number of targets and the precision in measuring the distance. After five years of nominal monitoring, the Gaia distance estimates are expected to be so precise that they can serve {as input\/} to improve the physics of stellar interiors, leading to model-independent radii and better ages than currently available as input for exoplanet studies and galactic archeology. Here we take a first step to compare asteroseismic distances with the astrometric values by considering the first Gaia data release \citep[Gaia DR1; e.g.][]{Brown2016, Prusti2016, Lindegren2016}. | Astrometric parallaxes in the Gaia DR1 TGAS catalogue have been compared with published asteroseismic distances of pulsating dwarfs and giants using an errors-in-variables approach. Deviation from the bisector would imply that either the models of stellar interiors combined with reddening models at the basis of the asteroseismic distances need revision, or that there is an unknown systematic uncertainty on the current version of the astrometric parallaxes, or both. Proper statististical analysis of the two parallax estimates is {a priori\/} more complicated than ordinary least-squares as we now have uncertainties on both estimates of one and the same quantity. We therefore set up an errors-in-variables model, on a linear scale and on a logarithmic scale, to do the fitting. For the 22 dwarfs and subgiants of \citet{SilvaAguirre2012}, the results reveal excellent agreement between the two distances, re-confirming the asteroseismic achievement for these stars now that we have more precise parallaxes from TGAS. For our sample of 938 giants taken from \citet{Rodrigues2014} and crossmatched with TGAS, the relative uncertainties on the astrometric parallaxes are much larger than on the seismic distances, turning the latter into a valuable instrument to validate the former. All the models we applied -- an errors-in-variables model, ordinary least-squares, and a logarithmic model -- lead to a significant difference between astrometric and asteroseismic parallaxes. There can be several underlying causes. The uncertainties of the TGAS parallaxes may be underestimated or could be subject to systematic errors. On the other hand, interstellar extinction corrections and/or too poorly known bulk metallicity may have introduced a systematic uncertainty for the asteroseismic parallax. Given that all stars in the Gaia DR1 TGAS catalogue were assumed to be single, binarity may also be part of the cause. Expectations are that the accuracy of the Gaia astrometric distance measurements will surpass the seismic measurements by the end of the mission. Gaia Data Release 2 (end 2017) will contain the astrometry of a billion stars, while Gaia Data Release 3 (end 2018) will deliver the orbital astrometric solutions for binaries with periods longer than 2 months. This will allow us to improve our current research and transfer it into a quantitative calibration of asteroseismic distances for a variety of stellar populations. | 16 | 9 | 1609.08945 |
1609 | 1609.08174_arXiv.txt | Most globular clusters (GCs) are now known to host multiple stellar populations with different light element abundances. Here we use narrow-band photometry and low-resolution spectroscopy for NGC~362 and NGC~6723 to investigate their chemical properties and radial distributions of subpopulations. We confirm that NGC~362 and NGC~6723 are among the GCs with multiple populations showing bimodal CN distribution and CN-CH anti-correlation without a significant spread in calcium abundance. These two GCs show more centrally concentrated CN-weak earlier generation stars compared to the later generation CN-strong stars. These trends are reversed with respect to those found in previous studies for many other GCs. Our findings, therefore, seem contradictory to the current scenario for the formation of multiple stellar populations, but mass segregation acting on the two subpopulations might be a possible solution to explain this reversed radial trend. | Recent observations suggest that most globular clusters (GCs) host multiple stellar populations showing star-to-star abundance variations in the light elements, such as C, N, O, Na and Al \citep[e.g.,][and references therein]{Car09,Gra12,Pio15}. Among several scenarios for the origin of these abundance variations, the most widely accepted one is the self-enrichment scenario, which explains these variations by the chemical pollution/enrichment from earlier generation stars, such as intermediate-mass asymptotic giant branch (IMAGB) stars \citep{VD08}, rotating AGB stars \citep{Dec09}, interacting binary stars \citep{de09}, and fast-rotating massive stars (FRMSs; \citealt{Dec07b}). In this scenario, later generation stars are expected to be formed by the gas ejected from earlier generation stars in the innermost region of a proto GC \citep{Dec07a,D'er08,Ves13}. Therefore, the later generation stars would be observed to be more centrally concentrated than earlier generation stars unless this radial distribution was seriously affected by dynamical evolution \citep[see, e.g.,][]{Mih15}. This radial trend is indeed observed in many GCs, including $\omega$~Cen, M13, NGC~3201, NGC~6752, and 47~Tuc \citep{Bel09,Kra10,Kra11,Lar11,Nat11,JP12,Mil12}. The incidence of spectroscopic binaries in different subpopulations, which is less affected by the dynamical evolution, also supports that later generation stars were formed in a denser environment where binaries are efficiently destroyed, resulting in a lower binary fraction for later generation stars \citep{D'Or10,Luc15}. \begin{figure*} \centering \includegraphics[width=0.9\textwidth]{fov.pdf} \figcaption{Our observed fields (blue squares) on the STScI images of NGC~362 and NGC~6723. The red circles indicate our spectroscopic target stars. \label{fig_fov} } \end{figure*} However, not every GC with multiple populations shows this general radial distribution trend. \citet{Car10} and \citet{jwlee15} report a central concentration of metal-poor earlier generation stars in NGC~1851\footnote{\citet{Mil09} showed that two stellar populations of sub-giant branch stars in NGC~1851 share the same radial distribution, which, however, is contradicted by \citet{Car11}.} and M22, respectively. These GCs are known as peculiar GCs showing intrinsic heavy elements dispersions \citep{jwlee09,Mar09,Car11,Lim15}, and therefore, their reversed radial trends may be considered to be the result of merging of two individual GCs \citep[see][]{Car11}. In the case of GCs without heavy elements spread, \citet{Dal14} have shown that two stellar populations in NGC~6362 share the same radial distribution, which is explained as a full spatial mixing accelerated by high mass loss rate of this GC (\citealt{Mih15}; see also \citealt{Muc16}). Furthermore, \citet{Lar15} recently suggested that first generation stars (primordial group) are more centrally concentrated than second generation stars (enriched group) in M15 using $Hubble~Space~Telescope~(HST)$ WFC3 photometry. This was not detected in the previous study for the same GC by \citet{Lar11} using the Sloan Digital Sky Survey (SDSS) data, which cover only the outer region of a cluster. Contrary to the cases of M22 and NGC~1851, the reversed radial trend in M15 is unlikely to be the result of merging, because this GC does not show Fe spread although variations in the light elements and some neutron-capture elements (Ba, Eu) were reported (\citealt{Sne97}; see also \citealt{BT16}). The spatial mixing due to dynamical evolution is also unlikely to explain this reversed radial trend \citep[see, e.g.,][]{Ves13}. Therefore, the presence of a reversed radial distribution trend in M15 casts some doubt on the current self-enrichment scenario, and thus a search for further instances of this radial characteristic is required. Investigating the radial distribution of multiple stellar populations, however, is not a simple task because spectroscopic observations are hard to secure a large enough number of samples and it is difficult to divide subpopulations using photometric observation alone. In this regard, our narrow-band photometry, combined with low-resolution spectroscopy, would be a useful tool for this investigation. Our previous studies have shown that narrow-band photometry using ``Ca'' and ``Ca+CN'' filters can efficiently detect multiple stellar populations with different chemical properties, and this is confirmed by low-resolution spectroscopy \citep{Lim15,Han15}. In this study, we have investigated the chemical properties and radial distributions of stars in NGC~362 and NGC~6723 by employing the same techniques and report that these two GCs show a central concentration of earlier generation stars, similarly to the case of M15. \begin{figure*} \centering \includegraphics[width=1.0\textwidth]{cmd.pdf} \figcaption{ CMDs for NGC~362 (left) and NGC~6723 (right) in ($y$, $hk_{Ca+CN}$) plane obtained with the Ca+CN filter set at CTIO. Spectroscopic target stars are also identified in these CMDs, where the blue and red circles are CN-weak and red CN-strong stars, respectively (see Section~\ref{spec}). Note that RGB spread and split are shown in NGC~362 and NGC~6723, respectively. \label{fig_cmd} } \end{figure*} \begin{deluxetable}{cccccc} \tabletypesize{\normalsize} \tablewidth{0pt} \tablecaption{Mask Descriptions and Spectroscopic Observation Log\label{tab_log}} \tablehead{ \colhead{Object} & \colhead{Mask} & \colhead{No of stars} & \colhead{Exposures (N$\times$s)} } \startdata NGC 362 & Bright & 23 & 2$\times$1200 \\ & & & 2$\times$1500 \\ & Faint & 26 & 4$\times$1500 \\ NGC 6723 & Bright & 19 & 4$\times$1200 \\ & & & 3$\times$1500 \\ & Faint & 31 & 5$\times$1500 \enddata \end{deluxetable} | \label{dis} We have shown that the RGB stars in NGC~362 and NGC~6723 are clearly divided into two subpopulations by CN index, but these two subpopulations show no difference in calcium abundance from our low-resolution spectroscopy. The well-known CN-CH anti-correlation between the two subpopulations is also shown in both GCs. These results, together with previous findings by other investigators \citep{WC10,Car13,Gra15}, suggest that these two GCs are ``normal'' GCs with multiple stellar populations. Furthermore, we found that the CN-weak earlier generation stars are significantly more centrally concentrated than CN-strong later generation stars in both GCs. These findings are important as the second and third cases that show such a reversed radial distribution trend, following M15 by \citet{Lar15}. We note, however, that the innermost region ($r$ $<$ 0.5$'$) of a cluster was not included in our analysis for NGC~362, and therefore further analysis with the recent HST UV survey \citep{Pio15} would be helpful to confirm the reversed spatial distribution discovered in this study. As described in Section~\ref{intro}, the self-enrichment scenario predicts more centrally concentrated stars belonging to the later generation \citep[see, e.g.,][]{Ves13}. Therefore, our finding of centrally concentrated CN-weak earlier generation stars in NGC~362 and NGC~6723 is not generally acceptable in this framework. The merging of two individual GCs might be a possible solution for different radial distributions of stellar populations \citep[see][]{Car11}, but these GCs do not show additional evidence of merging, such as metallicity difference. Moreover, their chemical properties, the bimodal CN distribution and CN-CH anti-correlation without Fe spread, are not intuitively understandable in the merging scenario. \citet{Lar15} suggested that mass segregation may explain the observed radial trend of M15. This scenario, however, requires extreme He enhancement of later generation stars ($Y$ $\geq$ 0.40) in order to produce a mass difference of 0.25$M_{\odot}$, which is not observed in this GC. In either case of NGC~362 and NGC~6723, He enhancement is not yet reported, although some $\Delta$$Y$ is expected. Therefore, the mass segregation by the difference in He abundance alone is not likely to explain the observed radial trends in M15, NGC~362, and NGC~6723. More recently, an alternative solution for the mass difference is proposed by \citet{Hen15}, which suggests that the different incidence of binary from primordial and enriched stellar populations may provoke this required mass difference \citep[see also][]{Hong15}. According to them, this effect is more efficient in more concentrated GCs. Interestingly, M15 has been known as a core collapsed GC, and NGC~362 and NGC~6723 are also classified as possible core collapsed GCs \citep{Har10}. The effects of mass segregation on the radial distribution of stellar generations require more detailed studies. To this end, we plan to run direct $N$-body simulations of NGC~362 and NGC~6723 (Pasquato et al. in prep.) under different assumptions for the relevant binary fractions, initial concentration of the second generation stars, and mass-difference between the subpopulations to assess to what extent mass-segregation can explain our findings in these GCs. At this stage, it is worth noting that the half-mass relaxation time of NGC~362 is relatively short (less than 1Gyr) while that of NGC~6723 is roughly of the same order of that of M15 ($\sim$2Gyr). Therefore, we expect that mass-segregation would play a more important role in the former cluster. One caveat in this analysis is that definitions of subpopulations are not exactly identical in many studies. We have divided stars in a GC into two subpopulations by the strength of the CN band, whereas many high-resolution spectroscopic studies based on Na and O abundances prefer to separate them into three subpopulations \citep{Car09,JP12}. In addition, \citet{Jang14} and \citet{JL15} suggest the presence of three subpopulations with different He abundances based on population models for RR Lyrae and HB stars. As shown in our previous study \citep{Lim15}, CN-weak stars in NGC~1851 could be further divided into two subpopulations, and therefore, it is possible that CN-weak stars in other GCs also might be further divided into two subgroups. This issue will be discussed in our forthcoming paper. | 16 | 9 | 1609.08174 |
1609 | 1609.06752.txt | We report the second complete molecular line data release from the {\em Census of High- and Medium-mass Protostars} (CHaMP), a large-scale, unbiased, uniform mapping survey at sub-parsec resolution, of mm-wave line emission from 303 massive, dense molecular clumps in the Milky Way. This release is for all \tco\ \joz\ emission associated with the dense gas, the first from Phase II of the survey, % which includes \tco, \ttco, and \ceto. % The observed clump emission traced by both \tco\ and \hcop\ (from Phase I) shows very similar morphology, indicating that, for dense molecular clouds and complexes {\em of all sizes}, parsec-scale clumps % contain $\Xi$ $\sim$ 75\% of the mass, while only 25\% of the mass lies in extended (\gapp10\,pc) or ``low density'' components in these same areas. The mass fraction of all gas above a density 10$^9$\,m$^{-3}$ is $\xi_9$ \gapp\ 50\%. This suggests that parsec-scale clumps may be the basic building blocks of the molecular ISM, rather than the standard GMC concept. Using \tco\ emission, we derive physical properties of these clumps in their entirety, and compare them to properties from \hcop, tracing their denser interiors. % We compare the standard $X$-factor converting \ico\ to \nhtwo\ with % alternative conversions, and show that only the latter give whole-clump properties that are physically consistent with those of their interiors. % We infer that the clump population is systematically closer to virial equilibrium than when considering only their interiors, with perhaps half being long-lived (10s of Myr), pressure-confined entities which only terminally engage in vigorous massive star formation, % supporting other evidence along these lines previously published. | The formation of massive stars and clusters from molecular clouds remains one of the major unsolved problems in astrophysics. Part of the reason for the ongoing debate about initial conditions, mechanisms, timescales, feedback, environmental factors, and other model parameters, is the complexity of massive star formation phenomenology and the relatively small number of wide-field, complete, uniform, unbiased, multi-wavelength studies, as explained by \citet{b11}. For example, it has not been established whether the parsec-scale massive clumps that form star clusters \citep{ll03} are long-lived entities (several 10s of Myr) that do not undergo vigorous massive star formation until the latter part of this time span \citep[e.g.,][]{kss09,b13}, or whether they are shorter-lived objects ($<$10\,Myr) that promptly form clusters and are then dissipated \citep[e.g.,][]{bsd10}. For long lifetimes, one must also argue for cluster-forming clumps to be either gravitationally bound or pressure-confined, as originally explained by \citet{bm92} based on stability arguments. On the other hand, even unbound clouds could form some stars \citep{wws14}. More recently, observational breakthroughs have provided new challenges to theory. The {\em Herschel} Observatory revealed widespread filamentary structures in the cold interstellar medium (ISM), and subsequent theoretical work is confirming how important filaments are to molecular clouds' overall physics and star formation activity \citep[see review by][]{a14}, including how gas flows assemble filaments and then clumps within them. Yet, many studies to date have concentrated on a few, typically nearby clouds. Similarly, there is a growing recognition that ``CO-dark gas'' \citep{gct05,p13} may contribute up to half the molecular mass of the Milky Way. Our direct knowledge of the distribution of this gas in the Galactic disk is currently limited to $\sim$0\fdeg2 scales \citep{L14}, yet new models suggest its distribution is related to the much finer scale of filaments \citep{s14}. In general, many questions about molecular cloud stability, dynamics, composition, and star formation activity would best be examined with high spatial dynamic range imaging of a wide sample of Giant Molecular Clouds (GMCs). We developed the Galactic {\em Census of High- and Medium-mass Protostars} (CHaMP) to address many of these issues. CHaMP was originally conceived to be a multi-wavelength survey of a statistically large but uniformly-selected sample of massive, dense molecular clumps, in order to analyse the massive star- and cluster-formation process with as little observational bias as possible. CHaMP began with a Mopra\footnote{The Mopra telescope is part of the Australia Telescope, funded by the Commonwealth of Australia for operation as a National Facility managed by CSIRO. The University of New South Wales Digital Filter Bank used for observations with the Mopra telescope was provided with support from the Australian Research Council.} molecular line mapping survey of a complete, flux-limited clump sample within a 20\degree$\times$6\degree\ area of the southern Milky Way in Vela, Carina, and Centaurus. Phase I of this mapping, of the emission from several ``dense\footnote{What exactly constitutes a dense clump is not always well-defined in the literature; we explore this in \S\ref{clumpcomps}.} gas tracers'' in the identified CHaMP clumps, took place over the period 2004--2007, with several follow-up studies at infrared wavelengths since then, using the AAT, CTIO, {\em Spitzer}, and archival spacecraft data. Early results from this work have been presented by \cite{y05}, \cite{z10}, \citet{b10}, and \citet{b13}, with more studies in preparation. In \citet[hereafter Paper I]{b11}, we described the overall survey strategy and reported the first mm-wave results for the ensemble of 303 massive molecular clumps in the southern Milky Way. We found that these clumps represent a vast population of subthermally-excited, yet massive clouds, 95\% of which are relatively quiescent and not currently engaged in vigorous massive star formation, suggesting long clump lifetimes. In \citet[hereafter Paper II]{mtb13}, we performed a global analysis of the spectral energy distributions (SEDs) of these clumps, and found a wide range of star-formation efficiencies (SFEs), consistent with such long lifetimes. At the same time, it was recognised that similarly-detailed information on the clump envelopes and embedding GMCs would be needed to provide critical comparisons with the denser gas, such as abundances, masses, and an environmental context for any star formation activity within the clumps. Thus, Phase II observations, primarily aimed at mapping the CO-isotopologue lines with Mopra, were conducted during 2009--2012. Because of its uniform and wide-area approach, CHaMP's strategy of fundamental cloud demographics was designed to access new discovery space in the pursuit of the science problems described above. In this sense, CHaMP differs strongly from several other molecular ISM/star formation projects. For example, detailed studies on individual massive star formation sites \citep[e.g.,][for OMC1, NGC\,6334, and DR21, resp.]{u97,r10,s10} find molecular gas properties which are more extreme than any members of the CHaMP sample. Surveys of massive star formation may also be selective: e.g., both \citet{wu10} and the MALT90 project \citep{j13,h13} examined more extreme samples of clouds than CHaMP, as discussed in Paper I. The former targeted $\sim$50 of the most luminous water masers in the Galaxy, while the latter mapped $\sim$2000 of the highest column density, pc-scale dust clumps in the southern Galactic Plane from the 870\,$\mu$m ATLASGAL \citep{s09} clump catalogue. In both cases, such projects specifically pick out the most extreme cloud population (e.g., in terms of luminosity, opacity, mass, or density), rather than the much more representative, and ultimately much larger and less biased, CHaMP cloud population, which probably numbers $>$10$^4$ across the Galaxy. Therefore, it should not be surprising that the cloud properties we find are different than in these other samples, even when observed with the same telescope. In this paper we describe the Phase II observing (\S2) and data reduction (\S3) procedures, paying particular attention to where these differ from Phase I, and give the first results from analysing the brightest Phase II line, \tco\ \joz, in \S4. We discuss these results in \S5, in the context of our previous results on the dense gas, IR emission, and star formation activity in these clumps, while also relating this to the current wider understanding of the cluster formation process. Our conclusions are summarised in \S6. %%%%%%%%%% % % % Section 2 % % % %%%%%%%%%% | As the second major mm-wave data release of the CHaMP project, we have presented new observations and analysis of the \tco\ line emission from a complete sample of $\sim$300 massive molecular clumps, originally defined by their \hcop\ emission in Paper I. The \tco\ emission traces the less dense molecular envelopes in which the denser \hcop-bearing interiors are embedded, and are the first results from Phase II of the Mopra mapping during 2009--12, which covered several molecular lines in the 107--115\,GHz range. The \tco\ observing and data processing include three significant advances over the techniques utilised in Paper I for \hcop: \hspace{1mm} \put(5,3){\circle*{3}} \hspace{2mm} We have used a sky-adaptive version of the standard on-the-fly (OTF) mapping technique, called ``Active Mapping'' or AM, which adjusts the mapping speed of each map for the sky conditions (whether elevation- or weather-dependent) at the time of observation. OTF-AM has the effect of making the rms noise levels in each map much more uniform than with regular OTF. \hspace{3mm} \put(5,3){\circle*{3}} \hspace{2mm} We have used data-screening techniques to filter out errors and correct for calibration problems. \hspace{3mm} \put(5,3){\circle*{3}} \hspace{2mm} We have used a smooth-and-mask (SAM) technique to dramatically improve the image fidelity in the moment maps derived from the data cubes, even for the intrinsically high S/N \tco\ data. With these new data and techniques, we have compiled the observed and derived physical properties of the identified ``massive dense clump'' envelopes, and compared these properties with those of the clump interiors as presented in Paper I. Our main results are as follows: %\begin{enumerate} \vspace{0mm}1.\ The observed clump sizes and linewidths are very similar in both \tco\ and \hcop, with only a $\sim$25\% contribution to the total mass of GMCs and molecular cloud complexes from extended ($>$5\,pc) cloud components. \vspace{0mm}2.\ This suggests that parsec-scale clumps ($<$$R$$>$ = 0.84\,pc, with a logarithmic dispersion equivalent to a factor of 1.9 around this size) comprise the basic building blocks %($\sim$75\% of the mass) of the molecular ISM, and that in the main, extended molecular structures such as GMCs are just collections of such clumps, since most of the mass, $\Xi\sim$ 75\%, is enclosed by the parsec-scale clumps we sample, with typical central densities 10$^{9-9.5}$\,m$^{-3}$. \vspace{0mm}3.\ We see a weak, but real, Larson-type size-linewidth relation for the envelopes, with $\sigma_{\rm V}$ = (1.67 $\pm$ 0.04\,\kms)\,$R_{\rm pc}^{0.24\pm0.04}$, whereas for \hcop\ we saw no statistically significant size-linewidth relation. %3.94$\pm$0.10 \vspace{0mm}4.\ When computed using a standard $X$-factor, the whole-clump properties including the envelopes give slightly lower central volume densities, and similar column densities, masses, virial-$\alpha$, and total internal pressures compared to the interiors alone. \vspace{0mm}5.\ When computed using new \ico\ to \nhtwo\ conversion formulae from \citet{bm15} and \citet{kl15}, the whole-clump properties including the envelopes give somewhat higher central volume densities and total internal pressures, systematically higher column densities and masses, and systematically lower virial-$\alpha$ compared to the interiors alone. \vspace{0mm}6.\ We interpret these results to mean that, including the envelope mass, $\sim$half the clumps detected in \hcop\ are gravitationally bound or near virial equilibrium, even when their interiors alone are not. \vspace{0mm}7.\ This suggests that about half the \hcop-defined ``dense-clump'' population are truly pressure-confined by their massive envelopes, as originally postulated by \citet{bm92}. \vspace{0mm}8.\ This in turn supports the view that a significant fraction of the observed dense clumps are long-lived structures, apparently in a state of low star formation rate/efficiency over several tens of Myr, until accumulating sufficient matter to pass a density threshold of $\sim$10$^{10}$ \htwo\ molecules m$^{-3}$, and only then engaging in vigorous massive star formation at the end of this time. \vspace{0mm}9.\ This is consistent with other studies that find a relatively small mass fraction $\xi_p<$ 10\% of molecular gas above a density $n=10^p$\,m$^{-3}$ for $p$=10, even while we find $\xi_p$ \gapp 50\% for $p$=9. It seems to be this smaller fraction of denser ``dense gas'' which is actively engaged in star formation, as found in several other studies. %\end{enumerate} \vspace{0mm}We look forward to further tests of the scenario argued for here, and proposed by \citet{b11}, that the long lifetime for massive dense clumps' quiescent phase is enabled by a pressure- and gravity-confined massive molecular envelope, that stabilises a less massive and still turbulent dense interior, which by itself is not dense enough to be gravitationally bound and efficiently form a star cluster until a sufficient density threshold is passed. \vspace{0.1mm} %% No more than seven \figcaption commands are allowed per page, %% so if you have more than seven captions, insert a \clearpage %% after every seventh one. \vspace{-8mm} | 16 | 9 | 1609.06752 |
1609 | 1609.02921_arXiv.txt | \noindent Most types of supernovae (SNe) have yet to be connected with their progenitor stellar systems. Here, we reanalyze the 10-year SN sample collected during 1998--2008 by the Lick Observatory Supernova Search (LOSS) in order to constrain the progenitors of SNe~Ia and stripped-envelope SNe (SE~SNe, i.e., SNe~IIb, Ib, Ic, and broad-lined Ic). We matched the LOSS galaxy sample with spectroscopy from the Sloan Digital Sky Survey and measured SN rates as a function of galaxy stellar mass, specific star formation rate, and oxygen abundance (metallicity). We find significant correlations between the SN rates and all three galaxy properties. The SN~Ia correlations are consistent with other measurements, as well as with our previous explanation of these measurements in the form of a combination of the SN~Ia delay-time distribution and the correlation between galaxy mass and age. The ratio between the SE~SN and SN~II rates declines significantly in low-mass galaxies. This rules out single stars as SE~SN progenitors, and is consistent with predictions from binary-system progenitor models. Using well-known galaxy scaling relations, any correlation between the rates and one of the galaxy properties examined here can be expressed as a correlation with the other two. These redundant correlations preclude us from establishing causality---that is, from ascertaining which of the galaxy properties (or their combination) is the physical driver for the difference between the SE~SN and SN~II rates. We outline several methods that have the potential to overcome this problem in future works. | \label{sec:intro} Of the various types of supernovae (SNe) we observe, only certain subtypes of SNe II have been conclusively connected to their progenitors. Pre-explosion imaging has shown that SNe II-plateau (SNe~IIP), for example, come from red supergiants (see review by \citealt{Smartt2009review}). Owing to their cosmological significance, the progenitor systems of SNe Ia have been pursued relentlessly over the last two decades, but their nature is still debated (see \citealt{2014ARA&A..52..107M} for a recent review). SNe II and stripped-envelope SNe (SE SNe, i.e., SNe~IIb, Ib, Ic, and broad-lined Ic; e.g., \citealt{1997ARA&A..35..309F,2001AJ....121.1648M,2014AJ....147...99M}) are attributed to the core collapse (CC) of stars more massive than $\sim 8~{\rm M}_\sun$. Spectroscopically, SE SNe are distinguished from SNe II by the lack of hydrogen features (either partial, as in SNe IIb, or nearly complete, as in SNe Ib; e.g., \citealt{1997ARA&A..35..309F,2016ApJ...827...90L}). SNe Ic also lack helium features. In order to explain this lack of hydrogen and helium, SE~SNe are thought to be the explosions of stars that have had their outer envelopes stripped away before the explosion (hence their name). Of all the processes suggested to explain this stripping, the leading models make use of either stellar winds (e.g., \citealt{2003ApJ...591..288H}), interaction with a binary companion (e.g., \citealt{1971ARA&A...9..183P,1992ApJ...391..246P,1998A&A...333..557D,1998A&ARv...9...63V}), or a combination of both (e.g., \citealt{2011MNRAS.412.1522S}). In the case of broad-lined SNe Ic connected to gamma-ray bursts, chemically homogeneous evolution (e.g., \citealt{2005A&A...443..643Y,2006ApJ...637..914W}) and explosive common-envelope ejection \citep{2010MNRAS.406..840P} have also been suggested. Pre-explosion imaging of the sites of these SNe has so far failed to reveal the nature of their progenitors conclusively (e.g., \citealt{2013MNRAS.436..774E,2016MNRAS.461L.117E,2016ApJ...825L..22F,2016ApJ...818...75V}), though the case for yellow supergiants in binary systems as the progenitors of SNe IIb is gaining traction (e.g., \citealt{2013ApJ...772L..32V,2014AJ....147...37V,2014AJ....148...68B,2014A&A...565A.114F,2015MNRAS.446.2689E}). New observational methods are required to address this question. SNe Ia are thought to be thermonuclear explosions of carbon--oxygen white-dwarf remnants of $<8$~${\rm M}_\sun$ stars. In order to disrupt the otherwise stable white dwarf, most models place it in a binary system where it can grow in mass, raising the temperature in the core until the carbon is ignited in a thermonuclear runaway. To grow in mass, the white dwarf can either siphon matter off of a main-sequence or evolved companion star (the so-called ``single-degenerate'' scenario; \citealt{Whelan1973}) or merge with a second carbon--oxygen white dwarf after the two spiral in owing to loss of energy and angular momentum to gravitational waves (the ``double degenerate'' scenario; \citealt{Iben1984,Webbink1984}). Recently, direct collisions of white dwarfs have also been suggested as a possible progenitor channel (e.g., \citealt{2012arXiv1211.4584K,2013ApJ...778L..37K,2015MNRAS.454L..61D}; but see also \citealt{2013MNRAS.430.2262H}). Many methods have been used to constrain these various SN progenitor models, including (but far from limited to) direct imaging of the explosion sites either before or long after the SN explosion (e.g., \citealt{2008MNRAS.388..421M,Li2011fe,2014MNRAS.442L..28G,2016ApJ...819...31G,2014ApJ...790....3K}), multiwavelength follow-up observations (e.g., \citealt{2012ApJ...746...21H,2013ApJ...767...71M,2014ApJ...790...52M,2016ApJ...821..119C}), and analyses of SN remnants (e.g., \citealt{2004Natur.431.1069R,Schaefer2012,2013ApJ...774...99K,2014MNRAS.439..354B}). Over the last few decades, studies have consistently shown that SNe Ia are more common in blue, star-forming, late-type galaxies than in red, passive, early-type galaxies (e.g., \citealt{1979AJ.....84..985O,1988A&A...190...10C,1989ApJ...345..752E,1990PASP..102.1318V,1991ARA&A..29..363V,1994ApJ...423L..31D,1997ApJ...483L..29W,cappellaro1999,2005ApJ...629..750D}). \citet{2006ApJ...648..868S} showed that SN Ia rates per unit mass decreased with increasing galaxy stellar mass. \citet[hereafter L11]{li2011rates} showed the same effect for all SN types, in all types of galaxies (but see Section~\ref{sec:SNIa}), and dubbed this the ``rate-size,'' or rate--mass, relation. We confirmed this trend for SNe~Ia in star-forming galaxies in \citet{GraurMaoz2013} and for SNe~II in \citet[hereafter G15]{2015MNRAS.450..905G}. Following \citet{Kistler2011}, we argued that the dependence of the SN Ia rates on stellar mass results from a combination of galaxy scaling relations (older galaxies are more massive, on average, than younger ones; \citealt{2005MNRAS.362...41G}) and the SN Ia delay-time distribution (DTD), which behaves as a power law with an index of $\sim -1$ (e.g., observations by \citealt{2008PASJ...60.1327T,Maoz2010clusters,Maoz2010magellan,Graur2011,Graur2014,2014AJ....148...13R}; and reviews by \citealt{2012NewAR..56..122W,2013FrPhy...8..116H,2014ARA&A..52..107M}). \citet{2005A&A...433..807M}, \citet{2006ApJ...648..868S}, \citet{2012ApJ...755...61S}, and G15 also measured SN Ia rates as a function of the galaxies' specific star-formation rate (sSFR). In G15, we showed that our explanation for the rate--mass correlation also explained the observed trend between SN Ia rates and sSFR, where the rates are constant in passive galaxies but rise with increasing sSFR in star-forming galaxies. In G15, we also measured SN rates as a function of stellar mass and sSFR for SNe II and claimed that their rate--mass relation was simply the result of their progenitors' short lifetimes: because SNe II come from massive ($>8~{\rm M}_\sun$) stars, their rates track the star-formation rates of their galaxies. Similar measurements of CC SN rates (i.e., combining SNe~II and SE~SNe) were made by \citet{botticella2012}. L11 was part of a series of papers that explored the SN sample collected by the Lick Observatory Supernova Search (LOSS; L11; \citealt{2011MNRAS.412.1419L,li2011LF,Maoz2010loss,2011MNRAS.412.1522S}). LOSS is an ongoing survey for SNe in local galaxies using the 0.76 m Katzman Automatic Imaging Telescope. For detailed descriptions of the survey, see \citet{2000AIPC..522..103L}, \citet{2001ASPC..246..121F}, and \citet{2003fthp.conf..171F,2005ASPC..332...33F}. Here, we use the LOSS sample to remeasure and reanalyze the SN rates originally published by L11. In Section~\ref{sec:galaxies}, we match between the LOSS sample and spectroscopy from the Sloan Digital Sky Survey (SDSS; \citealt{2000AJ....120.1579Y}) in order to go beyond L11 and measure the SN rates not only as a function of galaxy stellar mass but also of sSFR and metallicity, as expressed by the abundance of gas-phase oxygen in the centers of the LOSS galaxies. Throughout this work, we use the metallicity scale of \citet[hereafter T04]{2004ApJ...613..898T}. In Section~\ref{sec:addendum}, we make several addenda to the LOSS sample. We publish the control times necessary to measure SN rates with this sample, as well as updated tables of galaxy properties and SN rates. We deal with SNe Ia in Section~\ref{sec:SNIa}, and with CC SNe in Section~\ref{sec:CCSN}. We find significant correlations between the SN rates and the various galaxy properties. Most importantly, we find that the CC SN rates behave differently in different types of galaxies: the SE SN rates are shown to be depressed, relative to the SN II rates, in galaxies with low stellar mass, high sSFRs, and low metallicity values. Other studies have reported similar trends through measurements of correlations between fractions of SNe within a given sample and metallicity (e.g., \citealt{2008ApJ...673..999P,2009A&A...503..137B,2012ApJ...759..107K,2015PASA...32...19A}), or by splitting SN fractions between different types of galaxies, which encompass different metallicity regimes \citep{2010ApJ...721..777A,2014MNRAS.444.2428H}. We conduct an in-depth comparison of our results with these works in Section~\ref{sec:discuss}. We also show that our measurements rule out theoretical models based on single-star progenitors for SE SNe, but are consistent with models that assume binary-system progenitors. In Section~\ref{sec:discuss}, we additionally discuss how the various rate correlations are not independent. For all SN types, we show that a correlation with one galaxy property can be transformed into the measured correlations with any other of the galaxy properties studied here by using well-known galaxy scaling relations. This makes it impossible to distinguish causation from correlation, especially for the deficiency of SE SNe in lower-mass galaxies. However, we argue that the structure seen in the correlations (e.g., the way the ratio between the SE SN and SN II rates depends on galaxy stellar mass) can be incorporated into models and used to constrain progenitor models. We summarize our results in Section~\ref{sec:summary}. Paper II in this series \citep{2016arXiv160902923G} will use population fractions, as measured from the LOSS volume-limited subsample of SNe, to strengthen the results presented here and add further constraints on SN progenitor scenarios. For Paper II, we rely on a reclassification of the SNe in this subsample, as reported by \citet{2016arXiv160902922S}. | \label{sec:summary} This is the first of a series of papers in which we reanalyze the LOSS SN rates. Here, we matched the LOSS galaxy sample with SDSS and then remeasured the LOSS SN rates as a function of various global properties of galaxies: stellar mass, star-formation rates, and metallicity (in the form of nebular oxygen abundance). All of these measurements, including the control times necessary to compute them, are made public through the various tables in this work. We make the following observations. \begin{enumerate} \item The specific SN II rates are strongly correlated with all galaxy properties measured here. SE SN and SN Ia rates show strong correlations with stellar mass and sSFR, but not with nuclear metallicity. \item The SN Ia rate--mass correlation is statistically significant in star-forming galaxies, but not in passive ones, as also noted by \citet{2006ApJ...648..868S}. \item The SN Ia rates are well fit by a model that combines a $t^{-1}$ DTD with the galaxy mass--age scaling relation (Figures~\ref{fig:sSFR_Ia} and \ref{fig:OH_Ia}), as suggested by \citet{Kistler2011}, \citet{GraurMaoz2013}, and G15. \item The ratio between SE SN and SN II rates rises with galaxy stellar mass until it flattens in galaxies with $M_\star \ga 10^{10}~{\rm M}_\sun$ (or $12+{\rm log(O/H)}\ga9.2$) (Figures~\ref{fig:likelihood_ratio} and \ref{fig:ncomp}). This trend is statistically significant, at a $>3\sigma$ level. \item The rate ratio measurements rule out single stars as progenitors of SE SNe, but are consistent with models that assume binary-system progenitors, as suggested by earlier works (Figure~\ref{fig:ncomp}). \item Similar deficiencies in the SE SN rates relative to the SN~II rates are seen when correlating the rates with other galaxy properties, though those trends are not statistically significant (Figures~\ref{fig:rates_CC_sSFR} and \ref{fig:rates_OH}). \item The SN~II and SE~SN correlations do not exhibit significant breaks, which means that their underlying dependence on any of the galaxy properties studied here (or a combination of these properties) must be smooth within the dynamical range probed in this work. \item SE~SN host galaxies follow the same distribution in sSFR vs. metallicity space as SN~II host galaxies and LOSS galaxies that did not host SNe during the survey. \end{enumerate} Although the correlations shown here are broadly consistent with those shown in previous studies, the results of this work differ from previous studies by being based on absolute SN rates derived from a homogeneous, well-characterized SN sample. The LOSS sample, which is biased toward massive galaxies, has allowed us to sample a higher metallicity range than previous studies. Interestingly, in this range, the statistically significant correlation between the ratio of SE~SN to SN~II rates and galaxy stellar mass (or metallicity) levels off instead of continuing to increase monotonically. We have shown that, owing to the known galaxy scaling relations, any correlation between the SN rates---of any SN type---and a specific galaxy property can be transformed into the measured correlations with the other galaxy properties studied here. This precludes us from ascertaining which of the galaxy properties (or some combination of them) is responsible for the correlations we observe or for the deficiency of SE~SNe in low-mass galaxies. We have outlined several methods that might allow us to bypass this problem in future experiments. Finally, we have also enriched the LOSS sample with additional galaxy properties and the publication of the SN control times, so that further studies can be undertaken with this sample. | 16 | 9 | 1609.02921 |
1609 | 1609.07390_arXiv.txt | The 2014 March 29 X1 solar flare (SOL20140329T17:48) produced bright continuum emission in the far- and near-ultraviolet (NUV) and highly asymmetric chromospheric emission lines, providing long-sought constraints on the heating mechanisms of the lower atmosphere in solar flares. We analyze the continuum and emission line data from the Interface Region Imaging Spectrograph (IRIS) of the brightest flaring magnetic footpoints in this flare. We compare the NUV spectra of the brightest pixels to new radiative-hydrodynamic predictions calculated with the RADYN code using constraints on a nonthermal electron beam inferred from the collisional thick-target modeling of hard X-ray data from RHESSI. We show that the atmospheric response to a high beam flux density satisfactorily achieves the observed continuum brightness in the NUV. The NUV continuum emission in this flare is consistent with hydrogen (Balmer) recombination radiation that originates from low optical depth in a dense chromospheric condensation and from the stationary beam-heated layers just below the condensation. A model producing two flaring regions (a condensation and stationary layers) in the lower atmosphere is also consistent with the asymmetric Fe \textsc{ii} chromospheric emission line profiles observed in the impulsive phase. | The spectral energy distribution of the ultraviolet, optical, and infrared continuum (white-light) emission contains important information on the atmospheric response at the highest densities in the flare atmosphere but has remained largely unconstrained due to a lack of broad-wavelength spectral observations \citep{Fletcher2007}. The white-light emission provides important constraints on the strength and depth of flare heating resulting from magnetic energy release in the corona. White-light emission is thought to be produced by the energy deposition by nonthermal electrons due to the close spatial and temporal coincidence with hard X-ray emission \citep[e.g.,][]{RustHegwer1975, Hudson1992, Metcalf2003, Martinez2012}. The source of these nonthermal electrons is controversial; in the standard flare model, they are accelerated in the high corona as ``beams'' \citep[the collisional thick-target model;][]{Brown1971, Emslie1978}, but recently it has been suggested that limits on electron numbers and propagation effects (e.g., beam instabilities) require an alternative mode of energy transport, such as acceleration of particles in the lower corona or chromosphere by Alfv\'en waves \citep{Fletcher2008}. Continuum measurements from spectra are necessary to test heating models by providing constraints on the optical depths and electron densities that are attained in the deepest layers of the flare atmosphere. Spatially resolved flare spectra have been obtained at NUV/blue wavelengths from the ground and have shown a range of spectral properties around the expected location of the Balmer jump \citep{Neidig1983, Kowalski2015HSG}. The continuum emission from these spectra has been interpreted as optically thin hydrogen recombination radiation \citep{NeidigWiborg1984}, although an emission component from increased H$^{-}$ emission in the upper photosphere has also been suggested to explain some of the observed variation \citep{Hiei1982, Boyer1985}. However, the brightest regions of solar flares were rarely and poorly sampled in the past: the locations of the spectrographic slit relative to the small white-light footpoints were not precisely known \citep{NeidigWiborg1984, Neidig1989}, due to the rapid spatial and temporal evolution of the emission and also to variable and poor seeing during flares observed from the ground prior to the advent of adaptive optics \citep{Donati1984}. Therefore, the most extreme atmospheric conditions in the lower atmosphere during solar flares are not well constrained. Recently, Sun-as-a-star observations from SOHO/VIRGO's Sun PhotoMeter (SPM) have indicated the presence of a hotter blackbody emission component in the optical with a color temperature of $T\sim9000$ K \citep{Kretzschmar2011} which implies very large heating at high densities. High spatial resolution observations from Hinode during two X-class flares have shown a much lower color temperature in the optical of only $T\sim5000-6000$ K \citep{Watanabe2013, Kerr2014}, which is consistent with photospheric heating by several hundred K or optically thin hydrogen recombination radiation from heating of the mid-chromosphere. Blackbody fitting of NUV and optical spectra and broadband photometry of magnetically active M dwarf stars in the gradual and impulsive phases of flares results in larger color temperatures, $T\gtrsim 8000$ K but typically $T\sim9000-12,000$ K \citep{Hawley1992, Hawley2003, Zhilyaev2007, Fuhrmeister2008, Kowalski2010, Kowalski2013}. Clearly, a more thorough spatially resolved characterization of the brightest flare footpoints is necessary to determine the prevalence of hot blackbody-like emission in solar flares, and if any proposed heating model can self-consistently explain the implied heating requirements. Previous spectral observations of the NUV and optical have been interpreted using static isothermal slab models \citep[e.g.,][]{Donati1985} or semi-empirical static models \citep{Machado1980, Mauas1990, Kleint2016}, but the flare atmosphere is known to be highly dynamic and stratified \citep{Cauzzi1996, Falchi2002}. It has been proposed that the continuum emission may originate in impulsively-generated downflows in the upper chromosphere \citep{Livshits1981, Gan1992}, or chromospheric ``condensations'' \citep[hereafter CC;][]{Fisher1985b, Fisher1989}, which are also attributed to the formation of H$\alpha$ red-wing emission components that are often observed in solar flares \citep{Ichimoto1984, Canfield1987, Canfield1990}. \cite{Kowalski2015} recently found that an extremely large electron beam flux of $10^{13}$ \fluxunits\ could produce hot $T\sim10,000$ K blackbody-like emission in very dense CCs. With current computational facilities and constraints on electron beam fluxes from the Ramaty High Energy Solar Spectroscopic Imager \citep[RHESSI;][]{Lin2002} and high spatial resolution imagery, it is timely to critically examine the hydrodynamic and time-dependent radiative response of the models and compare to new spectral observations of solar flares. We have begun a large campaign to characterize the emission properties of the brightest flaring magnetic footpoints during Cycle 24 flares using new NUV and far-ultraviolet (FUV) spectra from the Interface Region Imaging Spectrograph \citep[IRIS;][]{DePontieu2014}. The high spatial resolution of IRIS allows improved intensity measurements of the continuum emission, which is observed in compact sources, or kernels, as small as 0.\arcsec3 \citep{Jess2008}. In this paper, we present radiative-hydrodynamic (RHD) modeling of the brightest continuum flaring pixels in the X1 flare of 2014-Mar-29, which has been extensively studied by \cite{Judge2014, Young2015, Liu2015, Battaglia2015, Kleint2015, Matthews2015, Kleint2016, Fatima2016}. In \cite{Heinzel2014} and \cite{Kleint2016}, the bright NUV continuum emission from this flare was identified and compared to static beam heated model atmospheres from \cite{Ricchiazzi1983} and to static phenomenological models with the RH code. \cite{Heinzel2014} concluded that the NUV continuum intensity was consistent with optically thin Balmer continuum emission. However, the time-dependent radiative transfer and the hydrodynamics, which can affect beam propagation through evaporation and condensation, have not yet been compared in detail to the continuum observations; \cite{Heinzel2015IAUS} and \cite{Fatima2016} present new RHD simulations with the \emph{Flarix} and RADYN codes, respectively, for relatively low beam fluxes compared to the flux inferred from imaging spectroscopy of the brightest source in the flare \citep{Kleint2016}. In this paper, we use the state-of-the-art Fokker-Planck treatment of energy deposition \citep{Mauas1997, Battaglia2012, Liu2009, Allred2015} from a high-flux electron beam, in order to understand the time evolution of the atmospheric stratification that self-consistently explains both the NUV continuum emission and chromospheric line profiles in the brightest flaring footpoints. By rigorously testing new RHD models guided by the combined information from RHESSI and new data of the white-light continuum and chromospheric lines, we seek answers to the following questions: \begin{itemize} \item Using electron beam parameters inferred from standard thick target modeling of RHESSI X-ray data, do electron beams produce an atmospheric and radiative response that is consistent with the high spatial and spectral line and continuum constraints from IRIS? \item Does a CC form that is hot and dense enough to explain the observed IRIS line and continuum emission? % \item Does the hydrodynamic response of the atmosphere to beam heating result in flare continuum emission in the IRIS NUV channel that is predominantly optically thin hydrogen recombination radiation? Is there evidence for hot ($T\gtrsim9000$ K) blackbody-like radiation from photospheric densities? \end{itemize} The paper is organized as follows. In Sections \ref{sec:calibration}-\ref{sec:area}, we describe the calibration of the IRIS observations and the high spatial resolution flare footpoint development; in Section \ref{sec:wl}, we discuss the constraints on the continuum intensity and emission line profiles in the spectra of the brightest flaring footpoints; in Section \ref{sec:radyn} we describe the radiative-hydrodynamic modeling of these spectra and the formation of the NUV continuum and Fe \textsc{ii} emission lines; in Section \ref{sec:sjibright}, we compare the modeling results to the brightest sources in the slit jaw images; in Section \ref{sec:limitations} we discuss several limitations of the modeling and future work; in Sections \ref{sec:summary} - \ref{sec:conclusions} we present our conclusions. Appendices A and B discuss broader wavelength model predictions for the optical continuum emission. | \label{sec:conclusions} The conclusions from our work are the following: \begin{itemize} \item The radiative-hydrodynamic modeling of the atmospheric response to a high flux (5F11) nonthermal electron beam is an improvement over static flare modeling because the dynamic effects on the emergent spectrum can be critically examined and compared to red-shifted emission line components. The NUV continuum brightness changes due to the atmosphere density evolution within a heated, chromospheric compression (condensation) that develops on short (several second) timescales in the 5F11 model. The 5F11 produces an electron density ($\sim4-5\times10^{13}$ cm$^{-3}$) in the early phase that is consistent with the values inferred from previous static, slab modeling of the Balmer continuum and lines \citep{Donati1985}, but the density in NUV continuum emitting layers increases by another order of magnitude as the condensation develops. \item In the brightest spectra of the \marflare, the excess NUV continuum excess can be explained by hydrogen Balmer recombination emission over several hundred km at chromospheric heights ($z\sim630-1020$ km) with low optical depth ($\tau_{2826} \le 0.2$). The excess NUV continuum emission originates over two flaring layers that are heated by the nonthermal electron beam: a chromospheric condensation with vertical downward velocities of $\sim20-55$ \kms\ and stationary flare layers just below the condensation. A variety of methods have inferred a low optical depth over the region producing the white-light continuum intensity \citep{Hudson1972, Neidig1983, Potts2010, Heinzel2014}. The high flux 5F11 RHD model demonstrates that a heated, downflowing compression increases the density to large values ($n_{\rm{H, max}} \sim 5\times10^{14}$ cm$^{-3}$) after 4~s when the continuum has nearly reached maximum brightness, but this density is not enough to produce a large optical depth. \item The low energy electrons ($E\sim25-50$ keV) in the beam heat the chromospheric condensation and higher energy electrons ($E\gtrsim 50$ keV initially, $E\gtrsim 80$ keV after the condensation becomes dense) heat the stationary flare layers. In the first few seconds of the beam heating the NUV continuum emission in the stationary flare layers contributes to the majority of the emergent intensity enhancement. Therefore, we expect less continuum intensity in the first few seconds of footpoint brightening (before the CC becomes bright and dominates the contribution to the emergent intensity) for flares with softer (higher $\delta$) time-averaged electron beams (but for same energy flux) because fewer $E>50$ keV electrons are available to heat the stationary flare layers. Although there has been no significant observational relationship established between hard X-ray spectral hardness and flare peak optical continuum intensity \citep{Matthews2003, Kuhar2016}, higher cadence observations than currently available for flare white-light emission are needed to critically test the models by constraining the properties of the continuum emission when it predominantly originates from the stationary flare layers, before the RWA line components develop. It has been proposed that the hardness of the electron beam for very high beam fluxes could explain the interflare variation in observed continuum flux ratios in dMe flares \citep{Kowalski2016}. \item We have developed a technique to include the Mg II $h+k$ wing opacity in the calculations of the excess continuum intensity for an accurate comparison to IRIS NUV observations. The far wing opacities are important for an accurate treatment of (sub-)photospheric continuum dimming as well as the amount of Mg II wing emission in the upper photosphere in response to a moderate temperature change. If these opacities are neglected, errors of $15-30$\% result for the excess continuum intensity in lower beam flux models such as the F11 model (Table \ref{table:model}). \item The \feone\ and \fetwo\ profiles provide new information on the dynamics in continuum-emitting layers at $T\sim8500-25,000$ K and electron densities of $\sim 5\times10^{13} - 5\times10^{14}$. The LTE Fe \textsc{ii} line profiles predicted by the 5F11 are qualitatively consistent with the spectral observations of the brightest flaring footpoints: the profiles are spectrally resolved and exhibit an emission component close to the rest wavelength that is produced in the stationary flare layers and a bright redshifted emission component that is produced in the chromospheric condensation. \item The physical depth range parameter is an important parameter for understanding how the emergent intensity varies as a function of wavelength for emission lines and continua because the physical depth range reflects the variation of $\tau_{\lambda}$. The \feone\ line is optically thin enough at line center to probe the conditions over a significant physical depth range in the flare chromosphere, but is more optically thick than the NUV continuum intensity. Interestingly, the spectrum of BFP1 (with an exposure time of 8~s) and the spectrum of BFP2 (with an exposure time of 2.4~s) exhibited a similar excess NUV continuum intensity and brightness in the rest-wavelength component of the chromospheric emission lines, but showed striking differences in the strength and redshift of the red-wing asymmetry (RWA) emission component in the chromospheric line profiles (Figures \ref{fig:feii_2814} and \ref{fig:feii_2832}). The range of relative brightnesses of the two emission components ($I_{\rm{RWA}} / I_{\lambda_{\rm{rest}}}$) for each Fe II line is adequately reproduced in the 5F11 model (using simulated exposure times of 2.4-8~s). The differences in $I_{\rm{RWA}} / I_{\lambda_{\rm{rest}}}$ between \feone\ and \fetwo\ are due to optical depth differences at \lamr: \fetwo\ is more optically thick at \lamr\ and has a larger value of $I_{\rm{RWA}} / I_{\lambda_{\rm{rest}}}$. Although there may be departures from LTE that are not accounted for in our models of Fe II, the LTE assumption identifies several important atmospheric parameters that lead to the line profile properties over an exposure time: 1) the brightness at \lamr\ is determined by the optical depth at line center, the beam energy deposition evolution, and the extent to which the stationary flare layers have been accrued by the chromospheric condensation; 2) the peak wavelength (and to some extent the broadening) of the red shifted emission line component is determined by the density and velocity evolution of the CC as it cools from $T=25,000$ K to $T=8500$ K. \item The Fe \textsc{ii} lines do not exhibit measurable thermal or pressure broadening, and a nonthermal broadening is required to account for the width of the observed line component at \lamr. We find that a time-variable nonthermal broadening given by a simple physical prescription for velocity broadening is appropriate for the profile shape, reproduces the observed bisector velocity, and accounts for some of the far red emission of the RWA emission component. \item The coronal heating model and 5F11 beam model produce redshifts of the H$\alpha$ line by $\sim30$ \kms\ and a relatively bright RWA emission component in the LTE model of the \feone\ line, but the conductive heating flux into the chromosphere does not produce bright NUV continuum intensity as observed in the spectra of BFP1 and BFP2. \item At the brightest times in the 5F11, the continuum intensity exceeds the spectroscopic constraints of the continuum intensity, and the predicted Fe \textsc{ii} profiles are very bright (a factor of $1.5-3$ brighter than the spectral observations). The goal of this study is not to precisely match all possible observables to the model intensity values, given the limits on spatial resolution and our simplified prescription of the electron beam heating function such as a constant heating flux and not including return current effects. Our goal is to quantify the range of continuum brightness values and line profile shapes achieved by the models and determine if they sweep through the regime of the observations as we change the beam flux to the highest values that are within reason. In future work, we intend to explore the range of fluxes between F11 and 5F11, which were chosen to bracket the value of the inferred flux for the brightest kernel \citep{Kleint2016}, and the range of possible values of $E_c$. A 2.5F11 $-$ 3.5F11 may produce a closer match to the observed intensity in the spectra. We speculate that a soft-hard-soft variation as observed in short hard X-ray flare bursts on the Sun \citep[e.g.,][]{Fletcher2002, Grigis2004} may also help reduce the rest-wavelength intensity averaged over an exposure time because less beam energy is deposited in the stationary flare layers for softer (higher $\delta$) beams. Nonetheless, we have tested the 5F11 against the slit jaw constraints of the brightest kernel in the flare and found that the very bright continuum and Fe II emission lines at $t\sim4$ sec are consistent with these constraints (Section \ref{sec:sjibright}). However, spectroscopic confirmation of these extremely bright continuum and emission line intensities is necessary. In a follow-up paper on the 5F11 model, we present optical predictions of the Balmer and Paschen jump region at $t=3.97$~s using the modeling techniques of the Balmer edge region in dMe flares (Kowalski et al. 2015b) in order to test the 5F11 model against future spectral observations of the brightest flaring kernels, such as with the Daniel K. Inoue Solar Telescope. \end{itemize} A high flux electron beam using the free-streaming, thick-target model reproduces several of the critical observations for the March 29th 2014 X1 flare, but the high nonthermal electron flux of 5F11 will require further modeling of the energy loss from the return current electric field as well as including the effects of beam instabilities. \emph{Despite the simplifications in the thick target electron beam model employed here, we conclude that flare heating with a high flux electron beam can be used as a powerful tool to interpret spectral phenomena and to understand important radiative-hydrodynamic processes in the brightest flaring footpoints.} The consistencies between the 5F11 model predictions and the IRIS observations will serve as a benchmark for models with additional physical processes that address observational challenges \citep{Battaglia2006, Krucker2011, Martinez2012, Dickson2013, Simoes2013} to the standard electron beam model for the brightest hard X-ray flare footpoints. Measurements of the Balmer jump ratio as a constraint on optical depth, the broadening of the hydrogen lines as a constraint on electron density, and modeling lines such as Si I \citep{Judge2014} for constraints on the heating of the upper photospheric layers will complement future comparisons of IRIS flare spectra and RHD models. | 16 | 9 | 1609.07390 |
1609 | 1609.05440_arXiv.txt | Classical Cepheids are among the most useful Galactic and nearby extragalactic distance tracers because of their well-defined period--luminosity relations (PLRs). Open cluster (OC) Cepheids are important objects to independently calibrate these PLRs. Based on Data Release 1 of the {\sl VISTA} Variables in the V\'ia L\'actea survey, we have discovered four new, faint and heavily reddened OC Cepheids, including the longest-period OC Cepheid known, ASAS J180342$-$2211.0 in Teutsch 14a. The other OC--Cepheid pairs include NGC 6334 and V0470 Sco, Majaess 170 and ASAS J160125$-$5150.3, and Teutsch 77 and BB Cen. ASAS J180342$-$2211.0, with a period of $\log P = 1.623$ [days] is important to constrain the slope of the PLR. The currently most complete $JHK_{\rm s}$ Galactic Cepheid PLRs are obtained based on a significantly increased sample of 31 OC Cepheids, with associated uncertainties that are improved by 40 per cent compared with previous determinations (in the $J$ band). The NIR PLRs are in good agreement with previous PLRs determined based on other methods. | Classical Cepheids are among the most useful Galactic and nearby extragalactic distance tracers because of their well-established period--luminosity relations \citep[PLRs; also known as the `Leavitt law':][]{Leavitt12}. Cepheids are commonly used to anchor extragalactic distances, constrain the Hubble constant and study Galactic structure. They also relate directly to secondary distance tracers, such as Type Ia supernovae, whose luminosities are calibrated using Cepheids in their host galaxies \citep{Riess11}. The Carnegie Hubble Program \citep{Freedman11} aims at determining an accurate Hubble constant based on the Cepheid distance scale, \citet{Matsunaga11, Matsunaga15}, \citet{Feast14} and \citet{Dekany15} have studied the Milky Way's structure in the Galactic bulge based on Cepheids, while \citet{Majaess09} and \citet{Dambis15} mapped the Galactic spiral arms using distances to classical Cepheids. Cepheid PLRs are not only important for measuring distances but also to constrain the pulsation physics and the evolution of Cepheids. However, independent Cepheid distances must be obtained, which can be achieved using (i) trigonometric parallaxes \citep{Feast97, Benedict07}, (ii) Baade--Wesselink-type methods/surface-brightness techniques \citep{Gieren97, Storm11} and (iii) main-sequence or isochrone fitting based on open cluster (OC) photometry \citep{An07, Turner10, Anderson13, Chen15}. The OC parameters derived based on the latter method can be used as independent constraints on the distances, ages and reddening values of Cepheids in these clusters. This method can also be used to study Cepheids in heavily reddened environments. The key to establishing the PLR using OC Cepheids is to select high-confidence cluster Cepheids. \citet{An07} obtained $BVI_{\rm c}JHK_{\rm s}$ PLRs based on seven Galactic OCs hosting Cepheids, while \citet{Anderson13} found five OC Cepheids using an eight-dimensional selection approach and obtained a $V$-band PLR for 18 OC Cepheids. \citet{Chen15} found six new OC Cepheids based on near-infrared (NIR) data and derived a $J$-band PLR for 19 OC Cepheids. As more accurate data are becoming available, the number of OC Cepheids is increasing, whereas the confidence of OC membership is improving simultaneously. \citet{Majaess11} provided new evidence to support the likely membership of the Cepheid TW Nor of the OC Lyng{\aa} 6. \citet{Turner12} discovered that SU Cas is a probable member of the OC Alessi 95; \citet{Majaess12a} conducted a detailed distance analysis of this latter OC and its member Cepheids based on new X-ray and $JHK_{\rm s}$ data. \citet{Majaess13} assessed the links between three Cepheids and NGC 6067 based on data from the {\sl VISTA} Variables in the V\'ia L\'actea (VVV) survey. However, the number of OC Cepheids that can presently be used to establish OC--Cepheid PLRs is still small, of order 20. One limitation to further progress is that approximately half of the OCs in the DAML02 catalogue \citep{Dias02} are poorly studied. Another limitation is a lack of availability of mean Cepheid intensities in multiple passbands; indeed, for many Cepheids we only have access to their $V$-band light curves. In this paper, we use VVV Data Release 1 (DR1) \citep{Minniti10,Saito12} to study faint OCs in the Galactic bulge and the Galactic midplane. We improve the quality of their age, reddening and distance determinations. We also convert single-epoch NIR photometry to mean intensities based on application of light-curve templates. We have collected the current largest sample of 31 high-confidence OC Cepheids based on NIR photometry, thus enabling us to obtain the most accurate $JHK_{\rm s}$ PLRs for Galactic OC Cepheids to date. In fact, in this paper we derive the first statistically meaningful OC--Cepheid $H$- and $K_{\rm s}$-band PLRs; our updated $J$-band PLR represents a significant improvement with respect to our earlier derivation \citep{Chen15}. In section 2 we present the data, the method used and our OC--Cepheid catalogue. Section 3 discusses the OC membership characteristics of our four newly found Cepheids. In Section 4 we discuss the Cepheid $JHK_{\rm s}$ PLRs, while Section 5 summarizes and concludes this paper. | The availability of VVV DR1 provides a unique opportunity to uncover and study faint and highly obscured OCs in the Galactic plane near the Galactic Centre. Using the DAML02 OC catalogue and other new OCs found in the VVV data, we have carefully examined 22 OCs and discovered four new OC Cepheids. Parameters such as distances, reddening values and ages were better determined compared with previous work based on isochrone fitting. By comparison of distances, apparent positions, proper motions and ages of the clusters and Cepheids, the newly found OC--Cepheid pairs include NGC 6334 and V0470 Sco, Teutsch 14a and ASAS J180342$-$2211.0, Majaess 170 and ASAS J160125$-$5150.3, and Teutsch 77 and BB Cen. ASAS J180342$-$2211.0 is the longest-period Cepheid -- $\log P = 1.623$ [days] -- thus far found in an OC, which is important in the context of constraining the slope of the PLR. The currently most complete NIR Cepheid PLRs based on 31 OC Cepheids were obtained (a significant improvement from the previous PLRs which were based on up to 19 OC Cepheids), i.e., $\langle M_J \rangle = (-3.077 \pm 0.090) \log P-(2.202 \pm 0.090)$, $\langle M_H \rangle = (-3.167 \pm 0.075) \log P-(2.434 \pm 0.074)$ and $\langle M_{K{\rm s}} \rangle = (-3.224 \pm 0.073) \log P-(2.469 \pm 0.072)$. The associated uncertainties have been improved by 40 per cent compared with previous Cepheid PLRs based on the OC-matching method. Our PLRs are in good agreement with the best NIR PLRs available for all Galactic Cepheids obtained using other methods. Particularly in the $K_{\rm s}$ band, the zero-point differences are very small; differences of less than 0.06 mag are expected for any Cepheid. Since there are more than 20,000 Cepheids and 20,000 OCs in the Galaxy, we expect that thousands of OC Cepheids will be discovered in the future. {\sl Gaia} will observe some 9000 Cepheids \citep{Windmark11}, which will improve the distances, proper motions and radial velocities of not only the Cepheids but also their host clusters. By combining the more accurate parallax distances from {\sl Gaia} with OC distances based on isochrone fitting, the Cepheid PLRs will be left with only small systematic errors. We thank the referees for their comments. We are grateful for support from the National Natural Science Foundation of China through grants 11633005, U1631102, 11373010 and 11473037. | 16 | 9 | 1609.05440 |
1609 | 1609.04266_arXiv.txt | We present a study of the magnetic field properties of NGC~4038/9 (the `Antennae' galaxies), the closest example of a late stage merger of two spiral galaxies. Wideband polarimetric observations were performed using the Karl G. Jansky Very Large Array between 2 and 4 GHz. Rotation measure synthesis and Faraday depolarization analysis was performed to probe the magnetic field strength and structure at spatial resolution of $\sim1$ kpc. Highly polarized emission from the southern tidal tail is detected with intrinsic fractional polarization close to the theoretical maximum ($0.62\pm0.18$), estimated by fitting the Faraday depolarization with a volume that is both synchrotron emitting and Faraday rotating containing random magnetic fields. Magnetic fields are well aligned along the tidal tail and the Faraday depths shows large-scale smooth variations preserving its sign. This suggests the field in the plane of the sky to be regular up to $\sim20$ kpc, which is the largest detected regular field structure on galactic scales. The equipartition field strength of $\sim8.5~\mu$G of the regular field in the tidal tail is reached within a few 100 Myr, likely generated by stretching of the galactic disc field by a factor of 4--9 during the tidal interaction. The regular field strength is greater than the turbulent fields in the tidal tail. Our study comprehensively demonstrates, although the magnetic fields within the merging bodies are dominated by strong turbulent magnetic fields of $\sim20~\mu$G in strength, tidal interactions can produce large-scale regular field structure in the outskirts. | Magnetic fields are pervasive in the Universe on all scales and they play crucial roles in various processes in the interstellar medium. The large-scale ordered magnetic fields ($\gtrsim1$ kpc) in galaxies are thought to be amplified via the $\alpha$--$\Omega$ dynamo mechanism---the buildup of initially weak seed fields ($<10^{-9}$ G; \citealt{ade15}) to microgauss fields via small-scale turbulence and differential rotation \citep{ruzma88, kulsr08}. The small-scale dynamo can efficiently amplify the magnetic fields on scales $\lesssim1$ kpc in $\sim10^6$ years \citep[][]{kandu99, feder11, chama13, schob13}. On the other hand, the conventional $\alpha$--$\Omega$ dynamo action requires $\sim10^9$ years to amplify the large-scale magnetic field in galaxies \citep{arsha09, pakmo14}, which is too long to explain the detection of coherent fields in young systems \citep[e.g.,][]{berne08, farne14}. This suggests that there must be other magnetic field amplification processes at work. \begin{figure*} \begin{tabular}{cc} {\mbox{\includegraphics[width=9.2cm, trim=5mm 10mm 2mm 0mm,clip]{figure1_LHS.eps}}} & {\mbox{\includegraphics[width=8.3cm, trim=40mm 10mm 105mm 5mm,clip]{figure1_RHS.eps}}} \\ \end{tabular} \caption{{\it Left-hand panel:} Total intensity image of the Antennae made by combined DnC+CnB array data at a centre frequency of 2.8 GHz having angular resolution $11\times9$ arcsec$^2$. Overlaid are the contours at [$-3, -2, 3, 6, 12, 24, 50, 75, 100, 150, 200, 350, 500$] $\times 40~\mu$Jy beam$^{-1}$ levels. The dashed contours represents negative total intensities. {\it Right-hand panel:} We label the various regions in NGC4038/9 that are studied in the paper on the HST-F658N narrow-band filter image tracing the H$\alpha$ emission of NGC 4038/9 with total radio continuum intensity contours shown same as the left-hand panel.} \label{totI} \end{figure*} In the current framework of hierarchical structure formation, galaxies build up their mass by merging. Galaxy encounters can compress, stretch and reshape fields in the progenitor galaxies, hence they provide a conducive environment for magnetic field amplification \citep{kotar10}. Since merger events were more frequent in the early Universe \citep[e.g.,][]{patto02}, it is important to assess how galaxy interactions affect the strength and geometry of galactic-scale magnetic fields in order to understand the overall evolution of cosmic magnetism. Gravitationally interacting galaxies possess a range of magnetic field properties. For example, despite their irregular appearances, the tidally interacting Magellanic Clouds have been shown to host large-scale ordered fields of microgauss strength \citep{gaens05, mao08, mao12}. On the other hand, \citet{drzaz11} suggested, based on study of 16 merger pairs, that interacting galaxies have lower field regularities and stronger total magnetic field strengths than non-interacting ones. Unfortunately, due to the lack of Faraday rotation measure (RM) information, the magnetic field coherency in these systems could not be probed. Moreover, the large distances to their sample galaxies and the limited angular resolution prevent one from studying magnetic field structures on scales $< 7$ kpc. Therefore, a high resolution mapping of the magnetic field of a prototypical merger event is much needed. To date, detailed high-resolution studies of galactic magnetic fields are of individual galaxies in isolation---only a few focus on magnetism in interacting galaxies \citep[e.g.,][]{brind92, humme95, chyzy04, rampa08}. The Antennae pair (NGC 4038/9) is the nearest merger \citep[22 Mpc;][]{schwe08} between two gas-rich spirals. The bodies of the colliding galaxies host sites of active star formation in the form of super starclusters, with a global star formation rate of 20 M$_\odot$\,yr$^{-1}$ \citep{zhang01}. The Antennae galaxies hosts tidal tails that measure over 100 kpc in H{\sc i} and likely originate from outskirts of the progenitors. The Antennae have been studied extensively from the radio to X-ray. They are also the subject of several numerical simulations \citep[e.g.,][]{mihos93, karl10, kotar10}. The latest work by \citet{karl10} proposed that the progenitors first encountered $\sim600$ Myr ago and they have just undergone the second passage. Our understanding of the pairs' interaction history, its multi-wavelength emission and its proximity make the Antennae an ideal candidate for a detailed magnetic field study. Magnetism in the Antennae was studied by \citet{chyzy04} with the Very Large Array at 1.49, 4.86 and 8.44 GHz. Enhanced polarized emission is found near the root of a tidal tail which is suggestive of a remnant spiral field. \citet{chyzy04} computed the RM of diffuse polarized emission at 8.44 GHz and 4.86 GHz at a resolution $17\times14$ arcsec$^2$ ($\sim2$ kpc linear scale). They pointed out that in the region where the galactic discs overlap, RMs are coherent on the scale of several synthesized beams, likely tracing the large-scale magnetic fields in the progenitors. There are several regions with consistent RM sign, which is suggestive of coherent magnetic fields. However, these RMs could suffer from the $n\pi$ ambiguity because they were computed using polarization angle measurements at only two bands, separated widely in frequency. A wideband study of the diffuse polarized emission from the Antennae is much needed to consistently derive RM to confirm the existence of coherent magnetic fields. In this paper, we present study of the magnetic field properties in the Antennae galaxies. In \textsection 2, we describe our observations and data analysis procedure. The results on Faraday depolarization and magnetic field strengths are presented in \textsection 3 followed by discussion in \textsection 4. Our results are summarized in \textsection 5. | \subsection{Regular field of $\sim20$ kpc in the tidal tail} \label{orderedtail} We detect highly polarized emission along the southern tidal tail of the Antennae with intrinsic polarization fraction close to the theoretical maximum (see \textsection\ref{depoltailmerger}). The magnetic field orientations in the plane of the sky corrected for Faraday rotation are well-aligned along the tail (see Figure~\ref{polimap}). Such a high degree of polarization can originate from a combination of anisotropic turbulent magnetic fields by compressing an initially random field, and regular magnetic fields in the plane of the sky \citep{beck16}. However, anisotropic fields dominating over the regular field will not contribute to Faraday depth \citep{jaffe10}. From our data, we find the Milky Way foreground-corrected Faraday depth along the tidal tail varies smoothly with positive sign throughout (see Figure~\ref{rmmap}) indicating the line-of-sight ordered field points toward the observer. Anisotropic random fields alone cannot give rise to a smooth large-scale variation of Faraday depth measured over several beams. We therefore conclude that the magnetic field in the plane of the sky along the tidal tail is regular (or coherent) and maintains its direction over $\sim20$ kpc. To our knowledge, this is the largest coherent magnetic field structure observed in galaxies. We discuss in detail the properties of the magneto-ionic medium in the tidal tail. \subsubsection{Turbulent cell size} The standard deviation of the Faraday depth within the 3D beam ($\sigmaRM$) in the tidal tail region is found to be $131$ rad m$^{-2}$ (see \textsection\ref{depoltailmerger}). This dispersion of the RM is caused by fluctuations of the field strength along the line of sight. The $\sigmaRM$ depends on the turbulent magnetic field along the line of sight and the properties of the magneto-ionic medium as \begin{equation} \sigma_{\rm RM} = 0.81 \left(\frac{\langle n_{\rm e}\rangle}{\rm cm^{-3}}\right) \left(\frac{B_{\rm turb, \parallel}}{\rm \mu G}\right) \left(\frac{L_{\rm ion} d}{f}\right)^{1/2}. \label{sigmaRM} \end{equation} Here, $\langle n_{\rm e}\rangle$ is the average thermal electron density along the line of sight, $B_{\rm turb, \parallel}$ is the strength of the random magnetic field along the line of sight, $L_{\rm ion}$ is the path length through the ionized medium (in pc), $d$ is the size of the turbulent cells (in pc) and $f$ is the volume filling factor of electrons along the line of sight. The RM dispersion in the plane of the sky ($\sigma_{\rm RM, sky}$) is related to the $\sigma_{\rm RM}$ as \citep{fletc11} \begin{equation} \sigma_{\rm RM, sky} \simeq N^{-1/2} \sigma_{\rm RM}, \end{equation} where, $N\approx(D/d)^2$ is the number of turbulent cells for the projected beam area in the sky of diameter $D$, for which $\sigma_{\rm RM, sky}$ is measured. Thus, we can compute the diameter of typical turbulent cell size $d$ using the observed $\sigma_{\rm RM, sky}$, $\sigma_{\rm RM}$, and the beam size $D\approx 1400$ pc. $d$ is found to be $\sim230$ pc, significantly larger than the typical turbulent cell size of $\sim50$ pc observed in the discs of galaxies \citep{fletc11, mao15, haver08}. Hence, assuming the line of sight extent of the tidal tail to be the same as the thickness in the plane of sky, i.e. $L_{\rm ion}=1.1$ kpc, the field along the line of sight must be regular. Therefore, the observed dispersion of the Faraday depth must be caused by systematic variations in the plane of the sky. The Faraday depth indeed varies smoothly in the tidal tail (see Figure~\ref{rmmap}). \subsubsection{Regular field strengths} The mean RM depends on the regular component of the magnetic field along the line of sight ($B_{\rm reg, \parallel}$) as \begin{equation} \langle {\rm RM}\rangle = 0.81 \langle n_{\rm e} \rangle B_{\rm reg, \parallel} L_{\rm ion \label{eqrm}}. \end{equation} Thus, the ratio of $\sigma_{\rm RM}$ to $\langle \rm RM \rangle$ can give us the estimate of $B_{\rm turb, \parallel}/B_{\rm reg, \parallel}$ and is given by \begin{equation} \frac{\sigmaRM}{|\langle \rm RM \rangle|} = \frac{B_{\rm turb, \parallel}}{|B_{\rm reg, \parallel}|}\left(\frac{d}{f L_{\rm ion}} \right)^{1/2}. \end{equation} From our estimated values of $\sigmaRM$, $\langle \rm RM\rangle$ and $d$, and assumed value for $L_{\rm ion}$, we find $B_{\rm turb, \parallel}/B_{\rm reg, \parallel}\sim 13 f^{1/2}$. Compared to the star-forming disc, the ionized medium in the tidal tail is more diffuse. Therefore, assuming typical $f$ in the range $0.2 - 0.5$ for a diffuse medium, $B_{\rm turb, \parallel}/B_{\rm reg, \parallel}$ lies in the range $5.8-9.2$. Assuming $B_{\rm turb}=5~\mu$G is isotropic (see \textsection\ref{bord}), we estimate $B_{\rm reg, \parallel}$ to be $\lesssim 1~\mu$G. Thus, the field strength in the plane of the sky ($B_{\rm reg, \perp} \sim8.5~\mu$G) is significantly larger than $B_{\rm reg, \parallel}$, perhaps caused due to stretching and twisting of the remnant spiral field in the disc of the galaxies during the tidal interaction. Here, we explore the degree of stretching required to amplify an initial regular field in the progenitor galaxy to the observed field strengths in the tidal tail. For this, we consider the scenario that an initially cylindrical spiral-shaped regular field generated in the progenitor galaxies by dynamo action has been stretched by the tidal interaction. Assuming the field is stretched keeping the cross-sectional area constant, then the ratio $B/\rho l$ is conserved if the magnetic flux is frozen in the tube \citep[see e.g.,][]{longa11}. Here, $\rho$ is the density of the gas and $l$ is the length of the cylinder. From observations of H{\sc i} gas \citep{hibba01}, the density in the tidal tail is about a factor of 3 lower than that in the remnant spiral arm. Therefore, a regular field of 8.5 $\mu$G in the tidal tail requires stretching by a factor of $\sim4-9$ for an initial regular field of $\sim3-6~\mu$G in strength (see \textsection\ref{n-arm}). This implies, the dynamo generated initial field was regular over $\sim2-5$ kpc, which is a typical length-scale observed in nearby spiral galaxies \citep{beck13book}. Thus, tidal stretching of field lines can also amplify large-scale regular fields within the dynamical time-scale of the merger event, i.e., few 100 Myr. A full treatment of the 3D magnetic field structure is beyond the scope of this paper and will be discussed in a forthcoming paper (A. Basu et al. 2016 in preparation). \subsubsection{Thermal electron densities} Using Equation~\ref{eqrm}, and the upper limit on $B_{\rm \parallel, reg}$, we constrain $\langle n_{\rm e}\rangle$ to be $\gtrsim0.02$ cm$^{-3}$ in the tidal tail. The column density of the H{\sc i} in the tidal tail is found to be $\sim6\times10^{20}$ cm$^{-2}$ having thickness $\sim 4.5$ kpc \citep{hibba01} which corresponds to $\langle n_{\rm H}\rangle \sim 0.04$ cm$^{-3}$. Thus, from our constraint on the $\langle n_{\rm e} \rangle$, the ionization fraction in the tidal tail is $\gtrsim30$ per cent. As pointed out in the study by \citet{hibba05}, the UV emission from the tidal tail predominantly arises from old stars, and there is little evidence of on-going star formation. This is insufficient to sustain the ionic medium and therefore the ionization of the tidal tail is maintained by the inter-galactic radiation field. \subsection{Turbulent fields in the merging disc} The polarized emission from the merging disc of the galaxies has low fractional polarization with a median value of only $\sim0.016$ at 2.8 GHz and the $\lambda-$dependent depolarization is best described by an external dispersion screen (see \textsection3.3). From the fitted value of $\sigmaRM\approx100$ rad m$^{-2}$ and observed $|\langle \rm RM \rangle|\approx40$ rad m$^{-2}$, we find $B_{\rm turb, \parallel}/B_{\rm ord, \parallel} \approx 16 f^{1/2}$, assuming path length $L_{\rm ion}=2$ kpc and a typical turbulent cell size of $\sim50$ pc observed in discs of galaxies. Hence, $B_{\rm turb, \parallel}/B_{\rm ord, \parallel}$ ranges from 4--5 for $f$ in the range 0.05--0.1. The measured low $\Pi_{\rm int}$ of 0.04 is likely caused by a turbulent magnetic field enhanced due to the merger event. For the estimated $\Pi_{\rm int}$ of 0.04, we find $B_{\rm turb, \perp}/B_{\rm ord, \perp}$ in the range 7--12, i.e., the turbulent field strength significantly dominates over the ordered field both in the plane of the sky and along the line of sight. For isotropic turbulence, $B_{\rm turb}$ is found to be $\sim19.5~\mu$G. Using this, we constrain the strength of the regular fields along the line of sight $B_{\rm ord, \perp}$ to be $\lesssim5~\mu$G, and in the plane of the sky, $B_{\rm ord, \parallel} \approx 1.2f^{-1/2} \lesssim 5~\mu$G for $f\gtrsim0.05$. \subsection{Ordered fields in the northern arm}\label{n-arm} \begin{figure} \begin{centering} \begin{tabular}{c} {\mbox{\includegraphics[height=9cm]{figure9.eps}}}\\ \end{tabular} \end{centering} \caption{Variation of the fractional polarization ($\Pi$) with $\lambda^2$ around the northern spiral arm region. The point with arrow is the 2$\sigma$ upper limit. The lines are the best fit models of different types.} \label{banddepolnorth} \end{figure} The polarized emission from the relic spiral arm of NGC 4038 in the north is weak with median $\Pi\sim0.08$ at 2.8 GHz. The $B_{\rm \perp, reg}$ is observed to follow a spiral pattern (see Figure~\ref{composite}). The polarized emission is depolarized in the optical spiral arm hosting sites of high star formation and the detected polarized emission is offset toward the outer parts. It is difficult to reliably measure the fractional polarization as the total intensity quickly falls off to the background noise level in that region. The mean Faraday depth of this feature is $\sim+9$ rad m$^{-2}$ and has a comparatively large dispersion of $\sim20$ rad m$^{-2}$ in the plane of the sky varying between $-30$ to $+40$ rad m$^{-2}$ (see Figure~\ref{rmmap}). We observe the Faraday depth to frequently change sign smoothly over a few synthesized beam, indicative of a less coherent regular component of the magnetic field in the plane of the sky than that in the tidal tail. In Figure~\ref{banddepolnorth}, we show the variation of $\Pi$ as a function of $\lambda^2$ in the northern arm region. Due to a limited $\lambda^2$ coverage of our data, it is not possible to select a best fit model from our data ($\chi^2=2.2, 1.6$ and $1.8$ for the DFR, IFD and EFD models, respectively). All the depolarization mechanisms indicate a $\Pi_{\rm int}$ between 0.08--0.1 and if we assume isotropic turbulent fields, $B_{\rm turb}$ is estimated to be $\sim19~\mu$G. Since, we do not have a estimate of $B_{\rm turb, \parallel} /B_{\rm ord, \parallel}$ due to lack of unambiguous fit to Faraday depolarization, we estimate an upper limit on the strength of $B_{\rm ord, \perp}$ as $6~\mu$G. \subsection{Extreme Faraday depolarization in southern star-forming regions} The southern part of the Antennae system around the dark cloud region shows strong depolarization, such that the polarized emission remains undetectable in our observations below $\sim3.6$ GHz, at 4.8 GHz \citep{chyzy04} and the fractional polarization at 8.44 GHz is extremely low \citep[$\sim1$ per cent;][]{chyzy04}. Although, such an effect could be caused by $\lambda-$independent beam depolarization due to the random component of the magnetic field, here we assess the possibility of extreme nature of Faraday depolarization. In \textsection\ref{sfreg} we showed, based on the estimated thermal emission, we expect $|\phi|$ $\sim10^4$ rad m$^{-2}$. Such high $\phi$ values are unlikely to cause bandwidth depolarization especially above $\sim3$ GHz within a 8 MHz channel. Assuming that the $\lambda-$dependent depolarization arises due to the same region being Faraday rotating and synchrotron emitting but with a regular magnetic field along the line of sight, i.e., depolarization due to differential Faraday rotation (DFR), the $\Pi(\lambda)$ is given by \begin{equation} \Pi(\lambda) = \Pi_{\rm int} \frac{\sin |\phi \lambda^2|}{|\phi \lambda^2|}. \end{equation} In this case, $\phi$ is related to the RM as $\phi = (1/2)\textrm{RM}$. Using this we find, for the estimated $\phi\sim10^4$ rad m$^{-2}$, the expected fractional polarization at 8.44 GHz with a bandwidth of 50 MHz \citep[same as the observations of][]{chyzy04} lies between $1-7$ per cent with occasional null values depending on the $|\phi|$. The variation of the fractional polarization at 8.44 GHz along this region was observed to be spatially smooth (1--2 per cent) and hence it is unlikely that an ionic medium with only a uniform magnetic field gives rise to such high depolarization. If the magneto-ionic medium is turbulent, driven by high star-formation activity and merger, then the internal and external Faraday dispersion models, give $\sigma_{\rm RM} \sim 10^3$ rad m$^{-2}$. Using Equation~\ref{sigmaRM} we infer the turbulent cells to be 10--50 pc in size for a turbulent field strength of $\sim20~\mu$G and the estimated $\langle n_{\rm e}\rangle \approx 3~\rm cm^{-3}$. The cell size is typical of what is observed in the spiral arms of normal star-forming galaxies \citep{fletc11, mao15}. Hence, it is difficult to disentangle whether the star-forming regions in the southern part are beam depolarized due to a turbulent magnetic field or Faraday depolarized due to extreme properties of the magneto-ionic media. To distinguish between these broadband properties require Stokes $Q$ and $U$ fitting of higher resolution (A- and B-array) and higher frequency (4--8 GHz) data. We have acquired data between 4 and 8 GHz using the VLA in the DnC, CnB and BnA array configurations and the data will be analyzed. \subsection{Implications on the buildup of galactic magnetic fields} {\it Implications on ISM of galaxies}: As a late stage merger, the Antennae is a classic example of a system at the peak of the cosmic star formation history. We, however, note that the early merging galaxies are believed to be different in their interstellar medium ISM properties as compared to the present day mergers. The present day merging systems, such as the Antennae galaxies, are predominantly between well-settled, dynamically cold discs with comparatively lower star formation rates as compared to early mergers which can be between more gas-rich, turbulent, and compact systems \citep{forst09, willi11, stott16}. Based on the bolometric far infrared luminosity, the Antennae pair is classified as a luminous infrared galaxy (LIRG). LIRGs contribute significantly to the comoving star formation density beyond redshift of 1 \citep{magne09}. Therefore, a detail understanding of the Antennae is essential to understand cosmic evolution of ISM in galaxies. Merger induced active star formation and turbulence \citep{veill02} is an essential ingredient in the evolution of hot interstellar gas through stellar feedback \citep{metz04}. It is crucial in the rapid amplification of magnetic fields via turbulent-dynamo mechanisms. Our study demonstrates the magnetic field strength in the Antennae is comparatively higher than that in isolated galaxies and is dominated significantly by turbulent fields within the merging bodies. The turbulent magnetic field---and its coupling with the ISM energy densities, especially the kinetic energy of the turbulent gas---is important in the origin and maintenance of the radio--far infrared correlation at higher redshifts \citep{schle13}. The stronger field strength in merging galaxies ($\sim20~\mu$G) compensate for the inverse-Compton losses due to the cosmic microwave background at higher redshifts and helps in maintaining the radio--far infrared correlation \citep{basu15b}. {\it Comparison to the pan-Magellanic field}: In the immediate neighbourhood of the Milky Way, the Magellanic bridge connecting the Large and Small Magellanic clouds (LMC and SMC, respectively) could potentially be an example of a system hosting regular magnetic fields of similar length-scale as detected in the tidal tail of Antennae galaxies. Through studies of starlight polarization and RMs inferred from background sources, it has been suggested that an aligned ``pan-Magellanic'' magnetic field possibly exists, connecting LMC and SMC \citep{mao08, wayte90}. \citet{deinz73} presented model of the H{\sc i} gas connecting the two clouds that spans $\sim20$ kpc, which supports the existence of pan-Magellanic magnetic field. Although, similar to the tidal tail of the Antennae galaxies, the Magellanic bridge is likely generated from a tidal event \citep{besla12, baghe13}, its progenitors are very different from those of the Antennae galaxies both in terms of morphologies and their initial field strengths. Detailed studies of magnetic field properties in LMC and SMC have revealed low total field strengths of $\sim3~\mu$G \citep{mao08, mao12}. Therefore, low progenitor magnetic field strengths along with significantly less tidal pressure because of shallow gravitational potential of the LMC--SMC system as compared to that of the Antennae galaxies, the strength of the pan-Magellanic ordered field could be low. For example, the model by \citet{deinz73}, predicts the strength of the pan-Magellanic ordered field to be sub-microgauss to a few microgauss. The pan-Magellanic ordered field could span about 20 degrees on the sky. Therefore, direct detection of the regular field from linearly polarized emission is difficult because of low surface brightness and systematic contribution from the Milky Way in the foreground. To firmly establish the existence of such a field, improved RM grid experiment with a large number of background polarized sources is essential as suggested by \citet{mao08}. A systematic comparison of the pan-Magellanic magnetic field with that of the tidal tail of the Antennae galaxies is not possible until we have firm observations about the strength and the structure of this pan-Magellanic magnetic field. {\it Implications on magnetic field measurements at high redshifts}: Usually, one studies magnetic field properties in high redshift intervening objects by measuring the excess RM towards quasar absorption line systems: Mg{\sc ii}, damped Lyman-$\alpha$ (DLA), sub-DLA \citep[see e.g.,][]{oren95, berne08, joshi13, farne14}. It has been suggested that the sub-DLAs can originate from neutral gas that lies $\gtrsim20$ kpc from the host galaxies and the absorbing gas is likely stripped via tidal interaction and/or ram pressure \citep{semba01, muzah16}. The tidal tail of the Antennae with $N_{\rm HI}$ in the range $\sim10^{20}-6\times10^{20}$ cm$^{-2}$, would be a classic example of a DLA or sub-DLA (depending on at what distance the from the host galaxies the line-of-sight intersects), when observed as a Lyman-$\alpha$ absorber against a quasar at high redshifts. Our study shows that such systems can host large scale regular magnetic fields and gives rise to rotation measure when observed at suitable viewing angles. It is therefore important to take into account that the inferred field could come from tidally stripped gas, and not only from the disc. {\it Implications to magnetization of the intergalactic medium}: Coherent magnetic fields in the outskirts of merging galaxies, extending in to the intergalactic medium, assists in propagation of cosmic ray particles along the field lines \citep[see e.g.,][]{cesar80, ptusk06}. Thus, apart from starburst driven galactic winds, galaxy mergers can also play an important role in magnetizing the intergalactic medium and enriching it with cosmic ray particles. {\it Implication to evolution of large-scale fields in galaxies}: In a study of cosmological evolution of large- and small-scale magnetic fields in galaxies, \citet{arsha09} predicted that major merger events dissipate the $10^{-6}$ G large-scale disc fields down to several $10^{-8}$ G. Our study comprehensively shows that, although the pre-merger regular magnetic fields in the galactic discs are mostly disrupted by the merger and are dominated by turbulent fields, they can assist in producing larger-scale coherent fields several microgauss in strength through field stretching. The detection of a $\sim20$ kpc coherent magnetic field in the tidal tail indicates that large-scale fields can be preserved even in advanced merging systems. Moreover, large-scale fields do not necessarily require Gyr timescales to develop as predicted by magneto-hydrodynamic (MHD) simulations \citep{arsha09, hanas09}. A detailed MHD simulation is essential to understand the nature of the coherent magnetic fields in the tidal tails and its implications in developing large-scale regular fields observed in isolated galaxies in the local Universe. | 16 | 9 | 1609.04266 |
1609 | 1609.06922_arXiv.txt | We explore the parameter space of the semi-analytic galaxy formation model \glf, studying the constraints imposed by measurements of the galaxy stellar mass function (GSMF) and its evolution. We use the Bayesian Emulator method to quickly eliminate vast implausible volumes of the parameter space and zoom in on the most interesting regions, allowing us to identify a set of models that match the observational data within model uncertainties. We find that the GSMF strongly constrains parameters related to quiescent star formation in discs, stellar and AGN feedback and threshold for disc instabilities, but weakly restricts other parameters. Constraining the model using local data alone does not usually select models that match the evolution of the GSMF well. Nevertheless, we show that a small subset of models provides acceptable match to GSMF data out to redshift 1.5. We explore the physical significance of the parameters of these models, in particular exploring whether the model provides a better description if the mass loading of the galactic winds generated by starbursts ($\beta_{0,\text{burst}}$) and quiescent disks ($\beta_{0,\text{disc}}$) is different. Performing a principal component analysis of the plausible volume of the parameter space, we write a set of relations between parameters obeyed by plausible models with respect to GSMF evolution. We find that while $\beta_{0,\text{disc}}$ is strongly constrained by GSMF evolution data, constraints on $\beta_{0,\text{burst}}$ are weak. \green{Although it is possible to find plausible models for which $\beta_{0,\text{burst}} = \beta_{0,\text{disc}}$, most plausible models have $\beta_{0,\text{burst}}>\beta_{0,\text{disc}}$, implying -- for these -- larger SN feedback efficiency at higher redshifts.} | Semi-analytic models of galaxy formation (SAMs) are well established tools for exploring galaxy formation scenarios in their cosmological context. The problem of how galaxies form and evolve is described by a set of coupled differential equations dealing with well-defined astrophysical processes. These are driven by dark matter halo merger trees that determine the source terms in the equation network \citep[for reviews see e.g.][]{BensonReview,Baugh2006,SomervilleDave2015}. Due to the approximate nature of the methods used in these simulations, and the uncertainties in the physical process that are modelled, these models include a large number of uncertain parameters. While order of magnitude estimates for these parameters can be made, their precise values must be determined by comparison to observational data. Traditionally, parameter values have been set through a trial-and-error approach, where the galaxy formation modeller varies an individual parameter developing intuition about its effects on the model predictions for a particular observable and then uses this understanding to select a parameter set that gives a good description of the observations. Despite its simplicity, and obvious limitations, this procedure has led to substantial progress in the field. Recently, however, several papers have employed more rigorous statistical methods to explore the high dimensional parameter space systematically \citep{Kampakoglou2008,Henriques2009,Bower2010,Henriques2013,Benson2014,Lu2014, Henriques2015}. Such approaches provide a richer analysis, and seek to identify the regions of parameter space that are in agreement with observational data, and not just to find optimal parameter values. This therefore informs as to the uniqueness of the parameter choices, and provides understanding of the degeneracies between different parameters. In this work we study which constraints are imposed on the semi-analytic model \glf by the observations of galaxy stellar mass function (GSMF). We first consider the constraints imposed by local observations and then investigate how the parameters are further constrained by the introduction of high redshift data. This makes powerful use of the iterative emulator technique described by \citet{Bower2010}, which provides an efficient way of probing a high dimensional parameter space. Importantly, the method allows additional constraints to be added in post-processing. Thus, we start by finding the region in the parameter space which contains models that produce a good match to the local Universe GSMF. This region is, then, further probed to check whether a match to higher redshift data is possible. By analysing 2D projections of the plausible models sub-volume and performing a principal component analysis of it, we are able to study the degeneracies and interactions between the most constrained parameters. We note that typical approaches to analysing comparable models using Bayesian MCMC require millions of model runs (at least), while the approach used here, which utilises Bayesian emulation, only required tens of thousands of model runs, representing a substantial improvement in efficiency. Closely reproducing the observed high-redshift galaxy mass function \citep{Cirasuolo2010, Henriques2013} is problematic for many galaxy formation models. \citet{Henriques2013}, for example, concludes that the effectiveness of galaxy feedback (specifically the re-incorporation time of expelled gas) must depend on \green{the virial mass of the dark matter halo on the basis of a Monte Carlo exploration of the parameter space of their model. This is not fully satisfactory, however, since one would expect the re-incorporation time to be physically related to the halo dynamical time and not the halo mass. } In this paper, we explore an appealing and well-motivated alternative. Observations of galaxy winds \citep[eg.,][]{heckman1990, martin2012} suggest that the effective mass loading is strongly dependent on the surface density of star formation. It appears that efficient outflows are more readily generated when star formation occurs in dense bursts than when the star formation occurs in a smooth and quiescent disk. These observations motivate a more careful exploration of the treatment of galaxy winds from starburst and quiescent disks, and in this paper we parametrize the mass loading of the wind independently in these two cases. This may naturally resolve the difficulty presented by observations of the high redshift GSMF since the cosmic star formation rate density may be more dominated by starbursts at high redshift, while it is dominated by quiescent star formation at low redshift \citep{Malbon2007}. This paper is organized as follows. In \S\ref{sec:model} we describe the galaxy formation model, specifying which parameters were varied and briefly reviewing the physical meaning of the most relevant of them. In \S\ref{sec:emulator} the iterative history matching methodology is reviewed. In \S\ref{sec:z0} we present our results for the matching to the local GSMF. In \S\ref{sec:high_z} we examine the effects of including higher redshift data. In \S\ref{sec:subspace}, 2D projections of the parameter space are analysed. In \S\ref{sec:pca}, the results of a principal component analysis of the non-implausible volume are shown. Finally, in \S\ref{sec:summary}, we summarize our conclusions. | 16 | 9 | 1609.06922 |
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1609 | 1609.01505_arXiv.txt | Parametric instability is an intrinsic risk in high power laser interferometer gravitational wave detectors, in which the optical cavity modes interact with the acoustic modes of the mirrors leading to exponential growth of the acoustic vibration. In this paper, we investigate the potential parametric instability for a proposed next generation gravitational wave detector based on cooled silicon test masses. It is shown that there would be about 2 unstable modes per test mass, with the highest parametric gain of $\sim 76$. The importance of developing suitable instability suppression schemes is emphasized. | Gravitational waves, predicted by Einstein a century ago, have finally been directly detected \cite{abbott2016observation} in 2015, enabling us to have a first glimpse of two black holes merging. The detections \cite{abbott2016observation} \cite{abbott2016gw151226} gravitational waves open the new windows on the universe and there are now strong reasons for improving detector sensitivity. At their design sensitivity, the kilometer scale interferometer detectors such as Advanced LIGO detectors \cite{aasi2015advanced} and Virgo detector \cite{accadia2012virgo} are expected to be able to detect $\sim10^3$ gravitational waves events per year \cite{belczynski2016first}. This data will be a critical resource for understanding the origin and distribution of stellar mass black holes in the universe. To reach high sensitivity, advanced laser interferometer detectors must use very high laser power inside the optical cavities to reduce quantum noise. Braginsky pointed out \cite{braginsky2001parametric} in the early 2000s that with extremely high laser power in the interferometer cavity, radiation pressure coupling between the acoustic modes and cavity optical modes can cause 3-mode parametric instability (PI), resulting in the exponential growth of many acoustic mode amplitudes and disruption of the operation of the interferometer. Zhao et al.\cite{zhao2005parametric} undertook a detailed 3D study of PI in a detector similar to the advanced LIGO design. The modeling considered interactions in individual arm cavities, and made predictions \cite{zhao2005parametric} that were in relatively close agreement to recent observations of PI on the Advanced LIGO detectors \cite{PhysRevLett.114.161102}. Zhao$'$s paper discussed the use of small variations of the mirror radius of curvature to tune the system to a local minimum of parametric gain. In fused silica the thermal distortion of test masses is used for this purpose. Subsequent detailed modeling \cite{gras2010parametric} showed that interactions between arm cavities, power recycling and signal recycling cavities could significantly modify the predictions, and further work identified many methods of suppressing PI, such as electrostatic feedback, optical feedback and the use of passive dampers \cite{degallaix2007thermal,ju2008strategies,gras2008test,gras2009suppression,fan2010testing,miller2011damping,zhao2015parametric,Gras2015damper}. Unfortunately the conditions that enable high sensitivity appear to always coincide with the conditions that enable PI. To date, no method has been demonstrated that can completely eliminate parametric instability. As advanced LIGO power is increased, it is anticipated that a combination of of thermal tuning, electrostatic feedback and passive dampers will be able to eliminate all instabilities. To increase the sensitivity of detectors beyond that achievable with the current advance LIGO detectors, various options are being investigated. Much effort is underway to design the next generation of detectors. Proposed designs include the Einstein Telescope (ET) \cite{punturo2010einstein} , several LIGO upgrades designs such as LIGO A+, LIGO Voyager \cite{instruments}, a 40km interferometer \cite{dwyer2015gravitational} and an 8km interferometer\cite{blair2015next}. All the future designs required the use of high laser power to overcome quantum shot noise for high frequency sensitivity. It is apparent that each of these designs will need to consider the PI problem. In 2012, Strigin studied the PI problem of the 10km Fabry-Parrot cavity in ET design with sapphire or silicon test masses \cite{strigin2012effect}. His result suggested that there would be $\sim 10$ unstable modes in such systems. Parametric instability arises through the coincidental matching of test mass acoustic mode frequencies and mode shapes with optical cavity transverse modes for which the frequency difference from the main TEM$_{00}$ pump laser mode is in the 1 - 100kHz range. Detailed predictions must take into account the details of the test mass shapes and optical cavity design. However in general, the risk of modal coincidences depends on the optical modes density and the acoustic mode density. The optical mode density depends on the cavity free spectral range, determined by the cavity length, and the transverse mode spacing called the mode gap, which is proportional to the inverse cosine of the cavity g-factor. The acoustic mode density depends on test mass sizes and inversely on the sound velocity. The particular design analyzed here replaces the 40kg fused silica test masses of Advanced LIGO with $\sim$ 200kg cooled silicon test masses. The high sound velocity and low acoustic loss of silicon could enable significant suppression of thermal noise which can be further improved through use of crystalline optical coatings\cite{cole2013tenfold}. The high sound velocity of silicon acts to reduce modal coincidences that lead to PI, but the increased mass acts in the other direction. This paper investigates this issue in detail for the specific design called the LIGO voyager blue design. While the results are specific to this design, they give an indication of the problems that could be faced by many future $3^{rd}$ generation gravitational wave detectors. For this analysis we use a single cavity model for PI modeling because of its general agreement with observations. Any more complex models depend on details of the power and signal recycling cavities, and do not change the general pattern of instability. Hence this model provides a good indication of the level of the PI. | The result presented here is the first step in analyzing the parametric instability in the laser interferometer gravitational wave detector with silicon test masses. The single cavity model was used for simplicity. In practical configurations, the power and signal recycling cavities will change the parametric gain depending on the resonance condition of the high order mode involved inside the recycling cavities. As reference\cite{gras2010parametric} pointed out that any unstable modes present in the single cavity will be potential threats in a full interferometer. The first detailed analysis of parametric instability of Advanced LIGO\cite{zhao2005parametric} using a single cavity model predicted that there would be 7 unstable acoustic modes per each fused silica test mass, with parametric gain $R$ up to 7, at circulating power of 800 kW. This means about a total of 28 unstable modes in 4 test masses of the two cavity arms. So far Advanced LIGO observed total $\sim$10 unstable modes, with the maximum parametric gain of the order of 10, operating at $\sim 25\%$ of designed power. This is consistent with the the single cavity simulation. A more complicated simulation \cite{gras2010parametric} that took into consideration the recycling cavities and losses gave deeper insight of the parametric instability of the real system. However, The simpler single cavity simulation is a very good indicator about the potential parametric instability problem of future design. Assuming a Q-factor of $10^7$ for the acoustic modes of the test mass, we estimated that there would be $\sim 2$ unstable modes with maximum parametric gain of 76 for the Voyager blue design. This means that the potential parametric instability risk is roughly at the same level as the advanced LIGO detectors, at a much higher optical power level. The choice of a high sound velocity test mass material is advantageous in regards to parametric instabilities. As designs are further developed, it will be necessary to consider all the effects mentioned above as well as details of test mass structures such as flats on the barrel for suspensions, and the diffraction losses of high order modes. | 16 | 9 | 1609.01505 |
1609 | 1609.07609_arXiv.txt | We report on C-band (5 - 7 GHz) observations of the galaxy, NGC~2992, from the CHANG-ES sample. This galaxy displays an embedded nuclear double-lobed radio morphology within its spiral disk, as revealed in linearly polarized emission but {\it not} in total intensity emission. The radio lobes are kpc-sized, similar to what has been observed in the past for other Seyfert galaxies, and show ordered magnetic fields. NGC~2992 has shown previous evidence for AGN-related activity, but not the linearly polarized radio features that we present here. We draw attention to this galaxy as the first clear example (and prototype) of bipolar radio outflow that is revealed in linearly polarized emission only. Such polarization observations, which are unobscured by dust, provide a new tool for uncovering hidden weak AGN activity which may otherwise be masked by brighter unpolarized emission within which it is embedded. The radio lobes observed in NGC~2992 are interacting with the surrounding interstellar medium and offer new opportunities to investigate the interactions between nuclear outflows and the ISM in nearby galaxies. We also compare the radio emission with a new CHANDRA X-ray image of this galaxy. A new CHANG-ES image of NGC~3079 is also briefly shown as another example as to how much more obvious radio lobes appear in linear polarization as opposed to total intensity. | \label{sec:introduction} The powerful jet and double-lobed character of radio emission around distant massive elliptical galaxies is well known and originates in the active galactic nuclei (AGNs) of their host galaxies. Indeed, since it is believed that the formation of a supermassive black hole (SMBH) is closely connected to the formation of the bulge mass, with both evolving together \citep{fer00,geb00,zhe09,kor13}, one would expect AGN-related activity to manifest itself in spirals as well. Yet, spiral galaxies do not generally show such impressive evidence for nuclear activity and observations of AGNs in spirals are relatively rare. The incidence of radio AGN in late type galaxies, for instance, is about an order of magnitude lower than in early-type galaxies \citep{kav15b}. Those spiral galaxies that do display nuclear outflows in the form of AGN-related jet and/or lobe-like features make for a short list. For example, Seyfert galaxies can harbour jets on pc to kpc scales \citep[][and references therein]{ho01}. Well-known examples also include NGC~3079 \citep{irw88,hum83}, the Circinus spiral galaxy \citep{har90,gre03}, NGC~7479 \citep{lai08}, NGC~4388 \citep{dam16}, NGC~4258 \citep{gre95,miy95,her98,kra04}, as well as the spectacular examples, 0313-192 \citep{led01,kee06}, `Speca' \citep{hot11}, J2345−0449 \citep{bag14}, and J1649+2635 \citep{mao15}. The reason for the apparent relative scarcity of such activity in spirals is not entirely clear, though a variety of possibilities have been raised. For example, advection-dominated accretion flow may result in very little radiation \citep{nar95,yua14}, the SMBHs in spirals may not be massive enough to produce the large-scale jets seen in ellipticals \citep{lao00}, accreting SMBHs in spirals may not be spinning rapidly enough \citep{wil95,sik07,tch10,dot13} or the magnetic flux could be insufficient to launch jets \citep{sik13}. Other issues such as the effect of the dense surrounding interstellar medium (ISM) on the outflow \citep{das15}, the difficulty in accreting high angular momentum gas, the possibility of episodic outflows \citep{mar03}, and environmental effects such as perturbations from nearby galaxies, all have yet to be fully understood and quantified. An additional complicating issue is that nuclear outflows can also originate from a central starburst \citep[][and references therein]{vei05} and both may be occurring in the same galaxy \citep[e.g. NGC~3079][and references therein]{sha15}. It has thus been difficult to identify the origin of an observed outflow as being due to star formation (SF), or an AGN, or some combination thereof. An active low-luminosity AGN (LLAGN), for example, could be embedded in a nuclear region with abundant SF activity. Since all spirals host star formation, the issue is really whether or not an AGN is also present and how best to identify it. Hard X-ray emission is an indicator of AGN activity, provided the luminosity exceeds what is expected from ultra-luminous X-ray sources ($10^{39-42}$ ergs/s) \citep{mus04}. Optical emission line ratios that indicate the presence of a hard ionizing continuum is another \citep{bal81} although dust obscuration can be problematic, especially for edge-on galaxies. The mid-IR PAH\footnote{Policyclic Aromatic Hydrocarbon} to continuum ratio has been used as a SF/AGN diagnostic for distant galaxies \citep{kir15} and other mid-IR line ratios \citep[see][]{per10} have also been shown to be effective. Evidence for an AGN, however, is not necessarily evidence for AGN-related outflows. As indicated above, it is in the radio continuum that jets and lobes have historically been detected, since synchrotron emission is a strong component of such outflow. AGN-related outflows have been identified by the presence of linear structures in Seyfert galaxies \citep{ulv84} or by small-scale jets in late-type spirals \citep{kav15a}, but only in in {\it total intensity} radio continuum \citep{ulv84}. In order to be so identified, such emission must dominate any SF-related emission in the vicinity. In this paper, we provide a new diagnostic for identifying AGN-related outflow via radio polarization, specifically {\it bipolar radio features or other evidence for AGN-related outflows that are masked in total intensity but are revealed in linear polarization.} In other words, LLAGNs may not, in fact, be rare, but require that our observations be targeted at the polarized rather than the total intensity radio emission to be clearly seen. Polarization images are well-suited to uncovering previously hidden AGNs because of their sensitivity to ordered magnetic fields that are prevalent in the outflows and also because the radio emission is insensitive to dust obscuration. As part of the CHANG-ES program \citep[`Continuum Halos in Nearby Galaxies -- an EVLA Survey',][]{irw12a}, we focus on one such galaxy at C-band (6 GHz), i.e. NGC~2992, which we take to be a prototype for such activity. This galaxy was previously known to harbour an active nucleus (Sect.~\ref{sec:n2992}) but now the polarized radio emission clearly shows the outflows. Note that NGC~2992 is {\it not} the only AGN in the CHANG-ES sample. A more complete census of AGNs in the CHANG-ES sample will be presented in a future paper. For a list of radio polarization observations and magnetic fields that have been observed in other mostly non-edge-on nearby galaxies, see the appendices of \citet{bec13}. | \label{sec:conclusions} The polarized emission from NGC~2992, uniquely provided by the CHANG-ES survey, has offered us a new glimpse into the nuclear activity in this galaxy. Although radio lobes are {\it not} observed in total intensity radio emission, they {\it are} observed in polarized emission. Thus, we are observing weak radio lobes which were not strong enough to be distinguished from other non-polarized emission seen in total intensity. For NGC~2992, at high resolution, a single discrete radio lobe is observed on the western side of the galaxy's center and at low resolution, two radio lobes are seen on either side of the nucleus, extending to the east and west. The eastern radio lobe has emerged from the galaxy whereas the western lobe is within the disk. We argue that the radio lobes are consistent with AGN-like outflow. The orientation of the discrete western radio lobe is consistent with outflow that has been redirected, via pressure-gradients in the ISM, along the minor axis. The outflow and ISM pressures are consistent with this picture, to within measurable factors of a few (assuming that the minimum energy criterion holds). We suggest a possible (though not unique) evolutionary scenario for the observed radio features such that the variety of radio activity seen in this galaxy can be simply explained by episodic outflow from the AGN, possibly embedded within more steady lower level flows. The AGN is at the radio core position which should designate the center of the galaxy, rather than the normally quoted optical center. Such polarized radio lobes provide a unique opportunity to study the physics of outflow/ISM interactions in spiral galaxies with LLAGNs. | 16 | 9 | 1609.07609 |
1609 | 1609.00818_arXiv.txt | { Catalogue cross-correlation is essential to building large sets of multi-wavelength data, whether it be to study the properties of populations of astrophysical objects or to build reference catalogues (or timeseries) from survey observations. Nevertheless, resorting to automated processes with limited sets of information available on large numbers of sources detected at different epochs with various filters and instruments inevitably leads to spurious associations. We need both statistical criteria to select detections to be merged as unique sources, and statistical indicators helping in achieving compromises between completeness and reliability of selected associations. } { We lay the foundations of a statistical framework for multi-catalogue cross-correlation and cross-identification based on explicit simplified catalogue models. A proper identification process should rely on both astrometric and photometric data. Under some conditions, the astrometric part and the photometric part can be processed separately and merged a posteriori to provide a single global probability of identification. The present paper addresses almost exclusively the astrometrical part and specifies the proper probabilities to be merged with photometric likelihoods. } { To select matching candidates in $n$ catalogues, we used the Chi (or, indifferently, the Chi-square) test with $2(n-1)$ degrees of freedom. We thus call this cross-match a $\rchi$-match. In order to use Bayes' formula, we considered exhaustive sets of hypotheses based on combinatorial analysis. The volume of the $\rchi$-test domain of acceptance -- a $2(n-1)$-dimensional acceptance ellipsoid -- is used to estimate the expected numbers of spurious associations. We derived priors for those numbers using a frequentist approach relying on simple geometrical considerations. Likelihoods are based on standard Rayleigh, $\rchi$ and Poisson distributions that we normalized over the $\rchi$-test acceptance domain. We validated our theoretical results by generating and cross-matching synthetic catalogues. } { The results we obtain do not depend on the order used to cross-correlate the catalogues. We applied the formalism described in the present paper to build the multi-wavelength catalogues used for the science cases of the ARCHES (Astronomical Resource Cross-matching for High Energy Studies) project. Our cross-matching engine is publicly available through a multi-purpose web interface. In a longer term, we plan to integrate this tool into the CDS XMatch Service. } {} | The development of new detectors with high throughput over large areas has revolutionized observational astronomy during recent decades. These technological advances, aided by a considerable increase of computing power, have opened the way to outstanding ground-based and space-borne all-sky or very large area imaging projects (e.g. the 2MASS \citep{Skrutskie2006,Vizier2MASS}, SDSS \citep{SDSS9,VizierSDSS9} and WISE \citep{Wright2010,CatAllWISE} surveys). These surveys have provided an essential astrometric and photometric reference frame and the first true digital maps of the entire sky. \newline As an illustration of this flood of data, the number of catalogue entries in the VizieR service at the Centre de Donn\'ees astronomiques de Strasbourg (CDS) which was about 500 million in 1999 has reached almost 18 billion as on February 2016. At the 2020 horizon, European space missions such as GAIA and EUCLID together with the Large Synoptic Survey Telescope (LSST) will provide a several-fold increase in the number of catalogued optical objects while providing measurements of exquisite astrometric and photometric quality. \newline This exponentially increasing flow of high quality multi-wavelength data has radically altered the way astronomers now design observing strategies and tackle scientific issues. The former paradigm, mostly focusing on a single wavelength range, has in many cases evolved towards a systematic fully multi-wavelength study. In fact, modelling the spectral energy distributions over the widest range of frequencies, spanning from radio to the highest energy gamma-rays has been instrumental in understanding the physics of stars and galaxies. \newline Many well designed and useful tools have been developed worldwide concurrently with the emergence of the virtual observatory. Most if not all of these tools can handle and process multi-band images and catalogues. When assembling spectral energy distributions using surveys obtained at very different wavelengths and with discrepant spatial resolution, one of the most acute problems is to find the correct counterpart across the various bands. Several tools such as TOPCAT \citep{Taylor2005} or the CDS XMatch Service \citep{Pineau2011b, Boch2012} offer basic cross-matching facilities. However, none of the publicly available tools handles the statistics inherent to the cross-matching process in a fully coherent manner. A standard method for a dependable and robust association of a physical source to instances of it in different catalogues (cross-identification) and in diverse spectral ranges is still absent. \newline The pressing need for a multi-catalogue probabilistic cross-matching tool was one of the strong motivations of the FP7-Space European program ARCHES \citep{Motch2016}\footnote{http://www.arches-fp7.eu/}. Designing a cross-matching tool able to process, in a single pass, a theoretically unlimited number of catalogues, while computing probabilities of associations for all catalogue configurations, using the background of sources, positional errors and eventually introducing priors on the expected shape of the spectral energy distribution is one of the most important outcomes of the project. A preliminary description of this algorithm was presented in \cite{Pineau2015}. Although ARCHES was originally focusing on the cross-matching of XMM-Newton sources, the algorithms developed in this context are clearly applicable to any combination of catalogues and energy bands \citep[see for example][]{Mingo2016}. | } In this paper we developed a comprehensive framework for performing the cross-correlation of multiple astronomical catalogues, in one pass. The approach employs a classical $\rchi^2$-test to select candidates. We computed two sets of likelihoods based on positions, individual elliptical positional errors and the $\rchi^2$-test region of acceptance: one that can be mixed without any caution with other parameters such as photometric values; and one for which the naive hypothesis of independence between positional uncertainties and magnitudes has to be tested. We also presented a way to estimate ``priors'' from the region of acceptance of the $\rchi^2$-test. Probabilities for each possible hypothesis can thus be computed from those likelihoods and ``priors''.\\ In practice the number of hypotheses, and thus the number of ``priors'', increases dramatically with the number of catalogues. To be able to cross-match more than six or seven catalogues, it is necessary to simplify the problem. One possibility consists of merging two catalogues of similar astrometric accuracy and similar wavelength range, considering all matches as non-spurious matches. Doing so we would effectively reduce the number of input catalogues by one. A large part of the statistical work carried out here depends on the simplifying assumptions made in \S \ref{sec:assumptions}: perfect astrometrical calibration (no systematic offsets), no proper motions, no clustering and no blending. In real life, the ``normalized'' distance between two detections of a same source present in two distinct catalogues hardly follows a Rayleigh distribution. The ``actual distribution'' ({in practice it is not easy to build such a distribution since it requires secure identifications) often has a broader tail \citep[see for example][Fig. 5]{Rosen2016} and a log-normal distribution may better fit it than the Rayleigh distribution. This is probably due to a combination of causes like small proper motions, imperfect reduction, systematics or bias from the calibration process, under or overestimated errors, etc.\\ In practice this means that the number of associations missed by the candidate selection criteria (based on Rayleigh) is larger that the chosen theoretical value ($\gamma$). We could for example add larger systematics to positional errors. The risk is then to distort (even more) the Rayleigh distribution. We could also try to re-calibrate locally the set of catalogues we want to cross-match, but we need secure identifications to do it properly; for each catalogue, all sources in the local area must have been calibrated at once (to possibly correct for a locally uniform systematic using four simple parameters $\Delta\alpha$, $\Delta\delta$, $scale$, $\theta$). Those two constraints (having secure identifications and at once calibration) are in practice quite hard to satisfy.\\ If we re-calibrate using a ``secure'' population (i.e. a population of objects having no proper motions like QSOs) we introduce a bias since QSOs are fainter than most stars in the optical and thus have errors larger than the global population of objects. And adding stars, we introduce noise due to proper motions. For these reasons, we believe that in case of ``old'' optical surveys based on photographic plates, a classical fixed radius cross-match may be more efficient that the $\rchi$-match to select candidates. We are nonetheless conviced that the equations we derived in this paper can help in building new catalogues, based for example on both multi-band and multi-epoch observations, and can be used to assess and improve the quality of coming surveys. We generated and processed synthetical catalogues, which meet the simplifying assumptions, in the tool we developed for the ARCHES project. The consistency between the theoretical results derived in this paper -- completeness of the candidate selection criterion, likelihoods and priors -- and the outputs of the tool has allowed us to cross-validate both the method and its implementation. The tool has also been used to generate ARCHES products which were used in the scientific work packages of the project. Currently the CDS XMatch Service \citep{Pineau2011b, Boch2012, Pineau2015} provides a basic but very efficient facility to cross-correlate two possibly large (> 1 billion sources) catalogues. It is planned to include the ARCHES tool into the CDS XMatch. This paper will be the basic reference for the extension of the latter to multi-catalogue statistical $\rchi$-match. | 16 | 9 | 1609.00818 |
1609 | 1609.09108_arXiv.txt | We investigate rapid TeV flaring in markarian 501 in the frame of a time-dependent one-zone synchrotron self-Compton (SSC) model. In this model, electrons are accelerated to extra-relativistic energy through the stochastic particle acceleration and evolve with the time, nonthermal photons are produced by both synchrotron and inverse Comtpon scattering off synchrotron photons. Moreover, nonthermal photons during a pre-flare are produced by the relativistic electrons in the steady state and those during a flare are produced by the electrons whose injection rate is changed at some time interval. We apply the model to the rapid flare of Markarian 501 on July 9, 2005 and obtain the multi-wavelength spectra during the pre-flare and during the flare. Our results show that the time-dependent properties of flares can be reproduced by varying the injection rate of electrons and a clear canonical anti-clockwise-loop can be given. | \label{sec:intro} Variability, which is found from radio to TeV $\gamma$-ray bands, is one of the major characteristics of blazars. The variability timescales from a few minutes to days in the optical band have been extensively investigated (e.g. Sillanp$\ddot{a}\ddot{a}$ et al. 1991; Wagner $\&$ Witzel 1996; Lainela et al. 1999). Particularly in the X-ray and TeV regimes in which photons are produced by radiation of ultra-relativistic electrons close to their maximum energy, the observations of variability timescales constrain on particle acceleration mechanism in TeV blazars. For examples, Kataoka et al. (2001) reported the observations of the X-ray flares with timescales of from hours to days for three TeV blazars (Markarian 421, Markarian 501, and PKS 2155-304), Albert et al. (2007) obtained a rapid TeV variability with of several ten minutes for Markarian 501 by MAGIC. Observed short timescales indicate that the variability is associated with small regions in the relativistic jet, which is located on a distance in excess of one hundred Schwarzschild radii ($r_{\rm s}$) with a central black hole mass $M=10^{9}M_{\odot}$, rather than the center region (Begelman et al. 2008). Relativistic particles maybe responsible for the emission flare. These particles are ejected from the central region alone with the subsistent jet structure, and radiate away their energy at $100r_{\rm s}$ quickly, or the particles are accelerated within the jet, close to the emission region. Generally, soft lags can be interpreted as due to electron cooling (Kirk et al. 1998; Kirk $\&$ Mastichiadis 1999; Kusunose et al. 2000). However, with the fast TeV $\gamma$-ray flare in Markain 501 on July 9, 2005, the evidence that hard $\gamma$-ray lagged the soft ones by $4\pm1$ minutes was discovered (Albert et al. 2007). For explaining these abnormal phenomenon, Bednarek $\&$ Wagner (2008) proposed the radiating blob accelerating during the flare, but in their model the particles would only undergo cooling processes without any accelerating around the high blob Lorentz factor plasma flow. Mastichiadis $\&$ Moraitis (2008) showed that allowing the particles to accelerate gradually can explain the observed feature, and reach the particles energy to $\gamma\sim10^{6}$, the acceleration timescales is of the order of hours. Following their model, Tammi $\&$ Duffy (2009) compared four different acceleration mechanisms, and pointed out that the timescale may be too long for first-order Fermi acceleration, so the stochastic acceleration may be as a promising candidate for the energy dependent time delays. Motivated by above arguments, we study the time-dependent one-zone synchrotron self-Compton (SSC) model in the presence of stochastic particle acceleration, and then apply the model to Markarian 501 for explaining its flare and time delay properties, especially the rapid flare of Markarian 501 on July 9, 2005. Throughout the paper, we assume the Hubble constant $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, the matter energy density $\Omega_{\rm M}=0.27$, the radiation energy density$\Omega_{\rm r}=0$, and the dimensionless cosmological constant $\Omega_{\Lambda}=0.73$. | \label{sec:discussion} In this paper, we have tried to explain the TeV $\gamma$-ray flare of markarian 501 observed by MAGIC telescope on July 9, 2005 in the context of the time-dependent one-zone SSC model which includes stochastic particle acceleration. In this model, particles with low energy are assumed to be injected and then are accelerated to higher energy by second-order Fermi acceleration mechanism (Fermi 1949), the most important photon targets for inverse Compton scattering by relativistic electrons are the synchrotron photons. We have studied time-dependent properties of flares by reproducing the pre-burst spectrum of the source and varying the injection rate. In our results, the behaviour of the mean multi-frequencies spectra before and during the flare is a little different, the peaks of both synchrotron and IC emissions move to lower frequencies, we argue that this can be explained by the energy loss of the electrons during the outburst. In this scenario, the hard lag flare can be obtained and during the flare, it shows a clear canonical anti-clockwise-loop. It should be noted that hard lags require some sort of particle acceleration. If the variability timescale is faster than the cooling timescale, the radiation from accelerated particles would show a hard lag (Albert et al. 2008). Kirk et al. (1998) argued that the hard lag from the acceleration process induces to the anti-clockwise-loop pattern. In this view, Albert et al. (2007) concluded that the acceleration process of low energy particles probably dominate over the TeV $\gamma$-ray flare. Assuming a low energy electrons injection and stochastic acceleration, our calculations predicted hard lags dependent flaring activity and showed a anti-clock-loop evolution of the hardness ratio with the flux. These are in agreement with the observational results on July 9, 2005, and imply that, during the flare, the dynamics of the system is dominated by the acceleration, rather than by the cooling processes. However, a detailed investigation of electron acceleration in the presence of losses has so far been performed only by a few investigators (e.g. Mastichiadis \& Moraits 2008). Given the complexity of the flaring activity of high energy radiation, this requires more detailed observations and the issue to be open. The magnetohydrodynamic turbulence will is generated if standing shocks form in the neighborhood of the central object, which amplify any incoming upstream turbulence in the downstream accretion shock magnetosheath (Campeanu \& Schlickeiser 1992). These magnetohydrodynamic plasma waves are the free energy and lead to stochastic acceleration of charged particles. Actually, stochastic acceleration occurs wherever there are turbulent magnetic fields and can spread to an extended region, the size is determined by the turbulence generation and decay rates. Virtanen \& Vainio (2005) simulated the stochastic acceleration in relativistic shocks and shows, when the particles were accelerated behind the discontinuity, a gradual shift of the whole particle spectrum to higher energy. Some recent observations of particles spectra with hard power law spectral indices, $N(\gamma)\propto\gamma^{-n}$ with $n <2$, suggests that the stochastic acceleration is seen in the observations (Katarzynski et al. 2006; B$\ddot{o}$ttcher et al. 2008). The model presented here contains the stochastic acceleration process. For simplicity, we introduced a constant acceleration term, which is associated with the momentum diffusion coefficient $D(p,t)$. The form of the diffusion coefficient due to interactions with magnetohydrodynamic waves has been discussed in details (e.g. Kulsrud \& Ferrari 1971; Schlickeiser 2002). In our model, both constant acceleration and escape times are assumed, leading to $D(\gamma,t)=\gamma^{2}/2t_{\rm acc}\propto\gamma^{2}$. The form of the diffusion coefficient corresponds to the hard-sphere approximation, in which the mean free path for article-wave interaction is independent of particle energy, and probably induces to a complicated spectrum. Furthermore, since the basic shock acceleration models postulate that $t_{\rm acc}\simeq t_{\rm esc}$ (e.g., Katarzynski et al. 2006), we adopt the shorter acceleration and escape timescales ($t_{\rm acc}=t_{\rm esc}=\frac{R}{c}=t_{\rm cr}$) than other investigators (generally, $t_{\rm acc}>t_{\rm cr}$, and $t_{\rm esc}>t_{\rm cr}$, see e.g., Kirk et al. 1998; Mastichiadis \& Moraitis 2008). These assumptions can lead to higher acceleration rate and lower escape rate, and make more particles acceleration up to high energy rapidly. There are two scenarios for explaining the intrinsic variability. The first scenario assumes that the observed variations origin from the geometry of emitting sources (e.g., Camenzind \& Krockenberger 1992; Gopal-Krishna \& Wiita 1992). The second scenario assumes that the variability is generated by change of the emission condition. A typical example is that fresh particles are injected into acceleration region and then are accelerated (e.g., Blandford \& Konigl 1979; Marscher \& Gear 1985; Celotti et al. 1991; Kirk et al. 1998). In order to reproduce both high energy radiation and variability of markarian 501, we change the injection rate of the low energy particles. It should be noted that when the shock front overruns a region in the jet in which the local plasma density is enhanced. The number of particles increase as an avalanche occurring in the jet, the injection rate can be expected to change. | 16 | 9 | 1609.09108 |
1609 | 1609.09422_arXiv.txt | Using numerical simulations, we show that jets with features of type II spicules and cold coronal jets corresponding to temperatures $10^{4}$ K can be formed due to magnetic reconnection in a scenario in presence of magnetic resistivity. For this we model the low chromosphere-corona region using the C7 equilibrium solar atmosphere model and assuming Resistive MHD rules the dynamics of the plasma. The magnetic filed configurations we analyze correspond to two neighboring loops with opposite polarity. The separation of the loops' feet determines the thickness of a current sheet that triggers a magnetic reconnection process, and the further formation of a high speed and sharp structure. We analyze the cases where the magnetic filed strength of the two loops is equal and different. In the first case, with a symmetric configuration the spicules raise vertically whereas in an asymmetric configuration the structure shows an inclination. With a number of simulations carried out under a 2.5D approach, we explore various properties of excited jets, namely, the morphology, inclination and velocity. The parameter space involves magnetic field strength between 20 and 40 G, and the resistivity is assumed to be uniform with a constant value of the order $10^{-2}\Omega\cdot m$. | \label{sec:Introduction} Magnetic reconnection is a topological reconfiguration of the magnetic field caused by changes in the connectivity of its field lines \citep{Priest_1984,Priest_et_al_2000}. It is also a mechanism of conversion of magnetic energy into thermal and kinetic energy of plasma when two antiparallel magnetic fields encounter and reconnect with each other. Magnetic reconnection can occur in the chromosphere, photosphere and even in the convection zone. In particular, the chromosphere has a very dynamical environment where magnetic features such as H$_\alpha$ upward flow events \citep{Chae_et_al_1998} and erupting mini-filaments \citep{Wang_et_al_2000} happen. The dynamics of the chromosphere at the limb region is dominated by spicules \citep{Beckers_1968} and related flows such as mottles and fibrils on the disk \citep{Hansteen_et_al_2006,De_Pontieu_et_al_2007a}. Spicular structures are also visible at the limb in many spectral lines at the transition region temperatures \citep{Mariska_1992,Wilhelm_2000}, and some observations suggest that coronal dynamics are linked to spicule-like jets \citep{McIntosh_et_al_2007,Tsiropoula_et_al_2012,Cheung_et_al_2015,Skogsrud_et_al_2015,Tavabi_et_al_2015a,Narang_et_al_2016}. With the large improvement in spatiotemporal stability and resolution given by the Hinode satellite \citep{Kosugi_et_al_2007}, and with the Swedish 1 m Solar Telescope (SST) \citep{Scharmer_et_al_2008}, two classes of spicules were defined in terms of their different dynamics and timescales \citep{De_Pontieu_et_al_2007c}. The so-called type I spicules have lifetimes of 3 to 10 minutes, achieve speeds of 10-30 km s$^{-1}$, and reach heights of 2-9 Mm \citep{Beckers_1968,Suematsu_et_al_1995}, and typically involve upward motion followed by downward motion. In \citep{Shibata&Suematsu_1982,Shibata_et_al_1982}, the authors studied in detail the propagating shocks using simplified one-dimensional models. In \citep{Hansteen_et_al_2006,De_Pontieu_et_al_2007a,Heggland_et_al_2007}, the authors study the propagation of shocks moving upwards passing through the upper chromosphere and transition region toward the corona. They also describe how the spicule-driving shocks can be generated by a variety of processes, such as collapsing granules, p-modes and dissipation of magnetic energy in the photosphere and lower chromosphere. In \citep{Matsumoto&Shibata_2010} the authors state that spicules can be driven by resonant Alfv\'en waves generated in the photosphere and confined in a cavity between the photosphere and the transition region. Another studies in this direction, like \citep{Murawski&Zaqarashvili_2010,Murawski_et_al_2011}, where they use ideal MHD and perturb the velocity field in order to stimulate the formation of type I spicules and macro-spicules. Furthermore in \citep{Scullion_et_al_2011}, the authors simulate the formation of wave-driven type I spicules phenomena in 3D trough a Transition Region Quake (TRQ) and the transmission of acoustic waves from the lower chromosphere to the corona. Type II spicules are observed in Ca II and H$\alpha$, these spicules have lifetimes typically less than 100s in contrast with type I spicules that have lifetimes of 3 to 10 min, are more violent, with upward velocities of order 50-100 km s$^{-1}$ and reach greater heights. They usually exhibit only upward motion \citep{De_Pontieu_et_al_2007b}, followed by a fast fading in chromospheric lines without observed downfall. Spicules of type II seen in the Ca II band of Hinode fade within timescales of the order of a few tenths of seconds \citep{De_Pontieu_et_al_2007a}. The type II spicules observed on the solar disk are dubbed ``Rapid Blueshifted Events'' (RBEs) \citep{Langangen_et_al_2008,Rouppe_et_al_2009}. These show strong Doppler blue shifted lines in the region from the middle to the upper chromosphere. The RBEs are linked with asymmetries in the transition region and coronal spectral line profiles \citep{De_Pontieu_et_al_2009}. In addition the lifetime of RBEs suggests that they are heated with at least transition region temperatures \citep{De_Pontieu_et_al_2007c,Rouppe_et_al_2009}. Type II spicules also show transverse motions with amplitudes of 10-30 km s$^{-1}$ and periods of 100-500 s \citep{Tomczyk_et_al_2007,McIntosh_et_al_2011,Zaqarashvili&Erdelyi_2009}, which are interpreted as upward or downward propagating Alfvenic waves \citep{Okamoto&De_Pontieu_2011,Tavabi_et_al_2015b}, or MHD kink mode waves \citep{He_et_al_2009,McLaughlin_et_al_2012,Kuridze_et_al_2012}. As mentioned above, there are several theoretical and observational results about type II spicules, but there is little consensus about the origin of type II spicules and the source of their transverse oscillations. Some possibilities discussed suggest that type II spicules are due to the magnetic reconnection process \citep{Isobe_et_al_2008,De_Pontieu_et_al_2007c,Archontis_et_al_2010}, oscillatory reconnection process \citep{McLaughlin_et_al_2012}, strong Lorentz force \citep{Martinez-Sykora_et_al_2011} or propagation of p-modes \citep{de_Wijn_et_al_2009}. Moreover, type II spicules could be warps in 2D sheet like structures \citep{Judge_et_al_2011}. A more recent study suggests another mechanism, for instance in \citep{Sterling&Moore_2016}, the authors suggested that solar spicules result from the eruptions of small-scale chromospheric filaments. The limited resolution in observations and the complexity of the chromosphere make difficult the interpretation of the structures, and even question the existence of type II spicules as a particular class, for instance in \citep{Zhang_et_al_2012}. In consequence, these difficulties spoil the potential importance of magnetic reconnection as a transcendent mechanism in the solar surface. Nevertheless, there is evidence that magnetic reconnection is a good explanation of chromospheric anemone jets \citep{Singh_et_al_2012}, which are observed to be much smaller and much more frequent than surges \citep{Shibata_et_al_2007}. A statistical study performed by \citep{Nishizuka_et_al_2011} showed that the chromospheric anemone jets have typical lengths of 1.0-4.0 Mm, widths of 100-400 km, and cusp size of 700-2000 km. Their lifetimes is about 100-500 s and their velocity is about 5-20 km s$^{-1}$. Other types of coronal jets can be generated by magnetic reconnection, for example in \citep{Yokoyama&Shibata_1995,Yokoyama&Shibata_1996} the authors using two-dimensional numerical simulations study the jet formation, or in \citep{Nishizuka_et_al_2008}, it is shown that emerging magnetic flux reconnects with an open ambient magnetic field and such reconnection produces the acceleration of material and thus a jet structure. The reconnection seems to trigger the jet formation in a horizontally magnetized atmosphere, with the flux emergence as a mechanism \citep{Archontis_et_al_2005,Galsgaard_et_al_2007}. Another approach uses a process that produces a magnetic reconnection using numerical dissipation of the ideal MHD equations, and the atmosphere model is limited to have a constant density and pressure profiles and assumes there is no gravity \citep{Pariat_et_al_2009,Pariat_et_al_2010,Pariat_et_al_2015,Rachmeler_et_al_2010}. In this paper we show that magnetic reconnection can be responsible for the formation of jets with some characteristics of Type II spicules and cold coronal jets \citep{Nishizuka_et_al_2008}, for that i) we solve the system of equations of the Resistive MHD subject the solar gravitational field, ii) we assume a completely ionized solar atmosphere consistent with the C7 model. The resulting magnetic reconnection accelerates the plasma upwards by itself and produces the jet. The paper is organized as follows, in Section \ref{sec:model_numerical_methods} we describe the resistive MHD equations, the model of solar atmosphere, the magnetic field configuration used the numerical simulations and the numerical methods we use. In Section \ref{sec:Results}, we present the results of the numerical simulations for various experiments. Section \ref{sec:conclusions} contains final comments and conclusions. | \label{sec:conclusions} In this paper we present the numerical solution of the equations of the resistive MHD submitted to the solar constant gravitational field, and simulate the formation of narrow jet structures on the interface low-chromosphere and corona. For this we use a magnetic field configuration of two superposed loops in a way that a current sheet is formed that allows the magnetic reconnection process, which in turn accelerates the plasma. An ingredient of our simulations is that we use a realistic atmospheric model that includes the transition region. The rarefaction of the environment above the transition region helps the acceleration of the plasma. We can summarize our findings in the schematic picture shown in Fig.\ref{fig:process}. This rarefied atmosphere allows the formation of a bulb at the top of the jet due to Kelvin-Helmholtz instability \citep{Kuridze_et_al_2016}, which is contained and stabilized by the magnetic field as found in \citep{Flint_2014,Zaqarashvili_2014}. We consider symmetric and asymmetric magnetic field configurations. In the symmetric case different jet properties where found in terms of separation and magnetic field strength of the loops. The magnetic field used ranges from 20 to 40 G, leading to the conclusion that the stronger the magnetic field the higher and faster the jet. The separation of the loops' feet was also found to be important, because it determines the thickness of the current sheet that later on produces the magnetic reconnection. With our parameters, a separation with $l_0=4$Mm is so large that no jet is formed anymore. The Temperature within the jet structure is of the order of $10^{4}$ K, which is within the observed range of a cool jet \citep{Nishizuka_et_al_2008}. An illustrative example is that of $B_{01}=B_{02}=40$G and $l_0=3.5$Mm, which shows a height of 7Mm measured from the transition region, and a maximum vertical velocity of $v_z\approx 34$ km s${}^{-1}$, parameters similar to those of Type II spicules \citep{De_Pontieu_et_al_2007c}. The evolution of these structures indicate that they last about 200 s, which is a lifetime similar to type II spicules. In the case of asymmetric magnetic field configurations we also simulated the formation of jets with similar properties of Temperature, velocity and height. These jets show a considerable inclination toward the loop with the weaker magnetic field. We found that the inclination of the jet depends on the magnetic field ratio of the two loops. According to the results of this paper, a good model for the formation of realistic jets mimicking type II spicules is to have two magnetic loops close together with opposite polarity. This produces a current sheet at chromospheric level capable to trigger magnetic reconnection. A key ingredient in the process is the inclusion of magnetic resistivity, which is a mechanism consistent with the magnetic reconnection process. Another feature of these jets is that they are based at the level of the transition region, which is characterized by a sharp gradient in density and temperature. | 16 | 9 | 1609.09422 |
1609 | 1609.00673_arXiv.txt | \noindent A model for strongly writhing confined solar eruptions suggests an origin in the helical kink instability of a coronal flux rope which remains stable against the torus instability. This model is tested against the well observed filament eruption on 2002 May~27 in a parametric MHD simulation study which comprises all phases of the event. Good agreement with the essential observed properties is obtained. These include the confinement, terminal height, writhing, distortion, and dissolution of the filament, and the flare loops. The agreement is robust against variations in a representative range of parameter space. Careful comparisons with the observation data constrain the ratio of the external toroidal and poloidal field components to $\Bet/\Bep\approx1$ and the initial flux rope twist to $\Phi\approx4\pi$. Different from ejective eruptions, two distinct phases of strong magnetic reconnection can occur. First, the erupting flux is cut by reconnection with overlying flux in the helical current sheet formed by the instability. If the resulting flux bundles are linked as a consequence of the erupting rope's strong writhing, they subsequently reconnect in the vertical current sheet between them. This reforms the overlying flux and a far less twisted flux rope, offering a pathway to homologous eruptions. | \label{s:intro} Erupting magnetic flux on the Sun often remains confined in the corona without evolving into a coronal mass ejection (CME). The rise of the flux then halts and any embedded filament or prominence material slides back to the bottom of the corona along the magnetic field lines \cite[e.g.,][]{HJi&al2003}. Confined and ejective eruptions begin similarly \citep{Moore&al2001}; both forms are usually associated with a flare. Their initially accelerating rise indicates the onset of an instability. Confined (or ``failed'') eruptions present an important testbed for theories of solar eruptions. Understanding what prevents an evolving eruption from becoming ejective is also relevant for the study of the space weather and its terrestrial effects \cite[e.g.,][]{Gosling1993, Webb&al2000}. In eruption models based on ideal MHD instability \citep{vanTend&Kuperus1978}, confinement results when the condition for the torus instability (in general terms: a sufficiently rapid decrease of the coronal field with height) is not met at or above the eruption site. This is possible if the eruption is caused by the onset of the helical kink instability in the stability domain of the torus instability \citep{Torok&Kliem2005}. If the helical kink saturates before the rising flux reaches the height range where the torus instability can act, then the eruption remains confined; otherwise a CME results. Another possibility arises if the coronal field is structured such that the condition for the torus instability is fulfilled in two separate height ranges enclosing a stable height range, as has been found, e.g., by \citet{YGuo&al2010} and \citet{ZXue&al2016a}. If slowly rising current-carrying flux reaches the lower unstable range, it erupts due to the torus instability, but is halted in the stable height range. Another model for confined eruptions suggests that the reconnection of two magnetic loops may yield two stable new loops, while producing a flare due to the release of magnetic energy during the reconnection \citep{Nishio&al1997, Hanaoka1997}. Observational support for this model was based on low-resolution data, which did not clearly reveal the nature of the interacting loop-shaped structures \cite[e.g.,][]{Green&al2002}. Coronal magnetic loops are no longer considered to contain sufficient free magnetic energy to power an eruption, rather the much larger amount of flux and free energy typically contained in a filament channel appears to be required. The suggested scenario indeed occurred in an event that showed the reconnection between two filaments \citep{Torok&al2011, Joshi&al2014a}. A CME was associated, but must have originated in the perturbed flux overlying the filaments which remained in place. \citet{YJiang&al2013, YJiang&al2014} reported partly similar (more complex) failed filament eruptions. Overall, events of this category are very rare, however. The confinement of eruptions may also be related to the existence of a coronal magnetic null point which spans a dome-shaped magnetic fan surface above the eruption site. In the case of an eruption, the footprint of the fan surface yields a circular flare ribbon \citep{Masson&al2009}. This configuration has been found to be associated with both confined eruptions \cite[e.g.,][]{HWang&CLiu2012, NDeng&al2013, Vemareddy&Wiegelmann2014, Kumar&al2015} and ejective ones \cite[e.g.,][]{CLiu&al2015, Joshi&al2015, Kumar&al2016}. Therefore, the fan surface of a coronal magnetic null does not appear to be the primary factor deciding the ejective vs. confined nature of eruptions launched under it. In fact, reconnection at the null may facilitate the removal of overlying, stabilizing flux \citep[e.g.,][]{XSun&al2013, CJiang&al2014}. Observational studies of confined eruptions have become quite frequent with the recent advances of observing capabilities \citep{YLiu&al2009, YShen&al2011, YShen&al2012, Netzel&al2012, Kuridze&al2013, HChen&al2013, HQSong&al2014, SYang&al2014, Kumar&Cho2014, RLiu&al2014, Joshi&al2014b, Kushwaha&al2014, Kushwaha&al2015, XCheng&al2015, TLi&JZhang2015, ZXue&al2016b}. Statistical studies of their association with source region structure and flare magnitude were also performed \citep{YWang&JZhang2007, XCheng&al2011}. Despite the large variety of these events, several trends are apparent. Strong overlying flux is typically observed, and for several events it was demonstrated to possess a height profile that prevents the torus instability. Naturally, this favors the central part of active regions above their periphery. Indications of the helical kink and the dissolution of embedded filaments are often found. Confinement also shows an association with the magnitude of energy release by the eruption, with only a minor fraction of the B-class flares but nearly all $>$X1-class flares being accompanied by CMEs. The underlying cause-effect relationship, i.e., whether the confinement limits the flare magnitude or an insufficient flare energy release prevents the eruption from developing into a CME, is not yet clear. However, since there are about three orders of magnitude between the weakest CME-associated flares (in the low-B-class range) and the strongest confined flare, the flare magnitude cannot be decisive by itself but must be considered in the context of the source region's magnetic structure, particularly the properties of the overlying flux. Occasionally, even eruptions producing X3 flares can remain confined, as in the exceptional active region (AR)~12192. Insufficient magnetic shear, twist or helicity in the AR's core field and overlying flux preventing the torus instability have been suggested to be the possible causes of confinement in this region \citep{Thalmann&al2015, XSun&al2015, HChen&al2015, JJing&al2015, Inoue&al2016, LLiu&al2016}. Since a strong eruption (the X3 flare) occurred in the first place, confinement by the overlying flux appears to be the most obvious explanation. Here we present a simulation study of the confined filament eruption on 2002 May~27, whose detailed and comprehensive observations were analyzed by \citet{HJi&al2003} and \citet{Alexander&al2006}. The initial phase of the eruption, up to the point the rising flux reached its terminal height, was already modeled by \citeauthor{Torok&Kliem2005} (\citeyear{Torok&Kliem2005}, henceforth TK05). Using a flux rope susceptible to the helical kink mode but not to the torus instability as the initial condition in their MHD simulations, the rise profile of the flux rope apex, the rope's developing helical shape, and its distortion during the deceleration showed close agreement with the observations. Thus, the helical kink instability appears to be the prime candidate mechanism for this event. Our simulations substantiate this model for confined eruptions in two ways. First, we extend the computations to model the whole event, and second, a parametric study suggests that the requirement on the initial twist can be relaxed to about $4\pi$, which is closer to twist estimates for other events than the estimate of $\approx5\pi$ by TK05. We also focus on the magnetic reconnection, demonstrating that it occurs in two distinct locations and phases which correspond to the observed brightenings and changes of topology, and consider the fate of the erupting flux, which can reform a (less twisted) flux rope. | \label{s:conclusions} (1) Using a kink-unstable force-free flux rope in equilibrium as the initial condition, the MHD simulations presented in this paper achieve good agreement with the essential properties of the confined filament eruption in AR~9957 on 2002 May~27, as observed by the \textsl{TRACE} satellite in the EUV. These include (i) the confined nature of the eruption, (ii) its terminal height, (iii) the writhing of the erupting flux according to the $m=1$ helical kink mode which yields the observed inverse-gamma shape, (iv) the dissolution of the erupting flux by reconnection with the overlying flux, and (v) the formation of the flare loop arcade, which shows indications of twist, by a second phase of reconnection. The agreement is obtained in a representative range of parameter space. This robustness supports the model for confined eruptions by TK05, which assumes a kink-unstable and torus-stable flux rope to exist at the onset of the eruption. (2) Through quantitative comparisons with the observations, the performed parametric study constrains the ratio of external toroidal (shear) field component and external poloidal (strapping) field component to $\Bet/\Bep\approx1$ and the average twist in the initial flux rope to $\Phi\approx4\pi$, in better agreement with recent twist estimates for other events than the estimate in TK05. (3) Different from ejective eruptions (CMEs), the confined eruption triggered by the helical kink is found to comprise two distinct phases of strong reconnection. The first phase occurs in the helical current sheet and destroys the rising flux rope through reconnection with overlying flux (opposite to standard ``flare reconnection'' in ejective events). In the whole range of parameters studied here, this reconnection occurs after a strong writhing of the erupted flux rope, such that the resulting two flux bundles are linked. This, in turn, causes the second phase of strong reconnection between the legs of the original flux rope in the vertical current sheet in the center of the system, which restores the overlying flux. Flare loops are formed in this flux above the vertical current sheet. (4) The second reconnection also results in the reformation of a flux rope of similar or only moderately reduced flux. Although the twist is strongly reduced (to $\Phi<\pi$ in our simulations), the reformation offers a pathway to homologous eruptions, preferably if the sequence starts with one or several confined events and includes only one CME at the end. | 16 | 9 | 1609.00673 |
1609 | 1609.09564_arXiv.txt | In this paper, the dependence of the broad \FeK line on the physical parameters of AGN, such as the black hole mass $\mobh$, accretion rate (equivalently represented by Eddington ratio $\ledd$), and optical classification, is investigated by applying the X-ray spectra stacking method to a large sample of AGN which have well measured optical parameters. A broad line feature is detected ($>3\sigma$) in the stacked spectra of the high $\ledd$ sub-sample ($\log\ledd>-0.9$). The profile of the broad line can be well fitted with relativistic broad line model, with the line energy consistent with highly ionized Fe K$\alpha$ line (i.e. \ion{Fe}{xxvi}). A model consisting of multiple narrow lines cannot be ruled out, however. We found hints that the Fe K line becomes broader as the $\ledd$ increases. No broad line feature is shown in the sub-sample of broad-line Seyfert 1 (BLS1) galaxies and in the full sample, while a broad line might be present, though at low significance, in the sub-sample of narrow-line Seyfert 1 (NLS1) galaxies. We find no strong dependence of the broad line on black hole masses. Our results indicate that the detection/properties of the broad \FeK line may strongly depend on $\ledd$, which can be explained if the ionization state and/or truncation radius of the accretion disc changes with $\ledd$. The non-detection of the broad line in the BLS1 sub-sample can be explained if the the average EW of the relativistic Fe K$\alpha$ line is weak or/and the fraction of sources with relativistic Fe K$\alpha$ line is small in BLS1 galaxies. | \label{sec:intro} Observational evidences of a broad \FeK line feature are found in the X-ray spectra of some active galactic nuclei (AGN), such as MCG-6-30-15 \citep{tanaka:1995, fabian:2002, miniutti:2007}, NGC 3516 \citep{turner:2002, markowitz:2006}, 1H 0707-495 \citep{fabian:2009} and others \citep{miller:2007, nandra:2007}. The broad \FeK line is generally believed to originate from the inner region of the accretion disc via the K-shell fluorescence process, and to be broadened due to the Doppler boosting, gravitational redshift and the transverse Doppler effect \citep{fabian:1989}. The energy at which the ``red'' wing of the broad line truncates is directly linked to the inner radius of the accretion disc that is commonly thought to be at the innermost stable circular orbit (ISCO). The spin of the black hole, which is related to the ISCO through a monotonic relation \citep{bardeen:1972}, can be inferred by modeling the broad \FeK line profile \citep{brenneman:2006, dauser:2010}. Besides its mass, an astrophysical black hole is completely characterized by its spin. The distribution of the black hole spin may yield important insights into the growth history and accretion process of supermassive black holes (SMBH). For instance, models in which SMBH growth is dominated by BH-BH mergers can lead to a bimodal distribution with one peaked at 0 and the other located at $\sim0.7$, whereas growth via gas accretion predicts a rapidly spinning or a slowly spinning population depending upon whether the BHs gain their masses via prolonged accretion or chaotic accretion \citep{moderski:1996, volonteri:2005}. Moreover, spin also determines the radiative efficiency, which is the mass-to-energy conversion efficiency, and thus influences how efficiently BHs accrete mass during the accretion phase. Black hole spin can also be a potent energy source, and may drive the powerful relativistic jets that are seen from many BH systems through the Blandford-Znajek mechanism \citep{blandford:1977}. There have been several studies for charactering the broad \FeK line and measuring the spin of SMBH based on relativistic reflection spectra in the literature \citep[e.g.][]{nandra:2007, de-la-calle-perez:2010, patrick:2012, walton:2013}. However, a significantly broad \FeK line is detected in only $<50$\, per cent of the sources in previous studies with different samples. In total, detections of the broad \FeK line are reported in the X-ray spectra of $\sim46$ AGNs. Among those, about 22 sources have reliable spin measurements \citep{reynolds:2013, brenneman:2013, reynolds:2014}, making it difficult to draw any robust statistical inferences on the distribution of BH spin. The reason for the lack of apparent relativistic broad \FeK line in the X-ray spectra of some AGN is still unclear. Observationally, one possible explanation is the low signal-to-noise (S/N) of the X-ray data for the majority of AGN. As demonstrated in \citet{mantovani:2014}, the broad \FeK line is revealed in the composite X-ray spectrum of IC 4329A observed by {\it Suzaku}, while it is absent in each individual observation due to the low S/N of the data. Indeed, to determine accurately the continuum spectrum and to reveal unambiguously the \FeK line profile, a large number of X-ray photons collected at high energies is required \citep{de-la-calle-perez:2010, nandra:2007, mantovani:2014}. Theoretically, the strength of the line is a function of the geometry of the accretion disc, which determines the solid angle subtended by the reflecting matter as seen by the X-ray source. It also depends on the elemental abundances of the reflecting matter, the inclination angle at which the reflecting surface is viewed, and the ionization state of the surface layers of the disc. \citet{bhayani:2011} analyzed a sample of 11 Seyfert galaxies observed by XMM-Newton that appear to be missing a broad \FeK line. They argued that the lack of apparent relativistic \FeK line can be explained if this feature becomes indistinguishable from the underlying continuum, as a result of a combination of several effects, e.g. blending and Comptonization in an ionized disc, strong relativistic effects and, in some cases, a high disc inclination. Calculations have also shown that the ionization state of the accretion disc, which will affect the observed line energy as well as the line equivalent width (EW) of the broad line \citep{matt:1993, ross:1993, nayakshin:2000}, depends strongly on the accretion rate. Thus a correlation between the properties of the broad \FeK line and the accretion rate is expected. Such a correlation has been reported in \citet{inoue:2007}, although a relativistic line is statistically not required by their data. In order to investigate the correlation between the properties of the broad \FeK line and the physical parameters of AGN, a homogeneously selected sample of AGN with well measured optical parameters is preferred. The virial BH mass $\mobh$ of AGN can be estimated through the empirical relations using the emission line widths and continuum luminosities \citep{kaspi:2005, shen:2013}. The Eddington ratio $\ledd$ is given by ${\ledd} = L_{\rm bol}/L_{\rm Edd}$, where the $L_{\rm Edd}$ is defined as $L_{\rm Edd} = 1.26\times10^{38}(\mobh/M_{\odot})$. The bolometric luminosity $L_{\rm bol}$ is usually estimated from the integration of the broadband spectral energy distribution (SED) or from a single band assuming a bolometric correction. The $\mobh$ and $\ledd$ for individual sources or small samples in previous studies (e.g. \citealt{inoue:2007}) were derived using different data analysis methods, BH mass estimation formalisms and/or heterogeneous data sets. Because of the lack of large homogeneously selected samples with high S/N X-ray data and well measured optical parameters, it is still unclear how is the appearance/property of the broad \FeK line dependent on the physical properties of AGN. X-ray spectral stacking is an effective way to obtain composite spectra with very high S/N. A broad relativistic Fe K$\alpha$ line is found in the stacked spectra of the Lockman Hole field using \xmm observations \citep{streblyanska:2005}. However, while a narrow line is significantly detected, the broad line is not clearly seen in the stack spectra of different samples in previous studies \citep{corral:2008, chaudhary:2010, chaudhary:2012, iwasawa:2012a, falocco:2013}. The relativistic Fe K$\alpha$ line is detect at $6\sigma$ using a sample with high S/N observed with \xmm in \citet{falocco:2014}. They pointed out that the low average S/N of the spectra, which make the continuum not to be well determined, can explain the low significance of the broad line in previous works. The non-detection of the broad line in the average spectra, which are obtained using samples including sources with diverse properties, can also be interpreted if the property of the relativistic line is highly dependent on one or more physical parameters of AGN. By applying the same rest-frame X-ray spectral stacking method \citep{corral:2008, falocco:2012} to a sample of NLS1 galaxies, \citet{liu:2015} found that there exists a prominent broad \FeK line in the composite X-ray spectrum. They suggested that broad \FeK line is perhaps common in AGN with high $\ledd$ (e.g. NLS1 galaxies). The average properties of the relativistic Fe K$\alpha$ line for a uniform selected BLS1 galaxies sample have not been explored yet, even though a broad \FeK line is detected in several BLS1 galaxies. In this paper, using a large homogeneously selected sample of AGN which have well measured $\mobh$ and $\ledd$, we investigate how the detection/property of the broad \FeK line depends on the physical parameters, such as the $\mobh$, $\ledd$, and optical classification of AGN, by means of X-ray spectral stacking. We use the cosmological parameters $H_\mathrm{0}=70\,\mathrm{km\,s^{-1}\,Mpc^{-1}}$, $\Omega_\mathrm{M}=0.27$, and $\Omega_\mathrm{\Lambda}=0.73$. All quoted errors correspond to the 90 per cent confidence level for one interesting parameter, unless specified otherwise. | \subsection{Dependence of the broad \FeK line on $\ledd$?}\label{subsec:edd_fe} A broad line feature, which is detected at $>3\sigma$ level ($\Delta\chi^2=18.14$ with one additional d.o.f, see Section \ref{subsubsec:edd_subs}) and can be well fitted with a relativistic \FeK line model (Case A, see Section \ref{subsubsec:edd_subs}), is revealed in the high $\ledd$ sub-sample, while no significant broad line features are found in the low and medium $\ledd$ sub-samples. The broad line feature in the high $\ledd$ sub-sample can also be fitted with multiple Gaussian components (Case B, see Section \ref{subsubsec:edd_subs}). The line energy of the broad line in Case A is consistent with a highly ionized \FeK line, assuming that the average inclination of the accretion disc is less than $45\degr$ for type-1 AGN (see Section \ref{subsubsec:edd_subs}). There is an indication that the line width of the \FeK line becomes broader as the $\ledd$ increases, as shown in Fig. \ref{fig:edd_subs} and Table \ref{tab:fit_subs}, consistent with the results presented in \citet{inoue:2007}. We should note that the relation between the $\ledd$ and the properties of the broad \FeK line may be even stronger considering that the uncertainties in the estimates of $\ledd$ can potentially weaken the relation. A highly ionized emission line is shown in the average spectrum of a sample with high $\ledd$ in \citet[][see also \citealt{liu:2015}]{iwasawa:2012a}. As suggested by \citet{iwasawa:2012a}, the narrow ionized emission line found in their sample can be produced by outflow/winds from accretion disc. This can also explain the highly ionized emission line shown in Case B of our high $\ledd$ sub-sample. However, the highly ionized line ($E_\mathrm{BL}\approx6.94\,\mathrm{keV}$, corresponding to \ion{Fe}{xxvi}) should be originated from the accretion disc if it is a relativistic broad line as in Case A. In this case the results can be interpreted by the fact that the ionization state of the disc will be higher when the $\ledd$ goes up \citep{inoue:2007}. Theoretical calculations have shown that the ionization state of the accretion disc becomes higher as the mass accretion rate increases, producing highly ionized Fe K emission line, e.g. \ion{Fe}{xxv/xxvi} \citep{matt:1993, ross:1993, nayakshin:2000}. The equivalent width of the fluorescence line will change with the ionization state of the accretion disc: due to the resonant trapping opacity, the equivalent width of the Fe K line will first decrease, then increase strongly when the iron is more ionized than \ion{Fe}{xxiii}, reaching a maximum value of $\sim500$\,eV for very high ionization state (e.g. $\xi\sim2000$). This is consistent with the measured equivalent width of the broad line in the stacked spectrum of the high $\ledd$ sub-sample. On the contrary, the equivalent width of the Fe K$\alpha$ line will be small, e.g. $<200\,$eV for a neutral or low-ionization accretion disc, which can explain the non-detection of the broad line in the low and medium $\ledd$ sub-samples (see Sec. \ref{subsec:nl_bl}). Alternatively, in case A the results can also be interpreted if the truncation radius of the accretion disc is strongly dependent on the $\ledd$. It has been suggested that the optically thick physically thin accretion disc, which exists at the low $\ledd$ regime, terminates beyond the ISCO and the inner region is filled with a hot advection-dominated accretion flow\citep[ADAF, e.g.][]{narayan:1994, narayan:1995}. This model has been widely used as a mechanism for explaining state changes in X-ray binaries \citep[XRB, e.g.][]{esin:1997, done:2007}. Spectral evidence for disc truncation has found in XRB \citep[e.g.][]{esin:2001, done:2010} and some radio-loud AGN \citep{marscher:2002, lohfink:2013}. The non-detection of the relativistic line is expected if indeed the disc is truncated at larger radius for AGN with low $\ledd$, as the case for our low and medium $\ledd$ sub-samples. On one hand, the profile of the emission line, as well as the reflection spectrum will not be significantly blurred by the relativistic effect at large radius. On the other hand, both the emissivity index, which is thought to be larger at the inner part of the accretion disc, and the region that is illuminated by the power-law continuum will decrease if the accretion disc is terminated at large radius, thus the strength and the equivalent width of the emission line will be weak. \subsection{Broad \FeK line in NLS1 and BLS1 galaxies}\label{subsec:nl_bl} \begin{figure*} \begin{center} \includegraphics[width=1.0\columnwidth]{bl_100} \includegraphics[width=1.0\columnwidth]{bl_50} \caption{The ratio of the stacked spectrum to a best-fitting continuum for the BLS1 sub-sample is shown as blue points. Different colours represent different disc inclinations (magenta: $i=15\degr$; red: $i=30\degr$; blue: $i=45\degr$). Left: the shadow areas show the $1\sigma$ confidence interval of the stacked spectra for the simulated sources (fraction $f=100$). The distribution of the EW for the broad line is estimated based on the $\sim44$ sources which have broad \FeK line measurements. Right: the EW of the broad \FeK line in the simulations is fixed at 400\,eV, while the fraction of sources with broad \FeK line is 50 per cent. \label{fig:bl_percent}} \end{center} \end{figure*} In addition to the neutral \FeK line, a highly ionized emission line feature is clearly shown in the average spectrum of the NLS1 sub-sample (Fig. \ref{fig:nl_bl}). The best-fitting line width of the highly ionized emission line is 156\,eV with a lower limit of 22\,eV. This may be an indication of a broad \FeK line, though with low significance ($\Delta\chi^2=2.9$ if the line width is fixed at 0). The broad relativistic \FeK line should be common in NLS1, as suggested by \citet{liu:2015}. They found a prominent broad \FeK line in the composite X-ray spectrum of a NLS1 sample consisting of 51 sources. As shown in the left panel of Fig. \ref{fig:mbh_edd}, there is a large overlap between the NLS1 and high $\ledd$ sub-samples, thus it is not surprising that a broad line may be detected in the NLS1s if the properties of the broad line indeed depends on $\ledd$ as suggested in Section \ref{subsec:edd_fe}. The stacked spectrum of our NLS1 sub-sample can be well fitted with the relativistic line model. The low significance of the broad line in our NLS1 sub-sample may be due to the much lower S/N of the data (see Table \ref{tab:sub_samples}). The EWs of the narrow ($48^{+56}_{-48}\,\mathrm{eV}$) and broad line ($495^{+288}_{-257}\,\mathrm{eV}$), as well as the BH spin ($a<0.53$), are in agreement with the results found in \citet{liu:2015}. No broad line feature is detected in the average spectrum of our BLS1 sub-sample and also the whole sample. \citet[][see also \citealt{corral:2008, falocco:2012, falocco:2013}]{falocco:2014} found a significant relativistic line in the stacked spectrum of a sample selected from the AGN catalogue built by \citet{veron-cetty:2010}, however. This may be simply due to the different samples used in different studies. The samples used in those previous works consist of different types of AGN, and often have very different properties (e.g. higher redshift and different luminosity range) comparing with our sample. Here we compared the results between the our BLS1 and NLS1 sub-samples, as well as the NLS1 sample presented in \citet{liu:2015}. The non-detection of the broad line in the BLS1 sub-sample can be explained if the fraction of sources with relativistic \FeK line is lower, or the average EW of the broad line is relatively smaller, than that in the NLS1 ($\sim 400$\,eV found in \citealt{liu:2015}). To roughly estimate the fraction of sources with broad \FeK line in BLS1 galaxies we carry out simulations. Using the method described in Sec.\,\ref{subsec:simulations}, we simulate a set of `observed' spectra for the sources in the BLS1 sub-sample. We assume that the fraction of sources with broad line is $f$ in the simulated spectra of the BLS1 sub-sample. Spectra without a relativistic component are generated using the best-fitting continua of the sources, while for the remaining ones a broad line model is added to the best-fitting continua. The spin parameter is randomly distributed in the range of 0-0.998. In order to compare the average broad line profile for different inclination angles, we simulate spectra for three disc inclinations, i.e. $i=15\degr$, $30\degr$, $45\degr$. Two different distributions for the EW of the relativistic component are assumed. In the first case (referred as Case 1 hereafter) the line EWs are randomly drawn from a probability distribution constructed on the basis of the $\sim 44$ sources \citep[][two sources with very strong broad \FeK, i.e. EW$>600$\,eV, are excluded]{brenneman:2013} which have broad \FeK line measurements. The typical values for the EW of the broad line component in Case 1 are smaller than the average line EW found in the stacked spectrum of our NLS1 sub-sample, e.g. $<300$\,eV. In another case (Case 2) the EW of the broad line component in the simulation is fixed at a comparable value to the observed average EW of the broad \FeK line in NLS1, e.g. $\mathrm{EW}=400$\,eV. The same stacking method is then applied to those simulated spectra. The results for Case 1 and Case 2 are shown in the left and right panel of Fig. \ref{fig:bl_percent}, respectively. The shadow areas represent the $1\sigma$ confidence intervals for the simulated composite spectra. The observed data of our BLS1 sub-sample are shown with blue points. The different colours represent the stacked line profiles for different disc inclinations (magenta: $i=15\degr$; red: $i=30\degr$; blue: $i=45\degr$). It is obvious that the non-detection of the broad \FeK line is not surprising in Case 1, since the observed data is consistent with the simulated results for all the different disc inclinations even when $f$ equals to 100 per cent. The fraction of sources with broad \FeK line in our BLS1 sub-sample should be much smaller than 50 per cent in Case 2 if the disc inclination is $i=15\degr$. For larger disc inclinations, e.g. $i=30\degr$ (red) and $45\degr$ (blue), the simulated results are consistent with the observed data if fraction $f\la50$\,per cent. This fraction is compatible with the fractions reported in previous studies, e.g. $\sim45\,$per cent in \citet{nandra:2007}; $36\,$per cent in \citet{de-la-calle-perez:2010} and $\sim50\,$per cent in \citet{patrick:2012}. \subsection{Non-detection of broad line in the $\mobh$ sub-samples} No significantly broad line feature is detected in the stacked spectra of the $\mobh$ sub-samples, though an indication of broad line in the MBH3 sub-sample is present. This may suggest that the fraction of sources with broad \FeK line or/and the equivalent width of the broad \FeK line do not strongly depend on the black hole mass, considering that a broad line is detected in the high $\ledd$ sub-sample with similar S/N. The non-detection of broad \FeK line in the $\mobh$ sub-samples can also be understood if the EW of the broad line is mainly driven by the $\ledd$. As shown in Fig. \ref{fig:mbh_subs}, although the $\mobh$ is slightly correlated with $\ledd$, there is a large dispersion of the $\ledd$ distribution in each $\mobh$ bin, which leads to a relatively small average EW of the broad line, thus make it difficult to be detected. | 16 | 9 | 1609.09564 |
1609 | 1609.05568_arXiv.txt | We consider localized soliton-like solutions in the presence of a stable scalar condensate background. By the analogy with classical mechanics, it can be shown that there may exist solutions of the nonlinear equations of motion that describe dips or rises in the spatially-uniform charge distribution. We also present explicit analytical solutions for some of such objects and examine their properties. | Spatially-homogeneous solutions in the complex scalar field theories with the global $U(1)$-invariance have been proven to be very useful in different branches of modern physics. Perhaps the most known example of their application to cosmology is the Affleck-Dine mechanism of baryogenesis \cite{Affleck:1984fy}. Evolution of the spatially-homogeneous condensate in the Early Universe, which is usually studied numerically, is subject to certain restrictions in order to yield a successful cosmological scenario \cite{Doddato:2011fz}. For instance, a possible spatial instability of the condensate results in its fragmentation into nontopological solitons --- Q-balls. The latter, in turn, can be a crucial ingredient in the solution of the dark matter problem \cite{Kusenko:1997si}. This makes inhomogeneous classical solutions also of considerable interest in cosmology. Their another application is related to the possibility of production of gravitational waves \cite{Zhou:2015yfa,Antusch:2016con,Katz:2016adq}. Emergence of localized stationary configurations was first discovered in the systems whose evolution is governed by the Nonlinear Schr\"{o}dinger Equation (NSE) \cite{Zakharov}. In nonlinear optics these solutions are known as bright solitons. Similar solutions in a theory of the complex scalar field in four dimensional space-time, possessing the global $U(1)$-charge, were called ``Q-balls'' by S. Coleman \cite{Coleman:1985ki}. NSE admits another interesting class of solutions corresponding to ``dark solitons'' in a stable medium \cite{Zakharov_Shabat}. They have the form of a dip in a homogeneous background. It is important to note that these solutions are of the topological nature. In particular, they cannot be deformed into the surrounding condensate by a finite amount of energy. Therefore, the question arises about the existence and properties of the analogs of dark solitons in the complex scalar field theory, where they presumably can be analyzed by the same methods as the ordinary Q-balls. The existence of the dip-in-charge-like solutions in scalar field theories is not a manifestation of some specific properties of these theories. In fact, such solutions exist for the usual ``Mexican hat'' scalar field potential. To see this, let us consider the complex scalar field $\phi$ with the Lagrangian density \begin{equation} \partial^{\mu}\phi^*\partial_{\mu}\phi-\frac{\lambda}{2}(\phi^*\phi-v^2)^2. \label{phi4} \end{equation} If $\lambda>0$, the theory admits the well-known real static solution --- the kink, which has the form \begin{equation}\label{RealKink} \phi=v\tanh\left(\sqrt{\frac{\lambda}{2}}v x\right). \end{equation} It can be generalized to a class of stationary but not static solutions as follows, \begin{equation} \phi=e^{i\omega t}f(x), \label{anz1_1} \end{equation} where $\omega$ is a constant parameter and \begin{equation}\label{chargedkink} f(x)=\sqrt{v^2+\frac{\omega^2}{\lambda}}\tanh\left(\sqrt{\frac{\lambda}{2}\left(v^2+\frac{\omega^2}{\lambda}\right)}\,x\right). \end{equation} Then, for the $U(1)$-charge density $\rho$ we get \[ \rho=2\omega f^2, \] which clearly has a dip around the origin $x=0$. The kink solution (\ref{RealKink}) is unstable in this model and can be interpreted as a sphaleron in the Abelian gauged version of (\ref{phi4}), see \cite{Bochkarev:1987wg} for details. Another distinctive feature of the model (\ref{phi4}) is the stability of the charged condensate as long as $\lambda>0$. We note that the solution (\ref{anz1_1}) requires an infinite amount of energy to be deformed into the spatially-homogeneous condensate of the same charge or frequency. In this paper we present the soliton-like localized solutions in a theory of the complex scalar field, which describe inhomogeneities in the charge distribution of the condensate and can be deformed into the spatially-homogeneous condensate of the same frequency using a finite amount of energy. We will refer to such solutions as ``Q-holes'' or ``Q-bulges'' in order to stress their similarity to the ordinary Q-balls and to the ``holes in the ghost condensate'' of \cite{Krotov:2004if}. In the next section we will argue in favor of existence of these solitons with the help of Coleman's overshoot-undershoot method and survey their general properties. In Section~\ref{Explicit examples of Q-holes} we will present and examine the explicit examples of Q-holes in one and three spatial dimensions. In Section~4 we will discuss the classical stability (in fact, instability) of Q-holes and Q-bulges. In Conclusion we will briefly discuss the obtained results. | In this paper we have presented Q-holes and Q-bulges --- two classes of localized configurations representing dips and rises in the spatially-homogeneous charged time-dependent scalar condensate. The important feature of these configurations is that they can be deformed into the condensate by a finite amount of energy. We expect that inhomogeneities of this type may be crucial for the nonlinear dynamics of the condensate in the Early Universe, in particular, for its fragmentation into Q-balls. We have also found the explicit solutions for Q-holes in the model with a simple piecewise-parabolic potential proposed in \cite{Theodorakis:2000bz}, and examined their properties. It has been shown that the renormalized energy $E_{ren}$ of Q-holes can take both positive, zero and negative values. In this paper, we did not address in detail the question of quantum stability of Q-holes and Q-bulges. Of course, if ``ordinary'' particles interact with Q-holes and Q-bulges through, say, the combination $\phi^{*}\phi$ (which is time-independent for these solutions and for the scalar condensate), Q-bulges and Q-holes with $E_{ren}>0$ can decay into such particles. Moreover, one may expect that Q-bulges and Q-holes can be created (even spontaneously) in processes involving these particles. So, this case is rather standard. However, the case of excitations of the scalar field $\phi$ above the condensate, which are supposed to form the corresponding scalar particles, is not so trivial. For ordinary Q-balls one can define a standard vacuum far from the core of the soliton and apply a standard quantization procedure to the perturbations above this vacuum. For the time-dependent scalar condensate, excitations on top of the background have nonstandard dispersion laws like the one in Eq.~(\ref{excitdispl}). Moreover, one can check that even the charge of the excitation with respect to the condensate charge also has a very nonstandard form, and the standard quantization procedure can not be applied to such excitations. It would be interesting to see what should be defined as ``particles'' related to the excitations of the form (\ref{pertback}) on top of the time-dependent background, and what must be the consistent quantization procedure, providing us with the correct definition of energy of the quantum excitations. These questions call for further detailed investigation. | 16 | 9 | 1609.05568 |
1609 | 1609.07047_arXiv.txt | Lifetime measurements using beam foil techniques for the $5s5p^{2}$ $^{2}D_{3/2,5/2}$ levels in Sn~{\sc ii} are presented. The resulting oscillator strengths for transitions at 1699.4, 1831.8 and 1811.2~\AA{} are reported. We also studied these levels with multi-configuration Dirac Hartree-Fock (MCDHF) calculations using a development version of the GRASP2K package. Our experimental and theoretical results are compared with other available studies. | Better understanding of radiative transitions of the Sn$^{+}$ ion has a number of applications including the determination of Sn abundances in the interstellar medium. The most recent studies on the Sn~{\sc ii} atomic structure for the configurations of our interest were performed by Oliver and Hibbert (2010) using large-scale Breit-Pauli configuration-interaction calculations and by Alonso-Medina \textit{et al} (2005) with relativistic Hartree-Fock (RHF) calculations. The discrepancies among the available calculations as well as the absence of experimental data especially on the $5s^{2}5d$ $^{2}D$ and $5s5p^{2}$ $^{2}D$ terms was the main motivation of our work. The transitions arising from the $5s^{2}5d$ $^{2}D$ multiplet are of great importance in astrophysics, and due to the strong mixing between the $5s^{2}5d$ $^{2}D_{3/2,5/2}$ and $5s5p^{2}$ $^{2}D_{3/2,5/2}$ levels, more experimental data will lead us to a better understanding of this ion's atomic structure. The $5s^{2}5d$ $^{2}D$ has already been experimentally studied by Schectman \textit{et al} (2000) at the University of Toledo using beam-foil techniques and therefore, the focus of the present work is on the $5s5p^{2}$ $^{2}D$ term. The results of our experimental work will be discussed in Section 2 of this paper. We have also performed a set of calculations with a fully relativistic approach using the GRASP2K package (J{\"{o}}nsson \textit{et al}~2014), with the main focus on the terms mentioned above as well as their radiative transitions to the ground state. These calculations will be discussed in Section 3. Our summary and conclusions are provided in Section 4 of this paper. | In this work we presented the results of our lifetime measurements using beam-foil techniques for the $5s5p^{2}$ $^{2}D_{3/2,5/2}$ levels in Sn~{\sc ii} which are presented in Tables 1 and 2. They are the only available experimental measurements for these levels to date. In addition to the values for lifetimes and oscillator strengths included in Tables 1 and 2 from previous work, there was another available calculation by Marcinek and Migdalek (1994). They used a relativistic Hartree-Fock method and included configuration interactions which yielded a multiplet \textit{f}-value of 0.0142 for the $5s^{2}5p$ $^{2}P$ $-$ $5s5p^{2}$ $^{2}D$ transitions and 1.33 for the $5s^{2}5p$ $^{2}P$ $-$ $5s^{2}5d$ $^{2}D$ transitions; from our experimental results we obtained 0.0247 and 1.17, respectively, for the corresponding transitions of these multiplets. We also reported the results of our fully relativistic MCDHF approach for this ion using the GRASP2K package for the \textit{ab-initio} energy eigenvalues of levels of interest in Table 3. Generally, these eigenvalues are in good agreement with the observed energy levels. There is a noticeable discrepancy for the fine structure splitting of the $^{2}P$ terms in our results. Using only the $5s^{2}5p$ $^{2}P$ configuration in the reference set with no substituitions, the value for the fine structure splitting was 4152.74 cm$^{-1}$ which is in good agreement with the observed energy levels; adding virtual orbitals up to the 7p orbital for the odd parity solution only reduced this value to 4079.53 cm$^{-1}$ and lastly by extending the calculations up to the 9p orbital, this value was 3903.18 cm$^{-1}$. However, this difference in the energy eigenvalue of the fine structure did not seem to affect the transition probabilities of levels of interest significantly. The results of the calculations for the relevant transitions are also reported in length and velocity gauges in Tables 1 and 2. Even though our \textit{ab-initio} calculations showed better consistency between the two gauges for the $5s^{2}5d$ $^{2}D$ term compared to the $5s5p^{2}$ $^{2}D$, they were overall in better agreement with the experimental results compared to other available theoretical calculations, especially for the $5s5p^{2}$ $^{2}D$ term. The velocity gauge results from the GRASP2K calculations are surprisingly in better agreement with the experimental values than the length gauge values. We have shown experimentally that the transitions to the ground state arising from the $5s^{2}5d$ $^{2}D$ term with the higher energies (71,406 and 72,048 cm$^{-1}$) are the more favorable transitions with shorter lifetimes, while the ones from the $5s5p^{2}$ $^{2}D$ term with the lower energies (58,844 and 59,463 cm$^{-1}$) are much longer lived. Previously this was known only theoretically. Our MCDHF calculations support these facts as well. \ack This work was supported in part by grant HST-AR-12123.001-A from the Space Telescope Science Institute. We thank the anonymous referees for helpful comments that improved the paper. | 16 | 9 | 1609.07047 |
1609 | 1609.09093_arXiv.txt | We present multi-epoch X-ray spectral observations of three Type IIn supernovae (SNe), SN~2005kd, SN~2006jd, and SN~2010jl, acquired with {\it Chandra}, {\it XMM-Newton}, {\it Suzaku}, and {\it Swift}. Previous extensive X-ray studies of SN~2010jl have revealed that X-ray spectra are dominated by thermal emission, which likely arises from a hot plasma heated by a forward shock propagating into a massive circumstellar medium (CSM). Interestingly, an additional soft X-ray component was required to reproduce the spectra at a period of $\sim$1--2 yr after the SN explosion. Although this component is likely associated with the SN, its origin remained an open question. We find a similar, additional soft X-ray component from the other two SNe IIn as well. Given this finding, we present a new interpretation for the origin of this component; it is thermal emission from a forward shock essentially identical to the hard X-ray component, but directly reaches us from a void of the dense CSM. Namely, the hard and soft components are responsible for the heavily- and moderately-absorbed components, respectively. The co-existence of the two components with distinct absorptions as well as the delayed emergence of the moderately-absorbed X-ray component would be evidence for asphericity of the CSM. We show that the X-ray spectral evolution can be qualitatively explained by considering a torus-like geometry for the dense CSM. Based on our X-ray spectral analyses, we estimate the radius of the torus-like CSM to be on the order of $\sim$5$\times10^{16}$\,cm. | \label{sec:intro} Mass loss from massive stars is directly linked to the massive stars' evolution, affecting a star's apparent temperature, and often luminosity and burning lifetime as well, for most of their lives. Therefore, it has been actively investigated from both observational and theoretical points of view \citep[e.g.,][]{2014ARA&A..52..487S}. In particular, the mass loss in the final stage toward SN explosions, which determines the type of a resulting supernova (SN) explosion, is one of the main issues in modern stellar astrophysics. One way to study the mass loss at the last stage of the massive star's evolution is to observe young SNe. Mass loss influences the stellar environments, forming the circumstellar medium (CSM) around the progenitor star. The CSM will be excited by the collision with the SN ejecta, and emits intense radiation in various wavelengths, allowing us to reveal the CSM properties and the mass-loss history of the progenitor. In fact, the mass-loss histories have been revealed for a number of various types of SNe, based on radio, optical, and X-ray observations \citep[e.g.,][]{2012MNRAS.419.1515D,2013MNRAS.435.1520M,2014ApJ...797....2K}. It is also important to measure the geometry of the CSM around a SN, because a direct comparison with a CSM directly imaged around an evolved massive star can help understand the nature of the progenitor (e.g., whether or not $\eta$ Carinae can explode immediately). However, little is known about the geometry of the CSM, since extragalactic SNe are spatially unresolvable during the early phase of their evolution. This situation has been gradually changing by recent detailed spectropolarimetry of young SNe, first pointed out by \citet{1982ApJ...263..902S}. However, these measurements are mostly for SN photospheric geometries (i.e., the SN ejecta), and the CSM geometries still remain uncertain \citep[][for a rewiew]{2008ARA&A..46..433W}. CSM geometries have been revealed/inferred for only a few SNe including a remarkable example, SN~1987A, for which a complex CSM ring was directly imaged \citep[e.g.,][]{1995ApJ...452..680B}. Others include SNe~1997eg, 1998S, and 2010jl, all of which are classified as Type IIn --- a rare class of SNe comprising of $\sim$9\% of all core-collapse SNe \citep{2011MNRAS.412.1522S}, characterized by an intense narrow H$\alpha$ line, and thought to have experienced the most drastic mass-loss episodes among all types of SNe. Polarization spectra of the three SNe IIn are strikingly similar with each other in that continuum emission is wavelength-independently polarized by $\sim$2\% and line emission is depolarized \citep{2000ApJ...536..239L,2001ApJ...550.1030W,2008ApJ...688.1186H,2011A&A...527L...6P}. This result led \citet{2000ApJ...536..239L} and \citet{2008ApJ...688.1186H} to suggest a dense, disk-like or ring-like CSM surrounding aspherical SN ejecta. We should however note that the interstellar polarization can change ``valleys" into ``peaks" (or vice versa) in the polarization spectrum, as clearly demonstrated by Figure~10 in \citet{2000ApJ...536..239L}. Since it is not easy to know a correct interstellar polarization, there remain large uncertainties on the CSM geometry. Therefore, additional observational information about the CSM geometry has been awaited. We here present new evidence for asphericity in the CSM for three Type IIn SNe, SN~2005kd, SN~2006jd, and SN~2010jl, based on multi-epoch X-ray spectral observations. Essentially, in the very early-phase, we can detect a heavily-absorbed X-ray component solely, but later ($\sim$1--2\,yr after explosions), an additional moderately-absorbed X-ray component emerges. The co-existence of the two components and the delayed emergence of the moderately-absorbed component argues for asphericity, presumably a torus-like geometry, in the CSM. In Section~\ref{sec:obs_ana}, we present information about observations as well as our analyses. We give our interpretation and conclusion in Sections~\ref{sec:discussion} and \ref{sec:conclusion}, respectively. | \label{sec:conclusion} Based on multi-epoch spectral analyses of three SNe IIn, SN~2005kd, SN~2006jd, and SN~2010jl, we found that their X-ray spectra were initially explained by a single heavily-absorbed component, and then became comprised of moderately- and heavily-absorbed components during $\sim$1--2 yr after explosions. Both of the two components are most likely high-temperature thermal emission associated with the forward shock propagating into the CSM, but have distinct absorptions. This X-ray spectral evolution requires a departure of a spherical symmetry in the CSM. Specifically, a torus-like geometry of the CSM would qualitatively explain the X-ray spectral evolution. We estimated that the radius of the torus-like CSM is on the order of $\sim$5$\times10^{16}$\,cm. | 16 | 9 | 1609.09093 |
1609 | 1609.05274_arXiv.txt | \noindent {The ex-nova RR Pic presents a periodic hump in its light curve which is considered to refer to its orbital period. Analyzing all available epochs of these hump maxima in the literature, and combining them with those from new light curves obtained in 2013 and 2014, we establish an unique cycle count scheme valid during the past 50 years, and derive an ephemeris with the orbital period 0.145025959(15) days. The O - C diagram of this linear ephemeris reveals systematic deviations which could have different causes. One of them could be a light-travel-time effect caused by the presence of a hypothetical third body near the star/brown dwarf mass limit, with an orbital period of the order of 70 years. We also examine the difficulty of the problematic of detecting sub-stellar or planetary companions of close red-dwarf white-dwarf binaries (including cataclysmic variables), and discuss other possible mechanisms responsible for the observed deviations in O - C. For RR Pic, we propose strategies in order to solve this question by new observations.} | \noindent Classical Novae comprise an important subclass of cataclysmic variables (CVs), characterized by a single observed outburst of large amplitude (typically between 9 a 15 mag) which is understood as a thermonuclear runaway explosion on the surface of the white dwarf component. RR Pic is one of the brightest representatives of this class, erupting in 1925 and leaving a stellar remnant of about magnitude V $\approx$ 12 mag. Since $\sim$1960 several observers noticed a modulation in the light curve of this remnant, with a period near 3.5 hours, which was confirmed by \citet{vogt75} who presented photometric observations and derived an ephemeris of this periodic ``hump'', based on all data available at that time. Later, this ephemeris was improved by \citet{kubiak84} based on additional data. Time-resolved spectroscopic observations were published by \citet{wyckoff77} and by \citet{haefner91}; their radial velocity curves based on the dominant He {\sc ii} ($\lambda$ 4686\AA) emission line confirm that the photometric hump refers to the orbital period. Additional tomographic studies of the spectral behaviour were published by \citet{linda2003} and \citet{ribeiro2006}. \noindent \citet{linda2008} presented new time-resolved photometric data, confirming the presence of the orbital hump and detecting, in addition, a superhump with a period excess of about 8.6\% over the orbital period.\\ Due to the large time gaps in the available observations previous authors were not able to determine a unique long-term ephemeris of the orbital period. Now, the main gap (between 1982 and 2005) could partly be filled by unpublished photometric observations in the CBA archive (kindly provided by J. Patterson and collaborators from the Centre for Backyard Astrophysics). In addition, we analyze here recent new observations obtained by one of us (F.-J.H.) in 2013 and 2014. The aim of our present paper is to derive, for the first time, a unique ephemeris for the orbital hump of RR Pic valid during the last five decades (section \ref{s2}). In this context some systematic deviations in the O - C diagram were detected which are discussed tentatively as light-travel-time effect in Section \ref{s3}. In Sections \ref{s4} and \ref{s5} we compare the third-body hypothesis of RR Pic with the actual situation of other known or suspected multiple stellar systems (including binary star-planet configurations), mention possible alternative interpretations and add some conclusions. | \label{s5} \noindent Based on a rather stable periodic orbital hump in the light curve of RR Pic between 1965 and 2014, we present a unique and precise ephemeris, valid for these five decades. In addition, we detected rather strong orbital period variations. In particular, there is a significant difference between the mean period valid in 1965--1982, and that afterwards, in the sense of an increasing period value during the total time interval of five decades covered. As one of the possible interpretations of these period variations we propose a light-travel-time effect and estimate that it could be caused by a hypothetic late M type companion with M$_3$ = 0.26 $\pm$ 0.02 solar masses, completing an eccentric orbit around the CV RR Pic in $\approx$ 67 years at a mean distance from the CV of $\approx$19 AU. Among the CVs, only for the three eclipsing polars planetary candidates have been reported: UZ For \citep{potter2011}, HU Aqr \citep{schwope2011} and DP Leo \citep{beuermann2011}. Reports on triple systems in other CVs, also involving stars and/or brown dwarf components, are very rare. To our knowledge, there are only two candidates for such triple systems with CV components: the first case is FS Aur, a dwarf nova and intermediate polar \citep{Neustroev2013}. For this star, a long-term modulation in its light curve with a period of $\approx$ 900 days was detected and interpreted by means of the interaction with a third body of a mass between 25 and 64 times that of Jupiter \citep{chavez2012}. All these four cases have ultra-short orbital periods ($<$0.088 d). The second case of a possible triple system is FH Leo, a wide visual binary consisting on two F/G type main sequence stars, at an angular separation of about 8 seconds of arc. HIPPARCOS had detected some possible flares without resolving the binary. \citet{vogt2006} suggested that these could be caused by a SU UMa type dwarf nova orbiting around one of the two double star components. The latter two cases, which do not involve light-travel-time effect, are rather hypothetic, without a stringent confirmation. We would like to emphasize that our interpretation of the long-term O - C diagram of RR Pic, if confirmed, would imply, for the first time, the detection of a cataclysmic binary with a third stellar companion, and also the first candidate for a triple system CV above the period gap between 2 and 3 hours. In addition, the case RR Pic would be the first application of the light-travel-time method in a non-eclipsing binary system. This was only possible because the amplitude of this effect, in case of a stellar companion, is about one order of magnitude larger than for planetary companions. Finally, RR Pic would be the first non-magnetic CV in a triple system. Therefore, we believe that it could be important to present these results in spite of the rather hypothetical character of our interpretation. It remains an open question why just RR Pic seems to be a record holder in so many respects, in spite of its rather normal behaviour as a non-magnetic CV. There are many other CVs with orbital hump light curves, and well defined eclipses, especially those with separate ingress and egress eclipse phases of the white dwarf, as present in Z Cha, OY Car and HT Cas among others. Does none of them possess a third body? Perhaps a re-analysis of the available data on eclipsing CVs might result in further discoveries. | 16 | 9 | 1609.05274 |
1609 | 1609.00384_arXiv.txt | The advantages of angular differential imaging (ADI) has been previously untested in imaging the host galaxies of damped Lyman $\alpha$ (DLA) systems. In this pilot study, we present the first application of ADI to directly imaging the host galaxy of the DLA seen towards the quasar J1431+3952. K-band imaging of the field surrounding J1431+3952 was obtained on the Gemini North telescope with the adaptive optics system and a laser guide star. We computed a sensitivity curve that demonstrates the sensitivity of our observations as a function of K-band magnitude, impact parameter and DLA angular size. For an impact parameter of 0.5" (3.4 kpc at the redshift of the absorber) our mass sensitivity is log (M$_{\star}$/M$_{\odot}) \sim$ 9.2 and drops to $\sim$ 9.0 at separations beyond $\sim$ 6 kpc for the smallest size model galaxy. Three candidate galaxies are identified within 5". Stellar masses were computed from the K-band photometry yielding values of log (M$_{\star}$/M$_{\odot}) \sim$ 9.9, 9.7 and 11.1 respectively. The likely identification of the absorbing galaxy is discussed, and we conclude that the galaxy with the largest impact parameter and highest stellar mass is unlikely to be the host, based on its inconsistency with the N(HI) impact parameter relation and inconsistent photometric redshift. Whilst we cannot distinguish between the remaining two candidates as the DLA host, we note that despite the low spin temperature and relatively high metallicity of the DLA, the host does not appear to be a particularly luminous (high mass) galaxy. | \label{sec:intro} When a quasar's line of sight is intercepted by intervening gas, the resulting absorption lines provide information about the chemical and physical properties in the early universe. The highest column density absorbers with log N(HI) $\ge$ 20.3 are the so-called damped Lyman $\alpha$ systems (DLAs), which have proven to be powerful probes of high redshift galaxies \citep{wolfe2}. Given the slope of the column density distribution function of high redshift absorption line systems \citep{tytler}, it is the DLAs, rather than the more numerous Ly$\alpha$ forest clouds, or Lyman limit systems, that contain the bulk of the neutral gas available for star formation. The abundance of the DLA atomic gas reservoir is little changed over many Gyrs at intermediate and high redshifts \citep[e.g.][]{chrighton,sanchez-ramirez,neeleman} indicating that DLAs offer a fertile location for star formation over a significant fraction of cosmic time. Combined with the chemical enrichment associated with DLAs, which manifest a wide range of metallicities from a less than 1/100 to in excess Z$_{\odot}$ metallicity \citep[e.g.][and references therein]{berg}, we expect these high column density absorbers to represent a broad cross section of galaxies and hence provide a window into galaxy evolution at these epochs. Indeed, there is a large body of research spanning the last two decades and beyond that have studied many aspects of DLAs, ranging from their elemental ratios \citep[e.g.][]{pettini00,prochaska,dessauges-zavadsky}, molecular fraction \citep[e.g.][]{ledoux03,noterdaeme}, kinematics \citep[e.g.][]{prochaskawolfe,ledoux06,neeleman13}, ionization properties \citep[e.g.][]{vladilo01,milutinovic} and dust depletion in the interstellar medium \citep[ISM, e.g.][]{pettini94,pettini97,vladilo11,murphy}. Despite these advances, some of the most fundamental properties of the absorbing galaxies, such as their luminosities, stellar masses and morphologies remain unknown for the vast majority of the DLA population. For many years, most DLAs with identified host galaxies were at relatively low-to-intermediate redshift \citep[e.g.][]{rao, bowen,chen,lebrun}, but there are now a growing number detected at $z>2$ \citep[e.g.][]{djorgovski,weatherley,fynbo,krogager,peroux12,krogager13,kashikawa,jorgenson14,hartoog,mawatari,peroux16}. Nonetheless, the identification of wholesale numbers of host galaxies for DLAs remains one of the outstanding challenges in the field. The fundamental challenge for the identification of a DLA host galaxy is the overwhelming brightness of the quasar relative to the galaxy, which is expected to be found at low impact parameter from the QSO \citep{rao03}. A few methods have been developed to circumvent the blinding quasar light, and to allow the observer to search for galaxy hosts close to the QSO line of sight. One such method is the 'Double-DLA' technique \citep{steidel92,fum,omeara06}, in which a QSO exhibits multiple high column density absorbers. The higher redshift system is used as a natural blocking filter to eliminate (rest-frame) far ultraviolet emission from the quasar so that the continuum emission from the lower redshift DLA can be directly measured. An alternative approach has been to search for absorbers in the spectra of the optical afterglow of gamma-ray bursts (GRBs) \citep[e.g.][]{vreeswijk,chen05,prochaska07}. GRB afterglows can shine brighter than quasars, but they fade rapidly \citep{kann}. Spectra taken along the sightline of the GRB can reveal intervining DLA, sub-DLA and other absorbing systems \citep{schulze}. After the afterglow has faded, follow up imaging and spectroscopy can be used to identify the host galaxy \citep[e.g.][]{masetti, pollack,vree,schulze,ellison06,chen09,chen10}. Again, this technique is limited to a relatively small number of absorbers. Finally, spectroscopy with integral field units (IFUs) has been used to successfully identify DLA hosts at z $\sim$ $1-2$ \citep[e.g.][]{peroux,bouche}. This technique uses narrow band images generated from IFU data cubes to search for H$\alpha$ or other emission lines along quasar sightlines, and is applicable when the DLA redshift optimally places emission lines between the bright sky lines in the IR. In the pursuit of identifying DLA hosts, adaptive optics (AO) seems a natural facilitator. AO has an obvious application in the search for DLA host galaxies, thanks to the improved diffraction-limited angular resolution and deep contrast close to the quasar that can be achieved. Despite its obvious benefits, relatively few AO-aided searches for DLA hosts outside IFU useage have been attempted in the past \citep{chun06,chun10}. Part of the historical challenge of AO observations has been the limited sky coverage of natural guide stars. Moreover, AO imaging techniques require fairly complex analysis methods, to deal with a highly time and space sensitive point spread function (PSF). Fortunately, significant progress has been made developing data reduction techniques, and the introduction of laser guide star technology opens the night sky to the application of AO. In this paper we consider the application of a particular AO technique called angular differential imaging (ADI), which was developed for directly imaging exoplanets \citep{marois06}. The main technical feature of the ADI technique is the disabling of the instrument rotator during observations, so that the field-of-view rotates around a central axis during the set of exposures. Given enough field-of-view rotation, these images can be combined to create a reference PSF and suppress quasi-static speckles by up to two orders of magnitude. The reference PSF removes off-axis light which increases the sensitivity, allowing for the detection of fainter objects at lower impact parameters. Previous AO imaging of DLAs have largely used multi-step methods of azimuthal PSF subtraction \citep{chun10,chun06}. Applying ADI to AO imaging of DLAs can greatly simplify the reduction process and potentially improve the limits of detection. This paper is organized as follows: Section~\ref{methods} describes the DLA that is the target of this pilot study, including details of14 the methods of observation (\ref{obsv}), data reduction (\ref{redu}) and PSF subtraction. In Section \ref{results} we describe the determination of candidate positions and magnitudes (\ref{simul}), our calculations of stellar mass (\ref{mass}), and the detection limits of this study (\ref{curves}). Section~\ref{disc} discusses candidate DLA host galaxies in the context of known scaling relations and directions for future work. Section~\ref{conc} summarizes our conclusions. Throughout this paper, a $\Lambda$ cold dark matter cosmology is assumed with $H_0$ = 68 km s$^{-1}$ Mpc$^{-1}$, $\Omega_\Lambda$ = 0.70 and $\Omega_M$ = 0.30. | \label{conc} The first application of ADI to direct imaging of DLA galaxies has resulted in three candidates (within 5 arcsec, or 30 kpc at the redshift of the absorber) for the $z_{abs}=0.602$ DLA seen in the quasar J1431+3952. Determination of the sensitivity curve for our observations indicates that we could have detected a galaxy whose stellar mass was as low as $10^{9.4}M_{\odot}$ at a separation of 3.4 kpc, or $\sim 10^{9.0}M_{\odot}$ beyond $\sim$ 6 kpc. Based on the K-band photometry of our NIRI observations, we determine stellar masses of log (M$_{\star}$/M$_{\odot}$) $= 9.9, 9.7 $ and $ 11.1$ for the three candidates, which are located at impact parameters of 15, 17 and 30 kpc respectively. The two galaxies at the lowest impact parameters are new detections in our NIRI data. Based on a photometric redshift of $z=0.08$ (Ellison et al. 2012), the unresolved nature of the object, and inconsistency with the N(HI) -- impact parameter relation (e.g. Krogager et al. 2012; Christensen et al. 2014), we conclude that the DLA is not associated with the highest mass, largest separation object of the three candidates. The remaining two galaxies are consistent with these scaling relations and therefore remain plausible candidates for the DLA host. Follow-up spectroscopy is required to confirm the redshifts of the remaining two candidates and observations in just one additional band would allow for additional mass-luminosity constraints \citep{bell01}. Our results indicate that despite its low spin temperature, the host galaxy of this DLA is unlikely to be of high stellar mass (or luminosity). | 16 | 9 | 1609.00384 |
1609 | 1609.00667_arXiv.txt | We use \nustar\ observations of the Galactic Center to search for X-ray lines from the radiative decay of sterile neutrino dark matter. Finding no evidence of unknown lines, we set limits on the sterile neutrino mass and mixing angle. In most of the mass range 10--50\,keV, these are now the strongest limits, at some masses improving upon previous limits by a factor of $\sim10$. In the $\nu$MSM framework, where additional constraints from dark matter production and structure formation apply, the allowed parameter space is reduced by more than half. Future \nustar\ observations may be able to cover much of the remaining parameter space. | Introduction} Is dark matter composed entirely of sterile neutrinos? In the neutrino minimal standard model ($\nu$MSM~\cite{Asaka:2005an, Boyarsky:2006jm, Asaka:2006nq, Canetti:2012vf, Canetti:2012kh}) --- an economical framework that can simultaneously explain neutrino mass, the baryon asymmetry of the Universe, and dark matter --- a definitive answer is possible. Sterile neutrino dark matter can be produced through mixing with active neutrinos. In the $\nu$MSM, where the lepton asymmetry is non-zero, dark matter is produced with resonant production, also called the Shi-Fuller mechanism~\cite{Shi:1998km}. (In the limit of zero lepton asymmetry, it corresponds to non-resonant production, also called the Dodelson-Widrow mechanism~\cite{Dodelson:1993je}.) When all constraints are considered, the allowed parameter space for sterile neutrino dark matter in the $\nu$MSM is finite~(see Refs.~\cite{Boyarsky:2009ix, Boyarsky:2012rt, Adhikari:2016bei} for reviews). In Fig.~\ref{fig:intro}, we summarize the current constraints and the improvements resulting from the work presented in this paper~(detailed in Sec.~\ref{sec:results}). Astrophysical X-ray constraints are model independent and provide \emph{upper} limits on the sterile neutrino mass~\cite{Dolgov:2000ew, Abazajian:2001vt}. If the $\nu$MSM is considered, structure-formation considerations provide \emph{ lower} limits on the mass~\cite{Dodelson:1993je, Shi:1998km,Dolgov:2000ew, Abazajian:2001nj}. At smaller masses~($\lesssim 10$\,keV), there are strong limits from X-ray telescopes such as \chandra, \suzaku, and \xmm, while at larger masses~($\gtrsim 50$\,keV), there are strong limits from \integral. However, until now, it has been particularly difficult to probe masses in the range 10--50 keV, which, since radiative decay produces an X-ray line at energy $E_\gamma = m_{\chi}/2$, corresponds to X-rays of energies 5--25 keV. This has been mostly due to the lack of new instruments sensitive to the relevant X-ray energy range. \begin{table*} \caption{\label{tab:obs} \nustar\ observations used for this analysis.} \begin{ruledtabular} \begin{tabular}{ccccccc} \multicolumn{1}{c}{Observation ID} & \multicolumn{2}{c}{Pointing (J2000)\footnote{Roll angle was 332$^\circ$ for all.} } & Effective Exposure\footnote{After all data cleaning.} & Detector Area\footnote{After stray light, ghost ray, and bad pixel removal.} & Avg. Solid Angle\footnote{Average solid angle of sky from which 0-bounce photons can be detected, after correcting for removal of stray light, ghost rays, and bad pixels, as well as efficiency due to vignetting effects.} \\ \multicolumn{1}{r}{} & RA (deg) & DEC (deg) & FPMA\,/\,FPMB (ks) & FPMA\,/\,FPMB (cm$^2$) & FPMA\,/\,FPMB (deg$^2$) \B \\ \hline \multicolumn{1}{r}{40032001002} & 265.8947 & $-29.5664$ & 39.7\,/\,39.6 & 9.89\,/\,11.10 & 3.73\,/\,4.09 \T \\ \multicolumn{1}{r}{40032002001} & 265.7969 & $-29.5139$ & 39.8\,/\,39.6 & 7.14\,/\,8.05 & 4.06\,/\,4.12 \\ \multicolumn{1}{r}{40032003001} & 265.6991 & $-29.4613$ & 39.8\,/\,39.6 & 8.18\,/\,8.92 & 3.47\,/\,4.01 \\ \multicolumn{1}{r}{40032004002} & 265.9550 & $-29.4812$ & 22.6\,/\,22.7 & 4.19\,/\,6.54 & 2.34\,/\,3.13 \\ \multicolumn{1}{r}{40032005002} & 265.8572 & $-29.4288$ & 25.6\,/\,25.8 & 9.78\,/\,7.85 & 3.80\,/\,3.85 \\ \multicolumn{1}{r}{40032006001} & 265.7595 & $-29.3762$ & 28.6\,/\,28.6 & 9.98\,/\,6.18 & 3.76\,/\,3.74 \\ \end{tabular} \end{ruledtabular} \end{table*} Launched in 2012, the \emph{Nuclear Spectroscopic Telescope Array} (\nustar)~\cite{Harrison:2013md} is the first focusing optic to cover the 3--79~keV energy range. Due to its combination of grazing-incidence design and multilayer-coated reflective optics, \nustar\ provides unprecedented sensitivity in this hard X-ray band, and its focal-plane detectors deliver energy resolution of 400~eV at $E_\gamma = 10$~keV. Moreover, \nustar\ has already completed (\emph{i.}) long exposures of the Galactic Center (GC), where the dark matter decay signal is expected to be bright, as well as (\emph{ii.}) extensive modeling of the astrophysical emission components, which form a significant background to sterile neutrino searches~\cite{NuSTARGRXE}. Due to the geometry of the \nustar\ instrument, photons arriving from several degrees away from the target of observation may directly enter the detectors without passing through the focusing optics. These ``0-bounce" photons (see Sec.~\ref{sec:obs}) normally constitute a background for pointed observations. However, an innovative use of these photons is to probe large-scale diffuse emission that extends over much larger scales than the field of view (FOV) of focused photons. We exploit the wide \nustar\ solid angle aperture for 0-bounce photons to perform a sensitive search for dark matter decay in the GC region. As show in Fig.~\ref{fig:intro}, this reduces the remaining parameter space for sterile neutrino dark matter in the $\nu$MSM by about half. In Sec.~\ref{sec:nustaranalysis}, we describe the \nustar\ instrument and the dataset used in this analysis (Sec.~\ref{sec:obs}), the particular analysis procedures necessary to utilize 0-bounce photons (Sec.~\ref{sec:zerobounce}), and the energy spectrum of the GC and corresponding line-search analysis (Sec.~\ref{sec:spec}). In Sec.~\ref{sec:jfactor}, we model the expected dark matter signal, which takes into account the non-trivial shape of the aperture for 0-bounce photons. In Sec.~\ref{sec:results}, we present our results in the mass-mixing plane and put them in the context of previous constraints. Conclusions and comments on future prospects are presented in Sec.~\ref{sec:conclusion}. \begin{figure} \includegraphics[width=0.75\linewidth]{0_bounce.pdf} \caption{\label{fig:0bounce} Illustration, from Ref.~\cite{Wik:2014boa}, of the \nustar\ observatory geometry. 0-bounce photons from far off-axis sources can bypass the aperture stops and shine directly on the detectors, though some of these rays are blocked by the optics bench.} \end{figure} | Conclusions} We search for dark matter that decays into monoenergetic keV-scale photon lines using a subset of the \nustar\ Galactic plane survey data. No obvious dark matter signals are found, and thanks to the novel use of 0-bounce photons, robust and stringent upper limits are placed on the decay rate of dark matter into X-rays. Our analysis has produced the strongest indirect detection limit on dark matter lines in the energy range $E_\gamma = 5-25$\,keV. This also allows us to place strong upper limits on the mixing angle for sterile neutrino dark matter. For the $\nu$MSM, where the sterile neutrino is produced via mixing in the Early Universe, only a small section of the original parameter space remained before out work. Our results significantly reduce the available parameter space, which is likely to be completely probed by future analyses of \nustar\ observations. In the case of a null detection, it would imply the minimalistic approach of $\nu$MSM is insufficient to explain neutrino mass, baryon asymmetry, and dark matter simultaneously. It would also further heighten interest in models where sterile neutrino dark matter is produced with different mechanisms~\cite{Shaposhnikov:2006xi, Kusenko:2006rh, Merle:2013wta, Frigerio:2014ifa, Lello:2014yha, Merle:2015oja, Patwardhan:2015kga}. \bigskip {\bf Note added:} As this paper was being completed, we learned of a sterile neutrino dark matter search that also considered the use of \nustar\ 0-bounce photons, but with a different data set~\cite{Neronov:2016wdd}. Our limit is comparable, and in some cases, more stringent, than theirs. | 16 | 9 | 1609.00667 |
1609 | 1609.01363.txt | %abstract %{}{}{}{}{} % 5 {} token are mandatory % \abstract %context {} %aims {The shaping mechanisms of old nova remnants are probes for several important and unexplained processes, such as dust formation and the structure of evolved star nebulae. To gain a more complete understanding of the dynamics of the GK Per (1901) remnant, an examination of symmetry of the nova shell is explored, followed by a kinematical analysis of the previously detected jet-like feature in the context of the surrounding fossil planetary nebula.} %methods {Faint-object high-resolution echelle spectroscopic observations and imaging were undertaken covering the knots which comprise the nova shell and the surrounding nebulosity. New imaging from the Aristarchos telescope in Greece and long-slit spectra from the Manchester Echelle Spectrometer instrument at the San Pedro M\'artir observatory in Mexico were obtained, supplemented with archival observations from several other optical telescopes. Position-velocity arrays are produced of the shell, and also individual knots, and are then used for morpho-kinematic modelling with the \textsc {shape} code. The overall structure of the old knotty nova shell of GK Per and the planetary nebula in which it is embedded is then analysed.} %results {Evidence is found for the interaction of knots with each other and with a wind component, most likely the periodic fast wind emanating from the central binary system. We find that a cylindrical shell with a lower velocity polar structure gives the best model fit to the spectroscopy and imaging. We show in this work that the previously seen jet-like feature is of low velocity.} %conclusions {The individual knots have irregular tail shapes; we propose here that they emanate from episodic winds from ongoing dwarf nova outbursts by the central system. The nova shell is cylindrical, not spherical, and the symmetry axis relates to the inclination of the central binary system. Furthermore, the cylinder axis is aligned with the long axis of the bipolar planetary nebula in which it is embedded. Thus, the central binary system is responsible for the bipolarity of the planetary nebula and the cylindrical nova shell. The gradual planetary nebula ejecta versus sudden nova ejecta is the reason for the different degrees of bipolarity. We propose that the $\lq$jet' feature is an illuminated lobe of the fossil planetary nebula that surrounds the nova shell.} %The results for the nova shell are in accordance with the inclination angle of the binary progenitor system lending credence to the binary shaping mechanisms thought to be at work. Also, the position angle matches that of the extended ancient planetary nebula perceived to be associated with the system, differences in the extent of the nova shell and the planetary nebula components hint at the efficiency of the shaping procedure.} % | With a proximity of 470 pc \citep{mclaughlin,Harrison:2013aa}, GK Per (1901) is a nearby, historic, and spectacular post-nova source. As the nearest and brightest of only two classical nova remnants observed within a planetary nebula to date, the other being V458 Vul \citep{Wesson458,Roy:2012aa}, it offers the best chance to study the evolution of nova explosion debris within a planetary nebula and thus potentially aids the understanding of both types of object. A classical nova event is the result of thermonuclear runaway on the surface of a white dwarf accreting from, typically, a main sequence or a late G- or K-type star \citep{warner}. The accreted shell is ejected, at velocities ranging from 5x$10^2$ to typically $<$ 5x$10^3$ km s$^{-1}$ \citep{BodeNova}, once a critical pressure is reached at the core-envelope interface. Before ejection, mixing occurs between the accreted envelope and the white dwarf core through convection, leading to a heavy-element enrichment of the roughly solar composition envelope \citep{Casanova:2010aa}. Dwarf novae, which are also exhibited by the GK Per system, result from an instability in the accretion disk surrounding the white dwarf star; the instability is caused by a disturbance in the magnetic field of the white dwarf leading to a brightening of 2-6 magnitudes \citep{osakiDN}. These events accelerate winds to the order of 1-6x$10^3$ km s$^{-1}$ \citep{dnvelone,dnvel3,dnvel2}, generally faster than classical nova winds. Dwarf novae are important in terms of accretion disk physics as during their rise and fall the majority of the emission is from the disk \citep{osakiDN}. The onset and progression of thermonuclear runaway in the surface envelope of post-enriched accreted hydrogen (and/or helium) is very difficult to explain \citep{Casanova:2010aa}. However, detailed observations of the outflow allow for estimates of the total mass and abundances of the heavy elements ejected, as is explored in \cite{helton2011}. This provides constraints on the mechanism and efficiency of dredge-up from the underlying white dwarf. A full understanding of this process would greatly improve inputs into models for nuclear reactions that follow the thermonuclear runaway, which would improve the assessment of ejecta masses and composition \citep{Casanova:2010aa}. As noted by \cite{ederoclitethesis}, and previously by others (e.g. %\cite{optspecoldnova} Ringwald et al. 1996 and references therein), that old novae are often neglected after their explosive lightcurves have reached quiescence. An important possible consequence of these systems is that they are candidates for type Ia supernova progenitors \citep{CNassn1a} along with their recurrent counterparts [e.g. T. Crb as discussed in \cite{growingWD}]. Old nova shells give insight into dust production, clumping, and ISM dispersal mechanisms, thus they should be understood at all evolutionary phases. There are many unknowns surrounding the origin of the morphology of nova shells such as whether their ejection is spherically uniform or intrinsically bipolar, (e.g. \cite{porterasphericity,lloydshaping}). There is recent evidence in the case of V339 Del that shaping of ejected material happens very early on; non-sphericity is evident in the first days after outburst \citep{Schaefer:2014aa}. The common envelope phase is thought to play a major role in the shaping of nova remnants, and planetary nebulae alike \citep{balick_shaping,Nord}. Slower nova events (regarding both photometric evolution and ejection velocities) are believed to have stronger deviations from spherical symmetry. This hints at the importance of the role the common envelope phase plays, when the time spent by the binary in the envelope is longer than their orbital period. Shaping is also expected from interaction with prior ejecta, which would also be non-spherical owing to various similar processes to the nova ejecta. There is observational evidence of dependence of axial ratio on speed class that can be seen in Fig. 8 of \cite{slavin}. %\subsection{GK Persei} As an old, bright, and close nova shell with published kinematics, GK Per is an ideal object to study and enrich our knowledge regarding the mechanisms at work in a nova system, see Fig. 1. The 1901 GK Per nova event was possibly its first \citep{Bode04} and it has emerged as a well-studied and peculiar object. The central system has changed state, in line with hibernation theory \citep{hibernation} although sooner than the theory predicts. It was the first classical nova remnant discovered in X-rays \citep{balogel} and non-thermal radio emission \citep{Seaquist89}, implying an interaction with pre-existing material surrounding the system. \cite{bode87} first observed that a probable planetary nebula surrounds GK Per and thus is a likely candidate for the preexisting material detected in radio and X-rays. The spectacular superluminal light-echos by Wolf, Perrine and Ritchey in the years directly following the nova event (explained by \citealp{Couderc} as the forward scattering of light along dust sheets) reveal an abundance of material in the vicinity of the system. The first direct image of the nebulosity associated with the classical nova event was taken in 1916 \citep{Barnard}. GK Per has the longest period classical nova progenitor binary known to date, see Table \ref{orbchar}, with a carbon deficient secondary star \citep{carboncompanion}. After undergoing several dwarf nova outbursts, observed since 1963, the object has been reclassified as an intermediate polar, which implies the presence of a strong magnetic field (1-10 megagauss, \citealp{watson85}). The dwarf nova outbursts on GK Per were first observed in 1963 and have a recurrence timescale of 3 years and a duration of 2 to 3 months. However, strong optical outbursts can be traced back to 1948 in the AAVSO data, once the central system had settled down to its quiescent state. Two more outbursts followed the 1948 explosion in quick succession, one in 1949 and another in 1950 \citep{sab83_dngk}. The central system is seen drifting through the local environment at about 45 km s$^{-1}$, a value derived from proper motion studies \citep{Bode04}. X-ray and radio observations reveal evidence of interaction between the SW quadrant of the shell with pre-existing material \citep{Anupama:aa,balogel,gkchandra15}. The apparent box-like appearance of the nebula has long been observed (e.g. \citealt{Seaquist89}). In the past arguments for flattening of the southern part of the shell through interaction with pre-existing material would not explain the flattened northern part of the shell. There have, however, been hints of an intrinsic symmetry to the nova shell such as the prolate structure proposed in \cite{Seittermorph}. They suggested a broken prolate structure, a missing southern cap, and an extended northern cap spanning the position angles (P.A.) 130-300$^{\circ}$. \cite{lawrencefp} conducted three-dimensional Fabry-Perot imaging spectroscopy of GK Per's nova shell, where they presented channel maps and a spatial model. It had been previously believed that the expanding nova remnant was decelerating at a faster rate, based on rigorous analysis of proper motions and radial velocities of individual knots \citep{Liimets:2012aa} (hereafter L12) came to the conclusion that the system is decelerating at a rate of at least 3.8 times slower than that derived by \cite{Duerbeck}. An eventual circularisation of the shell was proposed in L12. From considerations of the data presented in L12 as well as others, such as \cite{Seittermorph,lawrencefp,tweedy} and \cite{Shara:2012aa}, the opportunity to explore the structure and kinematics of the GK Per nova shell was undertaken. Knots associated with the shell are expanding almost radially away from the central system and have been followed over the decades; high-quality spectroscopy is also available. New data were collected to complement those found in the archives. GK Per exhibits a distinct arc of emission to the NE of the nova shell that is reminiscent of a jet but whose origin has not been settled \citep{Bode04,Shara:2012aa}. This study aims to determine the relationship between the expanding knotty debris, the dwarf nova winds, the nature of the jet-like feature and the planetary nebula. \begin{table}[htbp] \caption{GK Per, characteristics of the central binary system} %\centering \begin{tabular}[width=0.3\textwidth]{lc c c c cl} \toprule Orb T\tablefootmark{1} & WD T & e\tablefootmark{2} & inc($^{\rm \circ}$)\tablefootmark{3} & $M_{\rm 1}$\tablefootmark{4} & $M_{\rm 2}$\tablefootmark{5}\\ \midrule 1.997 d\tablefootmark{a} & 351 s\tablefootmark{b} & 0.4\tablefootmark{c} $\parallel$ 1\tablefootmark{a} &50-73\tablefootmark{d}&0 .9M$_{\rm \odot}$\tablefootmark{a} & 0.25M$_{\rm \odot}$\tablefootmark{a}\\ \bottomrule \label{orbchar} \end{tabular} \tablefoot{ \tablefootmark{1}{T = Period} \tablefootmark{2}{e = ellipticity} \tablefootmark{3}{inc = inclination} \tablefootmark{4}{$M_{\rm 1}$ = White dwarf mass} \tablefootmark{5}{$M_{\rm 2}$ = Mass of companion} \tablefootmark{a}{\citep{crampton_orbper}} \tablefootmark{b}{\citep{watson85}} \tablefootmark{c}{\citep{Kraft64}} \tablefootmark{d}{\citep{Morales-Rueda:2002aa}} } \end{table} %$\hfill \break %\hfill \break$ This paper follows the structure outlined here: Observations, both archived and new, are presented in Section 2; in Section 3 the analysis of the observations is explained; a discussion follows in Section 4; and conclusions in Section 5. %\citep{payne} %\citep{type1asn} %\citep{extragal} %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % | \label{conclusions} In this work different axisymmetric models were considered in order to put fusrther effort towards the understanding of the complex morphology of the old nova shell associated with GK Per, where we find a barrel equatorial feature with polar cones to give the best fit. In the light of our new echelle spectra gathered in November 2014, March 2015, and the WISE data archives, a new hypothesis on the origin of the mysterious jet-like feature has been put forward, i.e. that it is part of the ancient planetary nebula. We propose that the wavy tails of knots in the nova shell may be due to shaping by dwarf nova winds. This allows us to derive dwarf nova wind velocities of about 4400 km s$^{-1}$, although sophisticated hydrodynamical simulations would be needed to test the robustness of this hypothesis. Based on Doppler map profiles of the overall distribution of the nova shell knots it was found that a spherical shell, warped or otherwise, does not adequately explain the red-blue spread found in observations discussed here, although a spherical shell is a good first approximation to the overall shape of the knot distribution. Instead a cylindrical form with an axial ratio close to unity is found to fit best, and also shows a remarkable resemblance to the imaging data. After the application of the cylindrical shape to the main body of the nova shell, polar features were included to account for the large number of knots not explained by the barrel. The jet-like feature is most likely part of the surrounding planetary nebula owing to its low observed velocity and structure. Following the emission lines in our new observations through the shell, they are enhanced at a lower velocity at the area of interaction as observed in radio and X-ray observations. From the spatial modelling conducted it cannot be said whether the surrounding planetary nebula is indeed bipolar or cylindrical in structure with the polar over-densities (e.g. WISE band 3) attributable to both scenarios. Deep high-resolution echelle spectroscopy is needed to decipher between the two scenarios and to then test the efficiency of the shaping mechanisms. % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % This work was supported by the following: | 16 | 9 | 1609.01363 |
1609 | 1609.07930_arXiv.txt | In protoplanetary disks the aerodynamical friction between particles and gas induces a variety of instabilities that facilitate planet formation. Of these we examine the so called `secular gravitational instability' (SGI) in the two-fluid approximation, deriving analytical expressions for its stability criteria and growth rates. Concurrently, we present a physical explanation of the instability that shows how it manifests upon an intermediate range of lengthscales exhibiting geostrophic balance in the gas component. In contrast to a single-fluid treatment, the SGI is quenched within a critical disk radius, as large as 10 AU and 30 AU for cm and mm sized particles respectively, although establishing robust estimates is hampered by uncertainties in the parameters (especially the strength of turbulence) and deficiencies in the razor-thin disk model we employ. It is unlikely, however, that the SGI is relevant for well-coupled dust. We conclude by applying these results to the question of planetesimal formation and the provenance of large-scale dust rings. | The assembly of planets is a complex and multi-faceted phenomenon that spans a gulf of some 12 orders of magnitude in length: from micron-sized dust to $10^3$ km planetary cores. It draws on an equally wide range of physical processes: collisions, dust-gas aerodynamics, gravitational collapse, instabilities, and disk structures (e.g.\ vortices, dust traps), to name but a few (Papaloizou \& Terquem 2006, Chiang \& Youdin 2010, Armitage 2010). While it is relatively straightforward to grow cm sized particles from micron sizes, further growth is potentially halted by a number of `barriers' (bouncing, fragmentation, radial-drift; Johansen et al.~2014). Statistically a small number of `lucky' aggregates may hurdle these, but certain collective instabilities can aid aggregation through this difficult size range. These include classical gravitational instability (GI; Safronov 1969, Ward \& Goldreich 1973), streaming instability (Youdin \& Goodman 2005), and the secular gravitational instability (SGI; Ward 2000, Youdin 2005). It is to the last instability that this paper is devoted. One of the most attractive features of the single-fluid SGI is that its onset is unconditional; it should always be present. Unlike classical GI, which requires the Toomre parameter to be less than one, a single fluid analysis presents no analogous restriction: the SGI works no matter how thin or thick the particle sub-disk (Youdin 2005). The instability attacks longer scales preferentially, which ordinarily would be stabilised by the Coriolis force; but particles can shed (or gain) angular momentum via areodynamical drag, and hence are not obliged to undergo stabilising epicycles. As a consequence, rings that are radially drifting towards each other continue to do so unimpeded, and the instability can proceed. On small radial scales the SGI is suppressed by dust pressure or gas turbulence, and in fact, for well-coupled dust, turbulence decreases growth rates to potentially insignificant levels (Shariff \& Cuzzi 2011, Youdin 2011). Marginally coupled particles, however, could still be subject to respectable SGI growth rates at certain radii. The SGI has been thoroughly explored in single fluid models, which are applicable when the dust to gas density ratio is tiny (e.g.\ Ward 2000, Youdin 2005, Shariff \& Cuzzi 2011, Youdin 2011, Michikoshi et al.~2012). These models assume that the angular momentum bestowed onto, or removed from, the gas disk is negligible. On sufficiently long scales, however, both sides of this momentum transaction must be included and the gas dynamics explicitly calculated. An instability criterion then appears: in a two-fluid model the onset of SGI is no longer unconditional. Recently, Takahashi \& Inutsuka (2014, hereafter TI) made a start on this problem (see also Shadmehri 2016 and Takahashi \& Inutsuka 2016), but there is still much to be established. Putting aside the issue of growth rates, an especially important question is: at what radii and for what particle sizes should we expect SGI to exist at all? The first aim of this paper is to derive clean stability criteria for the SGI. In the limits of strongly-coupled and weakly-coupled particles these can be formulated analytically and involve a variety of parameters, including the gas's Toomre parameter and the dust-to-gas density ratio. Because they bypass the SGI's full 6th order dispersion relation, these criteria make it relatively easy to assess its prevalence. The criteria also motivate a straightforward physical picture of instability in a two-fluid system. In order for the instability to work, there must exist an intermediate range of lengthscales gas upon which (a) dust pressure or turbulent mass diffusion is subdominant, and (b) the gas is prevented from executing epicycles, despite its angular momentum transactions with the dust. Going to lengthscales longer than the dust pressure (or diffusion) scale takes care of the first restriction. But the second can only be satisfied if geostrophic balance holds in the gas fluid, and so we must simultaneously find shortish scales upon which gas pressure is dominant. The existence or not of this intermediate range furnishes us with the stability criterion. The formalism is applied to realistic disk models, where we find that it is unlikely that well-coupled dust is unstable to the SGI at any radius, unless the background turbulence is especially weak. Marginally coupled particles, however, can achieve appreciable growth rates in certain circumstances, emphasising that the SGI could help aggregation of solids of cm size. We conclude, however, that SGI is probably unrelated to the dust rings recently observed by ALMA (Brogan et al.~2015). The paper will be organised in the following way. First, in Section 2, we present the two-fluid razor-thin disk model that we employ, alongside a critical discussion of its shortcomings. The main parameters of the analysis will also be defined. In Section 3 we revisit the single-fluid model to fix some ideas and to provide context for the subsequent analysis, while in Section 4 we briefly treat a simple two-fluid system where the gas is regarded as incompressible. The main results of the paper are in Section 5, in which we derive analytic stability criteria in relevant limits that are then applied to realistic disk models in Section 6. We draw our conclusions in Section 7, where we discuss the relevance of the SGI in planet and structure formation in protostellar disks. | In this paper we have explored the secular gravitational instability (SGI) using a simple two-fluid model. Despite the complexity of its associated sixth order dispersion relation, analytic stability criteria and growth rates can be obtained in the two limits of weakly and strongly coupled particles. We find that on sufficiently long and short radial scales the SGI is stabilised; the existence of an unstable range of intermediate scales leads to an explicit instability condition involving the gas's Toomre parameter, a distinctive feature of the two-fluid SGI, as opposed to the single-fluid version. The mathematical analysis suggests a straightforward way to understand the instability mechanism. The SGI favours intermediate scales upon which stabilising dust pressure or turbulence is weak, but upon which the gas pressure is strong. The latter condition permits the gas to fall into geostrophic balance: hence when the gas is azimuthally accelerated by the dust drag, it will form a zonal flow rather than undergo epicycles that would disrupt the radially collapsing dust. An assessment of the prevalence of SGI in real disk models is handicapped by uncertainties in two parameters, the strength of the turbulence $\alpha_g$, and the mass ratio of a certain species of dust to the gas within the dust subdisk, $\delta$. Starting with a fiducial value of $\delta=0.01$, we find that a moderate level of turbulence $\alpha_g=10^{-5}$ prohibits the SGI on most radii, and when it does occur it grows too slowly $\sim 10^5$ years --- the timescale of the large-scale evolution of the disk, and of appreciable radial drift. Weaker turbulence $\alpha_g=10^{-6}$ permits growth for cm sized particles on radii $\gtrsim 10$ AU, with efolding times of a few $10^4$ years. Smaller sized particles may be subject to SGI but grow too slowly. It is only for $\alpha_g=10^{-7}$ that mm sized particles sustain growth at reasonable levels, and then only for $R> 10$ AU. Increasing $\delta$ improves the situation, of course, and $\delta>0.01$ might be the case for particularly well-settled and populous subclasses of particle, though further work is needed to better constrain this parameter. Even so, if $\alpha_g>10^{-6}$ it may be prove difficult for the SGI to meaningfully impose itself on the disk dynamics. We also discuss the various shortcomings of the razor-thin model we employ, which is especially a problem when the dust and gas disks exhibit different scale thicknesses. These issues no doubt impact quantitatively on our results, but the main qualitative conclusions and our picture of instability should be robust. They can be checked with a suitable vertically stratified analysis akin to Mamatsashvili \& Rice (2010) and Lin (2014), which will also provide more reliable quantitative estimates on the stability curves and growth rates. Our results extend previous analyses of the SGI, and for larger radii are in relative agreement with Youdin (2011) and Shariff \& Cuzzi (2011). A notable difference is that the two-fluid model prohibits SGI on radii less than a critical radius. As a result, the SGI is certainly unviable on radii $< 1 $ AU, and possibly absent on radii $<10$ AU, the expected regions of planet formation. The prospects for SGI in the cm class of particles on disk radii $\sim 10$ AU are reasonable as long as gas turbulence is not too efficient $\alpha_g \lesssim 10^{-6}$. The instability could then be an important route by which large aggregates could form further out, leapfrogging the entire range of difficult cm to km sizes. Note that our results are only for axisymmetric instability. It is likely, via analogy with classical GI, that non-axisymmetric SGI occurs for larger $Q_g$, in which case our stability curves may need some revision. It has been hypothesised that the SGI generates observed dust ring structures at larger radii in protoplanetary disks (TI). As discussed in Section 6, however, the SGI has great difficulty on radii $\gtrsim 10$ AU for small particle less than a cm in size. The dust-to-gas ratio $\delta$ needs to be increased, and $Q_g$ taken to levels approaching 1 in order to obtain instability. While it may be possible to justify increasing $\delta$, such a low $Q_g$ would mean the gas disk is marginally unstable to classical GI. Perhaps a more important point is that, while the linear phase of the SGI evolution is axisymmetric, its nonlinear phase will most likely involve a non-axisymmetric breadown into disordered flow, as in classical GI, not the formation of large-scale quasi-steady rings. Dedicated nonlinear simulations are required to test what dynamics the SGI exhibits once it reaches nonlinear amplitudes, and how readily it forms planetesimal clumps. This forms the basis of future work. | 16 | 9 | 1609.07930 |
1609 | 1609.05618_arXiv.txt | We present \textit{K}-band integral field spectroscopic observations towards 17 massive young stellar objects (YSOs) in the low metallicity Small Magellanic Cloud (SMC) and two YSO candidates in the compact H\,{\sc ii} regions N81 and N88\,A (also in the SMC). These sources, originally identified using \textit{Spitzer} photometry and/or spectroscopy, have been resolved into 29 \textit{K-}band continuum sources. By comparing Br$\gamma$ emission luminosities with those presented for a Galactic sample of massive YSOs, we find tentative evidence for increased accretion rates in the SMC. Around half of our targets exhibit emission line (Br$\gamma$, He\,{\sc i} and H$_2$) morphologies which extend significantly beyond the continuum source and we have mapped both the emission morphologies and the radial velocity fields. This analysis also reveals evidence for the existence of ionized low density regions in the centre outflows from massive YSOs. Additionally we present an analysis of optical spectra towards a similar sample of massive YSOs in the SMC, revealing that the optical emission is photo-excited and originates near the outer edges of molecular clouds, and is therefore consistent with a high mean-free path of UV photons in the interstellar medium (ISM) of the SMC. Finally, we discuss the sample of YSOs in an evolutionary context incorporating the results of previous infrared and radio observations, as well as the near-infrared and optical observations presented in this work. Our spectroscopic analysis in both the \textit{K-}band and the optical regimes, combined with previously obtained infrared and radio data, exposes differences between properties of massive YSOs in our own Galaxy and the SMC, including tracers of accretion, discs and YSO--ISM interactions. | \begin{figure*} \begin{center} \includegraphics{SMC_8um_300dpi_20160428_2} \caption{SAGE-SMC 8.0\,$\mu$m mosaic of the Small Magellanic Cloud \citep{Gordon2011} showing the positions of all 33 spectroscopic massive YSOs presented in \citet{Oliveira2013} with the addition of sources \#35 and 36. The sources observed with SINFONI are in green with the remaining sources marked in red. The H\,{\sc ii} regions N81 (\#36) and N88 (\#35) are also labelled.} \end{center} \end{figure*} The formation of massive stars relies on the balance between mass loss and accretion (see \citealt{Zinnecker2007} for a review). Both processes are heavily dependent on radiation pressure \citep{Krumholz2012}, the effects of which depend on the dust content of the circumstellar medium. The metallicity of massive star forming regions is an important parameter to consider; the heat dissipation necessary to develop dense pre-stellar cores relies on radiation via the fine-structure lines of carbon and oxygen and the rotational transitions of metallic molecules such as water and CO. Dust also plays a key role in massive star formation, driving the chemistry of molecular clouds, forming self-shielding dusty discs, and allowing continued accretion to form stars in excess of 10\,M$_{\sun}$ (Kuiper et al. 2010). It is therefore expected that massive star formation is influenced by the initial conditions of the natal molecular cloud. However the effects of metallicity on massive star formation remain poorly understood with the majority of studies being carried out in Galactic, approximately solar metallicity environments. As a nearby ($\sim$60 kpc) gas-rich galaxy, the Small Magellanic Cloud (SMC) provides a valuable opportunity to study the formation of stars at lower metallicity ($Z_{\text{SMC}}\approx 0.2$\,Z$_{\sun}$; \citealt{Peimbert2000}) than in the Milky Way on both the scale of an entire galaxy and on the scale of individual stars. The metallicity of a star forming region also has a profound effect on the structure and porosity of the interstellar medium (ISM). The lower dust abundance in low metallicity environments leads to a porous ISM and thus a larger mean free path length of UV photons \citep{Madden2006,Cormier2015,Dimaratos2015} meaning that UV photons can permeate over longer distances and excite gas which lies farther from the source of the excitation. This in turn will have a significant effect on the feedback mechanisms we observe towards massive star forming regions as well as the formation of the stars themselves. Furthermore, the CO emission in the SMC is thought to arise only from high density clumps \citep{Rubio2004} and the ISM in the similarly low metallicity galaxy NGC 1569 appears to be very clumpy \citep{Galliano2003}. The \textquotedblleft\textit{Spitzer} Survey of the Small Magellanic Cloud\textquotedblright (S$^{3}$MC, \citealt{Bolatto2007}) and the \textit{Spitzer} \textquotedblleft Surveying the Agents of Galaxy Evolution in the Tidally Stripped, Low Metallicity Small Magellanic Cloud\textquotedblright (SAGE--SMC, \citealt{Gordon2011}) survey have allowed the identification of a sample of YSO candidates in the SMC, \citet{Sewilo2013} identifying 742 high-reliability YSOs and 242 possible YSOs based on colour--magnitude cuts, visual inspection of multi-wavelength images and spectral energy distribution (SED) fits to near-IR and mid-IR photometry. A sample of 33 massive YSOs has been spectroscopically confirmed and studied by \citet{Oliveira2013}; a brief overview of this analysis is provided in Section 2. In order to determine whether massive YSO properties vary between the Milky Way and the Magellanic Clouds, we will compare our results for massive YSOs in the SMC with those of the Red MSX\footnote{Midcourse Space Experiment \citep{Egan2003}} Source Survey (RMS Survey; \citealt{Lumsden2013}). The RMS survey provides the most comprehensive catalogue of massive YSOs and Ultra-Compact H\,{\sc ii} regions (UCH{\sc ii}s) to date. \textit{H}- and \textit{K}-band infrared spectroscopy of a large sample of these objects has been carried out by \citet{Cooper2013} which we will use as a Galactic dataset to compare our results to. A similar analysis for a small sample of Large Magellanic Cloud YSOs in LHA120-N113 was performed in \citet{Ward2016} and is also included in our comparison. In this work we present the \textit{K-}band integral field spectroscopic observations of 17 previously confirmed YSOs and two YSO candidates in the compact H\,{\sc ii} regions N81 and N88\,A in the SMC, obtained with SINFONI (Spectrograph for INtegral Field Observations in the Near Infrared; \citealt{Eisenhauer2003}) at the European Southern Observatory's (ESO) Very Large Telescope (VLT). We also present an analysis of previously obtained optical long slit spectroscopic observations of the H\,{\sc ii} regions surrounding our \textit{K}-band targets (originally presented in \citealt{Oliveira2013}), obtained with the Double-Beam spectrograph (DBS) at the Australian National University Telescope. New spectroscopic data for N88\,A and N81 from the Robert Stobie Spectrograph (RSS) at the Southern African Large Telescope (SALT) are also presented. Section 2 summarises the results of previous observations towards our sources and Section 3 describes the SINFONI observations presented in this work and the data reduction procedure. The results of the SINFONI observations are presented in Section 4, with an analysis of the optical spectra in Section 5. Discussion and conclusions follow in Sections 6 and 7. | We have performed \textit{K}-band SINFONI observations for 17 of the 33 massive YSOs in the SMC presented in \citet{Oliveira2013}, as well as two additional targets in N81 and N88\,A, revealing a wide range of spectroscopic properties. We have also analysed optical spectra for 28 of the 33 massive YSO targets of \citet{Oliveira2013} plus sources 35 (N88\,A) and 36 (N81). \,We summarise our conclusions below. \begin{enumerate} \item 14 of the 17 \textit{Spitzer} YSOs in the SMC observed with SINFONI have been resolved into single \textit{K}-band continuum sources whilst three have been resolved into multiple components in the \textit{K}-band. Source \#35 (N88\,A) has been resolved into the previously known 4 continuum components. N81 (\#36), previously resolved into two components by \citet{Heydari-Malayeri2003}, has now been resolved into five \textit{K}-band continuum sources. \item Visual extinctions have been calculated towards all \textit{K}-band sources except source \#04. We find a median visual extinction towards the SMC sources of $A_V =$ 14 mag. The median for the Galactic sources from \citet{Cooper2013} is 44 mag and that for the three YSOs in N113 in the LMC is 20 mag \citep{Ward2016}. This seems to suggest that there is a correlation between extinction towards YSOs and metallicity as $Z_{\text{LMC}}\sim$0.5 Z$_{\sun}$ and $Z_{\text{SMC}}\sim$0.2 Z$_{\sun}$. \item Whilst consistent with the \citet{Cooper2013} sample, Br$\gamma$ luminosities are high compared with those in the Milky Way, suggesting accretion rates which are on average higher than for the Galaxy. \item He\,{\sc i} line luminosities are comparable to the Milky Way sample, indicating that the excess observed in the Br$\gamma$ emission is unlikely to be related to a stronger wind component. \item Average velocities have been measured towards the \textit{Spitzer} YSO sources of 169$\pm$5 km s$^{-1}$ and 173$\pm$4 km s$^{-1}$ for Br$\gamma$ emission and H$_2$ emission, respectively. The Br$\gamma$ emission line velocities are consistent with those of the H\,{\sc i} radio observations, indicating that the average motions of the YSOs are bound to the bulk motions of the ISM. \item The majority of the \textit{K}-band continuum sources (11/20; excluding observations towards the H\,{\sc ii} regions N81 and N88\,A) fall within the unresolved, ultra-compact regime. These sources exhibit a variety of spectral properties and include one probable UCH{\sc ii} region (\#26), a possible Herbig Be star (\#20) along with spectroscopically typical massive YSOs and two apparently featureless continuum sources (\#22\,B and 28\,B) which may not actually be YSOs. \item Seven sources with extended H$_2$ emission morphologies indicative of outflows have been identified, with velocity gradients measured towards four of these. Sources \#22\,A and 36\,B exhibit striking examples of bipolar outflow morphologies. Our observations of the well-studied H\,{\sc ii} region N81 are the first to identify the bipolar outflow originating from source \#36\,B. \item Source \#28\,A exhibits Br$\gamma$ and H$_2$ emission line morphologies indicative of an ionized cavity in the centre of the outflow. Source \#03 appears to be an extremely broadened (and possibly more evolved) example of the same structure. These structures are possibly the result of the broadening of outflows in massive YSOs predicted by the models of \citet{Kuiper2015} and \citet{Tanaka2015}. \item CO bandhead emission (commonly used as a tracer of discs) has only been detected towards one source (\#03), a detection rate of around one third of that towards massive YSOs in the Milky Way \citep{Cooper2013} for the same range of luminosities. This could be due to either the low gas-phase CO abundance of the SMC \citep{Leroy2007}, or conditions within protostellar discs which differ significantly (such as higher temperatures or less shielding from dust) from those in the Milky Way, leading to a higher rate of CO destruction. CO absorption is somewhat tentatively observed towards one source (\#02\,B) which could be indicative of a dusty circumstellar environment and possibly suggests a continuum source completely obscured by an edge--on disc. \item Optical emission towards all sources (where it is present) appears to originate from the outer edges of the molecular clouds as the average extinction measured towards the optical emission is significantly lower than that towards the \textit{K}-band emission for the same targets. The optical emission is photo-excited and is therefore unlikely to be produced by the interaction of outflows and winds with the ISM. This scenario is consistent with relatively large mean free path of high energy photons from the protostar and through the ISM, leading to a large number of UV photons leaking from the YSOs. This could occur due to a clumpy and/or torn-up circumstellar medium, consistent with expectations for a low metallicity ISM \citep{Madden2006,Cormier2015,Dimaratos2015}. \end{enumerate} To conclude, we have presented the first study of massive YSOs in the SMC using integral field spectroscopy, studying line emission and resolving a number of our targets into multiple components in the \textit{K}-band for the first time. Through comparison with existing data in the optical, infrared and radio regimes, we have been able to develop a greater understanding of these objects within an evolutionary context. Following from the work on ice chemistry by \citet{Oliveira2013}, we continue to expose differences between YSO properties in our own Galaxy and the low metallicity SMC, namely those which concern accretion, potential disc properties and the YSO--ISM interplay. | 16 | 9 | 1609.05618 |
1609 | 1609.08606_arXiv.txt | { I discuss the spectral energy distribution (SED) of all blazars with redshift detected by the {\it Fermi} satellite and listed in the 3LAC catalog. I will update the so called ``blazar sequence" from the phenomenological point of view, with no theory or modelling. I will show that: i) pure data show that jet and accretion power are related; ii) the updated blazar sequence maintains the properties of the old version, albeit with a less pronounced dominance of the $\gamma$--ray emission; iii) at low bolometric luminosities, two different type of objects have the same high energy power: low black hole mass flat spectrum radio quasars and high mass BL Lacs. Therefore, at low luminosities, there is a very large dispersion of SED shapes; iv) in low power BL Lacs, the contribution of the host galaxy is important. Remarkably, the luminosity distribution of the host galaxies of BL Lacs are spread in a very narrow range; v) a simple sum of two smoothly joining power laws can describe the blazar SEDs very well. } \keyword{galaxies: active; galaxies: jets; gamma--rays: general; radiation mechanisms: non--thermal; radiation mechanisms: thermal} \begin{document} | About 10\% of Active Galactic Nuclei have relativistic jets, whose emission is strongly boosted. When pointing at us, these jetted sources are called blazars. Blazars come in two flavours: BL Lacs, with weak or absent broad emission lines, and Flat Spectrum Radio Quasars, with strong broad emission lines. The non--thermal spectral energy distribution (SED) produced by the jet of blazars has two broad humps, peaking in the IR--X--ray band and in the MeV--TeV band. Often (but not always) fluxes in different bands vary in a coordinated way, suggesting that most of the SED is produced by the same electrons in a specific zone of the jet. Since this region must be compact, to account for the observed very fast variability, it cannot produce the radio emission, which is strongly self--absorbed, at all but the shortest radio wavelengths (sub--mm). Other, larger, regions must be responsible for the observed flat radio spectrum. | The left panel of Fig. \ref{anal} shows the analytical, phenomenological SED for the 5 luminosity bins. The different power law segments are labelled. The right panel compares these analytical approximations with the data. The detailed procedure will be fully described in a forthcoming paper, in the following I will only present briefly the main results and conclusions. \begin{itemize} \item There still is a blazar sequence, with the same overall properties of the ``1.0" version: the SED becomes redder, and the Compton dominance increases, as the total luminosity increases. \item In a sizeable fraction of FSRQs, the accretion disk becomes visible, through an upturn of the IR--optical spectrum. Pure data show that the accretion luminosity is related to the observed, beamed, non--thermal luminosity. \item On average, the Compton dominance in powerful blazars is slightly smaller than in F98. This is the main difference with the old sequence, and it is fully understood: the increased sensitivity of {\it Fermi}/LAT allows exploration of more ``normal" blazars, and not only the most luminous. As a consequence, the average $\gamma$--ray luminosity is less than in F98. This explains the puzzling result of Giommi \& Padovani (2015) when synthetizing the blazar contribution to the $\gamma$--ray background. They found that assuming the blazar radio luminosity function and the same $\gamma$--ray to radio luminosity ratio as in the original F98 blazar sequence, one overestimates the background, especially at high energies. With the new sequence, the problem is solved (Bonnoli et al. in prep). \item The $\gamma$--ray spectrum becomes steeper as the $\gamma$--ray luminosity increases. On the other hand, low power BL Lacs do not show, on average, a very hard high energy spectrum (rising in $\nu L\nu$), in the 0.1--100 GeV band, as was the case in F98 (based on only 3 sources...). \item At intermediate $\gamma$--ray luminosities, red FSRQs and blue BL Lacs share the same $\gamma$--ray luminosity. This is explained by a difference in black hole masses. \item One can define the average synchrotron ($\nu_{\rm S}$ ) and inverse Compton ($\nu_{\rm C}$) peak frequencies for each luminosity bin, but perhaps it is better to define them separately for BL Lacs and FSRQs. This is done in Fig. \ref{vpeak} (blue squares: BL Lacs, red circles: FSRQs) showing $\nu_{\rm S}$ and $\nu_{\rm C}$ as a function of $L_\gamma$. Both BL Lacs and FSRQs form a sequence, much more pronounced in the BL Lac case. \item The smallest $\nu_{\rm S}$ is $\sim$10$^{12}$ Hz, coincident with $\nu_{\rm t}$ of the most compact component. This suggests that there should be even more powerful blazars, with the real synchrotron peak hidden by self--absorption. These blazars should have $\nu_{\rm C}\sim$ 1 MeV (i.e. at or below $10^{20}$ Hz) with a steep spectrum above. In this case the flux in the 0.1--100 GeV band could become undetectable by {\it Fermi}/LAT, and thus be not represented in the new blazar sequence. However, they should be detectable in hard X--rays, and indeed in the {\it Swift}/BAT 3 year survey (Ajello et al. 2009) we find 10 very powerful blazars with $z>2$, and 5 of them are still not detected by {\it Fermi}/LAT. These are shown in the right panel of Fig. \ref{vpeak}. \end{itemize} | 16 | 9 | 1609.08606 |
1609 | 1609.01511_arXiv.txt | {% With the Auger Engineering Radio Array (AERA) located at the Pierre Auger Observatory, radio emission of extensive air showers is observed. To exploit the physics potential of AERA, electric-field amplitude measurements with the radio detector stations need to be well-calibrated on an absolute level. A convenient tool for far-field calibration campaigns is a flying drone. Here we make use of an octocopter to place a calibrated source at freely chosen positions above the radio detector array. Special emphasis is put on the reconstruction of the octocopter position and its accuracy during the flights. The antenna response pattern of the radio detector stations was measured in a recent calibration campaign. Results of these measurements are presented and compared to simulations. It is found that measurements and simulations are in good agreement. } | \label{intro} To do high precision measurements of ultra-high energy cosmic rays with the Auger Engineering Radio Array (AERA) \cite{SchulzProcICRC2015, GlaserArena2016} at the Pierre Auger Observatory \cite{Aab_OverviewAuger15}, the radio detector needs to be well-calibrated. The largest challenge in performing such a calibration is to place a signal source at every direction relative to the detector station and at least 20\,m away to guarantee far-field conditions. Different calibration methods have been developed and already performed at AERA \cite{Abreu_aeracalib2012} and other experiments \cite{2015ApelLOPES_improvedCalibration, TunkaRex_NIM_2015, NellesLOFAR_calibration2015}. In this contribution, we present a novel and very flexible method using a flying drone to carry the signal source. A drone is a very useful tool, since it can be placed at any point above the array and can carry any piece of equipment as long as it is not too heavy. With the presented method, a higher accuracy can be achieved than in the previous calibrations mostly due to a larger set of measurements and a better handling of environmental conditions. At AERA, two different antenna types are used. The core consists of a dense array of 24 stations equipped with log-periodic dipole antennas (LPDA). They consist of nine separate dipoles optimized for a frequency range of 30 to 80 MHz. The other stations are equipped with butterfly antennas. However, in this proceeding, only LPDA stations are considered. Extensive air showers produce short radio pulses. These pulses then induce a signal measurable as voltages at the antenna output. The value which relates the Fourier transformed electric field $\mathcal{\vec{E}}$ to the measured Fourier transformed voltage $\mathcal{U}$ is called the vector effective length (VEL) $\vec{H}$: \begin{equation} \mathcal{U}(f, \theta, \phi) =\vec{H}(f, \theta, \phi) \cdot \mathcal{\vec{E}}(f). \end{equation} $\vec{H}$ depends on the frequency of the signal $f$ and the incoming direction, described by the zenith angle $\theta$ and the azimuth angle $\phi$. In spherical coordinates, $\vec{H}$ is a superposition of a horizontal $H_\phi$ and a perpendicular meridional $H_\theta$ component. The absolute value of these components is determined in a transmission measurement using a calibrated transmitting antenna. For the absolute value of a VEL component at a specific frequency and direction the following holds \begin{equation} |H_{i}| = \sqrt{\frac{4 \cdot \pi \cdot Z_A}{Z_0}} \cdot R \cdot \sqrt{\frac{P_{r, i}}{G_t \cdot P_t}}, \label{eq:absH} \end{equation} where $i = \theta, \phi$ denotes the VEL component, $Z_A$ denotes the impedance of the antenna, \hbox{$Z_0 = 120 \,\pi\,\Omega$} is the vacuum impedance, $R$ describes the distance between the antenna under test (AUT) and the transmitting antenna, $P_r$ denotes the power received at the AUT, and $G_t$ and $P_t$ are the gain and power of the transmitting antenna \cite{WeidenhauptPHD2015}. | A calibration campaign of the LPDA antennas at AERA using an octocopter and a calibrated reference source has been performed. To determine the position of the reference source during the flight, a new optical reconstruction method has been developed and evaluated. The absolute vector effective length, which describes the antenna response pattern to an incoming electric field, has been measured with an uncertainty of 9.3\%. A good agreement between simulation and measurement was found. A journal article about the whole calibration process is currently in the works. This calibration is a crucial component in the determination of an absolute energy scale of cosmic rays from first principles. Using the new calibration presented here, the uncertainty of the cosmic-ray energy measurements can be lowered with regard to current values \cite{Aab_energyaera16, Aab_energyaera_prd2}. In the future, these results from AERA will be used to improve the measurements of the whole Pierre Auger Observatory \cite{GlaserArena2016}. | 16 | 9 | 1609.01511 |
1609 | 1609.08576_arXiv.txt | We use recently published redshift space distortion measurements of the cosmological growth rate, $f\sigma_8(z)$, to examine whether the linear evolution of perturbations in the $R_{\rm h}=ct$ cosmology is consistent with the observed development of large scale structure. We find that these observations favour $R_{\rm h}=ct$ over the version of $\Lambda$CDM optimized with the joint analysis of {\it Planck} and linear growth rate data, particularly in the redshift range $0< z < 1$, where a significant curvature in the functional form of $f\sigma_8(z)$ predicted by the standard model---but not by $R_{\rm h}$$=$$ct$---is absent in the data. When $\Lambda$CDM is optimized using solely the growth rate measurements, however, the two models fit the observations equally well though, in this case, the low-redshift measurements find a lower value for the fluctuation amplitude than is expected in {\it Planck} $\Lambda$CDM. Our results strongly affirm the need for more precise measurements of $f\sigma_8(z)$ at all redshifts, but especially at $z < 1$. | The large-scale structure revealed by the distribution of galaxies is believed to have formed through a process of gravitational instability, starting with primordial fluctuations in the early Universe. But while self-gravity amplifies perturbations in the cosmic fluid, cosmic expansion suppresses them. Their growth rate therefore depends rather sensitively on the dynamical expansion of the Universe and may be used to discriminate between different models. Measurements of the growth rate tend to focus on infall motions associated with condensing regions, with peculiar velocities largely correlated with the local gravitational potential. Galaxies trace these motions, carrying an imprint of the changing growth rate as the Universe evolves (Peacock et al. 2001; Viel et al. 2004; Jain and Zhang 2007; Ross et al. 2007; da Angela et al. 2008; Guzzo et al. 2008; Song and Koyama 2009; Song and Percival 2009; Davis et al. 2011; Hudson and Turnbull 2012; Macaulay et al. 2013; Alam et al. 2016). A principal statistical technique used to measure the growth rate is based on the redshift-space distortion (RSD) created by the galaxies' peculiar velocities (Kaiser 1987). Specifically, maps produced with distances inferred from redshifts in spectroscopic galaxy surveys show that the galaxy distributions are anisotropic due to the fact that the redshifts contain components from both the smooth Hubble flow and the peculiar velocities of the infalling matter. As long as one can reliably separate these two contributions to the redshift, one may thereby extract a history of the build-up of structure. The use of RSD, however, is practical primarily before non-linear effects begin to emerge, where measurements yield information on both the matter over-density and the peculiar velocities of galaxies. In the linear regime, the problem is typically reduced to solving a second-order differential equation for the time-dependent fluctuations, from which one may then infer their growth rate. Thus, although objects suitable for this work can in principle include individual galaxies, clusters of galaxies, and superclusters, their density contrasts ($\delta\rho/\rho$) today are, respectively, $\sim 10^6$, $\sim 10^3$, and $\sim O(1)$. These estimates assume a critical density $\rho_c\equiv 3c^2H_0^2/8\pi G\sim 10^{-29}$ g cm$^{-3}$, and that about $10^{12}$ $M_\odot$ of galactic matter are contained within a $\sim 30$ kpc region. Clusters typically contain fewer than $10^3$ galaxies, while superclusters have tens of thousands of galaxies. In $\Lambda$CDM, $\rho\sim (1+z)^3$ during the matter-dominated era, so galaxies and clusters presumably ceased growing linearly, i.e., grew with $\delta\rho< 1$, at redshifts $\sim 100$ and $\sim 10$, respectively. These structures are non-linear in the local neighborhood. The linear-growth analysis described in this paper therefore tends to address the formation of super-clusters, which could have grown linearly over the redshift range $z < 2-3$. This is the approach we shall follow in this paper, and therefore focus on surveys relevant to these large structures, including 2dFGRS (Peacock et al. 2001) and VVDS (Guzzo et al. 2008). By now, the development of linear perturbation theory is quite mature---at least within the context of the standard model ($\Lambda$CDM), in which dark energy corresponds to a cosmological constant $\Lambda$. Comparative tests using the growth of linear structure have been carried out between $\Lambda$CDM and an assortment of other cosmologies, principally those based on extensions to general relativity involving higher-order curvature terms or extra dimensions (Wetterich 1995; Amendola 2000; Dvali et al. 2000; Carroll et al. 2004; Capozziello et al. 2005). However, measurements of the linear growth rate have not yet been used to test the $R_{\rm h}=ct$ universe (Melia 2007; Melia and Shevchuk 2012; Melia 2016a, 2017), another FRW cosmology, which has thus far been shown to fit many other kinds of data better than $\Lambda$CDM. (A brief summary of these previously published results is provided in \S~3.1 below.) The principal aim of this paper is to address this deficiency. We are especially motivated to carry out this analysis by recent comparative studies using the Alcock-Paczy\'nski test (Alcock and Paczy\'nski 1979), based the changing ratio of angular to spatial/redshift size of (presumed) spherically-symmetric source distributions with distance (Melia and L\'opez-Corredoira 2016). The use of this diagnostic, with newly acquired measurements of the anisotropic distribution of BAO peaks from SDSS-III/BOSS-DR11 at average redshifts $\langle z\rangle=0.57$ and $\langle z\rangle=2.34$, has allowed us to determine the geometry of the Universe with unprecedented accuracy. Previous applications of the galaxy two-point correlation function to measure a redshift-dependent scale that could be used to determine the ratio of angular (i.e., transverse) size to redshift (i.e., radial) size were limited by the need to disentangle the acoustic scale in redshift space from redshift distortions from internal gravitational effects (L\'opez-Corredoira 2014). A major limitation of this process was that inevitably one had to pre-assume a particular cosmological model, or adopt prior parameter values, in order to estimate the possible confusion between the true cosmological redshift interval from one edge of the cluster to the other and the contribution to this redshift width from these internal gravitational effects. Unfortunately, the wide range of possible distortions for the same correlation-function shape resulted in very large errors associated with the BAO peak position and hence the inferred acoustic scale (often in the $\sim 20-30\%$ range). This situation has improved significantly over the past few years with (1) the use of reconstruction techniques (Eisenstein et al. 2007; Padmanabhan et al. 2012) that enhance the quality of the galaxy two-point correlation function, and (2) the use of Ly-$\alpha$ and quasar observations to more precisely determine their auto- and cross-correlation functions, allowing the measurement of BAO peak positions to better than $\sim 4\%$ accuracy (Cuesta et al. 2016). The most recent determination of $y(z)$ has been based on the use of three BAO peak positions: the measurement of the BAO peak position in the anisotropic distribution of SDSS-III/BOSS DR12 galaxies at $\langle z\rangle=0.32$ and $\langle z\rangle=0.57$ (Cuesta et al. 2016), in which a technique of reconstruction to improve the signal/noise ratio was applied; and the self-correlation of the BAO peak in the Ly-$\alpha$ forest in the SDSS-III/BOSS DR11 data at $\langle z\rangle=2.34$ (Padmanabhan et al. 2012), plus the cross-correlation of the BAO peak of QSOs and the Ly-$\alpha$ forest in the same survey (Font-Ribera et al. 2014). With these new measurements, the use of the Alcock-Paczy\'nski diagnostic (Melia and L\'opez-Corredoira 2016) has shown that the current concordance ($\Lambda$CDM) model is disfavoured by the BAO data at $2.6\sigma$. They instead show that the $R_{\rm h}=ct$ model has a probability $\sim 0.68$ (i.e., consistent with 1) of being correct. Measurements of the linear growth rate also critically depend on the Alcock-Paczy\'nski effect, so the observations considered in this paper provide an invaluable, complementary, set of data with which to test the $R_{\rm h}=ct$ cosmology. In \S~2 of this paper, we derive the necessary formalism for studying the time evolution of linear fluctuations in this model, which reduces to solving a second-order differential equation, though with some important differences compared with its counterpart in $\Lambda$CDM. A contextual background for $R_{\rm h}=ct$ is provided in \S~3, where we also solve the growth equation as a function of redshift, and describe the observables, specifically the volume-delimited variance $\sigma_8(z)$ of the fluctuations and its corresponding growth function. The standard model is analyzed in \S~4, and we end with a discussion and conclusion in \S\S~5 and 6. | In recent years, perturbation theory has matured to the point where the predictions of $\Lambda$CDM have been compared extensively to measurements of the growth rate $f\sigma_8(z)$, and to other models. The current consensus is that Planck $\Lambda$CDM (Planck Collaboration 2014a) is generally well matched to the data, and that no firm evidence exists for extensions to general relativity (Alam et al. 2016). Several recent analyses, however, have yielded some tension between the value of $\sigma_8(0)$ measured using redshift space distortions in galaxy surveys and that inferred by fitting anisotropies in the cosmic microwave radiation (Guzzo et al. 2008; Macaulay et al. 2013). \begin{figure} \centerline{\includegraphics[angle=0,scale=0.7]{fig3.eps}} \caption{Residuals relative to the best-fit (solid) curves in figures~1 and 2. Shaded sections correspond to 1 $\sigma$ confidence regions. For $\Lambda$CDM, only 4 of the 9 points lie within 1 $\sigma$ of the optimized model. More critically, all 5 of the remaining points lie below it. Together with the indication given by the $\chi^2$ values, these results suggest that the linear growth-rate measurements favour $R_{\rm h}=ct$ over Planck $\Lambda$CDM (Planck Collaboration 2014a).}\label{Residuals} \end{figure} The principal goal of this paper has been to ascertain whether or not the predictions of the $R_{\rm h}=ct$ universe are also consistent with the measured growth rate at redshifts $z < 2-3$. We have found that the current measurements, though still not sufficiently precise to clearly distinguish between $R_{\rm h}=ct$ and $\Lambda$CDM, nonetheless favour the former over the version of the latter optimized by the joint analysis of {\it Planck} and linear growth rate data. The two models are statistically indistinguishable when the optimization of the $\Lambda$CDM fit is based solely on the growth rate data in Table 2. Our results also suggest that the present consistency of the standard model with the growth-rate data may be an artifact of the relatively large errors associated with these measurements, which cannot yet clearly distinguish between functional forms of $f\sigma_8(z)$ with and without significant curvature. This work strongly affirms the need for more precise measurements of the growth rate in that critical redshift range ($0 < z < 1$) where differences in the growth function $f\sigma_8(z)$ between $\Lambda$CDM and $R_{\rm h}=ct$ are most pronounced. | 16 | 9 | 1609.08576 |
1609 | 1609.00792_arXiv.txt | Multiple outbursts of a Jupiter-family comet, 15P/Finlay, occurred from late 2014 to early 2015. We conducted an observation of the comet after the first outburst and subsequently witnessed another outburst on 2015 January 15.6--15.7. The gas, consisting mostly of C$_{2}$ and CN, and dust particles expanded at speeds of 1,110 $\pm$ 180 m s$^{-1}$ and 570 $\pm$ 40 m s$^{-1}$ at a heliocentric distance of 1.0 AU. We estimated the maximum ratio of solar radiation pressure with respect to the solar gravity $\beta_\mathrm{max}$ = 1.6 $\pm$ 0.2, which is consistent with porous dust particles composed of silicates and organics. We found that 10$^8$--10$^9$ kg of dust particles (assumed to be 0.3 \micron--1 mm) were ejected through each outburst. Although the total mass is three orders of magnitude smaller than that of the 17P/Holmes event observed in 2007, the kinetic energy per unit mass (10$^4$ J kg$^{-1}$) is equivalent to the estimated values of 17P/Holmes and 332P/2010 V1 (Ikeya--Murakami), suggesting that the outbursts were caused by a similar physical mechanism. From a survey of cometary outbursts on the basis of voluntary reports, we conjecture that 15P/Finlay--class outbursts occur $>$1.5 times annually and inject dust particles from Jupiter-family comets and Encke-type comets into interplanetary space at a rate of $\sim$10 kg s$^{-1}$ or more. | \label{sec:introduction} 15P/Finlay (hereafter 15P) was an undistinguished comet discovered by William Henry Finlay at Cape of Good Hope, South Africa, on 1886 September 26. This comet has a semi-major axis of $a$ = 3.488 AU, eccentricity of $e$ = 0.720, inclination of $i$ = 6.80\arcdeg, and Tisserand parameter with respect to Jupiter of $T_\mathrm{J}$ = 2.62, which are typical of Jupiter-family comets (JFCs). Since the discovery, it showed irregular magnitude light curves at different apparitions \citep{Sekanina1993}. The effective radius of 15P is estimated to be 0.92 $\pm$ 0.05 km \citep{Fernandez2013}, which is consistent with early results in \citet{Whipple1977} and \citet{Mendis1985}. It has maintained the perihelion around the Earth orbit at 0.98--1.10 AU for about a century and is sometimes linked to a meteor shower \citep{Beech1999,Terentjeva2011}. It is likely that the measured absolute magnitude reduced by a factor of $\sim$10 from 7.5 mag in 1886 to 10.1 mag in 1981 \citep{Kresak1989}, suggesting that 15P might have lost a fraction of volatile components near the surface while developing a dust mantle layer on the surface similar to other periodic comets \citep[e.g.,][]{Kwon2016,Hsieh2015}. The comet exhibited two large-scale outbursts around the perihelion passage in 2014--2015, with the first outburst occurring on 2014 December 16 \citep{Ye2015}. The image showed an envelope feature and near--nuclear tail, which is reminiscent of past cometary outbursts at 17P/Holmes and 332P/2010 V1 (Ikeya--Murakami), hereafter referred to as 17P and 332P, respectively \citep{Ishiguro2010,Ishiguro2014}. Soon after the report of the first outburst, we conducted an observation from 2014 December 23 to 2015 March 16 for deepening our understanding of cometary outbursts. We used six ground-based telescopes that constitute a portion of the Optical and Infrared Synergetic Telescopes for Education and Research (OISTER) inter-university observation network. As a result of frequent observation several times a week, we witnessed the second outburst on UT 2015 January 15. Such cometary outbursts have drawn the attention of researchers because they offer insight into the physical properties of comet nuclei \citep{Hughes1990}. The huge outburst of 17P could be explained by the crystallization of buried amorphous ice \citep{Li2011}. Although similar morphological features were found at 332P \citep{Ishiguro2014}, numerous of fragments were identified at its return in 2016 \citep{Weryk2016,Kleyna2016}. Motivated by a series of detections regarding cometary outbursts, we investigated the physical properties of the 15P multiple outbursts and estimated the frequency and mass production rate of outbursts on a scale similar to 15P. We describe our observations and data analysis in Section 2, and the photometric and polarimetric results in Section 3. We then discuss our findings considering the reports of recent outbursts in Section 4. | \label{sec:discussion} \subsection{Maximum Radiation Pressure Force and Constraint on Dust Physical Properties, $\beta_\mathrm{max}$} Cometary outbursts provide opportunities to investigate fresh cometary materials that are embedded in the surface processed layers but appear when a large amount of materials is ejected via explosions. In particular, the 15P event in 2015 January provided a unique opportunity to determine the maximum $\beta$ value with good precision for several reasons. Because we determined the precise onset time of the outburst, it was easy to follow the motion with respect to the reference time. In addition, the outburst ejecta were observed at the large phase angle of $\alpha$ = 45\arcdeg\ (c.f. $\alpha$ = 19\arcdeg\ for 332P and 17\arcdeg\ for 17P), which enabled observation of the dust motion toward the anti-solar direction more clearly. Moreover, the outburst occurred at a smaller heliocentric distance (1.0 AU), where the acceleration by radiation pressure was detected more effectively than those at distant locations such as 1.6 AU for 332P and 2.4 AU for 17P. It is well known that radiation pressure force depends on not only size but also and composition. The $\beta$ values have been investigated theoretically considering fluffy dust particles \citep{Mukai1992} and a variety of compositions \citep{Burns1979}. In section \ref{subsec:dynamics}, we adopted a simple model that $\beta$ is inversely proportional to the size under an assumption of a radiation pressure coefficient of $Q_\mathrm{pr}$ = 1. However, this assumption holds only when the particle size is larger than the optical wavelength ($\lambda$ = 0.5 \micron). $Q_\mathrm{pr}$ has almost constant value for $a>$0.1--0.3 \micron\ and significantly drops as the size of dust grains decreases \citep{Ishiguro2007}. For fluffy dust particles, $\beta$ is less dependent on the aggregate size but similar to that of each constituent \citep{Mukai1992}. \citet{Wilck1996} calculated the $\beta$ values of dust particles by using a core--mantle spherical particle model composed of silicate and amorphous carbon with different porosities and suggested that cometary dust particles have $\beta_\mathrm{max}$=1--1.8. Transparent materials such as silicates and water ice tend to have small $\beta$ maximum values (i.e., $\beta_\mathrm{max}<$1), whereas that of absorbing particles such as carbon is $\beta_\mathrm{max}>$1 \citep[see e.g.,][]{Kimura2016}. The $\beta_\mathrm{max}$ value determined in our measurement, 1.6 $\pm$ 0.2, is consistent with that in the porous absorbing particles model \citep{Wilck1996} and silicate-core, organic-coated grains model for dust aggregates \citep{Kimura2003}. However, our value is inconsistent with those of silicate spheres ($\beta_\mathrm{max}<$1) and organic spheres ($\beta_\mathrm{max}>$3). \subsection{Outburst Frequency} Cometary outbursts have been observed in a wide variety of comets. There are some references that analyzed historical observation of cometary outburst \citep[e.g.][]{Hughes1990}. However, the occurrence frequency and mass production rate have not been studied well because there is no coordinated observation system for monitoring comet magnitude. Some outbursts could have been missed because the magnitude contrast was too weak before and after to be noticed as an outburst or because the observation condition was not suitable for detection (i.e., too close to the Sun or Moon). Nevertheless, some outburst events have been reported in a voluntary manner by amateur observers. For example, the 17P event was first reported by A. Henriquez Santana (Spain), who noted sudden brightening to 8.4 mag by using a 0.2-m reflector \citep{Buzzi2007}. 332P was discovered in Japan during its outburst by S. Murakami and K. Ikeya. They observed this event by looking into the eyepieces of 0.46-m and 0.25-m telescopes when it reached $\sim$9 mag \citep{Murakami2010}. Moreover, 15P double outbursts with brightening to 8--9.4 mag were also reported by amateurs through a mailing list of comet observers \citep{Ye2015}. It has been suggested that 17P and 332P events were caused by a phase change of amorphous water ice \citep{Li2011,Ishiguro2013}, although other mechanisms such as rotational breakup of brittle cometary nuclei cannot be ruled out \citep{Li2015}. Excavations via impacts might create an environments favorable to outbursts, as suggested in \citet{Beech2002}. Because the energy per unit mass is similar to both events, we conjecture that 15P events were also caused by a phase change of amorphous water ice. Here, we consider the frequency of such outbursts in the inner ($r_\odot\lesssim5$ AU) solar system, where the physical mechanism of general cometary activity is confined to sublimation of water ice rather than that of super volatiles such as CO \citep{Jewitt2009}. For the reason, we excluded frequent outbursts at 29P/Schwassmann-Wachmann 1 beyond the Jupiter orbit. Moreover, we restricted our discussion to JFCs and Encke-type comets (ETCs) because significant observation samples are not available for Halley-type comets (HTCs) and long-period comets (LPCs). We applied a sample cut of the perihelion distance $q\lesssim5$ AU to the JPL Small-Body Database Search Engine\footnote{http://ssd.jpl.nasa.gov/sbdb_query.cgi} and retrieved 505 JFCs and ETCs in the list of known comets as of the end of 2015. We regarded disintegrated comets as a single object (e.g. 73P/Schwassmann-Wachmann 3 B, C, G... constituted a united body). Moreover, we excluded disappeared or dead comets (3D/Biela and 5D/Brorsen), designated as "D/," and main-belt comets, which are unlikely to contain amorphous ice \citep{Prialnik2009}. Among JFCs and ETCs, 357 objects passed their perihelion at least once between 2007 October and 2015 December, which make up $\sim$70\% of the entire population. This means that $\sim$70\% of JFCs and ETCs might have a chance to display outburst activities owing to the additional heat from the sun near their perihelia. We attempted to find outburst events through the Smithsonian Astrophysical Observatory/National Aeronautics and Space Administration (SAO/NASA) Astrophysics Data System (ADS) Astronomy Query Form by inputting "comet" and "outburst" as abstract terms and manually plumbing the results by reading the abstracts. In addition, we added one event at 205P which showed an outburst showing similar appearance to 15P (Seiichi Yoshida, private communication). We found that 15 outburst events occurred at 11 comets since 2007 (Table 4). We chose the arbitrary period beginning in 2007 because researchers have increased their consciousness toward outbursts since then, as motivated by the 17P event. Table 4 shows some minor outbursts detected by observers with larger telescopes because these comets drew their attentions due to space mission targets (81P/Wild 2 and 103P/Hartley 2) and repeated outbursts at 17P and 15P. Figure \ref{fig:cumM} shows the cumulative mass distribution of the outburst ejecta. We note that six outburst events, including 17P, 168P/Hergenrother, 332P, 217P/LINEAR, and two events at 15P, were brightened down to 10 mag. Such events would be detectable with observation through inexpensive equipment such as $\sim$10-cm class telescopes. Thus, we conjecture that outbursts $\lesssim$10 mag would be detected almost completely if observed conditions allowed ground-based observers to make observations. We fit the ejecta mass distribution with a power-law function for five objects (i.e., $\lesssim$10 mag events except 17P) and obtained a power index of $\gamma$ = 0.45 $\pm$ 0.09. Although there might be no physical basis to support the power-law function in Figure \ref{fig:cumM}, it is likely that the 17P outburst in 2007 was a very rare phenomenon to be detected within $\sim$8 years because the occurrence is one order of magnitude higher than that expected by the power function. In fact, the large discrepancy for 17P must be a bias of our analysis in which we counted the number of outbursts since the epoch-making event. In addition, the power-law function fits well in the mass range of 6 $\times$ 10$^7$--1 $\times$ 10$^{9}$ kg but deviates from the observed values in the mass range of $\lesssim$ 6 $\times$ 10$^7$ kg. Some outbursts might not be noticed because of the faintness or weak contrast before and after outbursts. For these reasons, we consider the unbiased outbursts rate in which the ejecta mass corresponds to 6 $\times$ 10$^7$--1 $\times$ 10$^{9}$, (hereafter referred to as "15P--class" outbursts). We applied the Poisson distribution following \citet{Sonnett2011}: \begin{eqnarray} P(n)=\frac{(fCN)^n \exp(-fCN)}{n!}, \label{eq:poisson} \end{eqnarray} \noindent where $n$ and $f$ denote the number of outburst detections and incidence of outbursts, $N$ is the number of observed samples, and $C$ ($\in$ [0, 1]) is the completeness of the survey. The 90\% upper confidence limit on the incidence, $f_{90\%}$, is given by \begin{eqnarray} 0.9=\frac{\int_0^{f_{90\%}} P(n)df}{\int_0^{1} P(n)df}~. \label{eq:poisson} \end{eqnarray} Assuming that observers could detect 15P--class outbursts completely at the solar elongation of $\epsilon_\odot>$60\arcdeg, we have $C$ = 0.67. Over eight years since October 2007, there were $n = 7$ 15P--class outbursts out of $N$ = 357 comets that passed their perihelion. By solving the implicit Eq. (\ref{eq:poisson}), we obtained $f_{90\%}$ = 0.05, which suggests an expected number $\langle n \rangle$ = $f_{90\%}CN$$\sim$12 in 8 years, or 1.5 times per year. Therefore, we can express the unbiased cumulative and differential frequency of 15P--class outbursts $f_{ub}$ and $f'_{ub}$ as \begin{eqnarray} f_{ub}\left(>M\right) = \int^{\infty}_{M} f'_{ub}\left(m\right)~dm = A \left(\frac{M}{M_0}\right)^{-\gamma}, \label{eq:funbias} \end{eqnarray} \noindent where $A$ = 1.5 yr$^{-1}$ and $M_0$ = 6 $\times$ 10$^7$ kg are constants. We obtained an effective mass production rate of 15P--class outbursts of $\sim$ 3 $\times$ 10$^8$ kg year$^{-1}$ or $\sim$10 kg sec$^{-1}$, from (see appendix A) \begin{eqnarray} \langle m \rangle = \int^{M_2}_{M_1} m f'_{ub}(m)~dm~, \end{eqnarray} \noindent where $M_1$=6$\times$10$^7$ kg and $M_2$=1$\times$10$^9$ kg are lower and upper bounds on the power-law behavior of 15P--class outbursts, respectively. We consider that the incidence would be underestimated because we optimistically assumed $C$ = 0.67. Some of the outbursts were missed owing to the crowded region of stars or the full lunar phase. Even considering these factors, which may decrease $C$ by several factor, our result would obtain one order-of-magnitude accuracy for the unbiased frequency. Therefore, it is reasonable to think that 15P--class outbursts inject dust particles into the interplanetary space at a rate of $\gtrsim$10 kg sec$^{-1}$. The mass is $\sim$2 orders of magnitude less than the mass required to sustain the interplanetary dust cloud. This implies that 15P--class outburst events may not compensate the mass eroded by Poynting--Robertson drag and other dynamical mechanisms onto the interplanetary dust. However, if we integrate $m f'_{ub}(m) dm$ up to the 17P-class ejecta mass (i.e., $M_2$=1$\times$10$^{12}$ kg), the mass production rate from 15P--17P class events would be $\sim$500 kg sec$^{-1}$, although we are not sure whether the frequency follows a simple power-law function up to the 17P class. These results suggest that large-scale cometary outbursts might contribute a significant fraction of the interplanetary dust source. We made an observation of 15P during the perihelion passage in 2015--2016, following a report of an outburst in the middle of 2015 December and detected the second outburst on UT 2015 January 15.6--15.7, which was equivalent to the first outburst. The results of our analysis are summarized in the following points. \begin{enumerate} \item{ Gas consisting mostly of C$_{2}$ and CN expanded at a speed of 1,110 $\pm$ 180 km s$^{-1}$, which is slightly faster than the speeds for other comets around 1 AU. The excess in speed can be explained by the large distance from the nucleus ($\approx 10^8$ km), where the gas flow velocity continues to increase. } \item{The dust ejecta accelerated up to a speed of 570 $\pm$ 40 km s$^{-1}$, which is comparable to the ejection speeds of 17P and 332P ejecta. These consistent speeds would have resulted in the similar appearances of these outburst ejecta. } \item{We derived the total mass of dust ejecta as 10$^8$--10$^9$ kg ($a$ = 0.3 \micron--1 mm was assumed). This mass is equivalent to that of the 2010 event at 332P but is three orders of magnitude smaller than the 2007 event at 17P.} \item{The polarization degree was measured to 6.8 $\pm$ 0.2\% at the phase angle $\alpha$ = 43\arcdeg, which fell into the common values of other comets. This similarity does not mean that outburst ejecta have similar polarimetric properties because it was diffused out at the time of our polarimetric measurement.} \item{Based on the immediate observation of dust ejecta for the second outburst, we derived a reliable estimate of $\beta_\mathrm{max}$ = 1.6 $\pm$ 0.2. This value is consistent with the theoretical prediction for porous absorbing particles, suggesting that such porous dust particles could remain inside the cometary nucleus and are released during the outburst. } \item{ The kinetic energy per unit mass (10$^4$ J kg$^{-1}$) is close to estimated values of 17P and 332P. In addition, the dust mass, speed, and kinetic energy are broadly comparable to the measured values of the 2010 outburst at 332P. This may suggest that these three outbursts occurred by a similar mechanism. } \item{ From a survey of cometary outbursts in publications in the SAO/NASA ADS, we estimated that 15P/Finlay--class outbursts occur annually, injecting cometary materials into interplanetary space at a rate of $\gtrsim$10 kg sec$^{-1}$ or more. } \end{enumerate} {\bf Acknowledgments}\\ This research, conducted at Seoul National University was supported by the Basic Science Research Program through the National Research Foundation of Korea (NRF) funded by the Ministry of Education (NRF-2015R1D1A1A01060025, No. 2012R1A4A1028713). The acquisition of observation data was supported by the Optical and Near-infrared Astronomy Inter-University Cooperation Program. We would like to thank Seiichi Yoshida for useful discussion and S. Goda for supporting observation at the Nayoro Observatory. Y. G. Kwon is supported by Global Ph.D Fellowship Program through the NRF funded by the Ministry of Education (NRF-2015H1A2A1034260). \appendix | 16 | 9 | 1609.00792 |
1609 | 1609.09811_arXiv.txt | We combine the decade long photometry of the Be/X-ray binary system \asource provided by the MACHO and OGLE~IV projects with high resolution SALT spectroscopy to provide detailed constraints on the orbital parameters and system properties. The $\sim$420d superorbital modulation is present throughout, but has reduced in amplitude in recent years. The well-defined 16.6409d orbital outbursts, which were a strong function of superorbital phase in the MACHO data (not occurring at all at superorbital maximum), are present throughout the OGLE~IV coverage. However, their amplitude reduces during superorbital maximum. We have refined the orbital period and ephemeris of the optical outburst based on $\sim$ 25 yrs light curves to HJD~=~2455674.48$\pm$0.03 + n*16.6409$\pm$0.0003d. Our SALT spectra reveal a B1~III star with \vsini of~285~\kms from which we have derived an orbital radial velocity curve which confirms the high eccentricity of $e$ = 0.72$\pm$0.14. Furthermore, the mass function indicates that, unless the neutron star far exceeds the canonical 1.44 \msun, the donor must be significantly undermassive for its spectral type. We discuss the implications of the geometry and our derived orbital solution on the observed behaviour of the system. | The recurrent X-ray transient \asource was first discovered in 1977 by \citet{White1978} when two X-ray outbursts were detected with the {\it Ariel V} satellite. Subsequently, several outbursts from the source which recur with a period of 16.7 d were observed with the {\it HEAO-1} modulator collimator \citep{Johnston1979, Johnston1980, Skinner1980}. The improved {\it HEAO-1} position enabled \citet{Johnston1980} to identify its optical counterpart with a bright (V$\sim$15.7) Be star. The radial velocity of its absorption lines confirms its physical association with the Large Magellanic Cloud (LMC), and hence a well-established distance of 50 kpc \citep{Alves2004}. At this distance the peak luminosity of the X-ray outbursts were estimated to be $\sim$ $\mathrm{10^{39} erg~s^{-1}}$. That makes \asource one of the most luminous X-ray sources known, and still the highest in the Be/X-ray binary (BeX) class \citep{Liu2005}. In one of the X-ray outbursts observed with the Einstein Observatory, \citet{Skinner1982} found an X-ray pulsation with a period of 69 ms which indicates the presence of a rapidly rotating NS in the system. However, for a distance of 50 kpc the observed X-ray luminosity is then substantially super-Eddington for a 1.44 $\mathrm{M_{\odot}}$ neutron star (NS). \citet{Skinner1981} derived an improved orbital period of 16.6515 d using archival plates taken with the UK Schmidt telescope. They also reported that the X-ray outbursts were accompanied by an optical brightening of 2 mag. \citet{Densham1983} observed four consecutive outbursts in late 1981, from which they found that the source can reach as bright as V$\sim$12, and that the optical brightening was accompanied by the sudden appearance of very strong \heII $\mathrm{\lambda4686}$ emission. Since late 1983, the scale of these periodic outbursts has reduced considerably, with only much weaker outbursts (few tenths of a magnitude) being seen in both optical and X-ray observations. Using optical and UV spectroscopy, \citet{Charles1983} classified the optical counterpart of \asource as a B2~III--V star. However, \citet{Hutchings1985} and \citet{Negueruela2002} suggested slightly earlier spectral types of B1 and B0.5~III, respectively, based on optical spectroscopic data taken during quiescence. The existence of long-term optical monitoring of the LMC by gravitational micro-lensing projects, such as the MACHO, has revolutionized our understanding of the long-term behaviour of \asource. \citet{Alcock2001} analysed the first $\sim$5 years MACHO light curve and found that the source displays large amplitude superorbital variability on a timescale of 420.8d. This was suggested to be a result of the formation and dissipation of the equatorial disc around the donor star. Analysis of the 70 years of archival Harvard and Schmidt plates by \citet{McGowan2003} also showed evidence of this long-term modulation. Within this 420.8d superorbital variation, the 16.6515d orbital outbursts were seen, but they were confined to only occur during optical minimum, never during the optical maximum (when the source remained quiescent). At optical maximum \asource is bluer, but remarkably when the disc forms it masks part of the hotter Be star which makes the source appear redder and fainter. This behaviour is typical for an equatorial disc viewed at high inclination angle \citep[for details, see][]{Rajoelimanana2011}. Based on the brightness changes between the optical maxima and minima, \citet{McGowan2003} suggested a lower limit for the inclination of the Be star of $i=74.9\degr$. Given its relative brightness and now well-established orbital period, there have been attempts to obtain spectroscopic radial velocity curves in order to constrain the system's kinematic properties \citep{Corbet1985, Hutchings1985}. However, with a likely rather high eccentricity proposed by \citet{Charles1983} (on the basis of the orbital optical/X-ray light curves) and a neutron star compact object, significant velocity changes were only expected close to periastron. Unfortunately, the radial velocity curves of \citet{Corbet1985} and \citet{Hutchings1985} are inconsistent, for which \citet{Smale1989} suggested that the spectra from that time may have been contaminated by residual emission components. This paper presents an analysis of the complete MACHO \citep{Alcock1996} and OGLE~IV \citep{Udalski2015} optical light curves of \asource as well as the results of optical spectroscopy obtained from the Southern African Large Telescope (SALT) and the SAAO 1.9 m telescope. The spectra are of sufficient resolution and S/N to obtain a new radial velocity curve, which is similar to that of \citet{Hutchings1985}. | \subsection{Long-term optical variations} BeX systems are known to exhibit large amplitude long-term photometric variability on timescale of years to decades. Such quasi-periodic superorbital variations have previously been reported in a number of BeX systems \citep[see][]{McGowan2008, Rajoelimanana2011}. \citet{Alcock2001} suggested that the 420.8 d superorbital modulations in the first 5 years of MACHO optical light curves of \asource are related to the variation in the size of the Be equatorial disc. As described in section~\ref{optical_lightcurve}, this superorbital variation appeared to have shortened in the second half of the MACHO. Hovever, in spite of a significantly reduced amplitude, the periodogram of the recent OGLE~IV data still shows a highest peak at 423.74d, very close to that previously reported, as well as additional power around 200d. These shorter period peaks are due to the change in the overall shape of the superorbital modulation in the OGLE~IV light curve compared to the MACHO era. \subsubsection{Disc-less phase (optical high state)} \label{discless_phase} The circumstellar disc is the main reservoir of material available for accretion by the NS. We would not, therefore, expect to see any outburst when the disc is absent or very weak (i.e. during the disc-less phase). The MACHO $V$ and $R$ light curves show flat optical maxima that last $\sim$ 200 d, during which the colour is bluest and does not change. This suggests that the system has entered its disc-less, normal B star phase. We therefore infer that the brightness ($V$ = 14.885, $R$ = 15.025) and colour ($V - R$ = --0.14) measured at those times correspond to the true intrinsic magnitudes and colour of the underlying B star. This was confirmed by our optical maximum spectra which are dominated by the photospheric spectrum of the underlying B star. The optical high state in the OGLE~IV observations occurred near 2012 Oct 27 and the observed \halpha line profile at that time appeared to be in absorption. However, some weak orbital outbursts are still seen in the light curve which indicates that the Be equatorial disc is not dissipating completely, and the NS can still influence it. This residual disc material at the optical maxima in the OGLE~IV data still partially obscures part of the B star, resulting to a decrease in the overall amplitude of the 420.8d compared to the disc-less phase during the MACHO observations. \subsubsection{Be phase (optical low state)} \label{bephase} We are viewing the Be equatorial disc in \asource nearly edge-on (see section~\ref{geometry}) which means when it forms it will mask part of the hotter (brighter) B star, causing the optical brightness to decline. The presence of the circumstellar disc makes the system redder and gives rise to Balmer and sometimes \heI emission lines. The passage of the NS near periastron has an affect on the structure of the Be circumstellar disc and its emission properties. If the disc is large enough, then the NS interacts and accretes from it near periastron which gives rise to periodic outbursts seen in X-ray and optical. During the orbital outburst the colour of the system reddens slightly, which suggests that the additional brightening originates from the cooler equatorial disc. The strength of outburst also varies significantly throughout the superorbital cycle, which means that the observed effect of the periastron passage of the NS depends strongly on the density and extent of the equatorial disc. The outbursts are much stronger when the disc is denser and larger (at optical minima) and do not occur when the disc is very weak or absent (disc-less phase). Furthermore, the light curve of \asource folded on the period of 16.6409~d shows a double peak profile with two peaks occurring at phase~0.0 and 0.071 (see Fig.~\ref{fold_all}). This suggests that there is a misalignment between the spin axis of the Be equatorial disc and the orbital plane of the NS. A spin-orbit misalignment has previously been seen in some BeX systems \citep[e.g. SXP327, ][]{Coe2008}. The phase of the two optical peaks ($\phi$=0.0 and 0.071) as well as the time of periastron passage ($\phi$=0.04) were plotted in the Fig.~\ref{plot_geom}. The presence of two peaks in \asource has also been reported by \citet{Corbet1997} in which they suggest that the increase in relative velocity near periastron may cause the dip in brightness between the peaks. We note that they observed the dip on 1997 Feb 02 which corresponds to phase~0.05 (very near periastron) in our new ephemeris, and so we interpret this phenomenon as two superposed peaks resulting from the misaligned equatorial disc. The appearance of the \heII~\lambdaf 4686 emission line in our spectra indicates the presence of hard ionizing radiation with energies above 54 eV which means \asource was in an X-ray "on" state at that time. The observed \heII~\lambdaf 4686 emission line has a double-peaked profile and is probably due to radiative recombination in a photoionized accretion disc around the NS rotating with Keplerian velocity. The \heII~emission line appears only in the spectrum taken near periastron and is strongest at phase~0 (optical outburst, EW$\sim$15.17 \AA~on 2013 Dec 13) where the NS is predicted to interact and accrete from the circumstellar disc. This suggests that the accretion disc in these systems is transient and is only present around periastron. The \heII~emission line has been previously detected in \asource, all of which occurred near phase~0 \citep[e.g.][]{Charles1983, Corbet1985}. \subsection{\halpha emission} The \halpha emission line in BeX systems originates from the geometrically thin, nearly Keplerian circumstellar disc around the Be star \citep{Porter2003}. Therefore, studying the evolution of its profiles will allow us to estimate the long-term structural changes that occur within the circumstellar disc. The values of EW (\halpha) show considerable variability throughout our SALT observations. Because of the short orbital period and large eccentricity in \asource, the periastron passage of the NS has a significant effect on the evolution of the Be circumstellar disc. The periodic gravitational influence of the NS near periastron in such a narrow orbit prevents the Be circumstellar disc from growing larger and causes truncation \citep{Reig1997, Okazaki2001}. The \halpha emission profile in \asource is a strong function of orbital phase. Far from periastron, it shows a symmetric shell profile, characteristic of an unperturbed equatorial disc viewed nearly edge-on and rotating with Keplerian velocity distribution. However, near periastron the line profile is double-peaked and highly asymmetric, becoming very strong and single-peaked at periastron. Near periastron, the circumstellar disc may be distorted by the NS' gravity, and this will make it highly non-axisymmetric with an enhanced density toward the compact object. Based on our orbital solution of $\omega$=183$\degr$ and $e$ = 0.72, the NS (and hence the high density region of the equatorial disc) is moving away from us near periastron (see Fig.~\ref{plot_geom}). Therefore, we would expect an increase in the strength of the red peak of the \halpha profile at phase $\sim$ 0.0--0.1 (on the right of the Be star in Fig.~\ref{plot_geom}). However once the NS is far from periastron the \halpha line profile will return to a symmetric shell profile again. This is what we have seen in our spectra taken on 2012 Dec 25. The spectra obtained at periastron (see $\phi$=0.038 and 0.029) are single-peaked and also have a broad wing component. This suggests that at periastron we are viewing the circumstellar disc at a lower inclination angle than it was before (outside periastron). This change in the disc inclination implies that at periastron the tidal torque from a misaligned companion neutron star tends to warp the equatorial disc towards the binary orbital plane which has a lower inclination angle than the disc. However, as a result of warping, the inner region close to the Be star, where the velocities are largest, can be at a still higher inclination angle. This high velocity region contributes to the observed broad wings in the profile. Similar behaviour has been observed in 4U 0115+63/V635 \citep{Negueruela2001}. \citet{Martin2011} also reported that the tidal torque from the NS could warp the equatorial disc in BeX system with short orbital periods if the orbit is eccentric and misaligned with the equatorial disc. \subsection{Geometry of the system} \label{geometry} The behaviour of the brightness and colour of \asource (brighter when bluer) indicates that the inclination angle of the circumstellar disc is large ($i \geq90\degr-\alpha $, where $\alpha \leq 10\degr$ is the opening angle of the disc). We note that the majority of BeX systems show the opposite behaviour, where the source is redder when brighter \citep[see][]{Rajoelimanana2011}. In addition, the presence of the shell profile in the \halpha line confirms that we are viewing the Be equatorial disc nearly edge-on \citep{Hummel2000}. Using our values for the projected rotational velocity (\vsini~= 285 \kms) and measured peak separation, and assuming a Keplerian rotation of the disc, the outer radius of the \halpha emitting region can be estimated using the equation \citep{Huang1972}: \begin{equation} \left( \frac {R_{\rm out}}{R_{1}}\;\right) = \; \left( \frac{2\,v\,\sin{i}} {\Delta V} \;\right)^{2} , \label{Huang1972} \end{equation} \noindent where $R_{1}$ is the radius of the Be star and $\Delta V$ is the peak separation. The presence of the \heII~$\lambda$4686 line in the spectra taken at $\phi$ = 0.017 (2012 Dec 28) at the end of our SALT campaign suggests that mass transfer onto the NS is again occurring in the system. To estimate the outer radius around that period of time, we measured a peak separation of 540 \kms from spectra taken far from periastron and at the same orbital cycle (2012 Dec 25). This gives an outer radius of $R_{\rm out}=1.15~R_{1}$ (shown as the grey circle in Fig.~\ref{plot_geom}). We note that we could still see weak outbursts in the optical high state during the OGLE observations, therefore the value of our \vsini and hence the outer radius of the \halpha emitting region that we measured is slightly underestimated. \begin{figure} \scalebox{0.5}{\includegraphics{plot_geometry.eps}} \caption{Geometry of the system, showing the relative orbit of the NS around the Be star computed using our orbital solution. The coordinates are in units of the semimajor axis. phase~0 is defined as the time when the outburst reaches its maximum brightness. The phase of the periastron (square), the two optical outburst peaks (circle) and the SALT observations (filled circle) were marked. The black and grey circles represents the optical Be star with radius R$_1$ =10.0 \rsun and the equatorial disc with radius $R_{\rm out}$= 1.15 $R_{1}$, respectively. The dashed ellipse represents the Roche lobe radius at each orbital phase for masses of $M_{1}$ = 8.84 \msun and $M_X$ = 1.44 \msun. The NS is moving counter-clockwise around the Be star. We note that the figure is in scale.} \label{plot_geom} \end{figure} Our spectroscopic study is essentially ``single-lined" as, until \asource is recovered as an X-ray pulsar, the only kinematic constraints on the NS are provided by \heII~$\lambda$4686, which we assume to be emitted close to the accretion disc, and that only appears near periastron passage. Our radial velocity curve therefore provides only the mass function. Furthermore, the binary inclination, which clearly must be different from the inclination of the Be equatorial disc, is also unknown. In Fig.~\ref{mass_range} , we have used our mass function to plot the range of possible masses of the Be companion star and the NS for several values of the binary inclination. If we consider a neutron star mass range of 1.44 \msun $\le M_{\rm X} \le$ 3.0 \msun~and a corresponding range for the Be star mass of 10.0 \msun $\le M_{\rm 1} \le$ 15.0 \msun, then clearly we require $i_{\rm orb}$ > $34\degr$ (see Fig.~\ref{mass_range}). However, there are additional constraints. The absence of X-ray eclipses in the light curves of \asource \citep[see][]{Skinner1980a} implies that $i_{\rm orb} \le 75\degr$ for a B1~III star with radius $R_1$ = 10.0 \rsun. This is in good agreement with our interpretation that the double peaked outbursts are related to the disc-orbit misalignment in the system. Using the mass function, an upper limit of $i_{\rm orb} \le 75\degr$, and adopting a canonical NS mass of 1.44 \msun, we estimate an upper limit of 8.84 \msun~for the mass of the Be donor. Therefore, unless the NS is heavily over-massive, the Be star is under-massive for its spectral type of B1~III. However, this is highly typical of high mass X-ray binaries in general, where the evolutionary path has led to a donor that has significantly different properties of those of a normal main-sequence equivalent \citep[see e.g.][]{Rappaport1983}. \begin{figure} \scalebox{0.5}{\includegraphics{plot_massrange.eps}} \caption{Mass constraints for the two components of the binary system. The lines have been calculated using our mass function for the labelled inclination values. The range of possible masses for the B1~III optical star (10.0 \msun < $M_1$ < 15.0 \msun) are represented by the grey region. The dashed horizontal line represents the standard mass of the neutron star ($M_{X}$=1.44\msun).} \label{mass_range} \end{figure} The radius of the Roche lobe can be computed using \citet{Eggleton1983} formula \begin{equation} \label{Roche_radius} \frac{R_L}{A} = \frac{0.49 q^{2/3}}{0.6q^{2/3} + \ln (1+q^{1/3}) } ~~~~~,~~~~~~(0<q<\infty) \end{equation} \noindent where q = M$_1$/M$_X$ is the mass ratio and $A$ is the orbital separation. The Roche lobe radius as a function of the orbital phase using q = 8.84/1.44~$\sim$ 5.86 is plotted in Fig.~\ref{plot_geom}. We note that corotation is only satisfied close to periastron, and so our calculation for the rest of the orbit is only very approximate. With these parameters, the Be star (with radius $R_1$ = 10.0 \rsun) would fill its Roche lobe at phase~0.025 ($\sim$0.5 day before periastron) and would overfill it at periastron ($R_{\rm Lper}$ = 8.98 \rsun). However, the absence of optical and X-ray activities during the extended optical maxima indicates that the naked B star may not fill its Roche lobe even at periastron, but this requires larger masses for both components. For inclinations $i_{\rm orb} \le$ 75\degr, we derived a lower mass limit of 1.73 \msun for the NS and 11.8 \msun for the donor star in order to get them to be within the Roche lobe. | 16 | 9 | 1609.09811 |
1609 | 1609.00101_arXiv.txt | Cosmic voids are promising tools for cosmological tests due to their sensitivity to dark energy, modified gravity and alternative cosmological scenarios. Most previous studies in the literature of void properties use cosmological N-body simulations of dark matter (DM) particles that ignore the potential effect of baryonic physics. Using a spherical underdensity finder, we analyse voids using the mass field and subhalo tracers in the EAGLE simulations, which follow the evolution of galaxies in a $\rm{\Lambda}$ cold dark matter Universe with state-of-the-art subgrid models for baryonic processes in a $(100 \rm{cMpc})^3$ volume. We study the effect of baryons on void statistics by comparing results with DM-only simulations that use the same initial conditions as EAGLE. When identifying voids in the mass field, we find that a DM-only simulation produces 24 per cent more voids than a hydrodynamical one due to the action of galaxy feedback polluting void regions with hot gas, specially for small voids with $r_{\rm{void}} \le 10\ \rm{Mpc}$. We find that the way in which galaxy tracers are selected has a strong impact on the inferred void properties. Voids identified using galaxies selected by their stellar mass are larger and have cuspier density profiles than those identified by galaxies selected by their total mass. Overall, baryons have minimal effects on void statistics, as void properties are well captured by DM-only simulations, but it is important to account for how galaxies populate DM haloes to estimate the observational effect of different cosmological models on the statistics of voids.\ | \label{Section: Introduction} In the year 1998, two independent groups of astronomers reported evidence for an accelerated expansion of the Universe, using type Ia supernovae as distance indicators \citep{ Riess1998, 1999ApJ...517..565P}. These results strongly suggest the requirement of a non-zero cosmological constant $\Lambda$ in Einstein's field equations. Together with the fact that as much as 85 percent of the matter in the Universe appears to be dark (only interacting through gravity), the so-called Lambda Cold Dark Matter ($\mathrm{\Lambda CDM}$) scenario has remained popular for being one of the simplest cosmological models that provides a reasonably good account of many key properties of the Cosmos, such as the expansion history of the Universe \citep{1999ApJ...517..565P}, the large-scale structure in the distribution of galaxies (e.g. \citealt{2005MNRAS.362..505C, 2005ApJ...633..560E}) and the existence and structure of the cosmic microwave background \citep{2004ApJ...617L..99O, 2013ApJS..208...19H, 2014A&A...571A..16P}. The $\mathrm{\Lambda CDM}$ model assumes that general relativity is the correct theory of gravity on cosmological scales. However, alternative approaches that modify the standard theory of gravity have been proposed (e.g. \citealt{2000PhLB..485..208D, 2004LRR.....7....7M, Carroll2005}). One such example is $f(R)$, a family of gravity theories that modify Einstein's Theory of General Relativity by replacing the Ricci scalar $R$ in the Einstein-Hilbert action by an algebraic function of it \citep{Carroll2005}. Since the solar-system and cosmological tests set tight constraints on the feasibility of modifications to gravity (e.g. \citealt{2011IJMPD..20.1357G, 2014IJMPD..2350036G}), $f(R)$ models have incorporated ``chameleon'' mechanisms that have the effect of removing deviations with respect to general relativity in regions where the gravitational potential is deep enough, such as in our Galaxy \citep{2008PhRvD..78j4021B}. Modified gravity models can produce expansion histories very similar to $\mathrm{\Lambda CDM}$, so it becomes necessary to find alternative ways to test whether they can constitute a better match to our Universe. A promising way to constrain modified gravity models consists in studying regions of low density, where the chameleon mechanism is effectively suppressed. Cosmic voids are the most prominent under-dense regions in our Universe: these are vast volumes with very low galaxy and mass densities, surrounded by the walls and filaments of the large-scale cosmic web. Previous theoretical studies have shown that voids might occupy more than 50 percent of the total volume of the Universe \citep{El-Ad1997, 2002MNRAS.330..399P, 2014MNRAS.441.2923C}, and their low matter density makes them useful tools for cosmological tests due to their sensitivity to dark energy \citep{2011MNRAS.411.2615L, 2012MNRAS.426..440B, Pisani2015, 2016MNRAS.463..512D}, modified gravity \citep{2012MNRAS.421.3481L, 2013MNRAS.431..749C, 2015MNRAS.451.1036C, 2015MNRAS.451.4215Z, 2016PhRvD..94j3524A} and alternative cosmological scenarios \citep{2015JCAP...11..018M, 2016JCAP...11..015B, 2015JCAP...08..028B}. Most studies of voids in the literature are based on cosmological N-body simulations that follow the gravitational interaction of dark matter. The baryonic component of the Universe has been so far ignored in these works, and although dark matter accounts for a large fraction of the mass of the Universe, baryons play an important role in a cosmological context. \citet{2015MNRAS.448.2941S} showed that because of cosmic reionization, most haloes below $\rm{3 \times 10^9\ M_{\odot}}$ do not contain observable galaxies. This breaks the assumptions of the commonly used abundance matching method, which relies on the premise that structure formation can be represented by dark matter-only simulations and that every halo hosts a galaxy. \citet{2014MNRAS.442.2641V} found that gas expulsion and the associated dark matter expansion induced by supernova-driven winds are important for haloes with masses $\rm{M_{200} \le 10^{13}\ M_{\odot}}$, lowering their masses by up to 20 per cent relative to a DM-only model. \citet{2015MNRAS.451.1247S} found that the reduction in mass can be as large as 30 per cent of haloes with $\rm{M_{200} \le 10^{11}\ M_{\odot}}$. They also found that baryons can affect the inner density profile of dark matter haloes, leading to cuspier profiles in the centre due to the presence of stars. Feedback mechanisms triggered by baryons can expell gas from galaxies, even polluting voids with processed material, as suggested by \citet{2016MNRAS.457.3024H}. One would intuitively expect that if feedback processes are strong enough, voids would be more polluted with baryons and hence their properties might differ from their dark matter-only counterparts. It might also be the case that cooling alters the spatial distribution of gas around voids. This modification of the mass distribution could have consequences for the weak lensing signal measured around voids. In this work we study the effects of baryonic physics on the properties of cosmic voids. We search for voids using the mass field and subhalo tracers in the Evolution and Assembly of Galaxies and their Environment (EAGLE) simulations \citep{2015MNRAS.446..521S, Crain2015}, a suite of cosmological, hydrodynamical N-body simulations that follow the evolution of baryonic and dark matter particles in a $\Lambda \textrm{CDM}$ universe. EAGLE consists of multiple simulations that were run with different mass resolutions, volumes and physical models (the largest simulation consisting on a 100 Mpc on a side comoving box), being one of the first projects that allows the study of baryonic processes in a large simulated volume. The hydrodynamical simulations in the suite have counterparts that were run with the same initial conditions but only following the evolution of dark matter. By comparing these different models we can study the role of baryons in the context of the cosmic web. EAGLE presents a unique opportunity to explore the effect of baryons on void regions for various reasons. The simulations implement state-of-the-art subgrid models that follow star formation and feedback processes from stars and active galactic nuclei (AGN). These subgrid models, together with a high mass resolution, allow us to trace the distribution of gas and dark matter at different scales in detail, which is a key point in our study given the discussion above. Another advantage of the EAGLE simulations is that they reproduce the present day stellar mass function, galaxy sizes, and many other properties of galaxies and the intergalactic medium with very good precision (e.g. \citealt{2015MNRAS.446..521S, 2015MNRAS.450.4486F, 2016MNRAS.456.1235S, 2016MNRAS.456.1115B, 2016MNRAS.tmp..895T}). This good agreement with observations becomes important when trying to extrapolate results inferred from voids identified in these simulations to the real Universe. With EAGLE, we can also mimic some of the selections of galaxies in the observations that are used to find voids. In fact, we will show that voids found using galaxies selected by their stellar mass are different than those found if galaxies are selected by their total mass. We identify voids using a modified version of the algorithm presented in \citet{2005MNRAS.363..977P} (from hereon mP05), which searches for spherical under-dense regions in simulations, either using the mass field or halo tracers. Many void finders have been presented in the literature, which use different tracers to define voids (e.g. \citealt{2002MNRAS.330..399P, 2002ApJ...566..641H, 2002MNRAS.332..205A, 2005MNRAS.360..216C, 2007MNRAS.375..184B, 2007MNRAS.380..551P, 2008MNRAS.386.2101N}). Watershed based methods, such as the Watershed Void Finder \citep{2007MNRAS.380..551P} and ZOBOV \citep{2008MNRAS.386.2101N} are suitable when studying the spatial structure of the cosmic web in detail, as they are parameter free and they do not make assumptions about void shapes or topology. Voids identified with these finders usually exhibit a smooth transition from the under-dense region of a void to the average density. Other algorithms, such as the one we employ in this work, search for spherical regions that satisfy some density criteria. These void regions are usually characterised by a fast transition to the average density, and as a consequence, they are suitable for weak lensing studies (e.g. \citealt{2017MNRAS.465..746S}). In particular, mP05 has been shown to effectively capture differences between $\rm{\Lambda CDM}$ and $f(R)$ models using void statistics in \citet{2015MNRAS.451.1036C}. A comprehensive comparison of different void finding methods can be found in \cite{2008MNRAS.387..933C}. It is important to mention that here we adopt a cosmological model ($\rm{\Lambda CDM}$) and study the effects of baryons on that particular model. It is difficult to predict what the impact of baryon effects would be if we adopted a different cosmology, but future works could expand on this matter. The paper is organised as follows: In Section \ref{Section: Simulation and methods} we describe the EAGLE simulations, the void finding algorithm and our methodology to identify voids in the simulations. In Section \ref{Section: Visual inspection of baryon effects} we visually explore the effects on baryons on the large-scale distribution of matter in the simulations, emphasizing on the effects of galaxy feedback on void regions. We present the main results of void statistics in Section \ref{Section: Void abundance}, \ref{Section: Void profiles} and \ref{Section: Voids and their large-scale environment}. Conclusions and discussions about the work are presented in Section \ref{Section: Conclusions and discussion}. In the Appendix we show the evolution of some of the results found in previous sections at higher redshift, and we explain the methods employed for the calibration of the void finder and the error estimation in detail. | \label{Section: Conclusions and discussion} We have analysed cosmic voids the EAGLE simulations, using subhalo tracers and the mass field for void identification. We have studied the impact of baryonic physics on void statistics by comparing void populations in the main hydrodynamical EAGLE simulation (Ref-L0100N1504) and its DM-only counterpart (DM-L0100N1504), which was run using the same initial conditions but only following the evolution of dark matter. EAGLE provides a unique opportunity to test the effects of baryons on void properties due to its high resolution and detailed treatment of baryonic physics, using subgrid models that follow star formation, radiative cooling, stellar feedback from massive stars, black hole growth and AGN feedback. We define voids as spherical regions that have integrated density profiles equal to 20 per cent of the mean mass/subhalo density in the simulation. We have used a modified version of the spherical under-density void finder presented in \citet{2005MNRAS.363..977P}. When finding voids in the simulation with baryons using subhalo tracers, we used two different samples of subhaloes: The first one consists in a sample of subhaloes selected by stellar mass ($\rm{M_{stel} \ge 10^8\ M_{\odot}}$). The second sample consists in a set of subhaloes with the same number density, but selected by their total mass. In DM-L0100N1504, subhaloes are selected by their total mass, and the number density of subhaloes remains fixed and equal to the other two samples. We have calculated the void abundance, density and velocity profiles for all the samples described above. We find effects originating from feedback, which expels gas into voids and, more importantly, from the stochasticity in the stellar mass that manages to form in subhaloes, that is from the scatter of the stellar mass-dark matter mass relation. Our main results can be summarised as follows: \begin{itemize} \item DM-L0100N1504 produces about 24 per cent more voids than Ref-L0100N1504 when the mass field is used to identify voids. However, this difference is mainly comes from small voids with $1.5 < r_{\rm{void}} < 5\ \rm{Mpc}$ (Fig. \ref{Figure: Void abundance - mass field}). The action of feedback mechanisms from supernovae and AGN in Ref-L010015N04 pollutes void regions with hot gas, which causes some of these voids to shrink in size compared to their DM-only counterpart. If these regions shrink enough, they do not enter our sample of voids selected by our algorithm because their sizes fall below the minimum void size selection, which results in the hydrodynamical simulation having a lower overall abundance of voids than the DM-only simulation. This effect is also observed for voids with $r_{\rm{void}} \sim 10\ \rm{Mpc}$, although the associated statistical uncertainties are large. A bigger simulated volume is needed to test this effect with higher significance. \item There are no significant differences in the abundance of voids identified using subhaloes in Ref-L0100N1504 and DM-L0100N1504 when subhaloes are selected by their total mass in both simulations. However, selecting subhaloes by their stellar mass tends to produce slightly larger voids (Fig. \ref{Figure: Void abundance - subhaloes}). This discrepancy arises because a stellar mass cut tends to select subhaloes that have a stronger bias with respect to the underlying matter distribution. These subhaloes are located in higher density peaks of the cosmic web than those selected by their total mass, which in the end results in larger voids. These differences are also seen for snapshots of EAGLE corresponding to $z=0.5$ and $z=1.0$ \item When subhaloes are selected by their total mass, small differences are seen in the subhalo density profiles of voids in Ref-L0100N1504 and DM-L0100N1504 (Fig. \ref{Figure: Density profiles}). These differences are mainly caused by the effect of galaxy feedback on the subhalo mass function of the simulations. When taking this effect into account by computing the density profile with subhaloes of DM-L0100N1504 for voids identified in Ref-L0100N1504, the profiles are almost identical. When subhaloes are selected by their stellar mass, voids show density profiles with a more pronounced peak, reflecting an overabundance of subhaloes around void walls. This confirms the differences found in void abundance, namely voids traced by galaxies selected by their stellar mass are larger due to the fact that these tracers are more strongly clustered than those selected by their total mass. These differences are also seen at $z=0.5$ and $z=1.0$. \item The void mass density profiles show similar features as the profiles measured using subhaloes, meaning that subhaloes correctly capture the features of the mass distribution within and around voids with reasonable accuracy for a biased tracer of the density field. By measuring the void gas density profile we find an excess of gas near void walls, suggestive of the action of feedback mechanisms polluting voids with hot gas coming from galaxy winds. \item No significant differences are found in the velocity profiles of voids in the simulation with baryons and its DM-only counterpart, if the subhaloes are selected by their total mass. For the case where subhaloes are selected by their stellar mass, subhaloes evacuate void regions at slightly lower (5 - 10 km/s) velocities than in the total mass case. The void velocity profiles appear to be very sensitive to the overlap that is allowed between voids in our algorithm, while the associated statistical uncertainties are still too large to establish a robust conclusion about this trend. \item We are able to identify two different voids populations in EAGLE: voids embedded in under-dense large-scale environments that appear to be expanding, and voids in contracting dense environments. These resemble the void-in-void and void-in-cloud populations found in previous works. The effects of baryons appear to be more significant for the void-in-cloud, as these voids show distinct density and velocity profiles when the subhaloes by which they are traced are selected by their stellar or total mass. The fraction of void-in-void that we find is probably affected by the restricted number of large-scale modes that are present in these simulations. Nevertheless, this raises many questions about the relation between this void classification and the properties of the simulations in which these structures are identified. \end{itemize} It appears that void properties are well captured by DM-only simulations, with baryons only adding second-order effects, which are less important than those so far reported for alternative cosmologies. When identifying voids in the mass field, we find that a DM-only simulation tends to produce larger voids than a hydrodynamical one due to the effects of galaxy feedback. While this is an interesting result that sheds light on the impact of baryonic processes on the large-scale distribution of matter, it does not directly have an impact on observational studies of voids identified in galaxy surveys. We do find, however, that care must be taken with the galaxy tracers that are used to find and characterise voids, since they can have a strong impact on the properties that are inferred. The different results from stellar mass and subhalo mass cuts indicate that this check needs to be done for each void search, mimicking the number density of the tracer sample. Moreover, the bias of the galaxy sample needs to be reproduced by the haloes used to identify voids in a mock sample. The fact that these effects are also present at higher redshifts in the simulations suggests that they should be considered when studying voids in current and future large-scale galaxy surveys, especially in the context of studies that aim to constrain cosmological models from void statistics. The results of this work are relevant in particular to spherical-based void finders, which are more suitable in the case of weak-lensing studies. It is worth noting, however, that it is difficult to predict whether the differences that we find in this work can be extrapolated to other void finders, especially to those that are not spherical-based. Nevertheless, this opens the possibility for future works to expand on this matter. Given the size of the EAGLE simulation, we are not able to make predictions for voids larger than $\sim 25\ \rm{Mpc}$, which is comparable to voids identified in the main galaxy sample of SDSS. This is the best that can be done including all baryonic physics to date, but since we have found that the actual gas distribution is not as important as the galaxy population that is used to identify voids, in Paillas et al. (in prep) we will look at voids in different semi-analytic models of galaxy formation, which will be relevant to many more surveys. It is also worth noting that it is difficult to predict whether the differences that we find in this work can be extrapolated to other void finders. Considering the results obtained in this study, a suitable avenue for future works would be to combine DM-only simulations with abundance matching or semi-analytical models to populate the DM haloes with galaxies in order to characterise voids closely to what observers would do. | 16 | 9 | 1609.00101 |
1609 | 1609.06727_arXiv.txt | We use the high-resolution cosmological simulation \textit{Illustris} to investigate the clustering of supermassive black holes across cosmic time, the link between black hole clustering and host halo masses, and the implications for black hole duty cycles. Our predicted black hole correlation length and bias match the observational data very well across the full redshift range probed. Black hole clustering is strongly luminosity-dependent on small, 1-halo scales, with some moderate dependence on larger scales of a few Mpc at intermediate redshifts. We find black hole clustering to evolve only weakly with redshift, initially following the behaviour of their hosts. However below $z \sim 2$ black hole clustering increases faster than that of their hosts, which leads to a significant overestimate of the clustering-predicted host halo mass. The full distribution of host halo masses is very wide, including a low-mass tail extending up to an order of magnitude below the naive prediction for minimum host mass. Our black hole duty cycles, $f_{\rm{duty}}$, follow a power-law dependence on black hole mass and decrease with redshift, and we provide accurate analytic fits to these. The increase in clustering amplitude at late times, however, means that duty cycle estimates based on black hole clustering can overestimate $f_{\rm{duty}}$ substantially, by more than two orders of magnitude. We find the best agreement when the minimum host mass is assumed to be $10^{11.2} M_\odot$, which provides an accurate measure across all redshifts and luminosity ranges probed by our simulation. | \label{sec:intro} It is now widely understood that supermassive black holes are found at the centre of massive galaxies \citep{KormendyRichstone1995}, and that properties of the host galaxy strongly correlate with black hole mass \citep[e.g.][]{Magorrian1998, Gebhardt2000, Graham2001,Ferrarese2002, Tremaine2002, HaringRix2004, Gultekin2009, McConnellMa2013, KormendyHo2013}. One fundamental aspect of black hole studies is clustering behaviour, which provides a unique way of linking black holes to their host galaxies. Black hole clustering has been studied extensively in observations \citep[e.g.][]{LaFranca1998, Porciani2004, Croom2005, Shen2007, Myers2007, daAngela2008, Shen2009, Ross2009, White2012, Ikeda2015, Eftekharzadeh2015}, as well as simulations \citep[e.g.][]{Bonoli2009, Croton2009, DeGraf2010, DeGraf2012}. Since the emergence of large scale surveys capable of probing a range of redshifts, the general consensus is that the clustering signal decreases with time. At low redshift (below $z \sim 2$) the evidence for redshift evolution is generally weak, but at higher redshifts the evidence is much stronger, with correlation lengths approaching $10$ h$^{-1}$ Mpc \citep[e.g.][]{Myers2006, White2012} and bias factors (the clustering strength relative to that of the underlying dark matter density distribution) as large as $b = 5-10$ \citep[e.g.][]{Shen2009, Ikeda2015}. In addition to the redshift evolution, the possibility of luminosity dependent clustering has crucial implications for our understanding of the relation between black holes and their host halos. In particular, under the simplistic assumption that Active Galactic Nuclei (AGN) luminosity is proportional to host halo mass, one would expect brighter samples to be more strongly clustered (consistent with the stronger clustering of more massive halos). On the other hand, most models suggest a more widely varying black hole luminosity history, such that both bright and faint AGN can populate similar halos at different phases of their lifetimes, in which case clustering behaviour should only weakly depend on instantaneous luminosity. Many observations have found a lack of luminosity dependence \citep[e.g.][]{Croom2005, Myers2007, daAngela2008, White2012, Krolewski2015} or only a weak dependence \citep[e.g.][]{Shen2009, Eftekharzadeh2015}, supporting this model. Work by \citet{Bonoli2009} suggests, however, that even in the case of varying luminosity histories, a luminosity dependence could be found among lower-luminosity black holes, which simulations are well suited to investigate as observations being to push to lower flux limits. By matching quasar clustering to that of dark matter halos, the typical mass of the halos that host quasars can be estimated , providing a relatively simple means of estimating host properties for a range of black hole populations. By taking the expected number density of such halos and combining with the number density of AGN (via a luminosity function), one can estimate the active fraction of black holes or duty cycle \citep[see, e.g.][]{HaimanHui2001, MartiniWeinberg2001, Grazian2004, Shankar2010}. Such duty cycle estimates, however, rely upon several assumptions, most significantly the accuracy of the typical and/or minimum host halo mass calculated from clustering. Furthermore, these estimates rely upon the link between AGN luminosity and the host halo. Some models, however, suggest that only peak AGN luminosity correlates with host halo mass \citep[e.g.][]{Hopkins2005,Hopkins2005c}; in these models the scatter between low-luminosity lifetimes and host properties suggests that clustering should have a weaker luminosity dependence \citep[e.g.][]{Lidz2006} and that the assumptions used when estimating duty cycles may not hold. Large-scale cosmological simulations are well suited for investigating these aspects of black hole clustering, which we focus on in this paper. Here we use the state-of-the-art Illustris simulation \citep{Nelson2015} to study the clustering of supermassive black holes across cosmic time, taking advantage of the statistically representative sample provided by a large-volume simulation. The Illustris simulation is a (106.5 Mpc)$^3$ box, providing a sufficiently large sample to predict clustering behaviour in detail for $z = 0-4$, including dependence on black hole luminosity. The Illustris simulation has been shown to reproduce several key black hole properties, including black hole mass density, mass function, luminosity function, and scaling relations \citep{Sijacki2015}, making it ideally suited to investigate black hole clustering behaviour. In addition to showing the clustering via the black hole autocorrelation function, we compare the clustering amplitude via both the correlation length and bias parameters to observational results. Using the strength of the black hole clustering signal, we are able to estimate the typical mass of host halos similar to observational approaches, and directly compare to the actual distribution of host masses and explain the discrepancies therein. Similarly, clustering can be used to estimate the duty cycle of black holes \citep[see, e.g.][]{Eftekharzadeh2015}, which we compare directly to the actual duty cycle in the simulation, probing the accuracy of this estimate for several different definitions for duty cycle. The outline for our paper is as follows. In Section \ref{sec:method} we outline the numerical methods used, including the Illustris simulation project, and the clustering calculation used throughout the paper. In Section \ref{sec:results} we discuss the results of our investigation. Section \ref{sec:bhclustering} covers the clustering of black holes, their evolution with redshift, and dependence on black hole luminosity. In Section \ref{sec:hosthalo} we link the clustering behaviour to properties of host halos. In Section \ref{sec:dutycycle} we characterize the duty cycle of both black holes and halos, and the issues involved in estimated duty cycle from clustering behaviour. Finally, we summarize our conclusions in Section \ref{sec:conclusions}. | \label{sec:conclusions} In this work we have used the Illustris simulation to investigate the clustering of supermassive black holes across a range of redshifts and luminosities. In addition to general agreement with observations, we use the clustering information to link black hole properties to the host masses as well as the AGN duty cycle of black holes and galaxies in general. Our main conclusions are the following: \begin{itemize} \item AGN clustering is found to be luminosity dependent, but primarily at small (1-halo) scales. At larger scales, luminosity dependence primarily occurs at intermediate redshift, where black hole accretion tends to be strongest. \item Correlation length ($r_0$) can have significant luminosity dependence, especially at intermediate redshifts and when satellite black holes are included. $r_0$ reaches a minimum at $z \sim 1.5-2$, with higher-redshift evolution being strongest for low luminosity thresholds. Our $r_0$ estimates are generally lower than observational measures. This is largely due to the limited luminosity range in our simulation (imposed by the simulation volume), however, and adjusting observations to match our mean luminosities produces fully consistent results. \item Our estimated black hole bias matches observations very well at low-redshift. For $z > 2$ we predict a lower bias than \citet{Croom2005} and \citet{Shen2009}, but consistent with \citet{Eftekharzadeh2015}. \item AGN clustering tends to be stronger than the expected clustering of halos of comparable mass; as a result, AGN hosts tend to be less-massive than predictions made based on AGN clustering. Although strongest when satellite halos (found most commonly in the largest halos) are included, this effect remains even when only central black holes are considered. This suggests that typical host halo masses found based on clustering behaviour may be underestimated by a factor of $\sim 2$, especially at intermediate redshifts. \item The scatter in black hole-host scaling relations and typical black hole Eddington fractions results in a wide distribution of host halo masses. Although the distribution for any given $L_{\rm{BH,min}}$ does drop off at low halo mass, there does remain a low-mass tail, especially at low-redshifts. \item Due to AGN being more strongly clustered than halos matched to the typical hosts and both the wide range and low-end tail of the host halo distribution, estimating the minimum host halo mass from AGN clustering tends to substantially overestimate $M_{\rm{h,min}}$, which can have a strong impact on duty cycle estimates. \item At low redshift, the black hole duty cycle follows a power law in $M_{\rm{BH}}$, with a normalization set by the luminosity threshold. Higher redshifts also tend to follow a rough power law for $f_{\rm{duty}} < 0.5$, above which the curve flattens out. \item Black hole duty cycle decreases with time, well fit by a logistic function with lower $L_{\rm{BH,min}}$ thresholds decreasing more rapidly and at lower redshifts. \item Black hole duty cycle is well matched by the halo duty cycle for halos with $M_h > 10^{11.2} M_\odot$, representing a characteristic minimum mass for black hole occupation. \item Estimating the duty cycle from AGN number and expected halo number above a given $M_{\rm{h,min}}$ is very inaccurate. In addition to the mis-estimate of $M_{\rm{h,min}}$, the rapid growth of typical host halo masses at low-redshift produces a significant increase in the calculation of $f_{\rm{duty}}$ which is not found in the true black hole duty cycle. \item We used the AGN-galaxy cross-correlation function to look for a possible signature of AGN-induced quenching of satellite galaxies. Although $\xi_{\rm{QG}}$ does show $M_{\rm{BH}}$-dependent clustering of quenched galaxies, we find this signal is caused by the larger physical size of halos hosting massive black holes rather than a direct causal link. After controlling for halo size, we find no evidence for AGN inducing quenching in satellite galaxies. \end{itemize} Using the Illustris simulation, we have shown black hole and AGN clustering consistent with current observations, and characterized the luminosity dependence of AGN clustering, which is strongest at intermediate redshift ($z \sim 1.5-2$). One of the most important aspects of clustering analysis is the use of a clustering signal to characterize properties of the host halos, particularly the halo mass. We find that the typical approach taken (matching AGN clustering to analytic estimates for halo clustering) does very well at high-redshift, but can overestimate host mass by $\sim 50 \%$ at low-redshift, as low-redshift AGN are found to cluster more strongly than an equivalent-mass halo. Finally, we considered the use of AGN clustering as an estimator for black hole duty cycles. A typical method for this estimation is to assume a minimum host mass for a given AGN luminosity (see Equation \ref{eq:minmass}) and a constant duty cycle among halos above this threshold. Contrary to this assumption, however, we find a wide distribution of halo masses, including a low-mass tail. This scatter among host masses (as also found in other simulations) must be accounted for or the AGN duty cycle can be strongly overestimated, particularly at low-redshift. Overall, we find the black hole duty cycle to evolve smoothly with redshift, and we provide numerical fits characterizing this evolution as well as the dependence on black hole mass and AGN luminosity. In summary, our work highlights that while black hole clustering is a powerful probe of host halo properties, cosmological simulations, such as Illustris, are needed to fully characterize and account for a number of biases which would otherwise lead to systematically overestimated clustering-predicted host halo masses and black hole duty cycles. | 16 | 9 | 1609.06727 |
1609 | 1609.08817_arXiv.txt | Measurements of nuclear line emission of cosmic origins enable us to investigate where and how new interstellar nuclei are produced and released. \Al radioactivity with its characteristic $\gamma$-ray line at 1808.73~keV and decay time $\tau$ of $\sim$1~Myr shows ongoing nucleosynthesis in our Galaxy, and is ideal to trace how ejecta are recycled from their nucleosynthesis sources into next-generation stars. The \emph{CGRO} mission with the imaging Compton telescope instrument \emph{'COMPTEL'} obtained a sky map of \Al $\gamma$-ray emission \cite{Diehl:1993a,Pluschke:2001c}. The Galaxy's \Al content as a whole then can be considered as probably being in a steady state, as many individual and independent sources contribute, and star formation is uncorrelated among local star forming regions across the Galaxy. But already from the COMPTEL \Al image, showing a clumpy structure, it had been concluded that massive star groups are the dominating \Al producers in the current Galaxy, and may individually not be in a steady state; rather, the age of the specific stellar populations would determine the current amount of \Al in such a region \cite{Voss:2009}. Measurements of systematic Doppler shifts of the line with Galactic longitude \cite{Diehl:2006d,Kretschmer:2013} had shown that the observed \Al $\gamma$ rays originate from sources throughout the Galaxy, including its distant and otherwise occulted regions at and beyond the inner spiral arms and bulge. Moreover, the Doppler shifts of the \Al-line centroid energy were found larger than expected from large-scale Galactic rotation, and suggested that large cavities around massive star groups play a major role in guiding ejecta flows from massive-star and supernova nucleosynthesis \cite{Krause:2015,Krause:2016}. The study of \Al from specific regions first focused on Cygnus \cite{Knodlseder:2000,Pluschke:2000,Pluschke:2001a}, which stands out as an individual source region in the COMPTEL \Al skymap. Population synthesis allowed comparison of the predicted impacts of massive star groups onto their surroundings, including nucleosynthesis ejecta, and also kinetic energy from winds and explosions as well as ionising starlight, to observations in a variety of astronomical windows and tracers of such massive-star action \cite{Voss:2009}. Detailed population synthesis and multi-wavelength studies of the Cygnus region \cite{Martin:2008,Martin:2009,Martin:2010b} have been followed by studies of Carina \cite{Voss:2012}, Orion \cite{Voss:2010a}, and Scorpius-Centaurus \cite{Diehl:2010} regions. \begin{SCfigure} \centering \includegraphics[width=0.48\textwidth]{Fig_OriSpec_blueShift} \caption{\Al spectrum from the Orion region, as measured with SPI on INTEGRAL. Goodness of fit (through the $\chi^2$ value), detection significance, and line parameters are given in the legend. The blue/grey/light-grey-shaded regions show the 1$\sigma$/2$\sigma$/3$\sigma$ uncertainty range of the Gaussian fit to the line; the hatched region shows the range of blue shift that is indicated by this measurement, comparing to the laboratory value of \Al decay at rest (green dashed line). This suggests that \Al appears to stream within the Eridanus cavity from Orion OB1 stars towards the Sun.} \label{Fig_26Al_Orion}% \end{SCfigure} | 16 | 9 | 1609.08817 |
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1609 | 1609.01936_arXiv.txt | {Morphological characteristics of Boxy/Peanut bulges are studied, in particular whether most of the flux associated to bulges in galaxies with masses similar to those of the Milky Way at redshift z$\sim$0, could belong to vertically thick inner part of the bar, in a similar manner as in the Milky Way itself. At high galaxy inclinations such structures manifest as Boxy/Peanut/X-shape features, and near to face-on view as barlenses. We also study the possibility that bulges in some fraction of unbarred galaxies could form in a similar manner as the bulges in barred galaxies.} {We use the Spitzer Survey of Stellar Structure in Galaxies (S$^4$G) and the Near-IR S0 galaxy Survey (NIRS0S), to compile complete samples of galaxies with barlenses (N = 85), and X-shape features (N = 88). A sample of unbarred galaxies (N = 41) is also selected, based on similarity in their surface brightness profiles with those of barlens galaxies. Sizes and minor-to-major axis ratios (b/a) of these presumably vertically thick inner bar components are compared, and interpreted by means of synthetic images using N-body simulation models. Barlenses and their parent galaxies are divided into different sub-groups. Their possible parent galaxy counterparts in galaxies where the barlenses are manifested as X-shape features, are also identified.} {Unsharp mask images are created for all 214 sample galaxies. These images are used to recognize the X-shape features, and to measure their linear sizes, both along and perpendicular to the bar. For detecting possible boxy isophotes (using B$_4$ -parameter), isophotal analysis is also performed for the barlens galaxies. In the interpretation N-body simulations from \citet{salo2016} % are used: the models, exhibiting Boxy/Peanut/X/barlens morphologies, are viewed from isotropically chosen directions, covering the full range of galaxy inclinations in the sky. The created synthetic images are analyzed in a similar manner as the observations.} {This is the first time that the observed properties of barlenses and X-shape features are directly compared, over a large range of galaxy inclinations. A comparison with the simulation models shows that the differences in their apparent sizes , a/r$_{\rm bar} \gtrsim$ 0.5 for barlenses and a/r$_{\rm bar}$ $\lesssim$ 0.5 for X-shapes, can be explained by projection effects. Observations at various inclinations are consistent with intrinsic a$_{\rm bl} \approx$ a$_{\rm X} \approx$ 0.5 r$_{\rm bar}$: here intrinsic size means the face-on semimajor axis length for bars and barlenses, and the semilength of X-shape when the bar is viewed exactly edge on. While X-shapes are quite common at intermediate galaxy inclinations (for $i$ = 40$^\circ$ - 60$^\circ$ their frequency is $\sim$ half of barlenses), they are seldom observed at smaller inclinations. { This is consistent with our simulation models which have a small compact classical bulge producing a steep inner rotation slope, whereas bulgeless shallow rotation curve models predict that X-shapes should be visible even in face-on geometry. The steep rotation curve models are also consistent with the observed trend with B$_4$ being positive at low inclination, and getting negative values for $i$ $\gtrsim$ 40$^\circ$-60$^\circ$, thus implying boxy isophotes}. In total, only about one quarter of barlenses (with $i$ $\le$ 60$^\circ$) show boxy isophotes.} {Our analysis are consistent with the idea that barlenses and X-shape features are physically the same phenomenon. However, which of the two features is observed in a galaxy depends, not only on galaxy inclination, but also on its central flux concentration. The observed nearly round face-on barlens morphology is expected when at least a few percents of the disk mass is in a central component, within a region much smaller than the size of the barlens itself. Barlenses participate to secular evolution of galaxies, and might even act as a transition phase between barred and unbarred galaxies. We also discuss that the large range of stellar population ages obtained for the photometric bulges in the literature, are consistent with our interpretation. } | What is the amount of baryonic mass confined into the bulges of galaxies and how was that mass accumulated, is a critical question to answer while constructing models of galaxy formation and evolution. The answer to this question depends on how well the different bulge components can be recognized, and assigned to possible physical processes making those structures. Most of the bulge mass associated to photometric bulges (ie. flux above the disk) is generally assumed to reside in classical bulges. These are relaxed, velocity dispersion supported structures, presumably formed by galaxy mergers \citep{white1978,hopkins2009}, or by coalescence of massive star forming clumps at high redshifts, drifted towards the central regions of the galaxies (Bournaud et al. 2008; Elmegreen et al. 2009; see also review by Kormendy 2016). This picture has been challenged by the discovery that most of the bulge mass in the Milky Way actually resides in a Boxy/Peanut (B/P) bar, showing also evidence of an X-shape morphology, without any clear evidence of a classical bulge \citep{mcwilliam2010,nataf2010,wegg2013,ness2016}. Whether such bar-related inner structures could form most of the bulge mass also in external Milky Way mass galaxies is a topic of this study. Boxy/Peanut (B/P) bulges are easy to distinguish in the edge-on view and it has been shown that even 2/3 of all disk galaxies in S0-Sd types have B/Ps (L\"utticke, Dettmar $\&$ Pohlen 2000; Bureau et al. 2006; but see also Yoshino $\&$ Yamauchi 2014). Many B/P bulges also show cylindrical rotation \citep{kormendy1982,bureau1999,falco2006,falco2016,ianuzzi2015}, which generally confirms their bar origin. Verification of a galaxy as barred is difficult in the edge-on view, but it has been shown that, at an optimal range of viewing angles, B/Ps are visible even in less inclined galaxies, as revealed by their boxy isophotes (Beaton et al. 2007; Erwin $\&$ Debattista 2013, hereafter ED2013). A new morphological feature, a barlens (bl), was recognized by \citet{lauri2011}, and it has been suggested (Laurikainen et al. 2014, hereafter L+2014; Athanassoula et al. 2015, hereafter A+2015; see also Laurikainen et al. 2007) that they might be the face-on counterparts of B/P bulges. Association of a barlens to the Milky Way bulge has been recently made by \citet{gerhard2016}. Because of their fairly round appearance barlenses are often erroneously associated with classical bulges (see the review by Laurikainen $\&$ Salo 2016), but there is cumulative evidence showing that barlenses might indeed form part of the bar. Their optical colors are very similar to the colors of bars (Herrera-Endoqui et al. 2016, hereafter HE+2016), and in particular, their surface brightness profiles are very similar to those predicted for the B/P-bulges in hydrodynamical simulation models when viewed face-on (A+2015). The first indirect observational evidence connecting barlenses with B/P bulges (which often have X-shape features in unsharp mask images), was based on the axial ratio distribution of the combined sample of their parent galaxies, which appeared to be flat (L+2014). However, it remained unclear why barlenses concentrate on earlier Hubble types than the B/P/X-shape bulges (peak values are T = -1 and T = +1, respectively). Is this simply an observational bias when classifying galaxies at low and high inclinations, or could it indicate some intrinsic difference between the parent galaxies hosting barlenses and X-shape features? The latter possibility is suggested by the recent N-body simulations by Salo $\&$ Laurikainen (2016; submitted to ApJ), who demonstrated that a steep inner rotation curve leads to realistic-looking round barlens morphology, with no trace of an X-shape in the face-on geometry. However, reducing the central mass concentration, and thus shifting the galaxy to a later Hubble type, produced more elongated barlenses, which exhibited X-features at a much larger range of galaxy inclinations. As barred and unbarred galaxies presumably appear in similar galaxy environments (see Aquerri, M\'endez-Abreu $\&$ Corsini 2009), it is not plausible that bulges in barred galaxies form smoothly by secular evolution, and bulges in unbarred galaxies by some violent processes, like major galaxy mergers. Therefore, our hypothesis that many classical bulges are misclassified B/P/X features can be valid only if an explanation is found also for the bulges of unbarred galaxies, in the same line with the explanation for the barred galaxies. In fact, there is observational evidence which hints to that direction. Namely, the inner lenses (normalized to galaxy size) in unbarred galaxies are shown to have similar sizes as barlenses in barred galaxies (Laurikainen et al. 2013; Herrera-Endoqui et al. 2015, hereafter HE+2015). Inner lenses in unbarred galaxies might therefore represent evolved bars where the thin bar component has been completely dissolved, or the classical elongated bar never formed. However, whether those lenses are also vertically thick needs to be shown. In this study the properties of 85 barlenses and their parent galaxies are studied, and compared with the properties of 88 galaxies hosting bars with X-shape inner feature. An additional sample of 41 unbarred galaxies is also selected. As a database we use the Spitzer Survey of Stellar Structure in Galaxies \citep{sheth2010}, and the Near-IR S0 galaxy Survey \citep{lauri2011}. The properties of the analyzed features are compared with those obtained for synthetic images, created from simulation models taken from \citet{salo2016}. To obtain a fair comparison, the analysis for the synthetic images is done in a similar manner as for the observations. | we have shown evidence that barlenses at low galaxy inclinations are physically the same inner bar components as B/P/X-shape features in more inclined galaxies. Whether these structures are barlenses or show boxy/peanut/X-shape features depends, besides galaxy orientation, also on the central mass concentration of the parent galaxy. This is shown by comparing directly the properties of barlenses and X-shaped features, and is also verified by our simulation models. For two barlens galaxies detailed stellar populations and kinematics, given in the literature, are discussed in the context of the identified barlenses. The properties of these galaxies are also compared with those of the Milky Way bulge. We conclude that the stellar populations of barlenses in these galaxies are similar to those of the Milky Way bulge. | 16 | 9 | 1609.01936 |
1609 | 1609.03791_arXiv.txt | {Do some environments favor efficient conversion of molecular gas into stars? To answer this, we need to be able to estimate the \Ht\ mass. Traditionally, this is done using CO observations and a few assumptions but the Herschel observations which cover the Far-IR dust spectrum make it possible to estimate the molecular gas mass independently of CO and thus to investigate whether and how the CO traces \Ht. Previous attempts to derive gas masses from dust emission suffered from biases. Generally, dust surface densities, \Hi\ column densities, and CO intensities are used to derive a gas-to-dust ratio (\gdr) and the local CO intensity to \Ht\ column density ratio (\xco), sometimes allowing for an additional CO-dark gas component (\kdark). We tested earlier methods, revealing degeneracies among the parameters, and then used a sophisticated Bayesian formalism to derive the most likely values for each of the parameters mentioned above as a function of position in the nearby prototypical low metallicity ($12 + \log(O/H) \sim 8.4$) spiral galaxy M33. The data are from the IRAM Large Program mapping in the CO(2--1) line along with high-resolution \Hi\ and Herschel dust continuum observations. Solving for \gdr, $\xco$, and $\kdark$ in macropixels $500 \pc$ in size, each containing many individual measurements of the CO, \Hi, and dust emission, we find that ($i$) allowing for CO dark gas ($\kdark$) significantly improves fits; ($ii$) $\kdark$ decreases with galactocentric distance; ($iii$) \gdr\ is slightly higher than initially expected and increases with galactocentric distance; ($iv$) the total amount of dark gas closely follows the radially decreasing CO emission, as might be expected if the dark gas is \Ht\ where CO is photodissociated. The total amount of \Ht, including dark gas, yields an average $\xco$ of twice the galactic value of $\sciexp{2}{20}\Xunit$, with about 55\% of this traced directly through CO. The rather constant fraction of dark gas suggests that there is no large population of diffuse \Ht\ clouds (unrelated to GMCs) without CO emission. Unlike in large spirals, we detect no systematic radial trend in $\xco$, possibly linked to the absence of a radial decrease in CO line ratios. } | \TabNotation{} Recent work has shown that large-scale star formation in galaxies is strongly linked to the molecular gas reservoir, in particular the dense molecular gas, and less so to the total amount of neutral gas (\Ht\ + \Hi) \citep{Kennicutt.2012,Lada.2012}. If we are to understand what affects the relationship between molecular gas and star formation, we need to be able to measure the amount of molecular gas at all positions within the disk of galaxies, ideally down to the scale of individual star-forming regions. In low-metallicity objects, we are very far from such an understanding. The cosmic star-formation rate density rises rapidly with redshift \citep{Madau.2014}, suggesting that either or both the molecular gas content and the star-formation efficiency (mass of stars formed per unit time and unit \Ht\ mass) also increase while the fraction of metals decreases with redshift \citep{Combes.2013}. This is such that what we learn about local star formation at subsolar metallicities may be useful to better interpret observations of the young universe. The small Local Group spiral galaxy M33 has a half-solar metallicity and is near enough \citep[840\unit{kpc},][]{Galleti.2004} to resolve Giant Molecular Clouds (GMCs) and has an inclination ($i=56^{\circ}$) that makes the position of the clouds in the disk well defined (in contrast to e.g. M31). The whole bright stellar disk of M33 (up to a radius of $\sim 7 \kpc$) was recently observed in the CO\Jtwo\ line down to a very low noise level \citep{Druard.2014,Gratier.2010} using the IRAM 30 meter telescope on Pico Veleta. The single-dish CO\Jtwo\ data do not suffer from missing flux problems which is an essential asset to the understanding of the entire molecular phase in the galactic disk. M33 is a chemically young galaxy with a high gas mass fraction and as such represents a different environment in which to study cloud and star formation with respect to the Milky Way. As the average metallicity is subsolar by only a factor of two and the morphology remains that of a rotating disk, M33 represents a stepping stone towards lower metallicity and less regular objects. Measuring the link between CO and \Ht\ is particularly important given the evidence that the conversion of \Ht\ into stars becomes more efficient at lower metallicities \citep{Gardan.2007,Gratier.2010,Druard.2014,Hunt.2015}. With the advent of high resolution dust maps in the Herschel SPIRE and PACS, and Spitzer MIPS and IRAC bands it is possible to determine reliable dust column densities with spatial resolution close to the size of individual GMCs in M33 \citep[see][]{Kramer.2010, Braine.2010a,Xilouris.2012}. Under the assumption of local independence of the gas-to-dust ratio (\gdr) with respect to the \Ht/\Hi\ fraction, it is possible to determine the local CO intensity to \Ht\ column density ratio (\xco). A simplified global version of such an approach has been applied in \citet[Fig. 4 of ][]{Braine.2010a}. A more sophisticated method based on maximizing correlation between dust column density structure and that of the gas as derived from \Hi\ and CO through an optimal $\xco$ factor has recently been proposed and successfully demonstrated by \citet{Leroy.2011} and \citet{Sandstrom.2013}. However, these methods have biases and/or degeneracies which will be studied in Sects.~\ref{sec.NH2ICO} and \ref{sec.LSpresent}, in particular they often do not consider a possible contribution from CO dark molecular gas. In this work, the dust, CO, and \Hi\ data covering the disk of M33 are analyzed using existing these methods along with simulations to quantify bias and degeneracy. A new Bayesian approach is then used and tested in order to calculate the \gdr\ and $\xco$ for any position but also the amount of potential CO dark gas, unseen in \Hi\ or CO. All the methods take as a basic assumption that any gas not traced by CO, or potentially optically thick \Hi, contains dust with similar properties as in the gas traced by CO and \Hi. This is common to all other studies using dust emission. | In order to investigate how \gdr, \xco, and \kdark\ vary in M33, the first step was to take a published estimate of the gas column density $\NHtot_{dust}$ based on the Herschel dust observations and plot $\NHtot_{dust}-\NHi$ versus \ico. The systematically positive intercept (Fig.~\ref{fig.nht_ico_nocut_sub}) suggests that there is low-column density gas traced by dust but not CO or \Hi, which we refer to as $\kdark$ \citep{Tielens.1985, Planck-Collaboration.2011}. The next step is to construct a map of the dust surface density. Two methods were used -- the classical $\beta=2$ dust emissivity (Fig.~\ref{fig.dust}, left panel) and the variable-$\beta$ (same Fig., right panel) developed by \citet{Tabatabaei.2014}. We adopt the second method because in other subsolar metallicity galaxies \citep{Galliano.2011} the classical approach yields too large a dust mass, presumably due to a change in grain properties with respect to Milky Way dust. Using $\beta=2$ for M33 also yields a very high dust mass and \citet{Tabatabaei.2014} show that $\beta=2$ is a poor approximation for M33. We then look for optimal values of \gdr, \xco, and \kdark\ to relate the dust surface density to the \Hi\ and CO intensities. Except where the signal-to-noise ratio is high, major degeneracies are present between these parameters (Fig.~\ref{fig.scatter_mean_sub}) such that they all increase (or decrease) simultaneously with similar scatter in $\log(\gdr)$. Using simulated data with noise, a similar effect is seen in that the deduced solutions generally have lower \gdr, \xco, and $\kdark$ than the input values (Fig.~\ref{fig.LSCenter}--\ref{fig.LSOuterCut}). Setting \gdr\ to the correct (input) value yields reasonably accurate results. Solving only for \gdr\ and \xco, implicitly assuming $\kdark = 0$ when the input value was $\kdark = 5$ \Msunpsqpc, yields results for \gdr\ and \xco\ that strongly depend on the amount of CO with respect to \Hi. The degeneracies are illustrated by Figs~\ref{fig.LSnoise} and \ref{fig.LSimu}. An extremely computation-intensive simulation using the Bayesian errors-in-variables approach was used to obtain ``true'' values of the parameters. Fortunately, a very similar result can be obtained using the Bayesian formalism but without the errors-in-variables approach, as shown from the comparison in Fig~\ref{fig.compare}. The main difference is the slightly lower uncertainty with the errors-in-variables approach. The degeneracies present using the other methods are (almost) no longer an issue (Fig.~\ref{fig.Correlnew}). \FigCorrelNew{} There is a radial increase in \gdr\ from $\sim 200$ near the center to nearly 400 in the outer disk. The \xco\ ratio remains constant with galactocentric distance, as does the CO(2--1)/CO(1--0) line ratio \citep{Druard.2014} and CO(3--2)/CO(2--1) line ratio (in prep.), unlike what is observed in large spirals. The surface density of dark gas, $\kdark$, decreases from the center (10\Msunpsqpc) to the outer parts (roughly zero) in the same way as the CO emission such that the dark gas represents close to half of the \Ht\ assuming that the dark gas is in fact \Ht. As a result, the ratio of all \Ht\ (dark gas plus the \Ht\ traced directly by CO), is about twice the local value of $2 \times 10^{20}\Xunit$. Some traces of the degeneracies between $\kdark$ and \gdr\ are still present in that some macropixels with little CO find optimal values that are physically unrealistic (typically \gdr\ $\sim$ 5000 with a corresponding divergence of $\kdark$). Limiting the \gdr\ to values less than 500 (5 times the Milky Way value) avoids the problem. Overall, our results argue for a fairly high \gdr\ in M33 (\gdr\ $\ge 200$), a radially decreasing $\kdark$ roughly proportional to the amount of CO emission, and a fairly constant \xco\ conversion both of the \Ht\ directly traced by CO and the total \Ht\ content including the dark gas (whose radial distribution is similar to that of the CO). The results presented here on the link between CO and total molecular gas mass (and/or any optically thick \Hi) confirm the earlier estimates of the \Ht\ mass of M33. As a result, either the \Ht\ is converted into stars more quickly than in large spirals or the star-formation rate is overestimated due to for example a change in IMF in this environment. | 16 | 9 | 1609.03791 |
1609 | 1609.08160_arXiv.txt | We perform the first magnetohydrodynamical simulations of tidal disruptions of stars by supermassive black holes. We consider stars with both tangled and ordered magnetic fields, for both grazing and deeply disruptive encounters. When the star survives disruption, we find its magnetic field amplifies by a factor of up to twenty, but see no evidence for a self-sustaining dynamo that would yield arbitrary field growth. For stars that do not survive, and within the tidal debris streams produced in partial disruptions, we find that the component of the magnetic field parallel to the direction of stretching along the debris stream only decreases slightly with time, eventually resulting in a stream where the magnetic pressure is in equipartition with the gas. Our results suggest that the returning gas in most (if not all) stellar tidal disruptions is already highly magnetized by the time it returns to the black hole. | Stars of all kinds possess magnetic fields, thought to arise from an internal dynamo. These magnetic fields do not dominate the energy budget of stars: for example, the ratio of gas pressure to magnetic pressure is $\beta_{\rm M} \equiv 8\pi{P}/B^2 \sim 10^{6}$ throughout the bulk of the sun \citep{Dziembowski:1989a}, except for its corona where $\beta_{\rm M} \sim 1$ \citep{Babcock:1963a}. However, even relatively weak fields influence convection, mixing, and winds from stars. Moreover, in a tidal disruption event a star is severely distorted and twisted by the tidal field of a black hole \citep{Rees:1988a}; motions that may greatly affect the strength and configuration of the stellar magnetic field. Stellar mergers also spin up stars and produce streams of unbound material; in this way, they are closely analogous to tidal disruption events. Simulations of stellar mergers can find that the magnetic field amplifies by anything between a factor of $\sim 10$ to $\sim 10^{12}$, depending on the initial conditions and numerical techniques employed \citep[see e.g.][]{Price:2006a,Kiuchi:2014a,Zhu:2015a}. These results suggest that tidal disruption events could produce extremely strong magnetic fields, which could in turn influence their observational signatures. In this paper, we present the first simulations of the tidal disruptions of stars that include magnetic fields. In Section~\ref{sec:method} we outline our approach and initial conditions, followed by a presentation of our primary results in Section~\ref{sec:results} and a discussion in Section~\ref{sec:discussion}. | \label{sec:discussion} In this paper, we present the first MHD simulations of tidally disrupted stars. In the streams of unbound debris leaving the star, we find that field geometry straightens to lie parallel to the direction of stretching, and that the pressure of this field eventually dominates over both gas pressure and self-gravity. This breaks self-gravity in the streams, causing them to grow homologously after a time which depends on the initial field strength (equation~2). This may occur before hydrogen recombination, previously thought to be the only process to break self-gravity in the streams. This transition changes the interaction between the streams and their surroundings, with potentially observable consequences \citep{Guillochon:2016b,Chen:2016b,Romero-Canizales:2016a}. The field configuration of any disk-like structure that forms from the debris will likely be toroidal, with periodic reversals in direction (clockwise, then anti-clockwise, etc.) with each wrap-around of the stream about the black hole. Such a configuration is not optimal for powering jets \citep[although spinning black hole may offer a path for converting toroidal to poloidal flux, see ][]{McKinney:2013a}, and because the flux is not amplified by the tidal stretching process but merely preserved, it is still likely that another mechanism is required to yield the $\sim 10^{29}~{\rm G}~{\rm cm}^2$ of flux required to power a jet \citep{Kelley:2014a}. But while the total flux is not increased within the debris, the parity between magnetic and gas pressures suggests that magnetohydrodynamic effects are likely crucial for understanding the subsequent evolution of the debris streams. The growth in field strength could influence the exchange of energy and angular momentum at the stream-stream collision point, leading to faster circularization times \citep{Bonnerot:2017a}. The magnetic field also offers the stream some protection from disruption via fluid interactions with ambient medium. Heat conduction into the stream will be suppressed in directions perpendicular to the magnetic field direction \citep{Dursi:2008a,ZuHone:2013a}, as well as the Kelvin-Helmholtz instability at the stream's surface \citep{McCourt:2015b}, both of which may improve the ability of infalling clouds that may be produced in disruptions to survive through periapse and beyond \citep{Guillochon:2014b}. For the surviving core, the amplification of about an order of magnitude suggests that repeated stellar encounters with the black hole, which arise naturally after a partial disruption \citep{MacLeod:2013a}, may yield stars that are highly magnetized. Whereas a partially disrupted star without a magnetic field will rejoin the Hayashi track and remain bright for a Kelvin time \citep{Manukian:2013a}, $\sim 10^4~{\rm yr}$, the inclusion of a magnetic field may permit the star to remain bright for much longer as the magnetic field slowly unwinds within the star and deposits heat \citep{Spruit:2002a}, potentially tens of millions of years. If a dynamo process acts within a partially disrupted star, repeated encounters may not be required, which would suggest that many thousands of tidally magnetized stars could lurk near the centers of galaxies. One piece of evidence for a large population of highly magnetized stars in our own galactic center would be the excess of X-rays their coronae would produce \citep{Sazonov:2012a}. Our simulations show that the influence of the magnetic field on the stream evolution and the stellar evolution of any surviving core are critically important to understanding the resulting dynamics and observability of tidal disruption events. In the future, simulations of stream-stream collisions and tidal disruption disk formation that include the strength and geometry of the strong fields we find here should be performed, as well as high-resolution simulations of the partial disruptions of stars to try to resolve any potential dynamo process. \bigskip | 16 | 9 | 1609.08160 |
1609 | 1609.03875_arXiv.txt | The long quest for the firm disentanglement among leptonic and hadronic scenarios represents one of the most important challenges for the high-energy study of SNRs, being directly related to cosmic rays origin and acceleration models. Accurate radio images of SNRs are typically available at low frequencies (Castelletti G. et al. $2007$ and $2011$; Gao et al. $2011$). On the other hand, multi-wavelength data on SNRs are sparse and spatially-resolved spectra are rarely available in the 5-20 GHz range, critical for model assessment, also for the most studied and bright objects (see Green D.A. $2014$). Deep multi-frequency imaging of the complex SNRs IC443 and W44 with the Sardinia Radio Telescope (SRT) can disentangle different populations and spectra of radio/gamma-ray-emitting electrons in these SNRs, in order to better address models and then firmly constrain high-energy emission arising from hadrons. On the other hand, accurate radio spectral imaging allows us to distinguish between shock parameters and different physical processes taking place within SNRs. Recent constraints on cosmic rays emission from SNRs and related models (Giuliani et al. $2011$; Ackermann et al. $2013$; Cardillo et al. $2014$) are based on integrated radio fluxes only (no spatially resolved spectra) implying the simplistic ”single-zone” assumption of a single electron population for the whole SNR. In the aim of studying the local properties of W44 and IC443, we performed with SRT accurate on-the-fly scans of these SNRs at three frequencies (L, C, K bands) in order to obtain detailed radio images and spatial-resolved spectral-slope measurements (synchrotron breaks are possibly expected in this range). Indeed, spectral index maps provide evidence of a wide physical parameters scatter among different SNR regions. | 16 | 9 | 1609.03875 |
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