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1808 | 1808.07318.txt | The composition of ultra-high energy cosmic rays is still poorly known and constitutes a very important topic in the field of high-energy astrophysics. Detection of ultra-high energy cosmic rays is carried out via the extensive air showers they create after interacting with the atmosphere constituents. The secondary electrons and positrons within the showers emit a detectable electric field in the kHz-GHz range. It is possible to use this radio signal for the estimation of the atmospheric depth of maximal development of the showers \xmax, with a good accuracy and a duty cycle close to $100\%$. This value of \xmax\ is strongly correlated to the nature of the primary cosmic ray that initiated the shower. We show in this paper the importance of using a realistic atmospheric model in order to correct for systematic errors that can prevent a correct and unbiased estimation of~\xmax. | \label{sec:intro} Recently a lot of efforts have been put into determining the mass composition of cosmic rays using the radio signal \cite{Bezyazeekov:2015ica,2016Natur.531...70B,gatearena2016}. Several methods exist by now with different approaches but the goal is the same: reconstructing the atmospheric depth of the shower maximum, \xmax, where the number of particles is maximum. This atmospheric depth is highly correlated to the mass of the primary cosmic ray. To be competitive, the uncertainty on its estimation should be close to or better than that achieved with the fluorescence technique ($\sim 20~\gcm$, see~\cite{xmaxpao2010}). The composition of the highest-energy cosmic rays (above $1$~EeV) is still poorly known, since it is difficult to measure composition using a surface detector that only samples the shower at ground level. Besides, the fluorescence light technique, more apt for composition measurements, has a duty cycle of the order of $14$\%~\cite{fdpao2010}, making it difficult to provide \xmax\ measurements for a large number of showers at the highest energies. The radio technique, consisting in the measurement of the electric field induced by the extensive air showers created by cosmic rays, could be an excellent alternative to obtain the \xmax\ with an almost $100$\% duty cycle. Extracting the \xmax\ using the radio signal relies on an atmospheric model. The electric field emission is highly beamed towards the direction of propagation of the shower and the shape of its distribution at the ground level depends on the distance between the point of maximum emission and the shower core. This property can be exploited to reconstruct \xmax\ from the radio signal. However, to make this method accurate, one needs to know the atmospheric depth corresponding to a given distance with precision. The electric field measured by the antennas strongly depends on the characteristics of the atmosphere in which secondary shower particles evolve: air density, air refractive index at radio frequencies, temperature, pressure and humidity. For a long time, simulation codes computing this electric field assumed a standard atmosphere. Nowadays, with high precision measurements on large radio arrays running continuously such as AERA~\cite{glaserarena2016}, it has become important to refine this atmospheric model. Indeed, it is clear that the atmospheric characteristics vary significantly with time (day/night effect and seasonal variations) and these variations are responsible for systematic uncertainties that can prevent an accurate estimation of the \xmax. Ideally, we need to know the atmospheric state at the time a shower is detected. This is possible using the Global Data Assimilation System~\cite{gdas} (GDAS) data. In this paper, we show how we use these data together with a standard atmospheric model for the highest altitudes to compute an accurate air density model as a function of altitude at the time of the detection of the event. The knowledge of the air density and humidity ratio also allows to compute the realistic air refractive index which is needed for the amplitude and time structure of the signal. % and for the Cherenkov ring diameter. Several descriptions of the atmosphere are in use in different simulation codes such as SELFAS~\cite{selfas2011}, ZHAireS~\cite{zhaires2012} and CoREAS~\cite{coreas2013}. % These descriptions are summarized in \tab~\ref{tab:tab1}. We show that the choice of the atmospheric model induces uncertainties in the atmospheric depths up to some tens of $\gcm$ which is comparable to the uncertainty on the \xmax\ obtained with the fluorescence data. %{\bf A GERER LES FORMULES $N(h)=Ns \exp(-Kr*h)$, where $Ns=325$ and $Kr=1.1218$/km} \\ The paper is organized as follows. In section~\ref{sec:geo}, we briefly present the geometrical description of the shape of the Earth and its atmosphere and the atmospheric depths computations. In section~\ref{sec:PCasp} we describe the GDAS data and its use to build a realistic atmospheric model that we will use to calculate the atmospheric depths and the air refractive index. We compare the results with those obtained assuming the basic US Standard model \cite{Usstd}. In section~\ref{sec:influence} we quantify the influence of the air refractive index and air density profiles calculated with the GDAS data on the produced electric fields. Then, in section~\ref{sec:recoxmax} we study the case of a simulated shower which develops in the atmospheric conditions of a sample day. We show that using the US Standard model on the \xmax\ estimation leads to biased results, unless we use the same atmospheric conditions than those of the day and time of the detected (here simulated) event. In this paper, we will note $\mathbf{V}$ the shower axis and $\mathbf{B}$ the geomagnetic field. | In this work, we have studied the influence of the description of the atmosphere on the electric field emitted by air showers and its effect on the reconstruction of the properties of the primary cosmic ray using the radio technique. %The radio signal allows to reconstruct the arrival direction of the primary cosmic ray, its energy, and the shower \xmax, with a $\sim 100\%$ duty cycle. In order to reach the required accuracy to be a competitive technique, we need to describe the atmosphere in a very precise way. With this objective, we have demonstrated the need to use a spherical geometry for the Earth and its atmosphere: the flat approximation leads to systematic errors larger than $10~\gcm$ for zenith angles above~$60^\circ$. %The electric field amplitude at ground level is also enhanced with the spherical description with respect to the flat approximation. %After that, we have built an atmospherical model in order to get the air density as a function of altitude $\rho(z)$; for free, we also get the air refractivity for the radio waves domain that we use in SELFAS. We used for that the GDAS data between the ground level and an altitude of $\sim 26$~km together with the US Standard model above $26$~km up to $\sim 100$~km. After that, we have used the Global Data Assimilation System (GDAS), which provides information on the atmospheric pressure, temperature and humidity for a range of altitudes every three hours. These three quantities allow us to know the density of the atmosphere and its refractive index, both of which are crucial for a correct simulation of the development of an extensive air shower and the calculation of the electric field it produces. Since the data provided by the GDAS are available up to $26$~km of altitude, our atmospheric model is a mixture of GDAS data below $26$~km (which is the most important region for the development of air showers), and the usual US Standard atmosphere above $26$~km. The atmospheric refractivity has been calculated with two formulas: the usual Gladstone-Dale (GD) formula, that does not take humidity into account, and a high-frequency (HF) formula that is more suited for radio frequencies (MHz-GHz) and takes the relative humidity as an input. The differences in refractivity between using a US Standard atmosphere coupled with the GD formula on one side and the GDAS atmospheric data with the HF formula are of $~15\%$ on average at $1$~km of altitude, and it can reach up to $35\%$ in the lowest layers of the atmosphere. We have studied the influence of the refractivity on the time traces and the lateral distribution function (LDF) of the electric field produced by air showers. When considering the $[20;80]$~MHz band, differences in refractivity up to $20\%$ result in a relatively small difference in the amplitude of the electric field and hence the LDF, indicating that at these frequencies, an accurate knowledge of the refractivity is not the most important factor for reconstructing the properties of the primary cosmic ray. However these small differences in the LDF vary as a function of the axis distance (+2\% to +8\% when increasing the refractivity by 20\%) leading to a shift in the reconstructed \xmax\ value. Moreover, when inspecting the $[120;250]$~MHz band, the shape of the time traces for the electric fields and the LDF on the ground change appreciably with the refractivity, making the reconstruction at high frequency more dependent on the correct knowledge of the atmospheric refractivity. In turn, if we can provide a precise refractivity, the $[120;250]$~MHz band presents the advantage that the electric field footprint on the ground varies dramatically with the shower maximum. In particular, the Cherenkov ring is clearly visible at these frequencies and can help us discriminating the position of the shower maximum. Finally, we have compared the performance in the reconstruction of the shower maximum with the several atmospheric densities (US Standard and GDAS) and refractivities (US Standard coupled with GD, GDAS coupled with GD and GDAS with the HF formula with relative humidity) available. We have used test events simulated with the GDAS density and HF refractive index, in order to quantify the error induced with the US Standard atmosphere and the GD formula if we assume the GDAS data are closer to the actual atmosphere. We have found that the most important parameter for the reconstruction of the shower maximum is the air density, since even if we correctly reconstruct the altitude of the shower maximum, an incorrect air density will bias the atmospheric depth of the \xmax. The bias induced with a US Standard air density lies around $\sim 30$ g/cm$^2$ for $30^\circ$ showers and $\sim 100$ g/cm$^2$ for $60^\circ$ showers. The bias induced by the refractivity calculated with the US Standard atmosphere and the GD index ranges from $\sim 5$ g/cm$^2$ for $30^\circ$ showers to $\sim 32$ g/cm$^2$ for $60^\circ$ showers. These biases are not negligible and indicate the need for a correct description of the atmospheric properties. The theoretical accuracy of the method, using the GDAS data and without taking into account uncertainties in the modelling of the electric field of the shower or the atmospheric parameters, is $\sim 2.4$ g/cm$^2$ for $30^\circ$ showers and $\sim 10$ g/cm$^2$ for $60^\circ$ showers. These accuracies constitute a theoretical limit for the precision of the \xmax ~reconstruction using the method discussed in this paper. %For hadronic air showers, the influence of this refractivity model can reach $10~\gcm$ at various zenith angles compared to less refined models such as the Gladstone-Dale law. It could be more important for electromagnetic showers for the gamma astronomy community through its influence on the Cherenkov ring diameter. Using an example day together with a simulated shower in the actual atmospheric conditions, we show that using the US Standard model instead of the actual atmospheric conditions leads to an error on the reconstructed \xmax\ of $34~\gcm$. Note that the reconstructed distance to the point of maximum emission is the same, using either the actual atmospheric conditions or those of the US Standard model. Only the conversion to the \xmax\ is affected by the choice of the model. To sum up with, the results of this paper indicate that a description of the atmosphere using the US Standard model paired with the GD formula cause non-negligible biases when reconstructing the \xmax, and therefore an alternative description is needed. The most complete description of the atmosphere publically available is the GDAS data, from which we can trivially calculate the properties of the atmosphere relevant for the simulation of the electric field produced by air showers. In doing so, we guarantee the minimum possible bias in the simulation of the electric field and the reconstruction of the shower maximum. Currently, the only way of improving this method is to directly measure the atmospheric properties for a given experiment \emph{in situ}. %We have therefore a fully coherent way to treat the electric field computations: use the GDAS data and the US Standard model to build a model of the air density and refractivity from ground level up to $\sim 100$~km of altitude, in a spherical geometrical description of the Earth and its atmosphere. | 18 | 8 | 1808.07318 |
1808 | 1808.01488_arXiv.txt | The strong enhancement of the ultraviolet emission during solar flares is usually taken as an indication of plasma heating in the lower solar atmosphere caused by the deposition of the energy released during these events. Images taken with broadband ultraviolet filters by the {\em Transition Region and Coronal Explorer} (TRACE) and {\em Atmospheric Imaging Assembly} (AIA 1600 and 1700~\AA) have revealed the morphology and evolution of flare ribbons in great detail. However, the spectral content of these images is still largely unknown. Without the knowledge of the spectral contribution to these UV filters, the use of these rich imaging datasets is severely limited. Aiming to solve this issue, we estimate the spectral contributions of the AIA UV flare and plage images using high-resolution spectra in the range 1300 to 1900~\AA\ from the Skylab NRL SO82B spectrograph. We find that the flare excess emission in AIA 1600~\AA\ is { dominated by} the \ion{C}{4} 1550~\AA\ doublet (26\%), \ion{Si}{1} continua (20\%), with smaller contributions from many other chromospheric lines such as \ion{C}{1} 1561 and 1656~\AA\ multiplets, \ion{He}{2} 1640~\AA, \ion{Si}{2} 1526 and 1533~\AA. For the AIA 1700~\AA\ band, \ion{C}{1} 1656~\AA\ multiplet is the main contributor (38\%), followed by \ion{He}{2} 1640 (17\%), and accompanied by a multitude of other, { weaker} chromospheric lines, with minimal contribution from the continuum. Our results can be generalized to state that the AIA UV flare excess emission is of chromospheric origin, while plage emission is dominated by photospheric continuum emission in both channels. | \label{sec:intro} Uultraviolet (UV) emission {from solar flares in the 1000--2000~\AA\ range} is often regarded as evidence of plasma heating in the lower solar atmosphere, highlighting the location of the energy deposition, { which is often coincident with} hard X-ray emission \citep[e.g.][]{WarrenWarshall:2001,AlexanderCoyner:2006,FletcherHudson:2001,SimoesGrahamFletcher:2015a} and mapping the connectivity of the flaring magnetic loops \citep[e.g.][]{Fletcher:2009,JoshiVeronigCho:2009,ReidVilmerAulanier:2012}. UV imaging of solar flares by the {\em Transition Region and Coronal Explorer} \citep[TRACE;][]{HandyActonKankelborg:1999} and the {\em Atmospheric Imaging Assembly} \citep[AIA;][]{LemenTitleAkin:2012} on board the {\em Solar Dynamics Observatory} \cite[SDO;][]{PesnellThompsonChamberlin:2012} has been used extensively to investigate the spatial morphology and temporal evolution of flare ribbons \citep[e.g.][]{WarrenWarshall:2001,FletcherPollockPotts:2004,SimoesFletcherHudson:2013,KazachenkoLynchWelsch:2017}. { Due to a lack of knowledge of the spectral content of the AIA UV images, detailed qualitative analyses of flaring plasma using AIA UV images have been rare.} A few exceptions are the investigation of plasma cooling \citep{QiuLiuHill:2010,ChengKerrQiu:2012,QiuSturrockLongcope:2013} and estimates of the radiative output { from the flaring chromosphere} \citep{MilliganKerrDennis:2014}. In contrast, the spectral content in the AIA extreme ultraviolet (EUV) images has been extensively explored. The radiation of the vast majority of spectral lines in the EUV range can be assumed to be optically thin. This assumption greatly facilitates the analysis of EUV data by the use of synthetic spectra generated by CHIANTI \citep{DereLandiMason:1997,LandiYoungDere:2013}, which has led to significant advancements in understanding the evolution of hot plasmas in the Sun \citep[e.g.][]{YoungDel-ZannaMason:2007,MilliganChamberlinHudson:2012,Del-ZannaWoods:2013}, and identifying the response of the EUV filters on AIA \citep{ODwyerDel-ZannaMason:2010,MilliganMcElroy:2013}. This knowledge of the spectral response of the AIA EUV channels has led to the development of techniques to estimate the differential emission measure of the plasma \citep[e.g.][]{HannahKontar:2012,Del-Zanna:2013,CheungBoernerSchrijver:2015}, and several other studies of the physical characteristics of the plasma in active regions \citep[e.g.][]{SchmelzJenkinsWorley:2011,Del-Zanna:2013}, microflares \citep[e.g.][]{InglisChriste:2014,WrightHannahGrefenstette:2017} and flares \citep[e.g.][]{HannahKontar:2013,BattagliaFletcherSimoes:2014,SimoesGrahamFletcher:2015a}. In the UV range, however, the presence of recombination edges and spectral lines originating from neutral and weakly-ionized metals, which are unlikely to be optically thin, prevents the direct use of synthetic spectra for similar analysis. One of the purposes of the UV filters on TRACE ({ 1550, 1600 and 1700~\AA}) was to estimate the \ion{C}{4} emission \citep{HandyBrunerTarbell:1998}, given the expected radiative importance of this line \citep[e.g.][]{BrunerMcWhirter:1988,HawleyFisher:1994}, and the potential of using its flux as a transition region `pressure gauge' \citep{HawleyFisher:1992}. TRACE \ion{C}{4} images have been explored for studies of quiescent structures \citep{de-WijnDe-Pontieu:2006,De-PontieuTarbellErdelyi:2003}, but not for flare studies, to the best of our knowledge. {Similarly, TRACE UV filters were also used to deconvolve the \ion{H}{1} Lyman-$\alpha$ channel and the estimate the flare emission of this line \citep{Rubio-da-CostaFletcherLabrosse:2009}.} The spectral content in the { AIA} UV images has seldom been explored { due to the lack of high-resolution flare spectra in this wavelength range}. The few instruments that have covered this spectral range in the past include the SO82B spectrograph \citep{BartoeBruecknerPurcell:1977} on the Naval Research Laboratory/Apollo Telescope Mount (NRL/ATM) on {\em Skylab}, the {\em SOLar STellar Irradiance Comparison Experiment} (SOLSTICE) \citep{BrekkeRottmanFontenla:1996}, and {\em Solar Ultraviolet Measurements of Emitted Radiation} (SUMER), on the {\em Solar and Heliospheric Observatory} (SOHO) mission \citep{GontikakisWinebargerPatsourakos:2013}. In 1973--1974, the SO82B spectrograph made observations of a few flares that provided the means for remarkable advances in our understanding of the flaring atmosphere, the formation of spectral lines and continua, emission measure distribution, electron densities, and mass motions via Doppler-shift analysis \citep[e.g.][]{DoschekVanhoosierBartoe:1976,FeldmanDoschekRosenberg:1977,Kjeldseth-MoeNicolas:1977,DereHoranKreplin:1977,CanfieldCook:1978,Withbroe:1978,ChengKjeldseth-Moe:1978,Cheng:1978,LitesCook:1979,CookBrueckner:1979,DereCook:1979,DoyleWiding:1990,WidingDoyle:1990}. This knowledge of the UV flare spectrum has not been fully applied to the analysis of more recent UV imaging. The {\em Interface Region Imaging Spectrograph} \citep[{\em IRIS};][]{De-PontieuTitleLemen:2014} covers the UV ranges 1332--1407~\AA\ and 2783--2837~\AA, with high spatial and spectral resolution. While IRIS has allowed many advances in our comprehension of flare physics \citep[e.g.][]{GrahamCauzzi:2015, KerrSimoesQiu:2015}, there is no overlap with the wavelength range covered by AIA. { However, it is worth noting that efforts have been made to identify the spectral content in the IRIS broadband 2832~\AA\ slit-jaw images (SJI). \cite{KleintHeinzelKrucker:2017} convolved the observed IRIS flare spectra with the SJI 2832~\AA\ response function and found that spectral lines have a significant contribution to the images, along with a possible contribution from the \ion{H}{0} Balmer continuum. \cite{KowalskiAllredDaw:2017} also identified that many \ion{Fe}{2} lines account for a significant contribution during flares and their models suggest that excess intensity in SJI 2832~\AA\ have similar contributions from \ion{Fe}{2} lines and continuum emission. } In this paper, we use high-resolution spectra in the UV range 1400--1960~\AA\ provided by the Skylab NRL SO82B spectrograph to estimate the spectral contributions to the AIA UV passbands, both during flares { and in quiescent plage regions. Definitive identification of the dominant spectral features, and their respective formation temperatures, will provide valuable knowledge on where in the solar atmosphere the observed emission is formed.} This will { potentially allow} plasma diagnostic information to be obtained from these two filters { in the future}. The calibration method of the SO82B data is described in the Appendix \ref{ap:skylab}. | We have estimated the contributions of UV spectral lines and continua to the 1600~\AA\ and 1700~\AA\ passbands on AIA. Given the lack of current spectroscopic data in the 1300--1900~\AA\ range, we employed observations of the flare SOL1973-09-07 with high spectral resolution made by the NRL SO82B spectrograph on Skylab. We find that the flare excess emission in both AIA UV filters is dominated by spectral lines formed in the chromosphere and transition region, with a temperature range $4.2 < \log T < 5.1$. We also confirm that quiescent plage emission captured by these filters is dominated by the continuum formed in the photosphere. The calibrated spectrum of the flare and a plage region (used as a reference spectrum during quiescent times) were convolved with the AIA filter responses and the { relative} contributions of the main spectral lines and continua for the flare excess were obtained. We find that in the AIA 1600~\AA\ images the flare excess is composed of 26\% of \ion{C}{4} 1550~\AA, 20\% of \ion{Si}{1} continua, 11\% of \ion{C}{1} 1656~\AA, 9\% of \ion{He}{2} 1640~\AA, 7\% of \ion{C}{1} 1561~\AA, 5\% of \ion{Si}{2} 1526 \& 1533~\AA, plus 21\% from hundreds of weaker lines from neutral and weakly-ionized metals. For AIA 1700~\AA\ images, the flare excess is formed of 38\% \ion{C}{1} 1656~\AA, 17\% \ion{He}{2} 1640~\AA, 12\% from the combined contributions from \ion{Al}{2} 1671~\AA, \ion{Fe}{2} 1713~\AA, \ion{Si}{2} 1808~\AA, and 1817~\AA, 28\% from weaker lines from neutral and weakly-ionized metals, and only 6\% from \ion{Si}{1} $^1$D continuum. The plage emission is dominated by continuum emission: 67\% in the AIA 1600~\AA\ filter and 87\% in the 1700~\AA\ filter. This is expected given the typical photospheric features often seen in the AIA UV images of the quiet Sun. Our results should be taken as a guideline for the interpretation of the flare emission in the UV bands observed by AIA. These should by no means be taken as the definitive fractions of the flare spectral components in those filters, since flare emission varies dynamically during an event, and also from event to event, so the relative contributions are likely to change. However, these variations should not be extreme, i.e. the \ion{C}{4} 1550~\AA\ doublet and other chromospheric lines should be the main contributors to the flare excess emission detected by AIA 1600~\AA. At the same time, the \ion{C}{1} multiplet at 1656~\AA\ and \ion{He}{2} 1640~\AA\ should account for the largest fraction of the flare excess emission into the AIA 1700~\AA\ passband, with negligible contribution from continua. Since spectral observations with high-resolution in this UV range are rare and thus the evolution of the UV spectrum during flares is still largely unknown, radiative hydrodynamic simulations should be used to examine the validity of our results. However, most of the lines in this range are not considered to be energetically important in the solution of the radiative transfer of the flaring plasma and therefore tend to be ignored in the current solutions. The inclusion of \ion{C}{0} and \ion{Si}{0} atomic database in current tools for radiative hydrodynamic modeling should allow us to revisit previous works regarding the behavior of the main spectral lines such as \ion{C}{4} and \ion{C}{1}, and \ion{Si}{1} continua during flares. Lines of neutral or weakly ionized metals such as \ion{C}{1}, \ion{Si}{1}, \ion{Fe}{2}, \ion{Si}{2}, and \ion{C}{2} are abundant and strongly enhanced during flares \citep{DoyleCook:1992}, and have been useful in the construction of semi-empirical flare atmospheric models \citep[e.g.][]{LitesCook:1979,MachadoAvrettVernazza:1980}. Lines of some of these elements are currently being observed by IRIS \citep[e.g. \ion{Si}{4} 1393.755~\AA and 1402.770~\AA, \ion{C}{2} 1334.535~\AA\ and 1335.708~\AA; see][]{PolitoReepReeves:2016,WarrenReepCrump:2016,ReepWarrenCrump:2016}, and could potentially be used to estimate the intensity of the lines of the same elements at longer wavelengths, within the AIA UV range. Armed with this new information about the spectral composition of AIA UV images, joint analysis of IRIS and AIA UV data may help to reveal more about the UV spectrum of flares and quiescent solar features and also to explore the AIA data since 2010 for long-term variability of several solar structures \citep[e.g.][]{Oliveira-e-SilvaSelhorstSimoes:2016}. | 18 | 8 | 1808.01488 |
1808 | 1808.05990_arXiv.txt | {Radio continuum surveys of the Galactic plane can find and characterize \ion{H}{ii} regions, supernova remnants (SNRs), planetary nebulae (PNe), and extragalactic sources. A number of surveys at high angular resolution ($\leq$~25$^{\prime\prime}$) at different wavelengths exist to study the interstellar medium (ISM), but no comparable high-resolution and high-sensitivity survey exists at long radio wavelengths around 21~cm.} {Our goal is to investigate the 21~cm radio continuum emission in the northern Galactic plane at $<$25\arcsec\ resolution. } {We observed a large fraction of the Galactic plane in the first quadrant of the Milky Way ($l=14.0-67.4\degr$ and $|b| \leq 1.25\degr$) with the {\it Karl G. Jansky} Very Large Array (VLA) in the C-configuration covering six continuum spectral windows. These data provide a detailed view on the compact as well as extended radio emission of our Galaxy and thousands of extragalactic background sources.} {We used the BLOBCAT software and extracted 10916 sources. After removing spurious source detections caused by the sidelobes of the synthesised beam, we classified 10387 sources as reliable detections. We smoothed the images to a common resolution of 25$^{\prime\prime}$ and extracted the peak flux density of each source in each spectral window (SPW) to determine the spectral indices $\alpha$ (assuming $I(\nu)\propto\nu^\alpha$). By cross-matching with catalogs of \ion{H}{ii} regions, SNRs, PNe, and pulsars, we found radio counterparts for 840 \ion{H}{ii} regions, 52 SNRs, 164 PNe, and 38 pulsars. We found 79 continuum sources that are associated with X-ray sources. We identified 699 ultra-steep spectral sources ($\alpha < -1.3$) that could be high-redshift galaxies. Around 9000 of the sources we extracted are not classified specifically, but based on their spatial and spectral distribution, a large fraction of them is likely to be extragalactic background sources. More than 7750 sources do not have counterparts in the SIMBAD database, and more than 3760 sources do not have counterparts in the NED database. } {Studying the long wavelengths cm continuum emission and the associated spectral indices allows us to characaterize a large fraction of Galactic and extragalactic radio sources in the area of the northern inner Milky Way. This database will be extremely useful for future studies of a diverse set of astrophysical objects.} | A number of surveys at high angular resolution ($\leq$~20$^{\prime\prime}$) at different wavelengths exist to study the interstellar medium (ISM), from infrared (e.g., UKIDSS\footnote{UKIRT Infrared Deep Sky Survey}, \citealt{lucas2008}; {\it Spitzer}/GLIMPSE\footnote{Galactic Legacy Infrared Midplane Survey Extraordinaire }, \citealt{benjamin2003,churchwell2009}, {\it Spitzer}/MIPSGAL\footnote{A 24 and 70 Micron Survey of the Inner Galactic Disk with MIPS}, \citealt{carey2009} ), to (sub)mm (e.g., ATLASGAL\footnote{APEX Telescope Large Area Survey of the Galaxy} and BGPS\footnote{Bolocam Galactic Plane Survey}, \citealt{schuller2009,rosolowsky2010,aguirre2011,csengeri2014}) and radio (e.g. MAGPIS\footnote{Multi-Array Galactic Plane Imaging Survey}, CORNISH\footnote{the Co-Ordinated Radio `N' Infrared Survey for High-mass star formation}, \citealt{helfand2006,hoare2012}) wavelengths. Previously, the best 21~cm \ion{H}{i} line survey was the HI Very Large Array Galactic Plane Survey (VGPS, \citealt{stil2006}) which has a resolution of 60$^{\prime\prime}$, significantly more coarse than the resolution of the aforementioned surveys. This was one of the motivations for initiating ``The \ion{H}{i}, OH, Recombination line survey of the Milky Way (THOR) \footnote{\url{http://www.mpia.de/thor/Overview.html}}'' \citep{beuther2016}. Using the {\it Karl G. Jansky} Very Large Array (VLA) in C-configuration, we achieve a spatial resolution of $<25^{\prime\prime}$. The WIDAR correlator at the VLA allows us to observe many spectral lines simultaneously, in particular several molecular OH transitions, a series of H$n\alpha$ radio recombination lines (RRLs, $n=$151 to 186), as well as eight spectral windows (SPWs) to cover the continuum emission between 1 and 2 GHz \citep{bihr2015, beuther2016, bihr2016, rugel2018}. We observed a large fraction of the Galactic plane in the first quadrant of the Milky Way ($l=14.0-67.4^\circ$ and $|b| \leq 1.25^\circ$) in several semesters (from 2012 to 2014). The continuum data from the first half of the survey ($l=14.0-37.9^\circ$ and $l=47.1-51.2^\circ$) have been published by \citet{bihr2016}. In this paper, we combine all the continuum data and present the results from the full survey. The radio continuum emission from 1 to 2~GHz is dominated by thermal free-free emission from electrons and non-thermal synchrotron emission of the relativistic electrons in magnetic fields \citep[e.g.,][]{wilson2013}. These can be distinguished by the spectral index $\alpha$, assuming $I(\nu) \propto \nu^\alpha$, where $I(\nu)$ is the intensity at frequency $\nu$. The thermal free-free emission shows an almost flat ($\alpha=-0.1$) spectrum if it is optically thin, or positive spectral index if it is optically thick with values varying between --0.1 and 2 \citep[e.g.,][]{keto2003, wilson2013}. \ion{H}{ii} regions and planetary nebulae are often the sources for thermal free-free emission. In contrast to this, synchrotron emission shows a negative spectral index whose value depends, amongst other things, on the particle energy distribution. Towards extragalactic jets powered by an active galactic nucleus (AGN), one often finds the synchrotron emission with a spectral index $\alpha<-0.5$ \citep[e.g.,][]{hey1971,rybicki1979}. Galactic SNRs often show synchrotron emission with a spectral index around $-0.5$ \citep[e.g.,][]{bhatnagar2011, green2014, reynoso2015}. Thus the spectral index can help us to characterize the nature of the continuum sources we detected in the survey. This allows us to determine whether they are Galactic or extragalactic, which is crucial for \ion{H}{i} and OH absorption studies. Compact galactic radio sources associated with X-ray emission are usually Pulsar Wind Nebulae \citep[e.g.,][]{brisken2005, miller2005}, or X-ray binaries \citep[XRB or microquasars; ][]{mirabel1998, mirabel1999}. By investigating the X-ray and radio flux ratios, the spectral indices and observations at other (optical/infrared) wavelengths of the Galactic sources in detail, we can also constrain the type of the sources, i.e., low mass XRB, PN, pulsar etc. \citep[e.g.,][]{seaquist1993, maccarone2012, tetarenko2016}. With the high-angular resolution ($<25\arcsec$) of our THOR continuum data, we can not only derive the spectral indices of the sources, but also further study the variation in frequency and space. The paper is structured as follows: In Sect.~\ref{sect_obs}, we present the observations and data reduction. Sect.~\ref{sect_extract} presents the methods we used to extract sources and determine the spectral indices. Sect.~\ref{sect_cata} describes the continuum catalog, and the distribution of the continuum sources we extracted. The nature of continuum sources are discussed in Sect.~\ref{sect_discuss}. The conclusions and summary are presented in Sect~\ref{sect_con}. The appendix gives additional information of the continuum observations and tables. | \label{sect_con} We observed a large portion of the first Galactic quadrant ($l=14.0-67.4\degr$, $|b|\leq1.25\degr$) using the VLA in C-configuration and achieved a spatial resolution of $\sim10-25\arcsec$ at 1 to 2~GHz with the THOR Galactic plane survey. In this paper, we present the catalog of the continuum sources from the whole survey. We summarise numbers of different types of sources in the catalog in Table~\ref{table_sum}, and the main results below. \begin{table} \caption{Summary of different types of sources in the catalog} \label{table_sum} \centering \begin{tabular}{l c} \hline\hline Source Types & Number of continuum sources\\ \hline \ion{H}{ii} regions & 713 \\ SNRs & 92$(+21)$\\ PNe & 164 \\ Pulsars & 38 (+663)\\ X-ray sources & 79\\ extragalactic jets & 299 \\ USS & 699 \\ \hline \end{tabular} \tablefoot{USS stands for sources with Ultra Steep Spectra ($\alpha\leq-1.3$). The numbers of the candidates are in brackets.} \end{table} \begin{enumerate} \item The catalog contains 10387 sources we extracted with the BLOBCAT software after removing the obvious observational sidelobe artifacts. About 72\% (7521 sources) of the extracted sources are detected at a significance higher than 7$\sigma$, and $\sim79\%$ are unresolved. The catalog is complete to at least 94\% above the 7$\sigma$ detection limit. The noise of our data is dominated by the sidelobe noise and spatially varying, although more than 60\% of the observed area has a noise level of $7\sigma\lesssim2$~mJy~beam$^{-1}$. We extracted the peak intensity of the six usable SPWs between 1 to 2~GHz, and we were able to determine a reliable spectral index (spectral index fitted with at least 4 SPWs) for 5657 sources. \item We cross-matched the THOR catalog with the {\it WISE} \ion{H}{ii} region catalog and found 713 continuum sources are associated with \ion{H}{ii} regions. Among the matched \ion{H}{ii} regions, 16 are in the radio quiet group in the {\it WISE} catalog which means they did not previously have radio continuum detected. 231 continuum sources are associated with more than one \ion{H}{ii} region. The spectral index distribution shows a single peak around $\alpha=0$, indicating thermal free-free emission. For 168 sources we can fit the SED with a simple homogeneous \ion{H}{ii} region model and derive the emission measure (EM) and the electron density ($n_e$) where the distance information is available in the {\it WISE} catalog. \item Although the diffuse emission from many of the large scale SNRs is filtered out by our interferometric observations, we identify 92 continuum sources associated with 39 SNRs from the SNR catalog by \citet{green2014}. 13 of the new SNR candidates from \citet{anderson2017} are detected in our continuum catalog. \item By cross-matching the THOR catalog with the HASH database, we detect 164 PNe in our continuum catalog. As 90 of them do not have radio emission information at 20~cm in the database, our survey provides this important information in Table~\ref{table_pn}. The spectral index distribution is similar to the one of the \ion{H}{ii} regions and shows a single peak around $\alpha=0$, indicating thermal free-free emission. \item We cross-matched the THOR catalog with the ATNF Pulsar Catalog and found 38 counterparts. One extremely intermittent pulsar J1841-0500 is also detected in our catalog. 663 sources with a spectral index $\alpha < -1.6$ could be Galactic pulsar candidates. \item We cross-matched the THOR catalog with X-ray source catalogs 1SXPS, 3XMM- DR7 and CSC, and found 79 overlaps. 12 of the them have a spectral index steeper than $-1.3$ and could be galaxy clusters. 43 of them do not have previous known radio counterparts within a radius of 15\arcsec on SIMBAD or NED. \item About 300 sources show clear structure of bipolar jets, we mark them as ``jet'' in the catalog and construct spectral index maps. We identified 699 Ultra Steep Spectra (USS) sources and they could be high redshifted radio galaxies ($-1.3>\alpha>-1.6$). Further spectroscopic observations are needed to confirm this. \item About 9000 sources in our catalog are not classified specifically. They are likely to be extragalactic background sources. More than 7750 sources do not have counterparts in the SIMBAD Astronomical Database, and more than 3760 sources do not have counterparts in the NED. \end{enumerate} With the THOR continuum catalog, we provide a rich dataset to the community. All the fits images and catalogs can be downloaded from the project website \\footnote{\url{http://www.mpia.de/thor/Overview.html}}. With the follow up observations of particular sources, such as absorption study of the extragalactic sources, or combining with other existing Galactic plane surveys, we can study the the structure of the Milky Way, and the ISM in different phases. | 18 | 8 | 1808.05990 |
1808 | 1808.05959_arXiv.txt | Detached double white dwarf (DWD) binaries are one of the main science case for the Laser Interferometer Space Antenna (LISA). As the most numerous LISA sources, they will provide important contributions towards understanding binary evolution, Supernovae Type Ia (SNIa) formation channels and the structure of the Milky Way. So far only detection prospects for the Milky Way have been computed. In this letter we show that LISA has the potential to detect DWDs in neighboring galaxies up to the border of the Local Group. In particular, we compute quantitative estimates for the number of detections in M31. We expect between a dozen to several tens of DWDs above the nominal detection threshold, for a mission duration between 4 and 10$\,$yr. We show that extra-galactic DWDs detectable by LISA are those with the shortest orbital periods and with the highest chirp masses, that are candidates SNIa progenitor, virtually undetectable at those distances in optical. This implies that LISA will be the only instrument able to provide SNIa merger rates across the Local Group. | \label{sec:intro} Detached DWD binaries with orbital periods $< 1\,$h will be important GW sources for the LISA mission in many ways \citep{ama17}. Firstly, DWDs are guaranteed LISA sources. A number of short period DWDs have already been identified at optical wavelengths \citep[e.g.,][]{kup18}. Those with strongest signals can be used as calibration sources as they will be detectable already after one week of observations; over time their signal will increase improving the accuracy with which these sources can be used to monitor data quality as new data are acquired \citep[][]{lit18}. Secondly, DWDs will be the most numerous LISA sources. The total number of expected detections exceeds $10^5$ \citep[e.g.,][]{nel01,rui10,mar11,kor17}. Thus, for the first time LISA will provide a sizeable sample of short period DWD binaries to test binary formation theories and validate SNIa formation channels \citep[e.g.,][]{nel01,nel04,reb18}. Moreover, such a large number of individually resolved sources spread all over the Galaxy will allow us to map the Milky Way in GWs and precisely measure its structural parameters like scale radii of the bulge and the disc \citep{ada12,kor18}. When combining GW and optical measurements for DWDs with optical counterparts we will also be able to derive the mass of the bulge and the disc component of the Galaxy \citep{kor18}. Finally, these binaries are so common in the Milky Way that their unresolved signals will form a background for the LISA mission \citep[e.g.,][]{rob17a}. This background contains information on the overall population of DWDs in the Milky Way and can be also used to derive the Milky Way's parameters, like the disc scale height \citep{ben06}. Previously the detectability of DWDs has been exclusively assessed in the Milky Way, while extra-galactic DWDs were only considered as contribution to the background noise \citep[e.g.,][]{kos98,far03}. In this letter we focus for the first time on the properties of extra-galactic DWDs that can be resolved by LISA in the Local Group, and especially in the Large and Small Magellanic clouds (LMC and SMC), and M31 (the Andromeda galaxy). We show that LISA will detect binaries with the shortest periods and highest total masses, and therefore double degenerate SNIa progenitors. As discussed in \citet{reb18}, these are difficult to find with optical telescopes in the Milky Way. Essentially, they are too faint to be identified from H$\alpha$ double-lined profiles in spectra and their eclipses are too short when considering the typical cadence of observations of optical sky surveys, like {\sl Gaia}. Therefore LISA might be the best tool to allow statistical studies of these systems. In this letter, we forecast the parameter space of DWDs accessible through GW observation located at the distance of SMC, LMC and M31 (Section \ref{sec:2}). We use a synthetic population to quantify the number of detection for M31 (Section \ref{sec:3}). In Section \ref{sec:4} we present our conclusions. | \label{sec:4} In this letter we explored the detectability of DWDs outside the Milky Way. We proved that LISA has the potential to detect binaries in neighboring galaxies: LMC, SMC and M31. We find that in the LMC and SMC LISA can detect any binary with $P< 20\,$min, while in M31 LISA will be sensitive to DWDs with $P<10\,$min and ${\cal M} > 0.5\,$M$_{\odot}$. Using an example DWD with $P=5\,$min and ${\cal M}=0.9\,$M$_{\odot}$, we showed that binaries with such characteristics can be detected up to $1\,$Mpc distance, i.e. within the large volume of the Local Group. In the Andromeda galaxy we found a few, to several tens, of DWDs above the LISA detection threshold for 4 and 10$\,$yr mission. This gives an optimistic prospects for detecting other kind of stellar type GW sources like AM CVns and low-mass X-ray binaries, which will likely have an electromagnetic counterpart. A large fraction of extra-galactic DWDs detectable by LISA will have total mass exceeding the Chandrasekhar mass limit and will merge in less than $1\,$Myr, meaning that LISA has the potential to provide SNIa merger rates across the Local Group. | 18 | 8 | 1808.05959 |
1808 | 1808.00267_arXiv.txt | The prevalence and consequences of convection perpendicular to the plane of accretion discs have been discussed for several decades. Recent simulations combining convection and the magnetorotational instability have given fresh impetus to the debate, as the interplay of the two processes can enhance angular momentum transport, at least in the optically thick outburst stage of dwarf novae. In this paper we seek to isolate and understand the most generic features of disc convection, and so undertake its study in purely hydrodynamical models. First, we investigate the linear phase of the instability, obtaining estimates of the growth rates both semi-analytically, using one-dimensional spectral computations, as well as analytically, using WKBJ methods. Next we perform three-dimensional, vertically stratified, shearing box simulations with the conservative, finite-volume code \textsc{PLUTO}, both with and without explicit diffusion coefficients. We find that hydrodynamic convection can, in general, drive outward angular momentum transport, a result that we confirm with \textsc{ATHENA}, an alternative finite-volume code. Moreover, we establish that the sign of the angular momentum flux is sensitive to the diffusivity of the numerical scheme. Finally, in sustained convection, whereby the system is continuously forced to an unstable state, we observe the formation of various coherent structures, including large-scale and oscillatory convective cells, zonal flows, and small vortices. | \label{Introduction} Thermal convection -- the bulk motion of fluid due to an entropy gradient -- transports heat, mass, and angular momentum. There is compelling analytical, numerical, experimental, and observational evidence for convection both in the inner regions of stars and in the outer core and mantle of the Earth \citep{spruit1990solar, aurnou2015rotating}, but its application to astrophysical discs is less clear. In particular, convection perpendicular to the plane of a disc bears marked differences to convection in most stars and planets. For example, the gravitational acceleration reverses sign at the disc midplane, and the gas supports a strong background shear flow. A third difference concerns the heat source necessary to drive the unstable entropy gradient. In stars this heat is provided by nuclear fusion, while in the Earth's outer core convection is driven by both thermal and compositional gradients. In discs heating at the mid-plane is normally assumed to be supplied by accretion (and the associated viscous dissipation of heat). Turbulence arising from the magnetorotational instability and dissipation of large-scale density waves are two mechanisms that might convert orbital energy into the required heat. Note importantly that the fluid motions associated with these heating processes may in fact impede the onset of convection, and at the very least interact with it in non-trivial ways. It has long been speculated that convection perpendicular to the plane of the disc influences the optically thick, high-accreting outburst phase of dwarf novae (for a succinct review, see \cite{2011cannizzo}). More recently, simulations of magnetorotational turbulence in stratified discs have shown that the MRI is capable of generating a convectively unstable entropy gradient. Moreover, three-dimensional, shearing box simulations with zero-net vertical magnetic flux reveal an interplay between convection and the MRI that might enhance angular momentum transport, typically quantified by the dimensionless parameter $\alpha$ \citep{bodo2013fully, hirose2014convection}. A similar interplay between magnetorotational instability and convection might operate in the ionized inner regions of protoplanetary discs, or during FU Orionis outbursts \citep{bell427jun, hirose2015magnetic}. In addition to dwarf novae and the inner regions of protoplanetary discs, hydrodynamic convection might drive activity in the weakly ionized `dead zones' of protoplanetary discs where the magnetorotational instability is unlikely to operate \citep{armitage2011dynamics}. Although irradiation by the central star usually results in an isothermal vertical profile in protoplanetary discs \citep{chiang1997spectral, d1998accretion}, heating by spiral density waves driven by a planet might nevertheless generate regions exhibiting a convectively unstable background equilibrium \citep{boley2006hydraulic, lyra2016shocks}. So too might the heating from the dissipation of the strong zonal magnetic fields associated with the Hall effect (Lesur, Kunz and Fromang 2014). As mentioned, accretion may instigate convection. But convection might drive accretion itself. Motivated by considerations of the primordial solar nebula, \cite{lin1980structure} constructed a simple hydrodynamic model of a cooling young protoplanetary disk contracting towards the mid-plane. They showed that vertical convection arises if the opacity increases sufficiently fast with temperature, with the heating source initially provided by the gravitational contraction. But if convection is modeled via a mixing length theory, with an effective viscosity, the process might \emph{self-sustain}: convective eddies could extract energy from the background orbital shear, and viscous dissipation of that energy might replace the initial gravitational contraction as a source of heat for maintaining convection. In order to go beyond mixing length theory, \cite{ruden1988axisymmetric} analysed linear axisymmetric modes in a thin, polytropic disc in the shearing box approximation and estimated that mixing of gas within convective eddies could result in values of $\alpha \sim 10^{-3}-10^{-2}$. They cautioned, however, that axisymmetric convective cells could not by themselves exchange angular momentum: such an exchange would have to be facilitated either by non-axisymmetric modes, or by viscous dissipation of axisymmetric modes. Following these early investigations, a debate ensued that centered not so much on the size of $\alpha$ but rather on its sign. Dissipation of \textit{non-linear axisymmetric} convective cells was investigated by \cite{kley1993angular} who performed quasi-global simulations (spanning about 100 stellar radii) of axisymmetric viscous discs with radiative heating and cooling. They measured an inward flux of angular momentum, but warned that this might be due to the imposed axial symmetry and to their relatively high viscosity. Convective shearing waves were first investigated by \cite{ryu1992convective}, who concluded that linear \emph{non-axisymmetric} perturbations would result in a net \textit{inward} angular momentum flux at sufficiently large time. These results were questioned however by \cite{lin1993nonaxisymmetric}, who examined analytically and numerically a set of \textit{localised} linear non-axisymmetric disturbances in global geometry. They demonstrated that these modes could transport angular momentum outwards in some cases, and opined that the inward transport of angular momentum reported by \cite{ryu1992convective} was an artifact of the shearing box approximation. Interest in hydro convection waned after local non-linear 3D compressible simulations showed that it resulted in \textit{inward} rather than outward angular momentum transport. Using the finite-difference code \textsc{ZEUS} and rigid, isothermal vertical boundaries, Stone and Balbus (1996, heareafter SB96) initialized inviscid, fully compressible, and vertically stratified shearing box simulations with a convectively unstable vertical temperature profile, and measured a time-averaged value of $\alpha \sim -4.2\times10^{-5}$. Inward angular momentum transport was also reported by \cite{cabot1996numerical} who ran simulations similar to those of SB96 but included full radiative transfer and a relatively high explicit viscosity. Analytical arguments for the inward transport of angular momentum by convection were presented by SB96 which, crucially, assumed axisymmetry, especially in the pressure field. However, in a rarely cited paper \cite{klahr1999azimuthal} presented fully compressible, three-dimensional, global simulations including explicit viscosity and radiative transfer of hydrodynamic convection in discs that gave some indication that non-linear convection in discs actually assumes non-axisymmetric patterns. Some fifteen years later, the claims of SB96 and of \cite{cabot1996numerical} were called into question: fully local shearing box simulations of Boussinesq hydrodynamic convection in discs using the spectral code \textsc{SNOOPY} and employing explicit diffusion coefficients indicated that the sign of angular momentum transport due to vertical convection can be reversed provided the Rayleigh number (the ratio of buoyancy to kinematic viscosity and thermal diffusivity) is sufficiently large. It would appear then that vertical convection can possibly drive outward angular momentum transport after all \citep{lesur2010angular}. Our aim in this paper is to explore the competing results and claims concerning hydrodynamic convection in accretion discs, in particular to determine the sign and magnitude of angular momentum transport. We also isolate and characterize other generic features of convection that should be shared by multiple disc classes (dwarf novae, protoplanetary disks, etc) and by different driving mechanisms (MRI turbulence, spiral shock heating, etc). We do so both analytically and through numerical simulations, working in the fully compressible, vertically stratified shearing box approximation. We restrict ourselves to an idealised hydrodynamic set-up, omitting magnetic fields and complicated opacity transitions, and in so doing ensure our results are as general as possible. The structure of the paper is as follows: first, in Section \ref{methods} we introduce the basic equations and provide a brief overview of the numerical code, the numerical parameters and set-up, and our main diagnostics. In Section \ref {eigensolver} we investigate the linear behavior of the convective instability, employing both analytical WKBJ and semi-analytical spectral methods to calculate the growth rates and the eigenfunctions. In Section \ref{unforcedsims} we explore the non-linear regime through \textit{unforced} simulations, in which the convection is not sustained and is permitted to decay after non-linear saturation. Because this unforced convection is a transient phenomenon which might depend on the initial conditions, in Section \ref{forcedsims} we explore the non-linear regime through simulations of \textit{forced} convection, in which the internal energy is relaxed to its initial, convectively unstable, state to mimic, possibly, the action of background MRI turbulent dissipation and optically thick radiative cooling. Finally, Section \ref{conclusions} summarizes and discusses the results. | \label{conclusions} Motivated by recent radiation magnetohydrodynamic shearing box simulations that indicate that an interaction between convection and the magnetorotational instability in dwarf novae can enhance angular momentum transport, we have studied the simpler case of purely hydrodynamic convection, both analytically and through three-dimensional, fully-compressible simulations in \textsc{PLUTO}. For the linear phase of the instability, we find agreement between the growth rates of axisymmetric modes calculated theoretically and those measured in the simulations to within a percentage error of $<1\%$, thus providing a useful check on our \textsc{PLUTO} code. The linear eigenmodes are worth examining not only to help understand the physical nature of convection in disks, but also because they may appear in some form during the nonlinear phase of the evolution, especially on large-scales, and during intermittent or cyclical convection. We then explored the nonlinear saturation of the instability, both when convection is continually forced and when it is allowed to reshape the background gradients so that it ultimately dies out. We focussed especially on the old problem of whether hydrodynamic convection in a disc leads to inward ($\alpha < 0$) or outward ($\alpha > 0$) angular momentum transport. In both forced and unforced convection we found $\alpha>0$ in general in the nonlinear phase. These results were confirmed by a separate run using the code \textsc{ATHENA}, but contradict the classical simulations of SB96 who reported inward transport in both cases using the code \textsc{ZEUS}. This discrepancy reveals a set of unfortunate numerical difficulties that complicate the simulation of convection in disks. These, in large part, issue from the fact that the inviscid linear modes of convection grow fastest on the shortest possible scales. Thus, no matter the resolution, the nature of the code's grid dissipation will always impact on the system's evolution, certainly in the linear phase and possibly afterwards. We argue that a more diffusive numerical set-up, such as supplied by \textsc{ZEUS} at low resolution or Riemann solvers such as HLL, impose an axisymmetry on the flow which leads to generally inward transport of angular momentum. But, on the other hand, we suspect that less diffusive solvers such as Roe and HLLC artificially alias shearing waves in the linear phase of the evolution, leading to spurious non-axisymmetric flow early in a simulation. Though concerning, we believe this is only a problem in the low amplitude linear phase because physical mode-mode interactions will dominate once the perturbations achieve larger amplitudes. Nonetheless, shearing wave aliasing certainly deserves a separate study. To properly dispense with these numerical issues one must add explicit viscosisty (and thermal diffusivity), as this regularises the linear problem. We find that at a Richardson number of $\text{Ri} \sim 0.05$, onset of convection is observed for a critical Rayleigh number $10^{5} < \text{Ra}_c \leq 10^{6}$. Just above this value convection is largely axisymmetric and $\alpha<0$. At a larger second critical Ra between $10^6$ and $10^7$, the sign of $\alpha$ switches and the flow becomes more turbulent and nonaxisymmetric. This sequence of states mirrors that simulated by Lesur and Ogilvie (2010). At large (resolved) Rayleigh numbers, viscous simulations are initially controlled by the axisymmetric modes; these are then attacked by secondary shear instabilities in both the $xz$ and $xy$-planes, which break the axisymmetry and order of these structures, leading to a more chaotic state. At lower Ra, viscosity suppresses the non-axisymmetric shear instabilities and axisymmetry is never broken. (At even lower Ra, convection never begins, of course.) In forced convective runs, rather than maintaining the convection by a fixed heating source at the mid-plane, we instead allowed the vertical equilibrium to relax to its initial, convectively unstable, state. Our thermal relaxation is artificially imposed, but its overall effect is to mimic the heating of the mid-plane and cooling of the corona due to physical mechanisms that maintain the convectively unstable entropy profile, such as the MRI and radiative losses present in the simulations of \cite{hirose2014convection}. We observed in the non-linear stage now the formation of large-scale convective cells (similar in some respects to elevator flow) that emerge and break down cyclically, in addition to zonal flows and vortices. Although further checks are required, it is tempting to link this cyclical convection to the convective limit cycle observed in the radiative magnetohydrodynamic simulations of \cite{hirose2014convection}, since both processes rely on the build-up and rapid evacuation of heat. Despite our demonstration that hydrodynamic convection can lead to positive stress and outward transport of angular momentum, the fact remains that the stresses are small (typically we measured $\alpha \sim 10^{-6}-10^{-5}$). Having said that, the magnitude of $\alpha$ is sensitive to the depth of the buoyancy frequency profile, and a deeper profile could increase $\alpha$ by an order of magnitude or more. Finally we have not observed self-sustaining hydrodynamic convection in any of our simulations. By self-sustaining convection we mean that (when $\alpha > 0$) energy extracted from the shear by convection might itself cause convective motions, which in turn extract energy energy from the shear, closing the loop. It is more likely that if convection is to occur in disks it will be as a byproduct of other processes, such as heating by density waves, emitted in the presence of a planet, or by dissipation of magnetorotational turbulence. We intend to investigate both mechanisms and their instigation of convection in future work. | 18 | 8 | 1808.00267 |
1808 | 1808.02438_arXiv.txt | We present a new method to study the characteristic scales of collapse and fragmentation in galactic disks. Clump formation is seeded in simulations via controlled perturbations with a specified wavelength and velocity. These are applied to otherwise quiet gas disks ranging from analogues of present day spirals to gas-rich, high-redshift galaxies. The results are compared to linear theory, turbulently perturbed disks and observations. The results reflect the expectations of linear, non-axisymmetric theory with a finite window for growth into a bound clump. We identify two new modes of clump formation: rotation-driven fission and fragmentation of tidal tails, though both are expected to rarely contribute to clump formation in observed disks. We find that bound clumps are generally much smaller than the commonly used Toomre mass. The preferred scale for fragmentation increases with the disk gas mass but cannot produce bound objects larger than $\sim10^9$ M$_{\odot}$. The most likely bound clump mass increases from $3\times10^6$ in low mass disks up to $5\times10^8$ M$_{\odot}$. We conclude that observed massive stellar and gaseous clumps on 1 kpc scales at high redshift are most likely aggregates of many initially distinct bound clumps. | The typical size of star clusters is expected to depend on the galactic environment. The Jeans' mass, which changes based on environment, plays a role \citep[e.g.][]{hopkinse2012}. The pressure likely plays a role; high pressure environments such as Arp 220 have larger star clusters when compared to other local galaxies \citep[e.g.][]{wilson2006}. If we begin by considering low redshift galaxies, a preferred mass-scale for star cluster formation is apparent. In the Milky Way itself, star clusters have typical masses between $10^3 - 10^4$ M$_{\odot}$ \citep{fall2012}. A preferred scale for star clusters may in turn suggest a preferred scale for Giant Molecular Clouds (GMCs). In the Milky Way GMCs have typical masses between $10^5 - 10^6$ M$_{\odot}$ and typical sizes of 50 - 100 pc \citep{fukuiARAA, mckeeARAA}. If we consider galaxies at higher redshifts, these preferred scales appear to change. Stellar observations are able to identify large UV-bright star-forming regions called {\it clumps} \citep{elmegreen2007}. The CANDELS survey has provided extensive clump catalogues for galaxies between 0.5 $< z <$ 3.5. The properties of clumps identified in CANDELS galaxies suggest they are extremely large, with typical masses of $10^7 - 10^9$ M$_{\odot}$ and typical sizes $\sim 1$ kpc \citep{CANDELS, CANDELS2}. Other compilation studies find masses between $10^5 - 10^9$ M$_{\odot}$ \citep{dessauges2018}. Either way, these are orders of magnitude more massive than present-day clusters or star-forming regions. This picture is even more complex if we add in starbursts or merger-driven systems. As mentioned above, if we consider Arp 220, there are many active sites of star formation and the star-forming complexes may be much larger than those in the Milky Way \citep[][and references therein]{murray2010}. In these environments the masses of stellar clusters increase dramatically. \citet{wilson2006} find masses approaching $10^7$ M$_{\odot}$, which may make them candidates for young globular clusters. The conditions in star-forming regions should be imprinted in the properties of gas. At higher redshifts, and in starburst systems, the conditions both within and around galaxies were different than at low redshift. Specifically, at higher redshifts the galaxy interaction rate was higher and galaxies themselves are much more likely to be gas rich. It is no surprise that in such gas-rich, highly molecular environments the properties of star-forming regions are likely to be different. Indeed, the appearance of gas disks beyond z$\sim$0.5 are much more clumpy in nature \citep[e.g.][]{forster2009}. This highly molecular clumpy nature may suggest that star-forming regions may be larger in both mass and spatial extent. Wide beam observational studies suggest these objects have masses between $10^8$ - $10^{10}$ M$_{\odot}$, or approximately $1-10\%$ of the total disk mass \citep[e.g.][]{tacconi2010, swinbank2010, swinbank2011, genzel2011, hodge2012}. The corresponding physical sizes range from as small as 100 pc to as large as 2 kpc \citep[e.g.][]{swinbank2010, tacconi2010}. However, recent results from the SGASS lensing survey have shown a finer level of substructure, albeit in galaxies less massive than those typical in the CANDELS sample \citep{johnson_model}. These results show that stellar clump sizes can be consistent with present-day star clusters, which would originate from objects similar to present day GMCs with high star formation efficiency \citep{johnson2017, rigby2017}. Another such lensing study has been done in the Cosmic Snake \citep{cava2018}. Using data from the CLASH survey, the authors have obtained both a lensed arc and counterimage. In this way they can compare two spatial resolutions for the same object. They find that the lower resolution image (counterimage) produces clumps that are amplified by a factor of 2-5 on average, for an decrease in resolution of 10 times. Studies like these suggest two important points. First, the gas mass of a galaxy is important for determining the scale of star formation. Second, the resolution of earlier studies may not be sufficient to resolve clumps. With the typical physical resolution of instruments at high redshift being generally poorer, it may be that we are treating collections of GMCs \citep{tacconi2010} as a single entity. This idea has been lent credence by samples of lensed galaxies: while directly observed galaxies are often large due to selection effects, lensed galaxies are typically of lower mass (M$_{\star}$ $\sim 10^9$ M$_{\odot}$). They provide us with a better resolved picture of the molecular gas. For example, \cite{swinbank2010} find gas clumps of similar size to Milky Way GMCs, approximately 100 pc. They propose that these objects would be similar to present-day star-forming regions except with more star-forming cores with higher densities. \cite{hodge2012} infer typical internal densities of $\sim100$ cm$^{-3}$, in accordance with the typical density of low-redshift GMCs. If we infer masses from these sizes, we can assume that these objects would likely have masses similar to large present-day GMCs: maybe between $10^5-10^7$ M$_{\odot}$. Local starburst galaxies with enhanced star-formation and kinematic properties similar to galaxies at z $\sim$ 1.5 can also be used as a testbed for these theories. Their proximity offers better resolution studies and HST-DYNAMO has studied 13 such galaxies \citep{white2017}. Studies here confirm that clump clustering is likely to impact the measurement of clump properties at higher redshifts, where resolution degrades \citep{fisher2017b}. If we look to isolated galaxy simulations, we see results consistent with small star-forming regions. For example, \citep{tamburello2015} find smaller clump masses ($< 10^7$ M$_{\odot}$). The higher resolution available in isolated galaxy simulations also offers the perfect place to study the impacts of resolution on structure identification. Some work has suggested that the resolution of HST at the redshifts concerned is insufficient to fully resolve these clumpy objects. \cite{tamburello2017} and \cite{dessauges2017} have argued that at the spatial resolution of 1 kpc, it is not possible to fully resolve these stellar clumps. This supports the idea that we are just seeing collections of smaller stellar clumps, clustered closely together. Indeed, \cite{behrendt2016} have shown that closely clustered clumps could be confused for more massive objects with the resolution of surveys like CANDELS. At the other extreme, it has been suggested that in massive, gas-rich galaxies the physics of clump formation changes with Violent Disk instabilities (VDI) producing different outcomes \citep{dekel2009}. This is built on the idea that a Toomre instability sets the characteristic scale of clumps in gas-rich disks \citep{shlosman1993,noguchi1998,noguchi1999}. This idea has been invoked to explain certain cosmological zoom simulations exhibiting larger clumps in the range 10$^{7-9}$ M$_{\odot}$ \citep{mandelker2014, mandelker2017}. However, it has been theorized that not all gas rich disks host VDI, and that this behaviour is largely dependant on feedback strength rather than a qualitative difference in how clumps form \citep{fiacconi2017}. Regardless, other studies of cosmological zoom simulations also report clump masses in this higher range \citep{agertz2009,oklopcic2017}. An alternative way to study clump formation in isolated galaxies is by employing spherical halo collapse. Here, rather than begin with a disks, a halo of dark matter and hot gas are allowed to collapse, thus forming the disk \citep[e.g.][]{kaufmann2006, kaufmann2007, teyssier2013}. This is commonly used to study the interaction of the disk with the surrounding medium, to study cold flows for example. However, it has also been used to study clump formation \citep{noguchi1999, inoue2012}. For example, \citet{inoue2014} report clump masses above $10^8$ M$_{\odot}$ in isolated galaxies formed in this way. Resolving these differences is complicated by the fact that all of these studies use different methods. Some involve simulations on cosmological scales while others model isolated galaxies, the resolution of these two types of simulations can be quite different. Different studies use different numerical methods or different hydrodynamical schemes. Beyond that, perhaps the largest variable, is the type of feedback chosen. The type and strength of feedback chosen plays a large role in determining the structure of star-forming gas and, consequently, the structure of stellar clusters. All of these variables makes it incredibly confusing to compare results among different studies. Add on top of that the different types of observations we are considering, stellar versus gas, lensed versus un-lensed, and we are left with great difficulty in interpreting the results in the literature. We propose a new method to study clump formation in simulations. Our method avoids the problems associated with many of the algorithm-specific assumptions discussed above. We seed clump formation events by hand and study their growth in high resolution isothermal disks that do not include feedback. In this way, we can constrain the initial mass of clumps formed in a variety of disks. These are directly comparable to both observationally determined masses and masses from theories of fragmentation. The rest of the paper is laid out as follows. We begin by examining the predictions from linear theory in section ~\ref{sec:theory}. Linear theory is difficult to extrapolate to non-linear clumps properties. Instead, we use it to design simulations to explore clump formation in disks ranging from Milky Way-like cases to the heavy, turbulent disks expected at high redshifts. In section~\ref{sec:methods} we describe our controlled simulation approach which allows for high resolution and relatively easy interpretation of the results. We make first use of our quiet disks in section ~\ref{sec:turbdisk}, looking at isolated, turbulent disks. These simulations show typical outcomes for turbulent-type initial conditions. Their purpose in this paper is to illustrate the value of a more controlled approach. Finally, in order to study the key scales for fragmentation and clump formation in a controlled way, we take a new approach of seeding non-linear perturbations. We present details of the approach and results in section~\ref{sec:seed}. In section~\ref{sec:mass}, we extrapolate from our simulation results to estimate likely clump masses based on disk conditions. Finally, in section~\ref{sec:implications} we discuss the observational implications of this study. | In this work, we have introduced a new method for studying the formation of clumps, or bound structures, in galactic disks. We seed clump formation events in initially stable isothermal disks, without star formation or feedback. We design these conditions to be purposefully simple and thus offer maximum control. Our clump mass spectrums are not impacted by feedback recipe choices, providing a complementary approach to other recent work that employs a variety of feedback assumptions. Our results indicate that the characteristic length and mass scale of clumps is not a fixed fraction of the Toomre estimate but depends on the properties of the galaxy in question. This complicates efforts to analytically estimate a characteristic mass. By seeding turbulent clump formation events, we are able to study the exact conditions under which different clump masses form. In general, we find that our largest clump masses can be over an order of magnitude smaller than the Toomre mass. Our results suggest smaller initial masses for clumps than reported in some observational studies. This is consistent with the idea that those studies have large beams that encompass many bound objects. We stress that when making this comparison, the clump observational mass may be different than the initial mass: just because an object is observed to be massive, does not mean it formed at that mass through gravitational fragmentation. Our method provides a new way to approach the problem of studying clump formation in simulations; it offers a new way to compare to observations. The method can be completely tailored to specific galaxies. The only requirements is that the rotation curve and surface density distribution for the galaxy are known. Our method provides a promising new way to study the formation of these clumps in specific galaxies without the biases introduced by including different feedback methods. | 18 | 8 | 1808.02438 |
1808 | 1808.02112_arXiv.txt | \noindent Self-similar solution is obtained for propagation of a strong shock, in a flat expanding dusty Friedman universe. Approximate analytic solution was obtained earlier, using relation between self-similar variables, equivalent to the exact energy conservation integral, which was obtained by L.I. Sedov for the strong explosion in the static uniform medium. Numerical integration of self-similar equation is made here, giving an exact solution of the problem, which is rather close to the approximate analytic one. The differences between these solutions are most apparent in the vicinity of the shock. For polytropic equation of state, self-similar solutions exist in more narrow interval of the adiabatic power than in the static case. | At early stages of star and galaxy formation we may expect strong explosions at last stages of evolution of very massive primordial stars, which enrich the matter with heavy elements. Detection of heavy elements at red shifts up to $z\sim 10$ from GRB observations (GRB090423 at $z\approx 8.2$, GRB120923A at $z\approx 8.5$, GRB090429B with a photo-$z \approx 9.4$) \cite{lzgrb} make plausible this suggestion. Propagation of a strong shock in the static uniform media was investigated by many authors \cite{stanyuk}, \cite{taylor}, but the finite analytic self-similar solution was obtained in \cite{sedov}. Analytic self-similar solution for the strong shock propagating through the uniform expanding media was obtained in \cite{bk15}, describing approximately a strong shock propagation in the flat Friedman universe \cite{znuniv}. The analytic solution was obtained in neglecting the energy input from the kinetic and gravitational energy of the non-perturbed Friedman model, into heating of the matter behind the shock. Here the numerical solution of the same self-similar equation is obtained, which take into account all processes. It is shown that restrictions for the adiabatic power $\gamma$ obtained for validity of the analytic solution, with small corrections, remain also for the exact numerical one, and numerical difference between both solutions is not large. A qualitative behaviour of the density dependence between two solutions happens in the thin layer near the shock front. The problem of a strong shock propagation in the expanding medium was considered earlier in \cite{soy75,och78,be83,ito83,vob85,ks93,ek98} Propagation of a detonation wave in the flat expanding universe was studied in \cite{kazh86,be85}. Shock propagation in the outflowing stellar wind was considered in \cite{cd89}. Review of papers on this topic is given in \cite{omke88}. | The constant $\beta$ in the definition of the non-dimensional radius $\xi$ in (\ref{eq25a}) is obtained from the explosion energy integral $E$. Due to zero energy (kinetic + gravitational) in the non-perturbed solution, the conserving value of the explosion energy behind the shock, in the uniformly expanding medium, with velocity and density distribu\-ti\-ons (\ref{eq20a}), with account of the gravitational energy, is determined as \be \label{eq52a} E=\int_0^{R(t)} \rho\left[\frac{v^2}{2}+\frac{c^2}{\gamma(\gamma-1)}\right]4\pi r^2 dr- \int_0^{R(t)}\frac{G_g m dm}{r}. \ee In non-dimensional variables (\ref{eq24a}) this relation reduces to the equation for the constant $\beta$ \be \label{eq53a} \beta^{-5}=\frac{64\pi}{25}\int_0^1 G\left[\frac{V^2}{2}+\frac{Z}{\gamma(\gamma-1)}\right]\xi^4 d\xi-\frac{8}{3}\int_0^1 G\xi\left(\int_0^\xi G\eta^2 d\eta\right)d\xi. \ee \begin{table}[h] \caption{The values $\beta(\gamma)$ for the analytic and numerical solutions} \label{tabular:timesandtenses} \begin{center} \begin{tabular}{ |c |c |c|} \hline \textbf{$\gamma$} & \textbf{$\beta_{an}$} & \textbf{$\beta_{num}$} \\ \hline 1.05 & 3.2910 & 3.3512 \\ \hline 1.10 & 2.2268 & 2.5003 \\ \hline 1.12 & 2.0423 & 2.3713 \\ \hline 1.15 & 1.8522 & 2.2416 \\ \hline \end{tabular} \end{center} \end{table} The values $\beta(\gamma)$ for the analytic solution at $\gamma<\gamma_{**}$, and numerical one for $\gamma<\gamma_{**}$, are given in the Table. It follows from the self-similar solution, that in the expanding medium the velocity of shock from (\ref{eq22a}) decreases as $\sim t^{-1/5}$, what is much slower than the shock veloci\-ty in the static uniform medium $\sim t^{-3/5}$, according to Sedov solution \cite{sedov}. Correspondingly the radius of the shock wave in the expanding self-gravita\-t\-ing medium increases $\sim t^{4/5}$, more rapidly that the shock wave radius in the uniform non-gravitating medium $\sim t^{2/5}$. It means, that the shock propagates in the direction of decreasing density with larger speed, than in the static medium, due to accele\-ra\-t\-ing action of the decreasing density, even in the presence of a self-gravitation. | 18 | 8 | 1808.02112 |
1808 | 1808.10707_arXiv.txt | {One of the main scenarios of planet formation is the core accretion model where a massive core forms first and then accretes a gaseous envelope. This core forms by accreting solids, either planetesimals, or pebbles. A key constraint in this model is that the accretion of gas must proceed before the dissipation of the gas disc. Classical planetesimal accretion scenario predicts that the time needed to form a giant planet's core is much longer than the time needed to dissipate the disc. This difficulty led to the development of another accretion scenario, in which cores grow by accretion of pebbles, which are much smaller and thus more easily accreted, leading to a more rapid formation.} {The aim of this paper is to compare our updated pebble-based planet formation model with observations, in particular the well studied metallicity effect.} {We adopt the Bitsch et al. 2015a disc model and the Bitsch et al. 2015b pebble model and use a population synthesis approach to compare the formed planets with observations.} {We find that keeping the same parameters as in Bitsch et al. 2015b leads to no planet growth due to a computation mistake in the pebble flux (Bitsch et al. 2017). Indeed a large fraction of the heavy elements should be put into pebbles ($Z_{\rm peb} / Z_{\rm tot} = 0.9$) in order to form massive planets using this approach. The resulting mass functions show a huge amount of giants and a lack of Neptune mass planets, which are abundant according to observations. To overcome this issue we include the computation of the internal structure for the planetary atmosphere to our model. This leads to the formation of Neptune mass planets but no observable giants. Reducing the opacity of the planetary envelope finally matches observations better.} {We conclude that modeling the internal structure for the planetary atmosphere is necessary to reproduce observations.} | \label{introduction} One of the main scenarios of planet formation is the core accretion model, which aims to explain, among other things, the formation of gas giants, believed to be the first object to form. These planets are massive and thus can have a huge impact on the formation of other planets in the system. The idea of the core accretion model is that a massive core forms first and then accretes a gaseous envelope. A key constraint in this model is that the accretion of gas must proceed before the dissipation of the gas disc. This implies that the core must become massive enough to accrete gas in a few Myr (Haisch et al., 2001, Li \& Xiao, 2016). In the classical accretion scenarios the core grows by accreting planetesimals, which are kilometer-sized objects and the time needed to form a giant planet's core is much longer than the lifetime of the disc (Pollack et al. 1996). This difficulty led to the development of another accretion scenario, in which cores grow by accretion of pebbles, which are centimeter-sized bodies (Birnstiel et al. 2012). Pebbles themselves form by collisions of micro-meter sized dust embedded in the disc (Lambrechts \& Johansen 2014). Due to their small size, pebbles are more affected by gas drag and thus can be more easily accreted by the core (Lambrechts \& Johansen 2012), resulting in a more rapid formation. While an appealing idea, this scenario needs to be confronted to observations to check wether it reproduces the trends seen in the data. We use the well studied metallicity effect which is fully established observationally for exoplanets and is reproduced by planetesimal-based formation models (Mordasini et al. 2009; Coleman \& Nelson, 2016a) to aim at testing this model, and present an updated pebble-based planet formation model which outcomes are comparable to observations.\\ This work is structured as follows: in section \ref{model} we set the theoretical basis that is used to perform our simulations. We present the disc model and its evolution, core growth by pebble and gas accretion, and the migration of the bodies. In section \ref{Formation_tracks} we compute formation tracks to disentangle the effects of growth and migration. In section \ref{population_synth} we present the impact of different initial conditions on the formation of planets, such as the starting time for the embryo and the metallicity of the disc, using a population synthesis approach. Considering a single planet per disc we then compare the results of our model with observations of exoplanets through mass functions. In section \ref{Differentinternalstructure} we show the impact of having an internal structure for the planet's atmosphere and, adjusting some parameters, compare this updated model with observations. Finally, section \ref{metallicity} discuss the metallicity effect and section \ref{conclusion} is devoted to discussion and conclusions. | \label{conclusion} In this work we provide an updated pebble-based planet formation model. We use the disc model given by Bitsch et al. 2015a and first test our computations by intending to reproduce their results. Aiming at this we noticed that the equation used to compute the flux of pebbles in B15b is not correct, as confirmed later on by B. Bitsch (private communication, Bitsch et al. 2017). Using the correct equation translates in the formation of very small mass planets, at least when the fraction of the total amount of solids taking part in the formation of pebbles is the same as in B15b ($Z_{\rm peb} = 2/3 \cdot Z_{\rm tot}$). Adopting the correct pebble flux equation therefore requires assuming that a larger fraction of the heavy elements takes part in the formation of pebbles ($Z_{\rm peb} / Z_{\rm tot} = 0.9$) in order to form massive enough planets.\\ Using this approach we study the impact of the amount of pebbles ($Z_{\rm peb}$) as well as the amount of dust ($Z_{\rm dust}$). The latter impacts not only the disc structure but also the growth of the planets through the isolation mass. $Z_{\rm peb}$ influences the surface density of pebbles and thus the planet growth. We show that increasing the density of pebbles favours the formation of massive planets. However the amount of dust should not be too small otherwise migration is too efficient and planets fall into the star. We therefore see that the changes in the disc properties (through $Z_{\rm tot}$ and thus $Z_{\rm dust}$) lead to a large diversity of outcomes. This is similar to Ida et al. 2016 where the boundaries in the disc (affected by the disc properties) have a strong influence on the results. \\ Using a population synthesis approach we then compute simulations and compare our results with observations (Mayor et al. 2011). The latter indicates a high fraction of Neptune mass planets that we do not reproduce. Indeed our mass distribution is such that most of the detectable planets are giant planets. This significant amount of giants come from the fact that in the B15b model, there is no internal structure. Indeed once the isolation mass is reached a huge amount of gas is accreted. In reality, once the planet reaches $\rm M_{\rm iso}$, the envelope contracts, producing some luminosity that will slow down gas accretion.\\ Computing the internal structure for the planetary atmosphere we produce another full set of simulations. We obtain a high amount of Neptune mass planets and a lack of giant planets. To recover an acceptable number of giant planets, we consider two different parameters: the starting time of the seeds and the opacity of the planetary atmosphere.\\ The starting time may indeed have a considerable impact on the outcome of the simulation, as shown by Chamber 2016. Studying the outcomes of planetary systems, they show that the embryos should have a mass between $10^{-3}$ to $10^{-2}$ $M_{\oplus}$ at 1.5 Myr and 2 Myr of their evolution to accrete efficiently and eventually grow to giant planets. In our model, we do not compute the growth of the embryos but insert them at later times (1.5 Myr and 2 Myr) with masses that are consistant with their results. Therefore the masses the embryos have at a given time influence the final planetary architecture for multiple-planets systems (Chambers 2016, Levison et al. 2015) as well as for single planet cases (this work). We highlight that inserting our embryos earlier in the disc evolution, for instance at $\rm t_{\rm ini} = 1$ Myr, would give them more time to grow and help forming giant planets. Reducing the opacity also increases the amount of giants and the resulting type of formed planets is in good agreement with observations (Fig. \ref{massfunction_smalleropacity}).\\ We thus use this last model to discuss the metallicity effect. We show that the amount of giant planets increases with the metallicity and that the B15b approach also reproduces this effect.\\ The main difference between the B15b approach and our model (with the internal structure and a reduced opacity) is the fact that we compute an internal structure for the planetary atmosphere. We conclude that computing the internal structure and reducing the opacity is more realistic and reproduces better the mass distribution of the observed planets (Mayor et al. 2011). However the replenishment of the envelope with the surrounding gas disc is not taken into account here which may lead to a more sophisticated model (Lambrechts \& Lega 2017, Brouwers et al. 2017).\\ Finally we should also keep in mind that our simulations are performed with one planet per disc and thus the eventual effects of several planets growing in the same disc are not taken into account, which we leave for future work (Coleman et al. in prep.), as well as long-term evolution, which is essential in order to compute the present day radius of planets. In addition, for planets close enough to the star, and of intermediate mass, evaporation needs to be taken into account (Mordasini et al. 2012). | 18 | 8 | 1808.10707 |
1808 | 1808.05218_arXiv.txt | HIRMES is a far-infrared spectrometer that was chosen as the third generation instrument for NASA's SOFIA airborne observatory. HIRMES promises background limited performance in four modes that cover the wavelength range between 25 and 122 $\mu$m. The high-spectral resolution ($R \approx 10^5$) mode is matched to achieve maximum sensitivity on velocity-resolved lines to study the evolution of protoplanetary disks. The mid-resolution ($R \approx 12, 000$) mode is suitable for high sensitivity imaging of galactic star formation regions in, for example, the several far-infrared fine structure lines. The low-resolution ($R \approx 2000$) imaging mode is optimized for spectroscopic mapping of far-infrared fine structure lines from nearby galaxies, while the low resolution ($R \approx 600$) grating spectrometer mode is optimized for detecting dust and ice features in protostellar and protoplanetary disks. Several Transition Edge Sensed (TES) bolometer arrays will provide background limited sensitivity in each of these modes. To optimize performance in the various instrument modes, HIRMES employs eight unique fully-tunable cryogenic Fabry-Perot Interferometers (FPIs) and a grating spectrometer. Here we present the design requirements and the mechanical and optical characteristics and performance of these tunable FPI as well as the control electronics that sets the mirror separation and allows scanning of the FPIs. | \label{sec:intro} % The High Resolution Mid Infrared Spectrometer (HIRMES) is the third generation science instrument for NASA's SOFIA. The development of HIRMES is led by Goddard Space Flight Center with Cornell University developing and providing the series of cryogenic scanning Fabry-Perot Interferometers described here. HIRMES's primary science goal is to investigate the physics and chemistry of protoplanetary disks and planet formation, using velocity resolved spectroscopy in the mid- to far-infrared spectral lines of HD, H\textsubscript{2}O, and [OI] and low resolution spectroscopy of solid state water ice features. These lines trace important building blocks in the formation of planets and the development of life. \subsection{Protoplanetary Disks} Stars form from condensations in dense molecular cloud cores. During the final phases accretion disks form, enveloping and feeding the pre-main sequence stars. Some fraction of these disks also provide the raw material for the formation of planets. Fig. \ref{waterdist} illustrates the morphology of this early phase of planet formation. It is thought that most planets form within a few tens of AU of the star in a disk of cold gas, dust and ice. The largest disk ice/rock particles and most dense gas and dust will naturally accumulate near the disk mid-plane, enveloped by a gas/dust ice/rock envelope extending perhaps an AU in scale height above the radially flaring plane. Material in the disk is shielded from the radiation of the protostar by dust. Where the shielding is high, ices will form, but towards the surface of the disk, temperatures will rise so that ices become gases. The region where water ice becomes water vapor is call the snow line. It is determined by the thermal balance of lessening of the stellar radiation field through shielding by dust and geometric fall-off. Within the denser regions of the disk near its mid-plane, dust and ice grains will collide and occasionally stick together slowly growing larger as more matter falls onto the inner plane from the envelope. At size-scales greater than about 1 km in diameter, these solid bodies are called planetesimals - the seeds of future planets. Planetesimals will continue to grow through accretion and collisional coalescence until they reach size-scales of order 1000 km diameter where runaway gravitational accretion can take over forming \textit{proto-planets}. During the planet formation process, the young star will often go through stages of outflow and/or high stellar winds which will eventually disperse the protoplanetary disk limiting the mass of the planetary system. A key question is what fraction of the original protoplanetary gas and dust mass is incorporated into the planetary system, and what drives the elemental abundance ratios found in various types of planets such as terrestrial (rocky) planets and gas giants. \begin{figure} [ht] \begin{center} \includegraphics[height=5cm]{disk.png} \end{center} \caption[disk] {\label{waterdist} Model of water distribution in a protoplanetary disk. HIRMES will help determine the distribution of water looking for depletion in the hot layers above the snowline and investigate the settling of solid particles to the mid-plane.} \end{figure} Protostellar and protoplanetary disks are probed with a variety of tracers using state-of-the-art facilities across the electromagnetic spectrum including, for example optical/near-IR work on the Hubble Space telescope, mid- and far-IR work using the Spitzer and Herschel space telescopes, and submillimeter and millimeter wave interferometry with the ALMA interferometer. However these observations leave an incomplete story. Optical observations cannot probe deeply into the protoplanetary disks since optical radiation from the inner disk will be obscured by dust both in the outer disk and the enveloping molecular cloud. Mid and far-IR radiation can escape from the inner regions, but the spatial resolution of Spitzer and Hershel is insufficient to resolve the disk. This can be overcome through high spectral resolution, where one uses velocity resolved spectral profiles of lines to place the source of line emission in the disk using the fall-off in rotation velocity with radial distance from the star as predicted by Kepler's 3$^{rd}$ law. The Herschel HIFI spectrometer, which had sufficient spectral resolving power, did not have sensitivity sufficient to perform these velocity-radial distance inversions for more than a few stars. Even then it only operated at frequencies below 1.9 THz, or 157 $\mu$m wavelength while the most important tracers of the inner disk lie at significantly higher frequencies. The ALMA interferometer has both the sensitivity and spatial resolution to begin to resolve protoplanetary disks, but it only operates at frequencies lower than about 900 GHz where the preponderance of spectral lines probe low excitation gas. These low-excitation line profiles will be dominated by the outer, cooler and less dense regions of the disk and the parent molecular cloud. The best spectral lines to probe the inner regions of the protoplanetary disk will be probes of high density (n $\geq 10^{4}$ cm$^{-3}$) and intermediate temperature (T $\sim$ 100 to 500 K) gas. Such spectral lines are typically found in the mid- to far-infrared, and include the 440 $\rightarrow$ 313 28.914 $\mu$m, 643 $\rightarrow$ 616 32.313 $\mu$m, and 651 $\rightarrow$ 624 34.987 $\mu$m rotational transitions of H\textsubscript{2}O, the 1$\rightarrow$0 112.073 $\mu$m, 2$\rightarrow$1 56.230 $\mu$m, and 3$\rightarrow$2 37.702 $\mu$m rotational transitions of HD, and the $^{3}P_{0} \rightarrow^{3}$P$_{1}$ 63.184 $\mu$m fine-structure line of [OI]. HIRMES was designed to specifically address these transitions with velocity resolution up to 3 km s$^{-1}$ (spectral resolving power $R \approx 10^{5}$). Since orbital velocities in the disk will be of the order 3 to 30 km s$^{-1}$, this is sufficient to spectrally resolve these lines and make the link to emission from specific regions in the Keplerian disk. HIRMES also employs direct detection bolometers with sufficient sensitivity to be background limited at these high resolving powers so that a wide variety of sources can be investigated. Furthermore, HIRMES also contains a grating spectrometer with resolving power $\sim$ 900 that enables detection of the 43-47 $\mu$m solid-state water ice features. This combination of instrument parameters enables unique investigations of protoplanetary disks including: \begin{enumerate} \item \textbf{Tracing the total gas content in protoplanetary disks.} HIRMES can trace the total molecular hydrogen gas mass (H\textsubscript{2}) in the disk via velocity resolved studies of the HD isotopologue in protoplanetary disks. [The pure rotational lines of H\textsubscript{2} are very difficult to observe since they trace warmer gas than is in the disk, and have very weak line transition strengths.] Much of the HD emission will arise in the outer regions of the disk where gas is still accreting. The gas mass derived from HD reveals the efficiency of transferring gas into planets, especially the gas giants. HD is not observable from the ground, and HIRMES brings a unique combination of resolving power and sensitivity not available with space-based platforms. \item \textbf{Tracing the water abundance.} HIRMES can provide uniquely sensitive observations of water in the disk in regions near the snow line and combined with our HD observations we can derive its relative abundance. The water abundance is important, as it's extreme depletion relative to rock on Earth compared with its astrophysical abundance has led to the theory that water was delivered to a dry Earth long after its formation \cite{albarede} \cite{chyba}. However, it is not known if this process is common for exoplanets in the habitable zones around their stars. The water lines that best probe the habitable zones lie in the 28 to 35 $\mu$m regime with upper-level energies of 700 to 1000 K so that they trace gas at temperatures as low as T $\sim$ 200 to 300 K. These lines are unique to HIRMES and trace water vapor down to the snow line. \item \textbf{Oxygen abundances, photochemistry, and velocities in the disk} HIRMES will measure and spectrally resolve the [OI] 63 $\mu$m fine-structure line and several OH rotational lines probing the photochemistry and thermal balance in outer layers of the disk \cite{weiden}, as well as the abundance of these two important gas phase reservoirs for oxygen. In addition, the [OI] line is particularly strong | it is expected to be four times brighter than the water or HD lines | so that substantial improvements in orbital dynamics will be made though its measurement, which as for the water and HD lines can be inverted through Kepler's laws to reveal spatial structure that is not available from direct imaging. \item \textbf{Water ice abundance and processing}. HIRMES's grating spectrometer mode delivers a complete, low resolution ($R \approx 600$) spectrum for protostars from about 25 to 120 $\mu$m. Of particular interest are the water ice feature near 45 $\mu$m, where HIRMES can spectrally resolve the solid state profiles of crystalline ice (peaking near 43 $\mu$m) and amorphous ice (peaking near 47 $\mu$m). The water ice content of the disk is derived from these features and the relative crystalline/amorphous abundance ratio provides information about its thermal evolution. In addition, the HIRMES grating spectrum can reveal the ice/rock ratio in protostellar disks \cite{McClure2015}. Is ice the dominant solid mass reservoir as in solar system comets, or is rock more important in the core-accretion phase of giant planet formation? \end{enumerate} HIRMES will be capable of observing hundreds of protoplanetary disks from molecular clouds within 500 parsecs (pc) and over 100 in the three nearest associations at a distance of 140-160 pc from the solar system. With these spectral probes, HIRMES investigates the evolution of protoplanetary disks. We are particularly interested in the abundance and spatial distribution of water, and water ice at the time of planetesimal formation since it is these parameters that determine how much water can be delivered to worlds that can potentially support life. \subsection{Fine-structure line studies} As for the Hi-res, and Grating modes of HIRMES, the Mid-res and Low-res modes will have applications for a wide variety of science cases. In the Mid-res mode, a long (16 beam) slit is employed, and the Low-res modes use an imaging array of 16 $\times$ 16 pixels which greatly improve mapping speed. Here we focus on our primary drivers for these modes, which is mapping of the far-infrared line emission from both Galactic star formation regions, and over large scales in nearby galaxies. Our primary lines in these studies are the [OIII] 52 and 88 $\mu$m, the [NIII] 57 $\mu$m and [NII] 122 $\mu$m line which arise from ionized hydrogen (HII) regions formed by the ionizing radiation from nearby by O/B stars, and the [OI] 63 $\mu$m line which arises from the warm, dense photodissociation regions (PDRs) that form on the surfaces of molecular clouds exposed to far-UV (6 eV $\leq$ h$\nu$ $\leq$ 13.6 eV). These lines are quite bright and important coolants of the gas and therefore probe the sources of gas heating, which is in most cases stars. The lines and their ratios probe both the physical conditions of the gas \cite{staceyIEEE}. The [OIII] 52 $\mu$m/88 $\mu$m line ratio is sensitive to HII region gas density, the [NIII] 57 $\mu$m/[NII] 122 $\mu$m and the [OIII] 88 $\mu$m/[NII] 122 $\mu$m ratio are sensitive the hardness of the stellar radiation fields that formed the HII region, the [NIII] 57 $\mu$m/([OIII] 88 $\mu$m + [OIII] 52 $\mu$m) line ratio is sensitive to the gas phase N/O abundance. The [OI] 63 $\mu$m line is sensitive to PDR gas density, and together with the [CII] 158 $\mu$m line yields PDR gas density and the strength of the ambient far-UV radiation field. Together these lines constrain many of the important physical parameters associated with star formation in both the Galaxy and at larger scales in external galaxies. Key questions include: \begin{enumerate} \item $\textbf{What is the age of the current stellar population?}$ The age of the stellar population reveals the time since the last major star formation event. Star formation is likely stimulated by mechanisms that compress the parent molecular clouds such as gas orbit crowding at molecular bar/spiral arm interfaces, passage through spiral density waves, compression of gas by nearby supernova events, or galaxy-galaxy collisions. Therefore, to find the age of the stellar population is to probe the importance of such events to star formation. On galactic scales, the most massive star on the main sequence typically dominates the local radiation field. The most massive star will have the hottest photosphere (hardest radiation field) and shortest main-sequence lifetime, so that the [NIII]/[NII] and [OIII]/[NII] line ratios, which trace the hardness of the stellar radiation fields yield the time since the last star formation event. The former ratio has a residual gas density effect, typically removed with the [OIII] 52/88 line ratio, while the later ratio has a residual O/N abundance effect (see item 3). \item $\textbf{What is the intensity of current day star formation?}$ [OI] line imaging, together with [CII] line imaging (from Herschel/PACS or FIFI-LS on SOFIA) constrains both gas density and the strength of the ambient FUV radiation fields. The strength of the field is related to the numbers of stars per unit volume near the molecular clouds so that it is a measure of the intensity of the star formation event and the star formation efficiency, i.e. the fraction of a cloud mass that ends up in each generation of stars. \item $\textbf{How many star formation events have occured?}$ Oxygen is a primary element, formed by the fusion of carbon and helium in the cores of massive stars. Oxygen is released from the cores in the explosion of massive stars in supernovae events. Nitrogen is a secondary element, formed predominantly in the CNO cycle that fuses hydrogen into helium in stars more than about 1.3 times the mass of the sun. Nitrogen is released to the interstellar medium through dredge-up and gas outflow events that occur as these intermediate stars enter their post main sequence red giant phase of evolution. Therefore, nitrogen is enhanced every time a stellar population ages to the point that these type of stars become red giants. This time-frame is about 3-5 billion years. Each generation of stars will therefore add more N relative to O, so the N/O ratio is a measure of the numbers of generations of stars that have occurred within a star formation region, or section of a galaxy. This then relates to star formation triggers and efficiency, thereby constraining star formation models. \end{enumerate} \begin{figure} [ht] \begin{center} \includegraphics[height=6cm]{suite.png} \end{center} \caption[disk] { \label{suite} HIRMES provides SOFIA with access to new scientifically important wavelengths with high resolving power and sensitivity. For our science, HIRMES's sensitivity outperforms GREAT and FIFI-LS in their overlap region.} \end{figure} In addition, the HIRMES Mid-Res mode provides access to a variety of fine structure lines covering the wide range of ionizations states generated by fast interstellar shocks that may occur in fast molecular outflows, supernova events, or galaxy-galaxy collisions. These lines include the fine-structure lines of [SI] 25.25 $\mu$m, [FeII] 25.99 $\mu$m, [SIII] 33.48 $\mu$m, [SiII] 34.81 $\mu$m, and [NeIII] 36.0 $\mu$m and rotational transitions of OH and CO. Line ratios probe shocked gas abundances, ionization state, and density in the shocked gas yielding mass flow estimates to be tested against shock models, while the line profiles that will elucidate the kinematics of the shocked gas \cite{stacey}. \subsection{HIRMES niche on SOFIA} HIRMES enables SOFIA to open up new scientifically valuable regions of resolving power-wavelength phase space (Figure \ref{suite}). HIRMES provides new resolving power capabilities (R = 10$^{5}$, 1.2 $\times$ 10$^{4}$, 2000, and 600) at wavelengths from 28 to 63 $\mu$m, and continuous spectroscopy from 25 to 120 $\mu$m. Due to the use of bolometers and a 16 $\times$ 16 spatial array, HIRMES is significantly faster than FIFI-LS for mapping programs in the [OIII], [OI], [NIII], or [NII] lines. Due to the use of background limited bolometers, HIRMES is also more sensitive than the GREAT heterodyne receiver for high resolution spectroscopy of point sources. | HIRMES's primary goal is to study key constituents of protoplanetary disks: water vapor and ice (which play a critical role in forming giant planet cores and producing habitable conditions on terrestrial planets), and neutral oxygen, a tracer of disk chemistry and radial structure \cite{kominami}. An important step is to measure hydrogen deuteride (HD), a tracer of molecular hydrogen, which will determine the protoplanetary disk mass with an accuracy not achievable by existing instruments. The high sensitivity and resolving power we will obtain with HIRMES will resolve narrow emission lines and determine their origin spatially in the disk from velocity information allowing for detailed disk modeling. In particular, the critical transition region demarcated by the snowline at a few to 10 AU will be solely accessible by HIRMES. This will allow HIRMES to help us learn more about the delivery of water to habitable worlds. The Hi-res FPIs have already been delivered to NASA GSFC by Cornell (Aug. 2018). The Mid-res FPIs are scheduled to be delivered on Sept. 1, 2018, and the Lo-res FPIs are expected to be delivered by the end of September. First light is expected one year later, and HIRMES will become a SOFIA facility instrument in April 2020. Assembly of the complete optical bench is still on going at GSFC. The high-res FPIs have all been tested at Cornell at cryogenic temperatures (77 K) prior to delivery to GSFC using the (16 KHz) cap. bridge driven PZTs and have met all required specifications. | 18 | 8 | 1808.05218 |
1808 | 1808.01920_arXiv.txt | Our understanding of protoplanetary disks is rapidly departing from the classical view of a smooth, axisymmetric disk. This is in part thanks to the high angular resolution that (sub)mm observations can provide. Here we present the combined results of ALMA (0.9 mm) and VLA (7 mm) dust continuum observations toward the protoplanetary disk around the solar analogue GM Aur. Both images clearly resolve the $\sim$35 au inner cavity. The ALMA observations also reveal a fainter disk that extends up to $\sim250$ au. We model our observations using two approaches: an analytical fit to the observed deprojected visibilities, and a physical disk model that fits the SED as well as the VLA and ALMA observations. Despite not being evident in the deconvolved images, the VLA and ALMA visibilities can only be fitted with two bright rings of radii $\sim$40 and $\sim$80 au. Our physical model indicates that this morphology is the result of an accumulation or trapping of large dust grains, probably due to the presence of two pressure bumps in the disk. Even though alternative mechanisms cannot be discarded, the multiple rings suggest that forming planets may have cleared at least two gaps in the disk. Finally, our analysis suggests that the inner cavity might display different sizes at 0.9 mm and 7 mm. This discrepancy could be caused by the presence of free-free emission close to the star at 7 mm, or by a more compact accumulation of the large dust grains at the edge of the cavity. | \label{sec:intro} Planetary systems have long been known to form in protoplanetary disks. These objects were usually assumed to be smooth, axisymmetric structures. However, this paradigm has recently started to change; complex structures such as gaps, rings, vortices, or spiral arms are being imaged and resolved using the high angular resolution provided by current radiointerferometers, such as the Atacama Large Millimeter/submillimeter Array (ALMA; \citealp{alm15,and16,per16,ise16}) and the NSF's Karl G. Jansky Very Large Array (VLA; \citealp{oso14,car16,mac16}). The discovery of these disk substructures has sparked a lot of interest, partly because they could solve one of the fundamental problems of the core-accretion model, our current paradigm of planetary formation. According to this scenario, planets form through the sequential growth and aggregation of dust particles (\citealp{tes14}, and references therein). However, large grains in the disk suffer a gas drag force that tends to drift the particles toward the higher pressures found in the inner regions of the disk \citep{whi72}. Dust evolution models predict that the mm/cm-sized particles in the disk should drift very rapidly, hence falling onto the star before they can grow up to planetesimal sizes \citep{bra07}. Local maxima in the gas pressure of the disk can provide a solution for this problem since they can stop the migration of large particles, accumulating them and allowing them to grow even more efficiently \citep{chi10,pin12}. The disk substructures revealed by VLA and ALMA could be the result of these dust accumulations. Among the several physical processes that have been proposed to explain the observed disk substructures, non-smooth gas distributions, resulting in the onset of dust traps in the disk, are often involved. Some examples of these scenarios are the dynamic interactions with planets \citep{zhu14,bae17}, the magneto-rotational instabilities (MRI) in magnetized disks \citep{flo15}, or the back-reaction of the dust onto the gas \citep{gon15,gon17}. However, other physical processes, such as changes in dust properties related to condensation fronts or snowlines of volatiles \citep{oku16,pin17b}, can produce similar features without the need of pressure bumps. More observations are needed to better understand the commonality of disk substructures and their role in the dust evolution and planetary formation process. Here we focus on the transitional disk around the T Tauri star GM Aur, a young solar analogue (K5.5 spectral type, $L_{\star}\simeq0.9$ $L_{\odot}$, $M_{\star}\simeq1.1~M_{\odot}$; \citealp{ken95,esp11}) located at $160\pm2$ pc \citep{gai16,gai18} in the Taurus-Auriga cloud. The disk of GM Aur has been very well studied from UV to cm wavelengths. From SED modeling, it was shown to present an inner cavity in the optically thick dust \citep{cal05}, although filled with residual optically thin small particles \citep{esp10,esp11}. Furthermore, the star of GM Aur presents significant accretion of material (0.39--1.1$\times10^{-8}~M_{\odot}$ yr$^{-1}$; \citealp{ing13,ing15}), which suggests that the inner cavity is also partially filled with gas. Sub-mm observations later confirmed the presence of a disk cavity of $20-30$ au in radius, but could not unveil additional details of the disk due to the limited sensitivity and angular resolution \citep{hug09,and11}. Recently, VLA observations at cm wavelengths revealed that GM Aur contains a substantial amount of ionized gas, suggesting that the disk could be undergoing some degree of photoevaporation \citep{mac16}. Some studies have shown that photoevaporation can explain the presence of small, significantly depleted cavities \citep{ale14,owe16}. However, the size of GM Aur's cavity, combined with the presence of gas and small dust grains inside the cavity, points toward a dynamical clearing origin \citep{esp14}, indicating that planetary formation is probably taking place in GM Aur. In this paper we present new mm observations of the protoplanetary disk of GM Aur obtained with ALMA at 0.9 mm and with the VLA at 7 mm. We model the observations and find that the large dust grains distribution of the disk is mainly composed of two rings, where large dust grains are being accumulated, surrounded by a fainter and more extended outer disk that has been significantly depleted from large dust grains due to radial drift. | \label{sec:conclusion} We have presented ALMA (0.9 mm) and VLA (7 mm) observations of the dust continuum emission of the protoplanetary disk around GM Aur. The deconvolved images at both wavelengths clearly resolve the inner cavity of the disk. Additionally, the 0.9 mm image displays an extended and fainter disk component that is not evident at 7 mm. We have modeled the observed visibilities at these two wavelengths following two different approaches: an analytical model of the surface brightness profile, and a physical model of the disk that also reproduces the SED of GM Aur. The results of our model have revealed that the disk of GM Aur is composed of at least two bright rings of dust, supporting the idea that rings and other types of substructures are common in protoplanetary disks. Our physical model shows that the observed ringed morphology can be explained by an enhancement in the abundance of large dust grains in two concentric rings. These enhancements are probably caused by gas pressure bumps that are trapping the mm/cm-sized dust particles of the disk. The inner pressure bump is likely produced by the dynamical clearing of the inner cavity, while the second ring could be produced by different physical mechanisms. Finally, the inner cavity appears to be smaller at 7 mm than at 0.9 mm. This discrepancy could be indicating that the mm/cm-sized dust grains, traced at longer wavelengths, are more accumulated at the cavity edge than the submm-sized particles. Another possibility is that the free-free emission at 7 mm is producing an apparent shift of the cavity edge toward inner radii. New observations are needed to confirm this discrepancy and shed some light on its possible origin. | 18 | 8 | 1808.01920 |
1808 | 1808.04502_arXiv.txt | We present the survey design, data reduction, construction of images, and source catalog of the Atacama Large Millimeter/submillimeter Array (ALMA) twenty-six arcmin$^2$ survey of GOODS-S at one-millimeter (ASAGAO). ASAGAO is a deep ($1\sigma \sim 61$~$\mu$Jy~beam$^{-1}$ for a 250 k$\lambda$-tapered map with a synthesized beam size of $0\farcs51 \times 0\farcs45$) and wide area (26~arcmin$^2$) survey on a contiguous field at 1.2~mm. By combining with ALMA archival data in the GOODS-South field, we obtained a deeper map in the same region ($1\sigma \sim 30$~$\mu$Jy~beam$^{-1}$ for a deep region with a 250 k$\lambda$-taper, and a synthesized beam size of $0\farcs59 \times 0\farcs53$), providing the largest sample of sources (25 sources at $\ge$5.0$\sigma$, 45 sources at $\ge$4.5$\sigma$) among ALMA blank-field surveys to date. The number counts shows that $52^{+11}_{-8}$\% of the extragalactic background light at 1.2~mm is resolved into discrete sources at $S_{\rm 1.2mm} > 135$~$\mu$Jy. We create infrared (IR) luminosity functions (LFs) in the redshift range of $z =$ 1--3 from the ASAGAO sources with $K_S$-band counterparts, and constrain the faintest luminosity of the LF at $2.0 < z < 3.0$. The LFs are consistent with previous results based on other ALMA and SCUBA-2 observations, which suggest a positive luminosity evolution and negative density evolution with increasing redshift. We find that obscured star-formation of sources with IR luminosities of $\log{(L_{\rm IR}/L_{\odot})} \gtrsim 11.8$ account for $\approx$60\%--90\% of the $z \sim 2$ cosmic star-formation rate density. | \label{sec:introduction} Revealing cosmic star formation history is one of the biggest challenges in astronomy. Because a significant fraction of star formation is obscured by dust at high redshift (e.g., \cite{mada14}, for a review), infrared (IR)--submillimeter/millimeter (submm/mm) observations are required to understand the true star-forming activity. The intensity of the extragalactic background light (EBL) in the IR--submm/mm is known to be comparable to that of the EBL in the optical, also showing the importance of IR--submm/mm observations for revealing the dust-obscured activity in the Universe. Deep surveys at submm/mm (850~$\mu$m and 1~mm wavelengths) with ground-based telescopes uncovered a population of bright ($S_{\rm 1mm} \gtrsim 1$ mJy) submm/mm galaxies (SMGs; \cite{blai02, case14}, for reviews). SMGs are highly obscured by dust, and the resulting thermal dust emission dominates the bolometric luminosity. The energy source of submm/mm emission is primarily from intense star formation activity, with IR luminosities of $L_{\rm IR} \gtrsim$ a few $\times 10^{12}$~$L_{\odot}$ and star formation rates of SFRs $\gtrsim$ a few $\times 100 M_{\odot}$~yr$^{-1}$. The redshift distribution of SMGs is characterized by a median redshift of $z \sim 2$--3 (e.g., \cite{chap05, yun12, simp14, chen16, mich17, bris17}). The stellar masses and SFRs of SMGs show that they are located above or at the massive end of the main sequence of star-forming galaxies (e.g., \cite{dadd07, mich12, mich14, dacu15}). It is thought that SMGs are progenitors of massive elliptical galaxies in the present-day Universe observed during their formation phase (e.g., \cite{lill96, smai04}). The contribution of SMGs to the EBL is estimated by integrating the number counts. Blank field surveys with single-dish telescopes resolved $\sim$20\%--40\% of the EBL at 850~$\mu$m (e.g., \cite{barg99, eale00, bory03, copp06}) and $\sim$10\%--20\% at 1 mm (e.g., \cite{grev04, pere08, scot08, scot10, hats11}). It is expected that deeper submm/mm observations trace less dust-obscured star-forming galaxies, which may overlap galaxies detected in rest-frame ultraviolet (UV) and optical wavelengths. \citet{whit17} found a dependence of the fraction of obscured star formation (SFR$_{\rm IR}$) on stellar mass out to $z = 2.5$: 50\% of star formation is obscured for galaxies with $\log(M/M_{\odot}) = 9.4$, and $>$90\% for galaxies with $\log(M/M_{\odot}) > 10.5$. Deep surveys probing fainter submm objects ($S_{\rm 1mm} < 1$~mJy), which are expected to be more normal star-forming galaxies rather than ``classical'' SMGs, are essential to understand the cosmic star-formation history and the origin of EBL, however, such observations have been hampered by the confusion limit of observations with single-dish telescopes since they have large beam sizes ($\sim$$15''$--$30''$). Interferometric observations enable us to reveal faint submm sources by substantially reducing the confusion limit. The Atacama Large Millimeter/submillimeter Array (ALMA) is now detecting submm sources more than an order of magnitude fainter than ``classical'' SMGs. Because of its high sensitivity and high angular resolution, ALMA can collect serendipitous sources from a variety of data sets to probe the fainter end of the number counts (\cite{hats13, ono14, carn15, fuji16, oteo16}). These studies show that more than 50\% of the EBL at 1 mm is resolved into discrete sources at a flux limit of $\sim$0.1~mJy. These studies are based on serendipitous sources detected in fields where faint submm sources are not the main targets, which could introduce biases due to the clustering of sources around the targets or sidelobes caused by bright targets. It is necessary to conduct ``unbiased'' surveys in a contiguous field rather than collecting discrete fields in order to obtain a census on the population of faint submm sources. Surveys in a contiguous field are also beneficial for clustering analysis. During ALMA Cycle 1, the central 2 arcmin$^2$ area of the Subaru/XMM-Newton Deep Survey Field (SXDF) was observed as an ALMA deep blank field survey \citep{kohn16, tada15, hats16, wang16, yama16}. From Cycle 1 to present, the GOODS-S/Hubble Ultra Deep Field (HUDF) has been observed with ALMA in different surveys \citep{walt16, arav16, dunl17, fran18}. There are also deep surveys in overdense regions such as the ALMA deep field in the $z = 3.09$ protocluster SSA~22 field (ADF22; \cite{umeh15, umeh17, umeh18}) and the ALMA Frontier Fields Survey of gravitational lensing clusters \citep{gonz17}. The GOODS-S/HUDF field has the deepest multi-wavelengths data from X-ray to radio with ground-based telescopes and satellites such as {\sl Chandra} \citep{xue11, luo17}, {\sl XMM-Newton} \citep{coma11}, {\sl HST}/ACS/WFC3 (HUDF, CANDELS, XDF; \cite{beck06, grog11, koek11, elli13, illi13}), VLT/HAWK-I (HUGS; \cite{font14}), Magellan/FourStar (ZFOURGE; \cite{stra16}), {\sl Spitzer} (S-CANDELS; \cite{ashb15}), {\sl Herschel}/PACS (PEP; \cite{lutz11}) and SPIRE (HerMES; \cite{oliv12}), APEX/LABOCA (LESS; \cite{weis09}), ASTE/AzTEC \citep{scot10, yun12}, SCUBA-2/JCMT \citep{cowi17}, and VLA \citep{mill13, rujo16}. Spectroscopic observations have also been conducted extensively (e.g., \cite{lefe04, bram12, skel14}). The VLT/MUSE spectroscopic survey of HUDF (the $3' \times 3'$ deep region region and $1' \times 1'$ ultra-deep region) provides 3-D data cubes of this field \citep{baco15, baco17}. JWST will conduct deep multi-band imaging and spectroscopy, offering the ability to diagnose optically-faint galaxies which are difficult to study with existing optical/near-IR telescopes. \begin{figure} \begin{center} \includegraphics[width=\linewidth]{fig1.eps} \end{center} \caption{ ASAGAO region consisting of nine sub-regions (red) overlaid on the {\sl HST}/WFC3 F160W image. The orange, purple, and green regions represent the ALMA survey areas of ASPECS \citep{walt16, arav16} at 1.2~mm, HUDF \citep{dunl17} at 1.3~mm, and GOODS-ALMA \citep{fran18} at 1.1~mm, respectively. } \label{fig:region} \end{figure} The ALMA surveys of the GOODS-S field have been conducted with different survey strategies: a deep but narrow survey (4.5~arcmin$^2$, $1\sigma = 34$~$\mu$Jy~beam$^{-1}$) at 1.3 mm (HUDF; \cite{dunl17}), a shallower and wider survey (69~arcmin$^2$, $1\sigma \sim 180$~$\mu$Jy~beam$^{-1}$) at 1.1 mm (GOODS-ALMA; \cite{fran18}), and spectral scans in an area of 1~arcmin$^2$ (ALMA Spectroscopic Survey; ASPECS) at 3 mm and 1.2 mm \citep{walt16, arav16} (figure~\ref{fig:region}). The spectral scans cover the full window of the bands, offering the deepest continuum maps ($1\sigma_{\rm 3mm} = 3.8$~$\mu$Jy~beam$^{-1}$ and $1\sigma_{\rm 1.2mm} = 12.7$~$\mu$Jy~beam$^{-1}$). The faint submm sources detected in these studies are found to be on the main sequence, but located at higher stellar mass and SFR ranges (e.g., \cite{hats15, yama16, arav16, dunl17}) due to the survey detection limit. In addition, the numbers of sources studied in these surveys are still very limited, and the demand for deeper and wider surveys remains high. In this paper, we present the results of ALMA twenty-six arcmin$^2$ survey of GOODS-S at one-millimeter (ASAGAO). ASAGAO is a deep ($1\sigma \sim 61$~$\mu$Jy~beam$^{-1}$ for a 250 k$\lambda$-tapered map) and wide-area (26~arcmin$^2$) survey on a contiguous field at 1.2~mm. The observing area matches the deepest VLA C-band 5 cm (6 GHz) observations (\cite{rujo16}; Rujopakarn et al. in prep.) and the ultra-deep VLT/HAWK-I $K_S$-band images. The primary goal of this survey is to obtain a census of galaxies with $L_{\rm IR} \gtrsim 3 \times 10^{11}$~$L_{\odot}$ or SFR $\gtrsim 50$~$M_{\odot}$~yr$^{-1}$ for the understanding of the dust-obscured star-formation history of the Universe. The initial results based on the ASAGAO data have been reported by \citet{ueda18} for the X-ray active galactic nucleus (AGN) properties, and by \citet{fuji18} for morphological studies. The results of the multi-wavelength analysis are discussed in \citet{yama18}, and the clustering analysis is conducted by Yoshimura et al. (in prep.). The arrangement of this paper is as follows. Section~\ref{sec:data} outlines the ALMA observations, data reduction, and archival data used in this study, and shows the obtained images. Section~\ref{sec:source} describes the detected sources, and we list the source catalog. In Section~\ref{sec:counts}, we describe the method of creating number counts, and compare with previous studies. We present the method of constructing luminosity functions and compare with previous studies in Section~\ref{sec:lf}. The conclusions are presented in Section~\ref{sec:conclusions}. Throughout the paper, we adopt a cosmology with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_{\rm{M}}=0.3$, and $\Omega_{\Lambda}=0.7$, and a \citet{chab03} IMF. All magnitudes are given in the AB system. | \label{sec:conclusions} We performed the ALMA twenty-six arcmin$^2$ survey of GOODS-S at one-millimeter (ASAGAO). The central 26~arcmin$^2$ area of the GOODS-S field was observed at 1.2~mm, providing a map with $1\sigma \sim 61$~$\mu$Jy~beam$^{-1}$ (250 k$\lambda$-taper) and a synthesized beam size of $0\farcs51 \times 0\farcs45$. By combining the ALMA archival data available in the GOODS-S field (HUDF by \cite{dunl17} and GOODS-ALMA by \cite{fran18}), we obtained a deeper map for the 26~arcmin$^2$ area, which has a rms noise level of $1\sigma \sim 30$~$\mu$Jy~beam$^{-1}$ for the central region with a 250 k$\lambda$-taper and a synthesized beam size of $0\farcs59 \times 0\farcs53$. We find 25 sources at 5$\sigma$ and 45 sources at 4.5$\sigma$ in the combined ASAGAO map, providing the largest source catalog among ALMA blank field surveys. The flux densities are consistent with those estimated in the other ALMA GOODS-S surveys by considering the difference in observing wavelength. The larger sample allow us to construct 1.2~mm number counts with smaller uncertainties from Poisson statistics. The flux coverage of the number counts connects the fainter range probed by ALMA deep observations and the brighter range constrained by ALMA follow-up observations of single-dish detected sources. We find that our number counts are consistent with previous ALMA studies. By integrating the derived differential number counts, we find that $52^{+11}_{-8}$\% of the EBL at 1.2~mm is revolved into the discrete sources. The integration of the best-fitting function reaches 100\% at $S_{\rm 1.2mm} \sim 20$~$\mu$Jy, although there is a large uncertainty to extend the function to the fainter flux range. Deeper surveys are required to individually detect faint submm sources, which significantly contribute to the EBL. By using the 5$\sigma$ sources, we construct IR LFs in the redshift ranges of $1.0 < z < 2.0$, $1.5 < z < 2.5$, and $2.0 < z < 3.0$. Our study constrains the faintest luminosity end of the LF at $2.0 < z < 3.0$ among other studies. We find that the ASAGAO LFs are consistent with those of \citet{kopr17}, supporting the evolution of LFs (positive luminosity evolution and negative density evolution with increasing redshift) derived in \citet{kopr17}. The integration of the best-fitting LF down to the lowest luminosity of the sources ($\log{(L_{\rm IR}/L_{\odot})} = 11.78$) gives a SFRD of $7.2^{+3.0}_{-1.9} \times 10^{-2}$~$M_{\odot}$~yr$^{-1}$~Mpc$^{-3}$. We find that the IR-based star formation of ASAGAO sources contribute to $\approx$60--90\% of the SFRD at $z \sim 2$ derived from UV--IR observation, indicating that the major portion of $z \sim 2$ SFRD is composed of sources with $\log{(L_{\rm IR}/L_{\odot})} \gtrsim 11.8$. \begin{ack} We are grateful to Maciej Koprowski for providing the scaling factor of their LFs. BH, KK, YT, HU, and YU are supported by JSPS KAKENHI Grant Number 15K17616, 17H06130, and 17K14252. RJI acknowledges support from ERC in the form of Advanced Investigator Programme, COSMICISM, 321302. This study is supported by the NAOJ ALMA Scientific Research Grant Number 2017-06B and 2018-09B, and by the ALMA Japan Research Grant of NAOJ Chile Observatory, NAOJ-ALMA-190. This paper makes use of the following ALMA data: ADS/JAO.ALMA\#2015.1.00098.S, \#2012.1.00173.S, and \#2015.1.00543.S. ALMA is a partnership of ESO (representing its member states), NSF (USA) and NINS (Japan), together with NRC (Canada), 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{ack} | 18 | 8 | 1808.04502 |
1808 | 1808.04391_arXiv.txt | The observation of the x-ray pulse profile emitted by hot spots on the surface of neutron stars offers a unique tool to measure the bulk properties of these objects, including their masses and radii. The x-ray emission takes place at the star's surface, where the gravitational field is strong, making these observations an incise probe to examine the curvature of spacetime generated by these stars. Motivated by this and the upcoming data releases by x-ray missions, such as NICER (Neutron star Interior Composition Explorer), we present a complete toolkit to model pulse profiles of rotating neutron stars in scalar-tensor gravity. We find that in this class of theories the presence of the scalar field affects the pulse profile's overall shape, producing strong deviations from the general relativity expectation. This finding opens the possibility of potentially using x-ray pulse profile data to obtain new constraints on scalar-tensor gravity, if the pulse profile is found to be in agreement with general relativity. | \label{sec:intro} Neutron stars are some of the most extreme objects in the Universe, and thus, they serve as a unique laboratory to probe fundamental physics. Their large masses ($m \approx 1.4\, M_{\odot}$) combined with their small radii ($R \approx 12$ km) result in supranuclear densities at their cores, whose description challenges our current understanding of matter. The latter is encoded in the star's equation of state, whose determination is an outstanding problem in nuclear astrophysics~\cite{Lattimer:2015nhk}. Moreover, the strong gravitational fields produced by neutron stars result in gravitational potentials $\sim G m /(R c^2)$ that are nine orders of magnitude larger than what we can probe on Earth's surface~\cite{Will:2014kxa,Baker:2014zba}. Therefore, to correctly describe these stars, we must rely on a relativistic theory of gravity~\cite{Bonolis:2017fdf}. The leading theory is of course Einstein's theory of general relativity. During its centennial existence, the theory has shown a remarkable predictive power, being consistent within all experimental tests carried out so far, ranging from local, Solar System experiments~\cite{Will:2014kxa}, to the spectacular detection of gravitational waves by the LIGO/Virgo collaboration~\cite{Abbott:2016blz,TheLIGOScientific:2016src}. This consistency with observation is even more striking when one notes that the theory (unlike most of its alternatives) does not possess any free parameters that can be tuned to make its predictions agree with Nature. Given the success of general relativity, why should we even consider modifications to it and examine their observational consequences? The reasons are many, but they can be organized in two main classes~\cite{Clifton:2011jh,Berti:2015itd,Alexander:2009tp}. On the observational front, the late time expansion of the Universe~\cite{Riess:1998cb,Perlmutter:1998np}, the rotation curve of galaxies~\cite{Sofue:2000jx,Bertone:2016nfn}, the baryon-antibaryon asymmetry~\cite{Spergel:2003cb,Canetti:2012zc}, and other cosmological observations seem to point at either exotic dark fluids or modifications to general relativity. On the theoretical front, the incompatibility of general relativity with quantum mechanics has prompted many attempts at extensions, from string theory to loop quantum gravity and other variations. Can neutron star observations\footnote{We here focus on properties of {\it individual} stars. Neutron stars have already proven invaluable tools to constrain deviations of general relativity when either in binaries~\cite{Damour:2007uf,Wex:2014nva,Kramer:2016kwa} (or triple~\cite{Archibald:2018oxs}) systems or more recently binary neutron stars mergers~\cite{TheLIGOScientific:2017qsa} (see e.g.~\cite{Sakstein:2017xjx,Baker:2017hug,Ezquiaga:2017ekz,Creminelli:2017sry,Bettoni:2016mij}).} be used to learn about gravity in extreme environments? A general prediction of modified theories of gravity is that the bulk properties of the star (e.g. its radius and mass) are different from those predicted by general relativity. Tests of gravity in this direction are however limited by a strong degeneracy problem between the equation of state and the gravity theory: the modifications of the bulk properties of these stars caused by changes in the gravitational theory are (often) degenerate with modifications due to different equations of state~\cite{Glampedakis:2015sua}. One option to bypass this issue is to focus on electromagnetic phenomena in the vicinity of neutron stars~\cite{Psaltis:2008bb}. These phenomena include e.g.~atomic spectral lines~\cite{DeDeo:2003ju}, burst~\cite{Psaltis:2007rv,Glampedakis:2016pes} and quasiperiodic oscillations~\cite{DeDeo:2004kk,Doneva:2014uma,Pappas:2015npa,Glampedakis:2016pes}. One can argue that in these scenarios one can in principle probe the {\it exterior spacetime} of the star, offering a glimpse on possible deviations from general relativity, without worrying about the intricacies of the stellar interior. The observation of x-ray waveforms or pulse profiles from rotating neutron stars is another potentially interesting phenomenon to consider~\cite{Arzoumanian:2009qn}. In this scenario, a region of of the neutron star's surface becomes hot (relative to the rest of the star's surface) generating an x-ray flux modulated by the star's rotation. This {\it hot spot} can be formed in a number of situations (see~\cite{Poutanen:2008pg,Ozel:2012wu,Watts:2016uzu} for reviews). In accretion-powered pulsars, material is channeled through the magnetic field lines and heats the star up when it reaches its magnetic poles. In burst oscillations, a thermonuclear explosion caused by infalling matter results in a hot spot on the surface of the accreting neutron star. In all these cases, the modeling of the resulting waveform, combined with x-ray observations, allows for the extraction of a number of properties of the source, including the neutron star's mass and radius, see e.g.~\cite{Miller:1997if,Weinberg:2000ip,Muno:2002es,Bogdanov:2012md,Poutanen:2003yd,Lo:2013ava,Miller:2014mca,Stevens:2016xiw,Lo:2018hes,Salmi:2018gsn}. The ongoing NICER (Neutron star Interior Composition Explorer) mission~\cite{2012SPIE.8443E..13G,2014SPIE.9144E..20A,2017NatAs...1..895G} offers a substantial improvement over the preceding x-ray observatory, the Rossi X-ray Timing Explorer (RXTE), opening the path to measurements of stellar masses and radii with unprecedented accuracy -- with immediate implications to our understanding on the neutron star equation of star. Given, the scientific potential of NICER, it is natural to ask: {\it can we use its observations to probe the strong-field regime of gravity}? In this paper, we take the first necessary steps to find an answer to this question. We find hints that NICER can indeed be used to probe the strong-field regime of gravity, but a definite answer will require a plethora of theoretical and data analysis work that this paper now enables (see~\cite{Sotani:2017bho,Sotani:2017rrt} for independent recent work in this direction). \subsection{Executive summary} We present a complete toolkit to model the x-ray flux from radiating neutron stars in scalar-tensor gravity, one of the most well-studied and well-motivated extensions of general relativity~\cite{Chiba:1997ms,Sotiriou:2015lxa}. This class of theories extends general relativity by introducing a scalar field that couples to the metric nonminimally, thus violating the strong-equivalence principle. Our formalism and the resulting toolkit are completely model-independent within this class of scalar-tensor theories. Moreover, the resulting waveforms include Doppler shifts, relativistic aberration and time-delay effects, thus extending~\cite{Sotani:2017rrt}. All of these effects are critical ingredients that are necessary to develop an accurate pulse-profile model, and thus, our toolkit now enables a complete data analysis study that will be carried out in the future. Figure~\ref{fig:pp_example} shows a sample waveform that takes all the previously mentioned effects into account, illustrating the difference between the predictions of both theories. As will discuss in Sec~\ref{sec:pulseprofile}, we find that the presence of the scalar field influences the exterior spacetime of the neutron star, altering the bending of light and the time-delay experienced by photons emitted by the star. The net result of these contributions is a waveform that can be considerably different in comparison to general relativity's predictions. Such a model now enables, for the first time, a serious data analysis investigation of whether such waveform differences can be detected/constrained with data or whether they are degenerate with other waveform parameters. \begin{figure}[t] \includegraphics[width=0.48\textwidth]{fig_illustrative_pulseprofile.pdf} \caption{An illustrative bolometric flux for a single hot spot over one revolution of the rotating neutron star in general relativity and scalar-tensor gravity. The star has $u \equiv 2 G_{\ast} m / (R c^2) = 0.5$ and mass $m = 1.4\, M_{\odot}$ in both cases. The only difference between the two models is the presence of a nonzero scalar charge for the star in scalar-tensor gravity. The magnitude of the scalar charge is quantified by the scalar charge-to-mass ratio $Q$ (described in Sec.~\ref{sec:review}) which is zero in general relativity and controls by how much the spacetime is different from the usual Schwarzschild spacetime. This particular examples a sample star labeled $c_3$ as discussed in Sec.~\ref{sec:review} (cf. Table~\ref{tab:stellar_models}).} \label{fig:pp_example} \end{figure} The remainder of this paper is organized as follows. In Sec.~\ref{sec:review}, we briefly review the basic of scalar-tensor gravity, discussing the properties of the exterior spacetime of neutron stars in this theory. Next, in Sec.~\ref{sec:pulseprofile}, we present in detail how we construct our pulse profile model. Finally, in Sec.~\ref{sec:results} we show a few examples of the pulse profile of burst oscillations and discuss the modifications introduced by scalar-tensor gravity. We discuss in details the roles played by different effects on the overall shape of the resulting waveform. We close with Sec.~\ref{sec:conclusion}, summarizing our main findings and discussing some possible extensions and applications of this work. | \label{sec:conclusion} We introduced a Just-Doppler approximation for calculating pulse profiles of rotating neutron stars in scalar-tensor gravity. The main result, encapsulated in Eq.~\eqref{eq:main_result}, allows one to calculate the waveforms, including effects of the strong-gravity and special relativistic effects, generalizing the Schwarzschild plus Doppler approximation to scalar-tensor theories of gravity. We presented a selection of sample results for burst oscillation waveforms of an infinitesimal hot spot and discussed the implications of the presence of a nonzero scalar charge on it. The model independent character of the formalism opens the possibility of constraining a large class of scalar-tensor gravity models with upcoming x-ray timing data releases from NICER. In this regard, our work is close in spirit to~\cite{Horbatsch:2011nh}, except that that work focused on binary pulsars. In a forthcoming paper, we will present a statistical likelihood analysis discussing the strength of potential future constraints on scalar-tensor gravity with pulse profile observations. We emphasize, however, that whether future constraints can be placed on scalar-tensor gravity with NICER data will require a full data analysis study that varies over all model parameters in the presence of noise; such a study is now possible thanks to the pulse-profile model presented here, and it will be carried out in the future. Despite the generality of the present formalism, it is necessary to discuss its limitations and signal in which directions it can be improved further. At considerably large rotational frequencies ($\nu \gtrsim 300$ Hz), the Schwarzschild-Doppler approximation becomes inappropriate as discussed, e.g.~in~\cite{Cadeau:2006dc}, because of the rotation-induced quadrupolar deformation of star. To include this effect in the pulse profiles, one must first extend the Just metric to rotating stars. To leading-order in rotation, that is, including only frame-dragging effects while keeping a spherical geometry for the star, this calculation was done in~\cite{Damour:1996ke}. Using the Hartle-Thorne perturbative expansion, Berti and Pani~\cite{Pani:2014jra} extended this work to second-order in rotation, which includes the quadrupolar deformation of the star. Alternatively, one could also work entirely numerically and carry out ray-tracing calculations (as in~\cite{Cadeau:2004gm,Cadeau:2006dc,Nattila:2017hdb,Psaltis:2013zja}), evolving the photon geodesics in a numerically constructed spacetime for rotating neutron stars in scalar-tensor gravity~\cite{Doneva:2013qva}. A caveat of this numerical approach is that the spacetime can only be determined numerically {\it after} choosing a particular conformal factor $A(\vp)$ \emph{and} the equation of state, therefore inevitably making it model-dependent. Another approach would be to follow the work in~\cite{Pappas:2014gca,Pappas:2015npa,Pappas:2016sye} and use the Ernst formalism to obtain an axisymmetric spacetime in Weyl-Papapetrou coordinates in terms of an multipolar expansions of the multipole moments of the rotating neutron star and the scalar field. The multipole moments appear as free constants in the spacetime metric, thus allowing one to develop an extension of the model-independent pulse profile model introduced here. In this approach, and as an intermediate step, one would first have to {\it estimate} numerically the values of these moments in order to have an analytic spacetime that describes accurately the one constructed numerically, for example in~\cite{Doneva:2013qva}. We emphasize, however, that Cadeau et al.~\cite{Cadeau:2006dc} have shown that even in the presence of stellar oblateness, the Schwarzschild-Doppler approximation works quite well at these frequencies when $\theta_{\rm s}$ and $\iota_{\rm o}$ are {\it near the equator} as in the case of Fig.~\ref{fig:illustrative_90}. Geometrically, this is due to the small difference between the normal vector ${\bm n}$ on a spherical and a oblate surface near the equator, which suppresses the effect of the star's oblateness. Because of the purely geometrical nature of this argument, we expect the Just-Doppler approximation to accurately describe the pulse profile when ${\theta}_{\rm s} \approx \iota_{\rm o} \approx 90^{\circ}$ even for rapidly rotating neutron stars in scalar-tensor gravity. Finally, it could also be interesting to consider other theories of gravity or to use a parametrized framework -- to consider model-independent deformations of the Tolman-Oppenheimer-Volkoff equations and the Schwarzschild spacetime -- as introduced in~\cite{Glampedakis:2015sua,Glampedakis:2016pes}. Work in all of these directions is currently underway and will be reported in several future publications. | 18 | 8 | 1808.04391 |
1808 | 1808.06488_arXiv.txt | {The way massive pre-main sequence (pre-MS) binaries form and evolve is by large an open question. Here we systematically address this topic from the perspective of stellar structure and evolution models. We are looking at the pre-MS detached evolution of massive binaries that evolve, through accretion growth, from a small mass stellar cores to massive zero age main sequence binaries.} {We explore the initial conditions that lead to detached binary evolution of massive pre-MS binaries and ask how large a fraction of the observed binary systems may have been initially formed as low-mass protobinaries and later undergone a significant accretion phase while remaining detached.} {We develop a family of analytic models to describe the orbital separation, $a$, and mass ratio, $q$, evolution. For a given mass accretion rate onto the binary system, we define a recipe for distributing this mass between the two components. For this we introduce a parameter $\eta$, such that $\tfrac{\dot{M}_2}{\dot{M}_1}=q^\eta$ at any time, to determine the binary mass ratio evolution. Depending on the choice of $\eta$, any type of mass ratio evolution is possible. Furthermore, we use MESA, a detailed stellar structure code, to calculate an extensive grid of binary sequences where a protobinary undergoes accretion, and we identify the initial conditions that separates detached from non-detached pre-MS binary evolution.} {A value of $\eta$ around 2 allows accretion growth in detached systems to form close massive binaries on the Zero Age Main-Sequence with minimum orbital periods down to about 1.2 days for $M_{\rm 1,ZAMS}=20-30$\msun{} twin-binaries. $\eta=2$ can also reproduce the observed population of binary systems with primary stars above 6\msun. } { The whole observed range of massive close binaries can form via accretion growth in detached systems, making the binary formation channel of accretion growth a strong contender to explain the formation of massive close binaries, including progenitors of coalescing binary black holes.} | \label{sec:intro} The formation of massive stars, >8\msun{}, is by large an open question, and even more so, is the formation of massive binary stars. This is in part because massive stars are rare and short lived, but also because their formation process lies hidden inside optical thick natal clouds that are not removed until the massive stars launch strong winds or some of them undergo supernovae whereby the cloud is cleared. However, this only happens after the massive stars have formed. We would expect the formation of massive single and binary stars to be subjected to the same physical processes and environments. Therefore we begin by a brief discussion on massive single star formation. There are two proposed ways to form massive single stars, both of which involves lower mass protostars. These are the accretion and the merging scenarios, respectively. In the accretion scenario a less massive star can accrete material from its surroundings and grow to become massive. The alternative is the merging scenario where the merging of two intermediate mass stars can produce a single massive star. How massive a star can become via accretion is determined from its metallicity and the geometry of the infalling material \citep[e.g.][]{larson1971, kuiper2010}. The stellar structure of single stars undergoing accretion has been studied by many authors \citep[e.g.][]{palla1992, behrend2001, omukai2001, omukai2003, yorke2008, hosokawa2010, hosokawa2013, hosokawa2016,lee2016, haemmerle2016, haemmerle2017}. Yet, open questions remain in this picture of which we note challenges on relations between the accretion rate, metallicity and swelling due to internal restructuring and shell D-burning, and transport of angular momentum and its impact on stellar rotation. How is angular momentum from the infalling material coupled to the star and transported with the accretion flow into the star and what is the effect on the stellar structure? A potentially important effect is the modified Eddington effect due to rapid rotation, also known as the $\Omega\Gamma$-limit \citep{maeder2000}. A rapid rotating star has a lower effective gravity due to the centrifugal force and this gives rise to a lower modified (effective) Eddington luminosity limit. In accreting stars angular momentum is carried into the star which is then spun up, and may reach a critical rotation, thus stopping the accretion. \citet{lee2016} looked at the $\Omega\Gamma$-limit in the formation of massive population III stars via accretion growth. They concluded that protostars quickly reaches a critical Eddington limit, whereby the accretion stalls and as a result producing stars more massive than the 20-40\msun{} seems impossible due to the $\Omega\Gamma$-limit. \citet{lee2016} treated rotation of their models in post processing and assumed solid body rotation. \citet{haemmerle2017} used a new version of the Geneva Stellar Evolution code to follow the accretion growth of protostar, keeping track of its structure and rotation profile. This allowed \citet{haemmerle2017} to observe how differential rotation plays a crucial role in preventing the star from reaching the $\Omega\Gamma$-limit already at the swelling phase postulated by \citet{lee2016}. As is shown in \citet{haemmerle2017} the formation of stars more massive than 40\msun{} is possible, if the protostar can dissipate 2/3 or more of the angular momentum gained from the accretion. This is possible by a strong breaking mechanism which modulate the angular momentum accretion history of the star. The second possibility, is to grow low mass stars into high mass stars through merger events, which requires a high stellar number density \citep{bonnell2001,bonnell2005}. The stellar structure resulting from the merger of two stars was investigated by \citet{glebbeek2013}. They found the mass loss in the merger event for head-on collisions is less than 10 \%, which permit the production of massive stars via mergers. \citet{schneider2016} argued the merger event produces large scale magnetic fields in the new star, and if so, that about 10 \% of OBA pre-MS and MS which displays such magnetic fields are merger products. \citet{baumgardt2011} performed N-body simulation of young star forming embedded clusters, accounting for the process for accretion growth onto each star. They find, that increasing the stellar number density of the embedded cluster increases the number of interactions. However, in many cases, the close encounter does not produce a merger, but a runaway star instead. These conclusions were revised by \citet{railton2014} who used pre-MS isochrone tracks of stars from 1-8 \msun{} to describe the earliest stages of stellar evolution. They find that collisions were common and could produce high mass stars. However, using pre-MS isochrone tracks of constant mass ({\it i.e.} non-accreting) stars gives an unrealistically large cross-section for interactions. Observing massive stars in the making is very hard as they are surrounded by infalling material and outflows. If they are sitting in an embedded cluster, still surrounded by their natal cloud, observations are even harder. % Since stars, including binaries, are generally believed to form inside embedded clusters \citep{lada2003}, both accretion and dynamical interactions should play a role in forming binary systems: accretion, because the embedded cluster offers a reservoir of gas, and dynamical interactions due to the potentially high stellar number density. This naturally leads to ask if massive binary stars may also form from low-mass proto binaries and accretion growth without reaching the Roche limit or being scattered. % An interesting example is the observation of the HII outflow region S106, which contains a massive pre-MS binary, S106IR, with an orbital period of 5 days, primary mass of $\approx19$ \msun{} and mass ratio $q\approx0.17$ \citep{comeron2018}. The sparse number of other sources in close vicinity makes S106IR a prime candidate for massive binary formation via accretion growth. Another candidate is IRAS 04191+1523 which is a very low mass binary. It has a high mass ratio $q\approx0.85$ and orbital separation $a\approx860$ au \citep{lee2017}. \citet{lee2017} proposed that IRAS 04191+1523 formed from a process of fragmentation in a turbulent medium, as initially suggested by \citet{bonnell2005}. Massive binaries can result from dynamical interactions, when, e.g., two massive stars form a bound system as they pass near each other at sufficiently low velocities for their mutual gravitational attraction to govern their further dynamics. If dynamic interactions would be the means by which the majority of binary stars form, then the stellar structure due to accretion growth is not relevant. We think however that although dynamical interactions can be important, especially in dense stellar systems, it is probably not the only channel through which close massive binary stars are formed. Here we focus on the formation of binaries via accretion growth and we define a binary system to be composed of a primary star $M_1$ and a companion star $M_2$ separated by some orbital separation $a$ and eccentricity $e$, with an orbital period $P$. We further define the binary mass ration $q\equiv \tfrac{M_2}{M_1}$ and call a binary with $q\geq0.95$ a twin-binary \citep{lucy1979}. The accretion process requires that a low-mass protobinary is concepted out of the gas phase. Three theories has been proposed to describe the formation of protobinaries in molecular clouds. The oldest one is the fission theory, where a rapidly spinning proto-star splits into two \citep{jeans1919}. However, the compressible fluids from which stars form, work against this formation scenario, as the rapidly spinning protostar generates spiral arms that dissipate angular momentum of the protostar before it can be split \citep{tohline2002, bate2015}. Alternatively, a protobinary forms from the gravitational fragments of a collapsing cloud, due to local density fluctuations which produce areas of smaller free fall times, hence a faster collapse. These fragments are then bound once they fragment and start their contraction \citep{boss1979, boss1988}. Finally, it has been proposed that sufficiently cool and massive protostar discs may be dynamically gravitational unstable and fragment to form a protobinary \citep{kratter2006}. In the present paper, we will remain agnostic about the exact protobinary formation scenario and we will assume that protobinaries are formed at relative low mass and undergo a phase of accretion growth that brings them to their final ZAMS masses. During this phase we ignore any potential dynamical interactions with the rest of the cluster or merger events. Several other studies of protobinaries undergoing accretion growth have concentrated on the accretion of material onto a protobinary of some mass ratio using hydrodynamic simulations \citep[][]{bate1997ballistic, bate1997gas,ochi2005,young2015a,young2015b}, but little is known on their stellar structure evolution and its role in binary formation. To our knowledge, the earliest study of accreting binaries was by \citet{rayburn1976} who considered accretion onto a main sequence binary that encounters a interstellar cloud. \citet{artymowicz1983} studied for the first time the pre-main sequence binary stars undergoing accretion growth. He used a three-body approximation to describe the accretion onto each star. The first to include the stellar structure to study the formation of binaries was \citet{tutukov1983}. He noticed that the orbital period distribution featured a distribution with a clear break around orbital periods of 100 yr, indicative of two formation mechanisms. \citet{tutukov1983} assumed binary formation via the fission process and that wide binaries ($P_{\rm orb} \gtrsim 100$ yr) fissioned out the collapsing protocloud before a hydrostatic-equilibrium gas-dust core could form. The close binaries ($P_{\rm orb}\lesssim 100$ yr) would fission out of the protocloud after the gas-dust core has formed. The most recent study on the stellar structure of accreting proto-binaries, to our knowledge, is by \citet{krumholz2007}. They used a simple one-zone model by \citet{mckee2003}, under the assumption of constant accretion rates, for binaries in circular orbits, and with fixed mass ratio. \citet{krumholz2007} conclude that in case that mass transfer occurs, it is always the primary star that fills its Roche lobe and this is due to internal restructuring and subsequent shell D-burning. They also find that mass transfer is unstable unless the initial mass ratio $q \geq 0.8$ and the donor is in the D-shell burning. Overall they suggest that mass transfer can explain the observed excess of twin-binaries. Many open questions on the formation via accretion remains to be highlighted. There is no reason to suspect, that the binary mass ratio remains constant during the accretion phase. It is not clear how close of a massive binary can be formed from accretion, as this depends both on the initial orbital separation- and the mass ratio-evolution, as well as the protobinary's accretion history. We expect the orbital period- and mass ratio-evolution to also put limits on the fraction of the binary population that can be explained as originating through an accretion process. Finally, it is not clear that the primary star is the only star that can overflow its Roche lobe during the pre-MS accretion phase. In fact, we will show that even under the assumption of constant mass ratio, the secondary star may also overflows its Roche lobe while undergoing accretion. In this study, we will simulate the accretion phase of a pre-MS binary system using the 1D stellar evolution code Modules for Experiements in Stellar Astrophysics \citep[MESA][]{paxton2011,paxton2013,paxton2015,paxton2018} to evolve, in detail, the stellar structure of each star as they accrete. Further, the MESA binary module allows us to track the orbital separation as the stars accrete. Our main goals in this first paper are to study the detached evolution of binaries and investigate the relation between initial binary orbit with post-accretion binary orbits, primary mass and mass ratio. The pre-MS evolution of binary stars undergoing accretion is a complex study and we address here only detached evolution in circular orbits. Effects of mass transfer, eccentricity, and tidal evolution during the pre-MS of accreting binary stars is beyond the scope of the present paper. Here we develop an analytic model for a pre-MS binary system undergoing accretion, that account for the evolution of orbital separation $a$ and mass ratio $q$. Further, we simulate a large grid of binary accretion sequences generating binaries with primary star masses between 6-60\msun{} at zero age main sequence (ZAMS) given different initial conditions. From the grid we deduce the limit separating detached from non-detached binary formation via accretion growth at the ZAMS. This detached vs. non-detached limit at ZAMS, we compare to the observed ZAMS binary populations mass ratio- and orbital period distribution for the relevant range of primary masses, to see if any observed binary systems at ZAMS suggests a non-detached pre-MS evolution. Our paper is built up as follow. In sect. \ref{sec:theory} we develop our analytic model of pre-MS binaries undergoing accretion. First, we consider in sect. \ref{sec:orbit_evo}, the orbital separation evolution dependent on the change in pre-MS binaries' orbital angular momentum as the protostars accrete material from some external source. Secondly, in sect \ref{sec:mass_ratio_evo}, we present our model for the change in mass ratio in response to the accretion. In particular we introduce a new parameter $\eta$ that governs the change in mass ratio. In Sect. \ref{sec:eta} we briefly review studies using hydrodynamic simulations of protobinaries undergoing accretion and estimate $\eta$ from these sources. We also present analytic estimates of $\eta$. Section \ref{sec:single_star} describes our MESA model of an accreting pre-MS star and discusses its structure and evolution while accreting. Section \ref{sec:grids} presents our grids of pre-MS binary accretion sequences. Sections \ref{sec:result}, \ref{sec:discussion}, and \ref{sec:conclusion} are reserved to present our results, discussion and conclusion respectively. | \label{sec:conclusion} In this work we have explored how binary star systems might become massive and close through a process of accretion that started from a low mass proto binary < 2\msun. We have done this by formulating a family of analytic models to describe the change in orbital separation and mass ratio as the binary system undergoes accretion growth. The orbital separation evolution we described assumes a disc geometry under the assumption that the two binary stars accrete material that carry the specific orbital angular momentum of the accretring star. The mass accreted is shared between the two components according to a law $\tfrac{\dot{M}_2}{\dot{M}_1}=q^{\eta}$ at all times. We combined our model of mass ratio and orbital separation evolution with the stellar evolution code MESA, and computed grids of pre-MS binary stars undergoing accretion growth, varying the initial orbital separation and mass ratio. We computed such grids for a range of $\eta$-values from 0.5 to 3. With these grids we explored the stellar structure and evolution in HR diagram of protobinaries undergoing accretion growth. From the grids we inferred the limits on how massive a binary system could become, given its initial conditions, before either of the binary stars would fill their Roche lobe, at which point our computations stopped. This yielded the result that accretion into binary stars during pre-MS can lead to both the primary star and the secondary star overflowing their Roche lobe. A final step we took was to compare the inferred limits of our model grids for different values of $\eta$ with the observed population of binary stars from \citet{moe2017}. From our efforts we note that: \begin{itemize} \item Accretion via discs allows binary stars to accrete more mass before filling their Roche lobe compared to accretion from a spherical symmetric accretion process. This is similar to accretion growth of single stars, where disc accretion allows single stars to overcome the luminosity barrier (see eq. \eqref{eq:ai2af_spherical} and \eqref{eq:afai_disk}). \item The mass ratio evolution can be parameterised by $\eta$ such that $q^{\eta}=\tfrac{\dot{M}_2}{\dot{M}_1}$, and determines wether the mass ratio goes towards unity or zero, see eqs. \eqref{eq:qi2qf} and \eqref{eq:qf2qi} and Fig. \ref{fig:qi2qf}. Two special cases are $\eta=1$ in which case the mass ratio remains a constant and $\eta=0$ in which the binary total mass accretion is split equally by the two stars or $q=1$. If $\eta<1$ the mass ratio evolution is towards unity and this puts a limit on the mass ratio distribution that emerges after the accretion phase by introducing a minimum final/ZAMS mass ratio of binary stars forming from accretion even if the initial mass ratio is very small (see Fig. \ref{fig:eta_qi2qf}). \item If $\eta>1$ the mass ratio evolution is towards zero and there is no upper or lower limit to the final mass ratio after end accretion growth. \item Estimating $\eta$ from hydrodynamic simulations of protobinaries is difficult and depends on the initial assumption of gas sound speeds and specific orbital angular momentum of the infalling gas, see Fig. \ref{fig:eta_q_estimates}. Using an analytic approach assuming Bondi-Hoyle accretion or Roche lobe radius we found that $\eta$ lies in the range 1-3. \item Both the primary and secondary protostars can overflow their Roche lobe and whether the former or the latter occurs is strongly correlated with the swelling phase of either star, see Figs. \ref{fig:grid_min_aiM1zams_eta1p0} and \ref{fig:grid_min_aiM1zams_eta2p0}. \item The minimum orbital period at ZAMS decreases with increasing primary mass and increasing mass ratio. This allows for the formation via accretion growth of chemically homogeneously stars in close binaries (see Figs. \ref{fig:minLogP_0.5} and \ref{fig:minLogP_2.0}). Systems like VFTS352 may have formed from accretion growth and ending as a contact system. It also seems that this result is somewhat insensitive to the value of $\eta$. \item For $\eta=1$ or $\eta=2$, the mass ratio and orbital period distribution of most binary systems, 88 to 99\%, for a primary star mass range from 6 to 50 \msun{} could be reproduced (see Figs. \ref{fig:pdf_permit_logP_zams} and \ref{fig:pdf_permit_q_zams}). The massive pre-MS binary S106IR \citep{comeron2018} can also be produced from the accretion channel if $\eta$ is 1 or 2. \item From the ZAMS distribution of binary stars we infer the protobinary mass ratio distribution at the onset of accretion and find the orbital period range to coincide with pre-MS solar type stars \citep{mathieu1994}, that likely have not accreted significantly. See Figs. \ref{fig:pdf_permit_logP_proto} and \ref{fig:pdf_permit_q_proto}. \item The picture of massive binary stars forming at small mass and subsequent accretion growth is a likely explanation for the distribution of the observed range of massive binary stars, and could allow for the formation of progenitors of coalescing binary black holes. \end{itemize} There are still many open questions on how binary systems form and we have not here been able to account for all possible relevant physics. In brief we mention; eccentric orbits, tidal forces, and dynamic effects, which will affect the conclusion of this paper and we look forward to continuing this effort exploring more physics. | 18 | 8 | 1808.06488 |
1808 | 1808.04996_arXiv.txt | We report an orbital characterization of GJ1108Aab that is a low-mass binary system in pre-main-sequence phase. Via the combination of astrometry using adaptive optics and radial velocity measurements, an eccentric orbital solution of $e$=0.63 is obtained, which might be induced by the Kozai-Lidov mechanism with a widely separated GJ1108B system. Combined with several observed properties, we confirm the system is indeed young. Columba is the most probable moving group, to which the GJ1108A system belongs, although its membership to the group has not been established. If the age of Columba is assumed for GJ1108A, the dynamical masses of both GJ1108Aa and GJ1108Ab ($M_{\rm dynamical,GJ1108Aa}=0.72\pm0.04 M_{\odot}$ and $M_{\rm dynamical,GJ1108Ab}=0.30\pm0.03 M_{\odot}$) are more massive than what an evolutionary model predicts based on the age and luminosities. We consider the discrepancy in mass comparison can attribute to an age uncertainty; the system is likely older than stars in Columba, and effects that are not implemented in classical models such as accretion history and magnetic activity are not preferred to explain the mass discrepancy. We also discuss the performance of the evolutionary model by compiling similar low-mass objects in evolutionary state based on the literature. Consequently, it is suggested that the current model on average reproduces the mass of resolved low-mass binaries without any significant offsets. | \label{sec1} \indent \subsection{Evolutionary models and their observational constraints} In recent years, many new low-mass companions to stars have been detected and characterized following advances of various complementary methods. Several of these companions have inferred masses below the deuterium-burning limit ($\sim13M_{\rm jup}$), and are therefore commonly referred to as planets \citep[e.g.][]{chauvin04, marois08}. Such directly imaged exoplanets are in a very early stage of their evolution, since current high-contrast imaging searches focus on exoplanets around young stars, which benefits the detection of thermal emission \citep[e.g.][]{lafreniere07b}. Stellar evolutionary models (i.e., theoretical isochrones) are important for characterizing physical parameters of low-mass objects, and to understand their formation and evolution. However, such evolutionary models for low-mass objects are poorly constrained by observations, although several studies on benchmark stars have been performed to evaluate their performance. \cite{hillenbrandwhite04} assembled several benchmark stars, including mainly eclipsing and visual binaries in the PMS or MS\footnote{PMS: Pre-main sequence, MS: Main sequence} phase, for which both dynamical masses and the photometric information necessary for deriving model masses were available, and performed a range of compatibility tests. As a result, they found that evolutionary models under-predict the masses (relative to dynamical for lower-mass stars, $<0.5M_{\odot}$). \cite{stassun14} focused on eclipsing binaries in the PMS phase as benchmark stars, and suggested that many models tend to over-predict masses by 10--30$\%$, with 20--50$\%$ scatters. This implies that current evolutionary models cannot reproduce the properties of those cooling young stars, probably because the models still lack the inclusion of several effects induced by the magnetic field \citep[e.g.][]{feiden13a} and the early accretion history in the PMS phase \citep{baraffe09, baraffe10}. Furthermore, the eclipsing binary systems used in these studies contain very close stellar pairs, and may have experienced various forms of interaction in the early stages of their evolution \citep{feiden16a}. Recently, the orbits of several PMS objects were determined with high-resolution imaging using adaptive optics and radial velocity measurements \citep[e.g.][]{montet15}. These resolved binaries are important benchmark stars to test stellar evolution, primarily because of their separations between the stellar components of the systems. The separations should be wide enough to be free from the tidal interactions between the components, compared with the eclipsing binaries. If the systems are the members of nearby young moving groups, their well-determined ages significantly help to characterize the properties of the systems. Using the dynamical mass and broad-band luminosities, the age of these multiple systems have been obtained, and seem to be consistent with the ages of young moving groups inferred from theoretical isochrones \citep{montet15, nielsen16}. However, the number of resolved binaries suitable for calibration purposes is still small at this stage, and the performance of evolutionary models needs further testing to advance our understanding of lower-mass objects. \subsection{Objectives of this work} In order to characterize young low-mass objects down to brown dwarfs and giant planets, and to understand their evolution, the calibration of evolutionary models is essential. However, dynamical mass estimation through orbital characterization is difficult to achieve for imaged planets detected so far, since the number of imaged planet is still small and their orbital periods are too long. Meanwhile, M-dwarfs have sufficiently low masses for their evolutionary phase to be observable for up to about 100 Myr \citep[e.g.][]{baraffe98}, and they are the most abundant population of stars in the solar vicinity. Planetary mass objects and M-dwarfs have several relevant physical aspects in common, including convection and molecular opacities. Furthermore, the contracting and cooling evolution needs to be known in order to advance the understanding of evolutionary aspects relevant for lower-mass objects, such as cloud formation and dissipation \citep[e.g.][]{burrows06}. Hence, orbital characterization of young M-dwarf binaries is a step toward expanding the understanding of yet lower-mass objects. Here we present the orbital characterization of one such resolved M+M binary in the PMS phase: GJ1108A, which was observed as part of the SEEDS survey \citep[Strategic Exploration of Exoplanets and Disks with Subaru, ][]{tamura09}. The rest of the paper is organized as follows. The target properties, and in particular the age of the star, are presented in Sect. 2. Observations and analysis are detailed in Sect. 3. Section 4 outlines the orbital solution for the system. In Sect. 5, we compare the dynamical mass of the system with the model-derived mass, and also discuss the performance of the recent stellar evolutionary models. The results are summarized in Sect. 6. | \indent \subsection{Revisit of the kinematic age for the GJ1108A system} \label{subsec:age_gj1108A} Due to orbital motions of the system, their true galactic motions were poorly constrained. The $\gamma$ velocity of GJ1108Aa is estimated as a few km larger than its intrinsic value. We employed an offset velocity in the RV curve as true galactic motion along the line of sight, the $\gamma$ velocity of Table \ref{tab:orbparam}. The proper motions in RA and Dec direction were determined by $Gaia$ satellite \citep{gaiadr2}. The $Gaia$ data release 2 is based on the data collected between mid 25 July 2014 (10:30 UTC) and 23 May 2016 (11:35 UTC). Using ``Observation Forecast Tool of $Gaia$\footnote{https://gaia.esac.esa.int/gost/index.jsp}'', we expected observational epochs for GJ1108A. Based on the observational epochs and the orbital solution in Table \ref{tab:orbparam}, the primary orbits their common center of mass with 118 degree in position angle. We approximated the GJ1108A's velocities along the RA and Dec directions by fitting linear functions with the orbital motions during the $Gaia$'s observation duration, finding that $Gaia$'s proper motion has the contamination of 36.2$\pm$6.0 and -0.32$\pm$0.02 mas/yr along the RA and Dec, respectively. Then, the uncertainty is an average deviancy between the expected orbital motion and the fitted linear function. We subtracted the contamination from the $Gaia$'s proper motions for GJ1108A, determining the true proper motions of GJ1108A system are -49.1$\pm$6.1 and -191.8$\pm$0.3 mas/yr in RA and Dec, respectively. The corrected proper motions of GJ1108A are more consistent with those of GJ1108B, -48.72$\pm$0.16 and -208.85$\pm$0.13 mas/yr in RA and Dec, respectively. Using the BANYAN $\Sigma$ calculator\footnote{http://www.exoplanetes.umontreal.ca/banyan/banyansigma.php} \citep{gagne18} with the updated proper motions and a $\gamma$ velocity of GJ1108A (10.1$\pm$0.2 km s$^{-1}$), we obtained the membership probabilities of GJ1108A to young groups, which are 45.6 and 54.4 $\%$ for the Columba and field region, respectively. \subsection{Dynamical mass and evolutionary models for the GJ1108A system} The combination of astrometry and RV revealed the dynamical mass of each component, $0.72_{-0.04}^{+0.04} M_{\odot}$ and $0.30_{-0.03}^{+0.03} M_{\odot}$ for the primary and the companion, respectively. Whereas the recent evolutionary model \citep[][hereafter called BHAC15]{baraffe15} suggests the model-derived masses as $0.68\pm0.01 M_{\odot}$ and $0.23\pm0.02 M_{\odot}$ based on their NIR flux and the age of Columba, in which the mass of both components are under-predicted and the older age may be preferred for GJ1108A system (Figure \ref{fig:masscomp_gj1108A}). The components in the system are here considered as coeval in the mass derivation using the evolutionary model. The right panel of Figure \ref{fig:masscomp_gj1108A} shows model-derived age distributions of the primary and the companion. The primary may be nearly on main-sequence, and the age is not determined well. We have obtained the properties of GJ1108Aa: flux, dynamical mass, and effective temperature ($T_{\rm eff}=4100_{-400}^{+200}$, See Appendix \ref{sec:app}). Using effective temperature and $H$-band flux, as similar to age estimation via mass and flux based on the evolutionary model, we estimated the isochronal age of 28$\pm$20 Myr, which is inconsistent with the model-derived age of GJ1108Ab, 69$\pm$15 Myr (Figure \ref{fig:masscomp_gj1108A}). The model may predict a higher $T_{\rm eff}$ of the primary or lower luminosity of the companion. The over-prediction of $T_{\rm eff}$ has also recognized on eclipsing binaries \citep[e.g., Figure 11 of ][]{irwin11}. If the orbital inclination of GJ1108Ab is assumed for the stellar inclination of GJ1108Aa, combining with the rotational period, a stellar radius of the primary can be estimated to be, $v_{\rm rotation} \times P_{\rm rotation}$ $\sim$ 1.2 $R_{\odot}$ ($v_{\rm rotation}$ = $v$sin$i_{\rm star}$ / sin$i_{\rm orbit}$). This estimate is inconsistent with model prediction, which is 0.7 $R_{\odot}$ for a star with the mass of 0.7 $M_{\odot}$ at 40 Myr. Even if GJ1108Aa is at 10 Myr, BHAC15 predicts $\sim$1.0 $R_{\odot}$ for the radius of the primary, which may suggest that the orbital inclination is misaligned with respect to the stellar inclination of the primary due to a perturbation such as Kozai-Lidov mechanism. Consequently, the system is indeed young as following properties: short rotational period, high UV and X-ray luminosity, kinematics, and the mass-luminosity relation determined by the orbital characterization, indicating the system is younger than 220 Myr. The little lithium abundance of GJ1108Aa should put a constraint on lower limit for their age (20Myr). In summary with current knowledge, a membership to Columba or GJ1108A is a young system in a field region are probable. Then, we here tentatively employ the age of Columba moving group, $\ageCOL$ Myr for the GJ1108A system as the age less independently determined to evolutionary models. \begin{deluxetable*}{ccc ccc ccc} \centering \tablewidth{0pt} \tablecaption{Physical parameters for the comparison of the dynamical mass and the model-derived mass} \tablehead{ \colhead{Name} & \colhead{Reference\tablenotemark{a}} & \colhead{distance} & \colhead{age} & \colhead{mass$_{\rm total}$} & \colhead{mass$_{\rm prim}$} & \colhead{mass$_{\rm comp}$} & \colhead{separation} & \colhead{tertiary} \\ \colhead{} & \colhead{} & \colhead{[pc]} & \colhead{[Myr]} & \colhead{[$M_{\odot}$]} & \colhead{[$M_{\odot}$]} & \colhead{[$M_{\odot}$]} & \colhead{[AU]} & \colhead{} } \startdata AB DorA ab & C05 & 15.18$\pm$0.13 & $\ageABD$\tablenotemark{b} & - & 0.865$\pm$0.034 & 0.090$\pm$0.005 & 3.07$\pm$0.15 & AB DorB \\ AB DorB ab & A15 & 15.18$\pm$0.13 & $\ageABD$\tablenotemark{b} & - & 0.28$\pm$0.05 & 0.25$\pm$0.05 & 0.79$\pm$0.03 & AB DorA \\ TWA 5A a+b & K07 & 48.7$\pm$5.7\tablenotemark{c1} & $\ageTWH$\tablenotemark{b} & 0.96$\pm$0.19 & - & - & 3.21$\pm$0.45 & TWA 5B \\ TWA 22 a+b & B09 & 17.5$\pm$0.2\tablenotemark{c2} & $\ageBPC$\tablenotemark{b, c2} & 0.220$\pm$0.021 & - & - & 1.77$\pm$0.04 & - \\ HD 130948 b+c & D09 & 18.17$\pm$0.11 & 790$_{-150}^{+220}$ & 0.109$\pm$0.003 & - & - & 2.20$\pm$0.11 & HD 130948 \\ HR 7672 b & C12 & 17.77$\pm$0.11 & 2400$_{-700}^{+600}$ & - & 1.08$\pm$0.04\tablenotemark{d} & 0.069$_{-0.003}^{+0.002}$ & 18.3$_{-0.5}^{+0.9}$ & - \\ Gl 417 b+c & D14 & 21.93$\pm$0.21 & 750$_{-120}^{+150}$ & 0.099$\pm$0.003 & - & - & 2.85$\pm$0.03 & Gl 417 \\ GJ 3305 ab & M15 & 29.43$\pm$0.30 & $\ageBPC$\tablenotemark{b} & 1.10$\pm$0.04 & 0.65$\pm$0.05 & 0.44$\pm$0.05 & 9.78$\pm$0.14 & 51 Eri \\ V343 Normae ab & N16 & 38.54$\pm$1.69 & $\ageBPC$\tablenotemark{b} & 1.39$\pm$0.11 & 1.10$\pm$0.10 & 0.290$\pm$0.018 & 3.07$\pm$0.08 & HD 139084B \\ HD 4747 b & C16 & 18.69$\pm$0.19 & 3300$_{-1900}^{+2300}$ & - & 0.82$\pm$0.04\tablenotemark{d} & 0.060$\pm$0.003 & 16.4$_{-3.3}^{+3.9}$ & - \\ GJ 1108A ab & - & 24.83$_{-0.22}^{+0.22}$ & $\ageCOL$\tablenotemark{b} & - & $0.72_{ -0.04}^{+ 0.04}$ & $0.30_{ -0.03}^{+ 0.03}$ & $4.11_{ -0.05}^{+ 0.05}$ & GJ 1108B \enddata \label{tab:lite} \end{deluxetable*} \begin{deluxetable*}{ccc ccc c} \addtocounter{table}{-1} \tablecaption{Continued} \centering \tablewidth{0pt} \tablehead{ \colhead{} & \colhead{} & \colhead{} & \colhead{} & \colhead{} & \colhead{} & \colhead{} } \startdata Name & m$_{J, {\rm prim}}$ & m$_{H, {\rm prim}}$ & m$_{K, {\rm prim}}$ & m$_{J, {\rm comp}}$ & m$_{H, {\rm comp}}$ & m$_{K, {\rm comp}}$ \\ \\ \tableline AB DorA ab & 5.32$\pm$0.02\tablenotemark{e} & 4.85$\pm$0.03\tablenotemark{e} & 4.69$\pm$0.02\tablenotemark{e} & 10.76$_{-0.24}^{+0.19}$ & 10.04$_{-0.15}^{+0.13}$ & 9.45$_{-0.15}^{+0.12}$ \\ AB DorB ab & 8.17$\pm$0.02\tablenotemark{e} & 8.29$\pm$0.04\tablenotemark{f1} & 7.97$\pm$0.03\tablenotemark{f1} & - & 8.55$\pm$0.04\tablenotemark{f1} & 8.23$\pm$0.03\tablenotemark{f1} \\ TWA 5A a+b & 8.40$\pm$0.07 & 7.69$\pm$0.04 & 7.39$\pm$0.04 & 8.45$\pm$0.15 & 7.79$\pm$0.05 & 7.62$\pm$0.08 \\ TWA 22 a+b & 9.12$\pm$0.10 & 8.61$\pm$0.15 & 8.24$\pm$0.19 & 9.52$\pm$0.11 & 9.12$\pm$0.15 & 8.70$\pm$0.25 \\ HD 130948 b+c & 13.81$\pm$0.06 & 13.04$\pm$0.10 & 12.26$\pm$0.03 & 14.12$\pm$0.06 & 13.33$\pm$0.11 & 12.46$\pm$0.03 \\ HR 7672 b & 4.69\tablenotemark{e} & 4.43\tablenotemark{e} & 4.39$\pm$0.03\tablenotemark{e} & 14.39$\pm$0.20\tablenotemark{f2} & 14.04$\pm$0.14\tablenotemark{f2} & 13.04$\pm$0.10\tablenotemark{f2} \\ Gl 417 b+c & 15.05$\pm$0.04 & 14.19$\pm$0.05 & 13.29$\pm$0.03 & 15.49$\pm$0.04 & 14.45$\pm$0.06 & 13.63$\pm$0.03 \\ GJ 3305 ab & 7.67$\pm$0.02 & 7.01$\pm$0.05 & 6.80$\pm$0.02 & 8.64$\pm$0.02 & 8.00$\pm$0.05 & 7.73$\pm$0.02 \\ V343 Normae ab & 6.44$\pm$0.12 & 6.05$\pm$0.10 & 5.93$\pm$0.11 & 9.69$\pm$0.12 & 9.20$\pm$0.10 & 8.79$\pm$0.11 \\ HD 4747 b & 5.81$\pm$0.02\tablenotemark{e} & 5.43$\pm$0.05\tablenotemark{e} & 5.31$\pm$0.03\tablenotemark{e} & - & - & 14.36$\pm$0.14 \\ GJ 1108A ab & 7.37$\pm$0.02 & 6.74$\pm$0.02 & 6.55$\pm$0.03 & 9.34$\pm$0.05 & 8.74$\pm$0.04 & 8.55$\pm$0.03 \enddata \tablenotetext{a}{Reference papers to determine the dynamical mass are shown, in which many of the parameters including distance, age, and the luminosities are presented unless otherwise noted. Abbreviations indicate: C05 for \cite{close05}, A15 for \cite{azulay15}, K07 for \cite{konopacky07a}, B09 for \cite{bonnefoy09}, D09 for \cite{dupuy09b}, C12 for \cite{crepp12b}, D14 for \cite{dupuy14}, M15 for \cite{montet15}, C16 for \cite{crepp16}, and N16 for \cite{nielsen16}.} \tablenotetext{b}{The age of a moving group is employed, $\ageABD$ Myr for AB Doradus, $\ageTWH$ Myr for TW Hydrae, $\ageBPC$ Myr for $\beta$ Pictoris, and $\ageCOL$ Myr for Columba \citep{bell15}.} \tablenotetext{c1}{The distance of stars in the TW Hydrae association is referred from \cite{ducourant14}.} \tablenotetext{c2}{The distance and membership of TWA 22 are referred from \cite{teixeira09}.} \tablenotetext{d}{Non-dynamical mass determined by empirical tracks, or theoretically determined using spectroscopic results using tools such as SME \citep{valentipiskunov96, valentifischer05}.} \tablenotetext{e}{Unresolved 2MASS photometry for those systems \citep{skrutskie06}.} \tablenotetext{f1}{Photometric measurements are obtained from \cite{janson07a}.} \tablenotetext{f2}{Photometric measurements are obtained from \cite{boccaletti03}.} \end{deluxetable*} \begin{figure*} \centering \begin{tabular}{ccc} \begin{minipage}[htbp]{0.33\hsize} \includegraphics[width=\hsize]{\dirGJ1108A/resolved_binaries_j.ps} \end{minipage} \begin{minipage}[htbp]{0.33\hsize} \includegraphics[width=\hsize]{\dirGJ1108A/resolved_binaries_h.ps} \end{minipage} \begin{minipage}[htbp]{0.33\hsize} \includegraphics[width=\hsize]{\dirGJ1108A/resolved_binaries_k.ps} \end{minipage} \end{tabular} \caption{The dynamical mass and model-derived mass based on BHAC15 model are compared on the figures. Data points are obtained from \cite{close05, konopacky07a, bonnefoy09, dupuy09b, crepp12b, dupuy14, azulay15, montet15, crepp16, nielsen16}, and classified according to their age as symbols with different colors. To estimate the model-derived mass, $J$, $H$, and $K$-band photometry are used in each panel (a), (b), and (c), respectively. The black solid and dashed lines on those show a mean offset and scatter, and the grey lines are for only resolved PMS stars.} \label{fig:masscomp_b15} \end{figure*} \subsection{What causes the mass discrepancy ?} \subsubsection{Accretion history} Several studies have been indicated the importance of accretion histories before the quasi-static contraction phase \citep[e.g.][]{baraffe09, baraffe12}. Due to the different types of accretion mechanisms, conditions of young low-mass stars can become much different from those considered in classical models. One of the most important subjects regarding accretion histories is the balance between expansion and contraction for a protostar, which strongly depends on details of accretion processes. However the details are still remained very uncertain. To quantitatively investigate the accretion effect: the energy loss during the accretion onto protostar's surface, a free parameter $\alpha$ have been employed in \cite{hartmann97}. The increase of internal energy for protostar and the radiation at its surface are respectively written as, \begin{eqnarray} E_{\rm add} & = & \alpha \epsilon \frac{GM}{R} \dot{M} \\ E_{\rm rad} & = & (1-\alpha) \epsilon \frac{GM}{R} \dot{M} \end{eqnarray} where $G, \ M, \ R$, and $\dot{M}$ indicate gravitational constant, mass of protostar, radius of protostar, and mass of accreted matter. The $\epsilon$ presents the fraction of internal energy retained in an accretion disk, up to 0.5. In case of $E_{\rm add}$ is dominant, a protostar will expand due to the accreted energy and become brighter than non-accreted objects in the models, and hence such accretion is called ``hot accretion''. In the opposite case, a protostar becomes smaller and fainter due to ``cold accretion''. The threshold of $\alpha$ is obtained by the energy equation \citep[Eq.6 of ][]{hartmann97}, and $\alpha$ of 0.1--0.2 have been assumed \citep{hartmann97, baraffe09}. If we assume that GJ 1108A is a member of the Columba group, cold accretion is preferred in its accretion history to explain the mass under-prediction. we assume that GJ1108A is a member of the Columba group, cold accretion is preferred in its accretion history to explain the mass under-prediction. However young stars may forget the accretion history at the age of GJ1108A, 40--50 Myr \cite[Figure 1 of ][]{baraffe10}, and the reason why such significant events had occurred in the system is a matter to be further investigated. \subsubsection{Magnetic activity} Magnetic activities are also essential to understand the stellar structure and their observable properties. Similar discrepancy to evolutionary models has been recognized on eclipsing binaries in mass-radius relation \citep[e.g.][]{irwin11}, and theoretical investigations suggest that models considering magnetic activities may explain the discrepancy \citep{feiden13a}. The magnetic field works as inhibition of stellar convection, making spots at stellar surface and making stellar radius inflated compared with classical models. In other words, active low-mass stars tend to be observed to classical models as larger radius for given $T_{\rm eff}$ or lower $T_{\rm eff}$ for given radius \citep[e.g.][]{mullanmacdonald01}. \cite{feiden13b, feiden14} demonstrate that stellar evolutionary code can explain the discrepancy by involving magnetic activity for both stars with a radiative core and fully convective stars. This approach also better reproduces the observed color-magnitude diagram of young stellar association with single isochrone \citep{feiden16}. GJ1108Aa is a rapid rotating and indeed X-ray luminous young system. If we assume the membership to Columba for GJ1108A, BHAC15 model not including the magnetic activity predicts its flux brighter than observed flux. The observed $T_{\rm eff}=4100_{-400}^{+200}$ is marginally consistent with the model prediction (4084 K) for 0.7 $M_{\odot}$ at 40 Myr. To explain the observed deviation, the primary must be smaller than the classical model prediction, which is opposite to the behavior of magnetic activity, ``a larger radius for a given $T_{\rm eff}$''. Hence, a scenario that GJ1108A is older than the age of Columba is more preferred than the non-included effects to classical models such as magnetic inhibition of convection or accretion history (See the previous section). \begin{deluxetable*}{ccc ccc cc} \centering \tablewidth{0pt} \tablecaption{Results of the mass comparison} \tablehead{ \colhead{Model set} & \colhead{Mean offset\tablenotemark{a}} & \colhead{Scatter\tablenotemark{a}} & \colhead{Method} & \colhead{Mass range} & $N_{\rm sample}$ & Tertiary rate & \colhead{Reference} \\ \colhead{} & \colhead{} & \colhead{} & \colhead{} & \colhead{$M_{\odot}$} & \colhead{} & \colhead{} & \colhead{} } \startdata BHAC15 (All,$J$) & 0.07$\pm$0.04 & 0.19$\pm$0.05 & age and $L_{J}$ & 0.05--0.7 & 11 & 6/9 & this work \\ BHAC15 (All,$H$) & 0.03$\pm$0.04 & 0.18$\pm$0.05 & age and $L_{H}$ & 0.05--0.7 & 12 & 6/9 & this work \\ BHAC15 (All,$K$) & 0.08$\pm$0.04 & 0.21$\pm$0.04 & age and $L_{K}$ & 0.05--0.7 & 13 & 7/10& this work \\ BHAC15 (PMS,$J$) & -0.05$\pm$0.04 & 0.17$\pm$0.04 & age and $L_{J}$ & 0.09--0.65 & 6 & 4/4 & this work \\ BHAC15 (PMS,$H$) & -0.05$\pm$0.04 & 0.18$\pm$0.04 & age and $L_{H}$ & 0.09--0.65 & 8 & 5/5 & this work \\ BHAC15 (PMS.$K$) & 0.08$\pm$0.04 & 0.21$\pm$0.04 & age and $L_{K}$ & 0.09--0.65 & 8 & 5/5 & this work \\ BCAH98\tablenotemark{b} (All,$J$) & 0.01$\pm$0.06 & 0.17$\pm$0.05 \\ BCAH98\tablenotemark{b} (All,$H$) & 0.02$\pm$0.06 & 0.17$\pm$0.04 \\ BCAH98\tablenotemark{b} (All,$K$) & 0.03$\pm$0.04 & 0.21$\pm$0.03 \\ \tableline BCAH98\tablenotemark{b} & 0.342$\pm$0.471 & 0.439 & $T_{\rm eff}$ and $L$ & 0.03--1.4 & 16 & 5/8 & S14\tablenotemark{d} \\ BCAH98\tablenotemark{c} & 0.212$\pm$0.353 & 0.308 & \enddata \tablenotetext{a}{The mean offset and scatter in the mass comparison are determined by a typical value of $(M_{\rm model}-M_{\rm dynamical})/M_{\rm dynamical}$.} \tablenotetext{b}{[M/H]=0, $l_{\rm mix}=1.0H_{p}$, and Y=0.275} \tablenotetext{c}{[M/H]=0, $l_{\rm mix}=1.9H_{p}$, and Y=0.282} \tablenotetext{d}{\cite{stassun14}} \label{tab:offset} \end{deluxetable*} \begin{figure*} \centering \includegraphics[width=\hsize]{\dirGJ1108A/s14vsm18.ps} \caption{The number of object focused on properties: mass, separation, and age of systems are shown. Solid and dashed histograms indicate resolved PMS stars compiled in this work and all the low-mass systems of \cite{stassun14}, respectively.} \label{fig:samples} \end{figure*} \subsection{Mass comparison using compiled samples} \subsubsection{Sample selection} In order to verify the performance of present evolutionary models, especially the luminosity evolution observed in NIR wavelength as a function of time, we compare the dynamical masses of stars with the model-derived masses of those (Figure \ref{fig:masscomp_b15}), by collecting information from the literature. Physical parameters of the literature-based samples such as age, dynamical mass, and broad-band fluxes, are summarized in Table \ref{tab:lite}. We set three criteria for samples in the mass comparison. As the first criterion, we selected objects whose mass are dynamically estimated. The mass of each component has not separately measured for some of our sample; their orbital RV modulations have not been well determined. We then simply employ the total masses of two components in the systems, which were determined by astrometry, in our mass comparisons. Second, we select the objects that have been age-dated. At this step, we do not adopt age dating based on isochrone for a single stellar system. If a sample is associated with any stellar cluster or moving group, the sample' age can be represented by a well-determined isochronal age of the coeval population in groups; our sample therefore includes the objects that are associated to stellar clusters or moving groups. Third, we here focus on PMS stars, since the tests of evolutionary models for the PMS phase should be more incomplete than the MS phase. Then, we do not include the members of star-forming regions, for which the stellar properties such as distance, brightness, and age are relatively poorly determined compared with post T-tauri stars in the nearby stellar groups. Age-dated brown dwarfs\footnote{As the age of brown dwarfs in our mass comparison (Table \ref{tab:lite}), we employed the age of stellar primary (tertiary). } are also included in our sample, because this study is motivated for better understanding brown dwarfs and giant planets based on evolutionary models. It should be noted that the mass comparison is conducted for both only PMS stars and all the sample including PMS stars and brown dwarfs. \subsubsection{Dynamical mass vs Model-derived mass} To estimate the masses of the samples, we need to select the fiducial sets of evolutionary models. For low-mass objects, non-grey atmosphere is crucial to properly reproduce the stellar structure as already discussed in \cite{chabrierbaraffe97}. We provide the following two criteria; the models (1) cover a wide mass range from brown dwarfs to stars and (2) adopt non-grey atmosphere. \cite{stassun14} reviewed several theoretical isochrones made by different groups, from which we selected the latest model of the Lyon model \citep{baraffe98, baraffe15}. The dynamical masses and model-derived masses are compared, and we evaluated an offset and scatter between the two masses. We ran a Monte-Carlo simulation in order to derive the uncertainties of the offset and scatter value. This simulation randomly generated masses of all the samples based on their errors, to make a distribution of those values and calculates the typical values from the distribution. Furthermore, the number of resolved young binary is still small, which may make the mass comparison unreliable. To verify this problem, we randomly omit 10--50$\%$ of samples and carried out the same comparison as described above. The offsets and scatters with those partially-selected samples are the almost same as what was derived using all the samples. The mass comparison of Figure \ref{fig:masscomp_b15} respectively results with small offset and scatter, $\le10\%$ and $\sim20\%$, indicating the luminosity evolution for low-mass objects are well reproduced on average by the evolutionary model of \cite{baraffe15}. Even if we exclude brown dwarfs in the mass comparison, similar results could be obtained. The old version of the model \cite[][hereafter BCAH98]{baraffe98} is also adopted for the mass comparison. Although BCAH98 has three model grids with different mixing length parameter, $l/H_{p}$=1.0, 1.5, and 1.9, the latter two grids do not have the mass range for brown dwarfs, and we used only the grid of $l/H_{p}$=1.0 in this work. We used the same approach as for BHAC15, providing a similar result between these two models. These are summarized in Table \ref{tab:offset}. \subsection{Comparison with previous study} \cite{stassun14} investigated the performance of stellar evolutionary models including BCAH98 using literature-based PMS eclipsing binaries whose isochronal age is spanned in 1--20 Myr. Using luminosity and effective temperature, they estimated the model-derived mass for 16 objects of 0.03--1.4$M_{\odot}$ with the BCAH98 model, and then conducted the mass comparison. They found the model overestimate the dynamical mass by 34.2$\pm$47.1 and 21.2$\pm$35.3$\%$ that are dependent on the mixing length parameter. For the mass discrepancy, they considered the mainly three possibilities: activities, initial events in star formation, and tertiary. Whereas significant improvements were not obtained due to corrections of the magnetic activity and accretion history for evolutionary models, the discrepancy might be seen for systems possessing a tertiary. Tertiaries may occur the non-negligible deviations in the mass comparison as they suggested. On the other hand, the evolutionary model has better performance for our resolved binaries. Figure \ref{fig:samples} compares properties of eclipsing binaries of \cite{stassun14} and resolved binaries in this work, focusing on mass, separation, and age. The resolved binaries in this work typically are older and have the larger separation, which may attenuate the the discrepancy of mode-derived masses from dynamical masses. | 18 | 8 | 1808.04996 |
1808 | 1808.01043_arXiv.txt | The growing evidence for energy-conserving outflows in powerful and luminous AGN supports the idea that high-velocity winds launched from the accretion disc evolve systematically after undergoing a shock with the ambient medium and that they are capable to expel enough mass and energy so as to produce feedback. This talk will give an overview of recent results on AGN ultra fast outflows, with focus on grating X-ray spectra of bright sources. I will review how UFO work, their observational properties and their relation with AGN outflows in other bands, what is their impact on the host galaxies and their role in feedback processes. | Feedback from Active Galactic Nuclei (AGN) is generally thought to be an important ingredient for galaxies evolution. After large amount of gas is accreted during the earliest stage of a quasar life time, the accumulated energy can be released via ejection of powerful outflows driven by the AGN. If the outflow is as strong as 0.5--5\% of the Eddington luminosity of the AGN, it has a profound impact on the development of the host galaxy itself. The effect of these winds is to eventually expel the gas that would otherwise be available for forming new stars in the host galaxy therefore providing an effective mechanism of quenching star formation. It is in this sense that we refer to AGN feedback as a mechanism able to regulate the growth of the galaxy and the growth of the central black hole as well (Di Matteo et al. 2005, Hopkins et al. 2010). A widely accepted scenario for explaining AGN feedback postulates that a fast wind observable in the X-ray band is launched at accretion disk scale (Faucher-Gigu{\`e}re \& Quataert 2012). This highly ionized X-ray gas is currently observed in the form of Ultra Fast Outflows in some AGN spectra (Tombesi et al. 2012, Gofford et al. 2013). While traveling outward, the impact of the wind with the ISM (inter-stellar medium) gives rise to shock processes (King 2010). After shocking with the gas, deceleration and cooling processes lead to the production of a slower outflow with less ionized lines observable in the optical band and to the formation of a bubble of hot, tenuous gas (e.g. Zubovas \& King 2012). As a result of the cooling, the presence of molecular gas outflowing at a much lower velocity is expected. This latest phase is frequently observed in several ULIRGS and Quasars (Cicone et al. 2014, Feruglio et al. 2010, 2015). To date, only two cases of ULIRGS are reported where the observed X-ray and molecular phases of the outflow are physically related, IRAS~F11119+3257 (Tombesi et al. 2015) and Mrk~231 (Feruglio et al. 2015). Both results remarkably fit in with the prediction of the energy-conserving outflow model outlined above. \\ \begin{figure}[b] \begin{center} \includegraphics[width=12cm]{fig1.eps} \caption{{\it XMM-Newton}-RGS spectrum of the NLSy1 Galaxy Mrk~1044 fitted by four components of ultra fast outflows (Krongold et al. submitted). The fastest component (blue labels) of this system reaches an outflow velocity of $\sim$48,000~km~s$^{-1}$ and column density of N$_H$$\sim$10$^{23.32}$~cm$^{-2}$, while the other three are outflowing at 25,000~km~s$^{-1}$. } \label{fig_1044} \end{center} \end{figure} This fascinating picture, among several other variables, relies on the existence and the properties of the nuclear wind, the so-called ``Ultra Fast Outflow". The X-ray spectra of some AGNs show signature of gas outflowing at high speed ($v \ge 0.1$~c). This gas is so highly ionized by the nuclear radiation that the only dominant ions left are He-like and Hydrogen-like ionic species. These systems were christened as ``Ultra-Fast Outflows" (UFOs) (Tombesi et al. 2010) and at the beginning they were observed mainly in the Fe K band at E$\ge$~7~keV. Several papers reported on UFOs hosted in individual AGNs (Pounds et al. 2003, 2011, 2014, Chartas et al. 2009, Lanzuisi et al. 2012, Nardini et al. 2015), and statistical studies show that they are detected in 30-40\% of nearby AGN (Tombesi et al. 2010, 2012, Gofford et al. 2013). The approximate ranges of mass outflow rate (0.01-1M$_{\odot}$~yr$^{-1}$) and kinetic energy (10$^{42-45}$ erg~s$^{-1}$) are in good agreement with theoretical predictions for black hole winds (King 2010). Two launching mechanisms are envisaged for the production of such rapid outflows, and both locate the region of action in the accretion disk of the AGN. One way to produce a fast outflow is through radiation driving (Proga and Kallman 2004, Sim et al. 2010), and another one is via magnetic fields in radio-loud sources (``magnetically driven outflows", Fukumura et al. 2015). | 18 | 8 | 1808.01043 |
|
1808 | 1808.04516_arXiv.txt | We provide the first in situ measurements of antenna element (tile) beam shapes of the Murchison Widefield Array (MWA), a low radio-frequency interferometer and an SKA\thanks{\url{https://www.skatelescope.org/project/}} precursor. Most current MWA processing pipelines use an assumed beam shape, errors in which can cause absolute and relative flux density errors, as well as polarisation `leakage'. This makes understanding the primary beam of paramount importance, especially for sensitive experiments such as a measurement of the 21$\,$cm line from the epoch of reionisation (EoR). The calibration requirements for measuring the EoR 21$\,$cm line are so extreme that tile to tile beam variations may affect our ability to make a detection. Measuring the primary beam shape from visibilities alone is challenging, as multiple instrumental, atmospheric, and astrophysical factors contribute to uncertainties in the data. Building on the methods of~\citet{Neben2015}, we tap directly into the receiving elements of the MWA before any digitisation or correlation of the signal. Using ORBCOMM satellite passes we are able to produce all-sky maps for 4 separate tiles in the XX polarisation. We find good agreement with the cutting-edge `fully' embedded element (FEE) model of~\citet{Sokolowski2017}, and observe that the MWA beamformers consistently recreate beam shapes to within~$\sim1$dB in the reliable areas of our beam maps. We also clearly observe the effects of a missing dipole from a tile in one of our beam maps, and show that the FEE model is able to reproduce this modified beam shape. We end by motivating and outlining additional onsite experiments to further constrain the primary beam behaviour. | \label{sec:intro} The Murchison Widefield Array (MWA) is a low radio-frequency interferometer located in the western Australian outback, at the Murchison Radio-astronomy Observatory (MRO). One of the observational strengths of the MWA is its substantial field of view. Each receiving element (tile) in the interferometer consists of a $4\times4$ grid of bow-tie dipoles mounted on a $5\,\mathrm{m} \times 5\,\mathrm{m}$ reflective ground screen. Analogue beamformers are used to create and electronically steer the MWA primary beam, which has a full-width half-maximum of $\sim 25^\circ$ at 150$\,$MHz~\citep{Tingay2013}. The regular spacing of the dipoles means that the quantised beamformer delays are exactly correct for a set of pointings, reducing the complexity of the instrument. By using identical receiving elements, a number of computational simplifications can be made to calibration and imaging, as beam corrections can be made in image space, avoiding costly convolutions in visibility space. Many calibration schemes make assumptions on beams/receiving elements~\citep[e.g.][]{Kazemi2013,Tasse2013}, and others explicitly use the beam shape during calibration/imaging~\citep[][specifically for the MWA]{Mitchell2008,Sullivan2012}. However, for these assumptions to be valid, each tile must actually produce identical beam shapes, and that beam shape must be correctly modelled. The exact degree of precision required of a primary beam and model depends on the science case, but epoch of reionisation (EoR) science is perhaps in the greatest need of well understood beam shapes. Particularly, the spectral behaviour of the beam is a significant possible contaminant of a detection~\citep{Barry2016,Trott2016b}. \citet{Barry2016} state that the frequency response of an EoR experiment should have spectral features no larger than $10^{-5}$. This kind of spectral behaviour could be injected via calibration using an incorrect beam model. Estimating the degree of precision of a beam model required for EoR science is an ongoing area of research. The effects of incorrect beam modelling on output data products can be subtle, and will change according to the calibration scheme used. To fully estimate the effects on an MWA EoR power spectrum experiment\citep[e.g.][]{Jacobs2016}, one has to take the following points into consideration: \begin{itemize} \item The primary beam is often used as a gridding kernel. An incorrect model would bias the gridded data used to either image or create a power spectrum \item Errors on the beam model can be direction dependent. As the MWA primary beam is stationary during a 2 minute observation, primary calibrators will move through these errors during the observation. The MWA EoR observing strategy often employs the same pointing for $\sim30\,$minutes, potentially injecting time-dependent calibration errors \item If the primary beam varies from tile to tile, the spatial scales that are measured by each baseline are affected differently. The effect of this upon a measured power spectrum is hard to estimate. \end{itemize} We plan to investigate the points above in future work, using simulated observations and power spectrum measurements to directly quantify the effects on potential EoR science. As the accuracy of the beam model can affect science outputs, significant work has gone into both simulating and measuring the MWA primary beam. The MWA is a full Stokes instrument, and as such the beam model must accurately describe the instrument in all polarisations. Of particular interest is `leakage' from Stokes I to other polarisations, whereby flux density from Stokes I is transferred into other Stokes parameters. This effect is often elevation dependent, and can be attributed to an incorrect beam model~\citep{Lenc2016,Lenc2017}. The sophistication of the MWA's beam model has progressed through a number of stages. The first model was a simple analytical short-dipole radiation pattern multiplied by the tile array factor~\citep{Ord2010}. The tile array factor is a direction dependent function that describes the cumulative beam shape effects of the superposition of the individual dipoles, in this case generated purely through the geometrical $4\times4$ layout of the tile. The first advanced model was the `average' embedded element (AEE) model, presented by~\citet{Sutinjo2015a}. This model used numerical simulations generated using the commercially available FEKO\footnote{\url{https://www.feko.info/}} simulation package, including mutual coupling effects, to create an average dipole radiation pattern, which was assumed identical for all dipoles. In addition, the model includes mutual-coupling induced changes to dipole impedance, which affects the array factor depending on pointing direction. Whilst this model was shown to reduce leakage,~\citet{Sutinjo2015a} showed that leakage still remained, and suggested that a `fully' embedded element (FEE) model was required. This was verified as the AEE model was used in calibration and image-based primary beam correction on GaLactic and Extragalactic All-sky MWA survey~\citep[GLEAM,][]{Wayth2015} data. Frequency/elevation dependent errors in the model produced polarisation leakage and incorrect absolute flux density of sources in the images~\citep[sometimes $>10$\%; see][for details]{Hurley-Walker2017}. Given catalogues such as GLEAM are often used to calibrate observations, a correct flux scale is of paramount importance. As a response, the FEE model was presented in~\citet{Sokolowski2017}. This model more rigorously takes into account mutual coupling, amongst other improvements. \citet{Sokolowski2017} were able to use GLEAM observations to show reduced leakage compared to the AEE model. Using polarisation leakage to test beam performance is a valuable metric, but is a measure of average beam effects, rather than the beam pattern itself. To validate a beam model, one must directly measure the actual beam shape of individual tiles. The simplest way to do this is to have a radio emitter of known location and strength, and use the instrument to measure the flux density seen from the emitter from various locations. Recent advances have been made in using commercially available drones to fly a radio emitter to achieve this~\citep[e.g.][]{Ustuner2014,Chang2015,Picar2015,Pupillo2015,Paonessa2016}. These experiments have the advantage of being able to run at multiple frequencies and along a controlled flight path. The disadvantage is that these experiments happen in the near-field of the instrument; observations occur in the far-field. Even with this limitation, once the stability of the drones and repeatability of results improves, this approach could provide accurate beam measurements in the future. Another option lies in using transiting satellites, bright in the MWA frequency band, as locatable emitters. \citet{Neben2015} have successfully used ORBCOMM satellites\footnote{\url{https://www.orbcomm.com/en/networks/satellite}} to measure an MWA tile beam shape in an experimental setup at the National Radio Astronomy Observatory in Green Bank, West Virginia. The constellation of ORBCOMM satellites operate at 137$\,$MHz, which does limit any experiment to a single frequency band, but the sky coverage and intrinsic brightness of the satellites allows a beam map to be made across the entire MWA primary beam. Satellites have the advantage of naturally being in the far-field of the instrument. \citet{Neben2016} were able to apply this method to the Hydrogen Epoch of Reionisation Array\footnote{HERA - \url{http://reionization.org/}}, to both measure beam shapes, and comment on the implications to their science goals. The success of Neben et al. has motivated a comparison of ORBCOMM and drone-based beam mapping techniques~\citep{Jacobs2017}, aimed towards being able to conduct reliable onsite experiments. To date, no one has actually directly measured the individual MWA tile beam shapes onsite at the MRO. Given the discussion above, we suggest the ideal MWA beam measurement experiment should: \begin{itemize} \item occur onsite at the MRO using actual MWA tiles \item cover the whole observable sky, in the far-field \item measure multiple tiles simultaneously to establish similarities/differences \item span the entire frequency range covered by the MWA \item measure both XX and YY polarisations so Stokes parameters can be calculated \item cover as many pointings as possible to fully explore predictive powers of any model \item compare any measured beams to the FEE model. \end{itemize} For a single experiment, this is a daunting list. In this paper, we make a first step towards onsite beam measurements, and choose to limit ourselves to checking for variance in tile to tile beam patterns. The goals of this paper are: make an initial assessment of the accuracy of the new FEE model in describing the primary beam shape generated by onsite MWA tiles; quantitatively measure how similar primary beam shapes are from tile to tile; use the results to inform follow-up experiments covering the items on the list above. The paper is organised as follows. In Section~\ref{sec:exsetup}, we describe our experimental setup and data collection method. In Section~\ref{sec:analysis} we detail our data analysis, and present our results, along with some experimental biases, in Section~\ref{sec:results}. We discuss our results and outline a possible future direction in Section~\ref{sec:discussion}. | We have measured the XX polarisation primary beam shape for 4 onsite MWA tiles. The coverage of the ORBCOMM constellation above the MRO allows full sky coverage down to an elevation of $10^\circ$, if one can continually observe with a single pointing over a period of two weeks. The maps have a dynamic range of $>30\,$dB, and show good agreement with the cutting-edge FEE model of \citet{Sokolowski2017} in the east-west direction. Geographical and hardware limitations reduce that agreement, and the reliability of the maps, in the north-south direction. We observe that the FEE model is able to accurately describe the zenith beam pointing of an MWA tile, even when one dipole is missing. We also demonstrate the strong dependence of the primary beam shape on the number of contributing dipoles. This is important since the MRO is an inhospitable environment and it is common for dipoles to fail. We have developed a method for collecting raw RF data from the tiles using relatively inexpensive and commercially available RF Explorers and Raspberry Pis. Most importantly, not only can all of this equipment easily fit within a receiver box, it can be remotely controlled and scheduled. This will allow us to easily expand this method to measure both instrumental polarisations and multiple beam pointings in future work, enabling a better understanding of instrumental performance, which will benefit all scientific programs of the MWA. | 18 | 8 | 1808.04516 |
1808 | 1808.06808_arXiv.txt | We construct parameter sets of the relativistic mean-field model fitted to the recent constraints on the asymmetry energy $J$ and the slope parameter $L$ for pure neutron matter. We find cases of unphysical behaviour, i.e.\ the appearance of negative pressures, for stiff parameter sets with low values of the effective mass $m^*/m$. In some cases the equation of state of pure neutron matter turns out to be outside the allowed band given by chiral effective field theory. The mass-radius relations of neutron stars for all acceptable parameter sets shows a maximum mass in excess of $2M_\odot$ being compatible with pulsar mass measurements. Given the constraints on the model in the low-density regime coming from chiral effective theory, we find that the radius of a $1.4M_\odot$ neutron star is nearly independent on the value of $L$. This is in contrast to some previous claims for a strong connection of the slope parameter with the radius of a neutron star. In fact, the mass-radius relation turns out to depend only on the isoscalar parameters of symmetric matter. The constraints of GW170817 on the tidal deformability and on the radius are also discussed. | Neutron stars have received a lot of attention over the years \cite{Lattimer:2004pg,Lattimer:2006xb,Oertel:2016bki}, especially since the detection of gravitational waves from the neutron star binary merger GW170817 \cite{TheLIGOScientific:2017qsa}. The two most important properties of neutron stars, maximum mass and radius, are still a matter of intense analysis since they are linked to the physics of their interior, which is nowadays an open question. It is now established that neutron stars, usually observed as pulsars, can have masses of up to $2M_{\odot}$ \cite{Demorest:2010bx,Antoniadis:2013pzd,Fonseca:2016tux}. The precise determination of neutron star radii is still an ongoing process. Model dependent constraints on the radius have been derived by fits to low mass quiescent X-ray binary data and thermonuclear bursts, sometimes with conflicting results \cite{Verbiest:2008gy,Ozel:2010fw,Suleimanov:2010th,Lattimer:2012xj,Steiner:2012xt, Bogdanov:2012md,Guver:2013xa,Guillot:2013wu,Lattimer:2013hma,Poutanen:2014xqa, Heinke:2014xaa,Guillot:2014lla,Ozel:2015fia,Ozel:2015gia,Ozel:2016oaf, Lattimer:2015nhk,Steiner:2017vmg}. Recent analysis of tidal deformabilities from the neutron star merger seen in gravitational waves, GW170817, however, were able to set limits on the radius of a $M=1.4M_\odot$ neutron star in the range of $R=12$--$13.5$~km based on statistical approaches \cite{TheLIGOScientific:2017qsa,Annala:2017llu,Most:2018hfd,De:2018uhw,Abbott:2018exr}. Future high-precision X-ray space missions, such as the on-going NICER (Neutron star Interior Composition ExploreR) \cite{2014SPIE.9144E..20A} and the future eXTP (enhanced X-ray Timing and Polarimetry Mission) \cite{Zhang:2016ach}, will improve the situation by simultaneous measurements of masses and radii with higher accuracy \cite{Watts:2016uzu}. Limits on neutron star radii are also expected to be refined by future detections of gravitational-wave signals from neutron star mergers. The mass and radius of neutron stars strongly depend on the properties of matter in their interior, that is described by means of the equation of state (EoS). Indeed, the determination of the EoS is a field of extensive and active research. Among the different approaches to obtain the EoS, the relativistic mean field (RMF) models \cite{Serot:1984ey,Mueller:1996pm,Serot:1997xg,Glendenning:2000,Fattoyev:2010rx,Chen:2014sca} have been widely used for describing the interior of neutron stars based on fits to nuclear ground state properties and/or on fitting the parameters of the model directly to properties of nuclei, such as masses, charge radii and surface thickness. Yet it is far from a trivial task to generate an EoS that respects the properties of nuclear matter and nuclei as well as describes pure neutron star matter. Recall that it is an extrapolation of $\sim$18 orders of magnitude from the radius of a nucleus to the radius of a neutron star \cite{Weber:1989hr,Horowitz:2000xj}. For densities $\rho \leq 4\cdot 10^ {11}~\rm{g/cm^ 3}$ neutron stars are expected to have a outer crust consisting of a lattice of neutron-rich nuclei. Up to densities of about $\rho \sim 10^ {14}~\rm{g/cm^ 3}$, the inner crust consists of a lattice of nuclei immersed within a neutron liquid. At higher densities the outer core is liquid neutron rich matter, consisting of a liquid of neutrons with a small admixture of protons and electrons. For the inner core, probably starting at twice saturation density, i.e.\ at $\rho \geq 5\cdot 10^ {14}~\rm{g/cm^ 3}$, one may speculate about exotic phases of matter, such as hyperon matter (see Ref.~\cite{Chatterjee:2015pua} for a review) or quark matter (see e.g.\ \cite{Weber:2004kj,Alford:2006vz,Weissenborn:2011qu,Buballa:2014jta,Zacchi:2015lwa}). In this paper we present relativistic parameter sets for the EoS of the interior of neutron stars that fulfill the $M\geq 2M_\odot$ neutron star mass constraint from the observations of pulsars and the radius constraint from GW170817 of $R=12$--$13.5$~km, while fulfilling the saturation properties of nuclear matter. Moreover, we impose further constraints on the EoS for neutron matter coming from chiral effective field theory ($\chi$EFT) \cite{Drischler:2016djf}. These constraints are met by simultaneously fitting the isoscalar couplings to saturation properties, while allowing for variations of the isovector parameters and the effective nucleon mass ($m^*/m$) so as to reproduce the symmetry energy ($J$) and its slope ($L$) within reasonable theoretical and experimental limits \cite{Li:2013ola,Lattimer:2012xj,Roca-Maza:2015eza,Hagen:2015yea,Oertel:2016bki,Birkhan:2016qkr}. We find that the values of the symmetry energy and its slope that allow for a physical solution for the neutron matter EoS compatible with $\chi$EFT depend on the value of the nucleon effective masses at saturation density. We also observe that the behaviour of both the maximum mass and the radius of neutron stars is dominated by the effective nucleon mass. The restricted range of $L$ values coming from $\chi$EFT constraints does not allow for an appreciable variation of the radius, while being not relevant for the determination of the maximum mass. All parameters sets result in maximum neutron star masses in excess of the $2M_\odot$ limit. Large values of the effective nucleon mass induce small neutron star radii, so that effective nucleon masses of $m^*/m > 0.60$ are needed in order to have radii compatible with the recent upper limit of the tidal deformabilities and radii from GW170817. The article is organized as follows. In Section \ref{formalism} we present the model Lagrangian and derive the corresponding equations of motion. We determine the parameters of the model in Section \ref{parameters}, while presenting our results for the EoS, mass-radius relation and dimensionless tidal deformability in Section \ref{results}. Our conclusions are summarized in Section \ref{conclusions}. The tables with the isoscalar and isovector parameters of the model can be found in the Appendix \ref{app}. | \label{conclusions} We have studied the EoS in the inner core of neutron stars within the relativistic mean field theory with the aim of fulfilling several recent astrophysical constraints, such as, the $2M_\odot$ neutron star mass limit \cite{Demorest:2010bx,Antoniadis:2013pzd,Fonseca:2016tux} and the extraction of neutron star radii $\lesssim$ 13.5~km from the recent analysis on tidal deformabilities of the GW170817 neutron star merger event \cite{Abbott:2018exr,Most:2018hfd,Annala:2017llu,De:2018uhw,Kumar:2017wqp,Fattoyev:2017jql,Malik:2018zcf}. The phenomenological model satisfies the saturation properties of nuclear matter together with constraints on low-density neutron matter coming from $\chi$EFT ab-initio approaches \cite{Drischler:2016djf}. These constraints are fulfilled by simultaneously fitting the isoscalar couplings to saturation properties (saturation density, energy per particle and compressibility), while allowing for variations in the isovector parameters so as to reproduce the symmetry energy and its slope within reasonable theoretical and experimental limits \cite{Li:2013ola,Lattimer:2012xj,Roca-Maza:2015eza,Hagen:2015yea,Oertel:2016bki,Birkhan:2016qkr}. We have found that the values of the symmetry energy ($30 \leq J \,\, [{\rm MeV}] \leq 32$) and its slope ($40 \leq L \,\, [{\rm MeV}] \leq 60$) that allow for a physical solution for the neutron matter EoS compatible with the $\chi$EFT are determined by the value of the nucleon effective mass ($0.55 \leq m^*/m \leq 0.75$). It is indeed difficult to find a physical solution compatible with the $\chi$EFT results once the values for $m^*/m$ and $L$ are reduced (increased) simultaneously for a fixed value of $J$. A softening (hardening) of the EoS induced by a small (big) value of $L$ competes with the stiffening (softening) of the EoS as we reduce (increase) the effective nucleon mass, leading to either a solution outside the allowed area from $\chi$EFT or the appearance of unstable solutions below saturation density. With regards to the mass and radius of neutron stars, we have obtained that the effective nucleon mass turns out to be the dominant parameter controling both, the maximum mass and the radius of a neutron star. This is due to the fact that, on the one hand, the restricted range of $L$ values coming from $\chi$EFT constraints does not allow for noticeable variations on the radius and, on the other hand, the isovector parameters turn out to be not relevant for the determination of the maximum mass, as seen in Ref.~\cite{Tolos:2016hhl,Tolos:2017lgv}. Large values of $m^*/m$ induce small masses and radii, as expected from the Hugenholtz-van-Hove theorem. Thus, effective nucleon masses of $m^*/m > 0.6$ are needed in order to reproduce 2$M_{\odot}$ observations and have radii compatible with recent astrophysical determinations \cite{Guillot:2013wu,Guillot:2014lla,Guver:2013xa,Heinke:2014xaa,Lattimer:2014sga,Lattimer:2013hma,Nattila:2015jra,Ozel:2015fia,Ozel:2016oaf,Lattimer:2015nhk}, that are corroborated by the analysis on tidal deformabilities of the GW170817 event \cite{Abbott:2018exr,Most:2018hfd,Annala:2017llu,De:2018uhw,Kumar:2017wqp,Fattoyev:2017jql,Malik:2018zcf}. In fact, our values of the dimensionless tidal deformability are within the 90$\%$ confidence level for $m^*/m > 0.65$ at $J=32$ and $L=60$. Note, however, that the effective nucleon mass has to be reconciled with the binding energies and charge radius of atomic nuclei \cite{Reinhard:1989zi}. In the near future, apart from the expected detection of gravitational-wave events from other neutron-star binary systems, high-precision X-ray space missions, such as the on-going NICER \cite{2014SPIE.9144E..20A} and the eXTP \cite{Zhang:2016ach}, will shed some more light on the properties of matter inside neutron stars by offering simultaneous measurements of their masses and radii \cite{Watts:2016uzu}. \appendix | 18 | 8 | 1808.06808 |
1808 | 1808.03100_arXiv.txt | We briefly summarize the BRITE observations of classical Cepheids so far. | BRITE is a constellation of five satellites, each equipped with a 3-cm telescope, observing the brightest stars in the sky \citep[see e.g.,][]{weiss,pablo}. Its primary targets are stars brighter than $4$\thinspace mag in $V$-band, however fainter stars, with slow variability, such as classical Cepheids, can be observed at high enough precision for their normally large amplitudes. So far BRITE has observed eleven Cepheids. Initial results were reported in \cite{brite2} and here we provide a short update. Table~\ref{tab} provides basic data on the observed Cepheids, of which six pulsate in the fundamental mode (top section of the Table) and five pulsate in the first overtone (bottom section). The consecutive columns provide the star's name and HD number, period as determined from the analysis of BRITE data, mean $V$-band brightness and a summary of observational data, i.e. satellite ID (UBr -- UniBRITE, BTr -- BRITE Toronto, BHr -- BRITE Heweliusz) and field ID. A short series of test observations in the blue filter are not included. \begin{table} \begin{tabular}{lrrrl} star & HD & $P$\thinspace (d) & $\langle V\rangle$& summary of observational data\\ \hline \multicolumn{5}{l}{\it fundamental mode Cepheids:}\\ T~Vul & 198726 & 4.4355 & 5.75 & UBr, BTr (Cyg-I), BTr (Cyg-II)\\ $\delta$~Cep & 213306 & 5.3663 & 3.95 & BHr, BTr (CasCep-I)\\ X~Sgr & 161592 & 7.0128 & 4.55 & UBr (Sgr-I), BHr (Sgr-II)\\ W~Sgr & 164975 & 7.5949 & 4.67 & UBr (Sgr-I), BHr (Sgr-II)\\ l Car & 84810 & 35.601~ & 3.72 & UBr (Car-I) \\ U Car & 95109 & 38.868~ & 6.29 & BHr (Car-I) \\ \multicolumn{5}{l}{\it first overtone Cepheids:}\\ DT~Cyg & 201078 & 2.4991 & 5.77 & UBr, BTr (Cyg-I), BTr (Cyg-II)\\ V1334~Cyg & 203156 & 3.3328 & 5.87 & UBr, BTr (Cyg-I), BTr (Cyg-II)\\ BG Cru & 108968 & 3.3427 & 5.49 & BTr (CruCar-I) \\ AH Vel & 68808 & 4.2272 & 5.70 & BHr, BTr (VelPic-II) \\ MY~Pup & 61715 & 5.6953 & 5.68 & BHr (VelPic-I) BHr, BTr (VelPic-II)\\ \hline \end{tabular} \caption{Basic data about Cepheids observed with BRITE (pulsation period, mean $V$-band brightness) and indication of satellites and observing campaigns in which data were gathered (red filter only; see \textsf{http://brite.craq-astro.ca/} for more details). Stars are sorted by pulsation mode (fundamental mode stars in the top section of the table) and by increasing pulsation period.} \label{tab} \end{table} The data reduction procedure was briefly described in \cite{brite2}. In a nutshell, the variability was first modelled with a Fourier series. The residuals were used to decorrelate the data with parameters like CCD temperature, star's position on the CCD or the satellite's orbital phase. The decorrelated data were then orbit-averaged. The remaining slow residual trends were modelled with polynomials. Outliers were removed during the procedure. Further details on BRITE photometry and its reduction can be found in \cite{popowicz}. | 18 | 8 | 1808.03100 |
|
1808 | 1808.06039_arXiv.txt | We present results from our SMA observations and data analyses of the SMA archival data of the Class I protostar IRAS 04169+2702. The high-resolution ($\sim$0$\farcs$5) $^{13}$CO (3--2) image cube shows a compact ($r \lesssim 100$ au) structure with a northwest (blue) to southeast (red) velocity gradient, centered on the 0.9-mm dust-continuum emission. The direction of the velocity gradient is orthogonal to the axis of the molecular outflow as seen in the SMA $^{12}$CO (2--1) data. A similar gas component is seen in the SO (6$_5$--5$_4$) line. On the other hand, the C$^{18}$O (2--1) emission traces a more extended ($r \sim$400 au) component with the opposite, northwest (red) to southeast (blue) velocity gradient. Such opposite velocity gradients in the different molecular lines are also confirmed from direct fitting to the visibility data. We have constructed models of a forward-rotating and counter-rotating Keplerian disk and a protostellar envelope, including the SMA imaging simulations. The counter-rotating model could better reproduce the observed velocity channel maps, although we could not obtain statistically significant fitting results. The derived model parameters are; Keplerian radius of 200 au, central stellar mass of 0.1 $M_{\odot}$, and envelope rotational and infalling velocities of 0.20 km s$^{-1}$ and 0.16 km s$^{-1}$, respectively. One possible interpretation for these results is the effect of the magnetic field in the process of disk formation around protostars, $i.e.$, Hall effect. | \label{sec:intro} Recent observational efforts have been finding Keplerian disks around not only T-Tauri stars ($e.g.$, Simon et al. 2000; Williams \& Cieza 2011) but also protostars \cite{tob12,har14,lin14,aso15}. The disks around protostars are considered to be under the formation and growth process of “protoplanetary disks”, precursors of planetary systems. Studies of disk formation processes around protostars are thus important to understand the initial condition of planet formation. Theoretical studies have suggested that magnetic fields play a vital role in disk formation from protostellar envelopes \cite{li14,mac16}. The magnetic field connects the inner regions and the outer envelopes and efficiently transfers the angular momentum from the inner to outer regions. This process is known as magnetic braking \cite{gil74,mou85,all03}. If an ionization degree of the cloud cores is high enough and ideal magnetohydrodynamics (MHD) approximation is valid, the magnetic braking is so efficient and almost completely suppresses the circumstellar disk formation in cloud cores under a typical magnetic-field strength \cite{mel08,hen09}. The ionization degree of the real cloud cores is, however, very low \cite{ume90,nis91,cas98,nak02}, and non-ideal MHD effects (Ohmic diffusion, Hall effect, and ambipolar diffusion) caused by low conductivity of the gas must play a role during the cloud core collapse. It has been suggested that the Ohmic and ambipolar diffusion decouples the magnetic field and the gas at $\rho>10^{-12}~{\rm g~cm^{-3}}$ and circumstellar disk formation is enabled even at the very early phase of star formation \cite{ma11a,tsu15a,tom15,mas16}. Recent theoretical simulations show that the Hall effect imprints the characteristic velocity structure in the envelope and disk, $i.e.$, a flip of the rotation velocity, or counter rotation \cite{kra11,li11,tsu15b,wur16,tsu17}. Another important physical mechanism that controls disk formation is turbulence. Theoretical simulations show that turbulence can reduce the magnetic braking through magnetic diffusions and reconnections, and add additional angular momenta \cite{san12,joo13,sei13,mat17}. Such turbulent effects promote disk formation around protostars. Furthermore, the rotational axis of the formed disk can be misaligned from that of the surrounding protostellar envelope \cite{mat17}. It has been difficult, however, to observationally identify gas motions which are indeed controlled by magnetic fields or turbulence in disk-forming regions. Detailed observational comparisons of gas motions from dense cores, envelopes, to central disks are essential to tackle this problem \cite{tob11,har14}. In particular, radial rotational profiles from the envelope to the inner disk have been measured observationally \cite{yen13,har14,har15,aso15}. Those observations show that the radial rotational profiles in the envelopes and disks can be approximated to be $v_{rot} \sim r^{-1}$ and $\sim r^{-0.5}$, that is, rotation with the conserved specific angular momenta and Keplerian rotation, respectively. While these results apparently show that the magnetic or turbulent effect on the gas motions is not significant, the rotational profiles and the power-law indices derived from these observations are not accurate enough to be directly compared with those from theoretical simulations. A more straightforward observational signature of such effects is desirable. We consider that change of the rotational axes between the central disk and the outer envelope, or even the flip of the rotational vectors, is an intriguing tracer for the effect of magnetic field and turbulence. In this paper, we report SubMillimeter Array (SMA)\footnote{The SMA is a joint project between the Smithsonian Astrophysical Observatory and the Academia Sinica Institute of Astronomy and Astrophysics and is funded by the Smithsonian Institution and the Academia Sinica.} observations of IRAS 04169+2702 at 330 GHz as well as data analyses of the SMA archival data at 230 GHz. IRAS 04169+2702 (hereafter I04169) is a Class I protostar ($L_{bol}$ = 0.76 $L_{\odot}$; $T_{bol}$ = 133 K) \cite{ken93a,ken93b,you03} located in the molecular filament of the B213 / L1495 region at $d$ = 140 pc \cite{hac13,taf15}. The protostar is associated with a $\sim$20000 AU scale, $\sim$1 $M_{\odot}$ dense core as seen in the 1.3-mm and 850 $\micron$ dust-continuum \cite{mot01,you03} and N$_{2}$H$^{+}$ (1--0) emission \cite{tat04}. Previous millimeter interferometric observations of I04169 in the C$^{18}$O (1--0) \cite{oha97} and H$^{13}$CO$^{+}$ (1--0) lines \cite{sai01} have found a $r \sim$1000 au scale protostellar envelope elongated along the northwest to southeast direction (P.A. = 154$\degr$) with an inclination angle of $i \sim60\degr$. In the C$^{18}$O (1--0) line, the southeastern part of the envelope is blueshifted ($\sim$-0.8 km s$^{-1}$ from $v_{sys}$ = 6.8 km s$^{-1}$) and northwestern part redshifted ($\sim$+0.8 km s$^{-1}$), and this velocity gradient along the major axis is regarded as the rotation of the envelope \cite{oha97}. CARMA 1.3-mm continuum and Keck I-band imaging of I04169 exhibit a small, almost unresolved ($\lesssim 1\arcsec$) dusty disk without any scattered light \cite{eis12}. From the SMA data of I04169, we have found possible observational evidence for a counter rotation between the protostellar envelope and circumstellar disk, which will be shown in the rest of the present paper. In Section 2, we shall describe our new SMA observations and archival data, and calibrations and imaging of those data. In Section 3, the 0.9-mm and 1.3-mm continuum, $^{12}$CO (2--1), $^{13}$CO (2--1; 3--2), C$^{18}$O (2--1), and SO (6$_5$--5$_4$) results are presented and compared. Section 4 describes our modeling efforts to reproduce the observed velocity structures with the SMA. In Section 5.1 we discuss physical origin of the observed gas motions around I04169 and in section 5.2 implications of these results. | \label{sec:dis} \subsection{Nature of the Detected Velocity Features} \label{subsec:nature} The molecular-line data toward I04169 reveal distinct velocity features at different spatial scales. At $r \sim$400 -- 1000 au, the northwestern part is redshifted and the southeastern part blueshifted, as traced by the C$^{18}$O (2--1; 1--0) lines (see Table \ref{sumvel}). In the inner $r \lesssim200$ au the northwestern part becomes blueshifted and the southeastern part redshifted, as seen in the SO (6$_5$--5$_4$) and $^{13}$CO (3--2) emission. Our toy model shows that the observed velocity features can be reproduced with a system of the central disk plus the outer counter-rotating, infalling envelope, although our simple modeling effort cannot prove that the counter-rotating model is significantly better than the forward-rotating model in a quantitative way. A possible source to change directions of velocity gradients is turbulence. Recent theoretical simulations of magnetized and turbulent collapsing dense cores show that turbulences produce additional angular momenta and induce magnetic diffusions and reconnections \cite{san12,joo13,sei13,mat17}. These effects reduce relative strengths of magnetic braking and thus promote disk formation around the central stars. The axis of the formed disk (or the angular-momentum vector) can be misaligned from the that of the magnetic field, outflow, and the envelope \cite{mat17}. Observationally, Harsono et al. (2014) conducted PdBI observations of disks around low-mass protostars, and compared the rotational directions of the $r \sim$100 au scale disks to the directions of the velocity gradients of the dense cores as seen in the N$_{2}$H$^{+}$ emission observed with FCRAO \cite{cas02}. The comparison between the $r \sim$100 au scale disks and $\sim$10000 au scale cores shows that there are indeed differences of the directions of the velocity gradients between the disks and cores, and the differences range from $\sim$90$\degr$ to $\sim$230$\degr$. Such differences of the velocity gradients between the large-scale cores and disks can be attributed to the effect of the turbulences. It is also possible that the identified change of the direction of the velocity gradients in I04169 is due to the turbulence. Presence of multiple gas clumps, which have different motions and velocities, can also reproduce the observed, apparent flip of the velocity gradients in I04169. If the interpretation of the counter rotation is correct, the effect of magnetic fields is a promising source to realize such a velocity structure. The physical mechanism to produce such a counter rotation through magnetic fields is known as the Hall effect, one of the non-ideal MHD effects \cite{war04,war12,bra12a,bra12b}. In collapsing cloud cores, the Hall effect induces toroidal (or azimuthal component of) magnetic fields from poloidal magnetic fields at the mid-plane of the pseudo-disk or the flattened envelope. The induced toroidal field exerts the magnetic torque and the gas rotation with the left-handed screw direction of the global poloidal field in the case of the negative Hall resistivity. If the Hall induced magnetic torque is large enough and has an opposite direction to the initial rotation (this corresponds to the case in which the poloidal magnetic field direction is parallel to the angular momentum of the cloud core), the gas rotation can flip. The counter rotating structure also appears even when the direction of the Hall induced magnetic torque is the same as that of the initial rotation. In this case, the midplane of the envelope and the disk maintains the initial rotation direction, but the upper envelope exhibits the counter rotation due to the back reaction of the Hall induced forward rotation at the midplane. In both cases, % the flattened envelope or peudo-disk with the scale of $\sim 100$ AU can exhibit the counter rotation \cite{kra11,li11,tsu15b,wur16,tsu17}. Whether the Hall effect can induce the counter rotation or not depends on the magnetic field strength of the parent cloud core. Previous theoretical studies have shown that the counter rotation is realized in the cores with the mass-to-flux ratio $\lambda \sim$5 \cite{li11,tsu15b,tsu17}. Here $\lambda$ is normalized by its critical value. The Hall effect is more effective under a stronger magnetic field ($i.e.$, lower $\lambda$), and the previous observations of the magnetic fields in dense cores have found $\lambda \sim$2 \cite{tro08,cru12}. Thus, the flip of the rotational direction caused by the Hall effect is likely possible under a typical condition of cloud cores. The magnetic field strength can be estimated from given $\lambda$ as $B=7.6 \times10^{-21} N(H_2) \lambda^{-1}~\mu G$ \cite{tro08}. Toward I04169 $N(H_2)$ value is estimated to be 1.4$\times$10$^{22}$ cm$^{-2}$ \cite{mot01}, and $\lambda \sim$2 yields the field strength of 54 $\mu$G. The limited spatial resolution and dynamic range of the present SMA data, however, prevent us from discriminating these different interpretations. If further higher-dynamic range observations of I04169 unveil the consistency of the velocity gradient from $r \sim$1000 au down to 400 au scale and then the flip of the velocity gradient below 400 au, such results must be a strong evidence for the presence of counter rotation between the disk and envelope and the magnetic-field origin. In these spatial scales, theoretical simulations show that the possible range of the misalignment between the central disk and the outer envelope originated from turbulence is at most $\sim$30$\degr$ \cite{mat17}. Unless the core mass is as high as 100 $M_{\odot}$, the misalignment does not show a complete flip $i.e.$, 180$\degr$ \cite{sei13}. Thus, if the future higher-dynamic range observations unveil a systematic, consistent velocity structure and a flip of the velocity gradient simultaneously, that can rule out the origin of turbulence or multiple gas components. In addition, a more thorough theoretical modeling (not toy model) and statistical parameter search are required to prove that the counter rotation caused by the magnetic effect is the most probable and unique interpretation. \subsection{Implications of the Opposite Velocity Gradients} \label{subsec:impli} The present SMA observations of the Class I protostar I04169 have found that at different spatial scales the directions of the velocity gradients are opposite. With our simple model, we suggest that one of the intriguing interpretations is counter rotation between the protostellar envelope and disk caused by the magnetic fields, whereas we admit that at this stage we cannot exclude the other possible explanations. Furthermore, a marginal, slow infalling velocity ($\sim$0.16 km s$^{-1}$) in the envelope has been identified. The identified infalling velocity corresponds to the free-fall velocity toward the central mass of $\sim$0.01 $M_{\odot}$ or smaller at $r \lesssim$700 au. If the mass of the central protostar is 0.1 $M_{\odot}$ as inferred from the $^{13}$CO (3--2) and SO (6$_{5}$--5$_{4}$) results and our modeling, the observed infalling velocity is much smaller than the corresponding free-fall velocity. Such small infalling velocities around the central disks have also been seen in the other protostellar objects from our recent ALMA observations \cite{oha14,aso15}. These results have been interpreted as the transitions from the infalling envelopes to the central disks with the increasing centrifugal support. The physical origin of the slow infalling velocity may also be the effect of magnetic fields \cite{li11,ma11b}. Opposite velocity gradients at different spatial scales have also been seen in the other protostellar sources. In HL Tau, the $r \sim$100 au scale protoplanetary disk exhibits blueshifted emission to the southeast of the protostar and redshifted emission to northwest \cite{alm15}. On the contrary, the follow-up ALMA observations by Yen et al. (2017b) have found that the southeastern part of the $r \sim$1000 au envelope around the protoplanetary disk of HL Tau as seen in the $^{13}$CO (2--1) emission is redshifted and the northwestern part blueshifted. While Yen et al. (2017b) argued that a simple counter-rotating model is not sufficient to fully reproduce the observed gas motions, presence of opposite signs of the velocity gradients is identified with the ALMA observations. Among the protostellar sample investigated by Harsono et al. (2014), L1527 IRS shows almost a complete flip ($\sim$177$\degr$) of the velocity gradient between the core and the disk. Tobin et al. (2011) have also found that the northern part of the $r \sim$8000-au scale protostellar envelope around L1527 IRS is blueshifted and the southern part redshifted, and that in the inner $r \sim$1000-au scale the direction of the velocity gradient flips. The larger-scale velocity gradient is consistent with the result from the single-dish C$_{3}$H$_{2}$ (2$_{12}$--1$_{01}$; 2$_{02}$--1$_{11}$) observations \cite{tak01}, and interferometric observations of L1527 IRS have confirmed the presence of the $r \lesssim 100$ au scale Keplerian disk with the opposite velocity gradient \cite{tob12,sak14,oha14,aso17}. In these two cases the differences of the spatial scales are within an order of magnitude. These results imply that our SMA results of I04169 are not unique. While both turbulences and magnetic fields have been considered to play a vital role in star and circumstellar-disk formation out of cloud cores, % it has been difficult to observationally identify such effects in protostellar sources. Radial rotational profiles in the protostellar envelopes have been measured observationally to study gas motions into circumstellar-disk formation ($e.g.$, Harsono et al. 2014; Yen et al. 2017a). Whereas these observations show that the radial rotational profiles in the envelopes and disks can be approximated to be $v_{rot} \sim r^{-1}$ ($i.e.$, rotation with the conserved specific angular momenta) and $\sim r^{-0.5}$ (Keplerian rotation), the rotational profiles measured from these observations are not accurate enough to be directly compared with those from theoretical simulations including magnetic fields and turbulence. Thus, it is not straightforward to infer the impact of magnetic fields and turbulence from the observed rotational profiles, if the direction of the rotational vector is common in the envelopes and disks. By contrast, the flip of the velocity gradient, if present, is rather easy to identify observationally. We thus suggest that further (re-)investigation of the opposite velocity gradients of molecular gas around protostellar sources should shed new light on the studies of star and circumstellar-disk formation. | 18 | 8 | 1808.06039 |
1808 | 1808.02330_arXiv.txt | Key information about the progenitor system and the explosion mechanism of Type Ia supernovae (SNe~Ia) can be obtained from early observations, within a few days from explosion. iPTF16abc was discovered as a young SN~Ia with excellent early time data. Here, we present photometry and spectroscopy of the SN in the nebular phase. A comparison of the early time data with a sample of SNe~Ia shows distinct features, differing from normal SNe~Ia at early phases but similar to normal SNe~Ia at a few weeks after maximum light (i.e. the transitional phase) and well into the nebular phase. The transparency timescales ($t_0$) for this sample of SNe~Ia range between $\sim$ 25 and 41 days indicating a diversity in the ejecta masses. $t_0$ also weakly correlates with the peak bolometric luminosity, consistent with the interpretation that SNe with higher ejecta masses would produce more \Nif. Comparing the $t_0$ and the maximum luminosity, \lm\, distribution of a sample of SNe~Ia to predictions from a wide range of explosion models we find an indication that the sub-Chandrasekhar mass models span the range of observed values. However, the bright end of the distribution can be better explained by Chandrasekhar mass delayed detonation models, hinting at multiple progenitor channels to explain the observed bolometric properties of SNe~Ia. iPTF16abc appears to be consistent with the predictions from the \mch models. | Type Ia supernovae (SNe~Ia) have long been linked to the explosion of a C/O white dwarf (WD) in a binary system \citep{Hoyle1960}, which has been confirmed by observed limits on the progenitor \citep{Nugent2011,Bloom2012}. However, there is still heated debate about the fundamental physical properties of the system, e.g. the mass of the progenitor, the nature of the companion and the mechanism of the explosion \citep[see,][for a review]{Hillebrandt2013,Maoz2014}. Many attempts to answer these open questions regarding the physics of SNe~Ia have concentrated on observations near maximum light. These observations have been critical to derive global parameters for SNe~Ia, e.g. synthesized \Nif\, mass, total ejecta mass \citep{Stritzinger2006a,Scalzo2014,Dhawan2016,Dhawan2017a}. However, there is a wealth of information available in observations shortly after explosion as well as at late times \citep[$\sim$ a year after maximum, i.e. the nebular phase, e.g.][]{Maguire2016,Graham2017}. Early time observations can shed light on the interaction between the SN ejecta and its companion \citep{Kasen2010}, and can be used to constrain the size of the companion. The signature of such an interaction can be seen as a sharp excess in the UV and blue flux at early epochs \citep[e.g.][]{Cao2015}, however, these pulses can be interpreted differently as in \citet{Kromer2016} and \citet{Noebauer2017}. Very early time observations of SNe~Ia \citep{Zheng2013,Zheng2014,Goobar2014,Goobar2015,Marion2016,Hoss2017,Miller2017,Jiang2017} have shown a diversity in their behaviour shortly after explosion and are a rich source of information regarding the progenitor system and the explosion mechanism. iPTF16abc was discovered shortly after explosion and showed some interesting early time features. It has a linear rise for the first three days after the time of first light, blue colours at early times compared to normal SNe \citep[e.g. SN~2011fe][]{Nugent2011}, strong carbon features in early spectra that disappear after $\sim$ 7 days \citep{Miller2017} and the near absence of the Si II\, 6355\,\AA\, in the earliest spectrum. Here, we present a nebular spectrum of iPTF16abc and analyse the photospheric (i.e. pre-maximum), transitional (i.e. $\sim$ +30 to +100 days) and nebular ($\sim$ +300 days) phase properties of iPTF16abc in context of SNe in the literature. The observations of the early, photospheric phase mostly probe the outer layers of the ejecta, the late phase, when the $\gamma$-ray escape fraction increases \citep{Jeffrey1999,Stritzinger2006a}, is sensitive to the inner layers of the SN ejecta, which are dominated by iron group elements (IGEs). We, therefore, aim to answer if the features seen in the early phase also manifest at late epochs and hence, whether they are a result of only the composition of the outer layers of the ejecta or also due to the inner core. The structure of the paper is as follows. We present the dataset in Section~\ref{sec-data}. We compare the properties of iPTF16abc to a sample of SNe in Section~\ref{sec-res}. We discuss our results in Section~\ref{sec-disc} and conclude in Section~\ref{sec-conc}. | 18 | 8 | 1808.02330 |
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1808 | 1808.05660_arXiv.txt | In this work we perform an analysis on the recently proposed conjoined cosmic growth and cosmic expansion diagram~\cite{Linder:2016xer} to compare several dark energy models using the Figure of Merit showed in~\cite{Basilakos:2017rgc}, which consists in the inverse of the $1\sigma$ confidence region in the $f\sigma_8(z)-H(z)$ plot. Our analysis also consists of comparing the models by performing different statistical criteria: Bayes factor~\cite{Trotta:2005ar}, the Bayesian Information Criteria~\cite{Schwarz:1978tpv} and the Akaike Information Criterion~\cite{AIC2}. We also developed a 3-dimensional Figure of Merit to account simultaneously for the errors on the growth rate and the Hubble parameter. The main idea is to consider several cosmological models and compare them with the different statistical criteria in order to highlight the differences and the accuracies of each single criterion. | Recent observations~\cite{Riess:1998cb,Perlmutter:1998np} pointed out that the Universe seems to be in a phase of accelerated expansion. These evidences have led cosmologists to revise the theory of the expansion of the Universe either by introducing a new component called dark energy~\cite{Sapone:2010iz} or by modifying directly the theory of gravity~\cite{Tsujikawa:2010zza}. Within the framework of Friedmann-Lema$\text{\^i}$tre-Robertson-Walker (FLRW) cosmologies, such accelerated expansion can be generated by adding up a simple cosmological constant $\Lambda$ to the total budget of the Universe. Even though the latter gives rise to severe coincidence and fine-tuning problems, observations still confirm such an explanation~\cite{Betoule:2014frx, Ade:2015xua, Abbott:2017wau}. Over the years a series of dark energy models have been considered in order to solve, or at least alleviate, the theoretical problems related to dark energy. However, none of these explanations seem to be convincing. Alternative theories of gravity came naturally as a consequence of the incapability of having a self-consistent model of dark energy. This class of models intends to modify General Relativity (GR) and to explain the observed acceleration of the Universe as a pure weakening of gravity at very large scales. The important question here is whether the two scenarios can be distinguished. It is well known that any Hubble expansion can be generated by choosing an appropriate equation of state for the dark energy, see~\cite{Bonvin:2006en}. However, over the years there have been claims that it is possible to distinguish alternative theories of gravity from dark energy models by using growth data; the last assumption is not always true unless the expansion history is fixed, \cite{Kunz:2006ca}. Nonetheless, recent works have proposed to study the cosmic growth versus cosmic expansion history conjoined diagram, the $f\sigma_8-H$ plot, to put constraints on the parameter space of cosmological models, or to compare different models directly~\cite{Linder:2016xer}. Model comparison using this approach has already been investigated in~\cite{Moresco:2017hwt, Basilakos:2017rgc}. The advantage of the $f\sigma_8-H$ plot over other probes relies on the degeneracy break of the history curves when comparing different models or the parameter space, since it contrasts a geometrical observable, given by $H(z)$, to a pure gravitational effect, given by $f\sigma_8(z)$. Using this approach, dark energy models were compared using the $f\sigma_8-H$ plot~\cite{Basilakos:2017rgc} through the FoM defined as the inverse area of the $1\sigma$ confidence region in the conjoined diagram. In this work, we follow a similar approach and we also compare the models using different statistical tools: the standard Bayesian evidence~\cite{Trotta:2005ar}, the Bayesian Information Criteria~\cite{Schwarz:1978tpv} (BIC), the Akaike Information Criterion \cite{AIC} (AIC) and the FoM. Furthermore, we considered an extension to the FoM which we define 3-FoM, which considers both errors on $f\sigma_8(z)$ and $H(z)$. Anticipating the results, we find that the FoM is a fairly good estimator of the errors, however, its extension, the 3-FoM, captures simultaneously the growth of matter and the expansion history making it more stable over different models. The criteria BIC and AIC$_c$ penalize substantially models with extra parameters. The paper is structured as follows: in Section \ref{basic-eqs} we report the basic equations that will be used in our work, whereas in Section \ref{models} we list the cosmological models that will be compared, and the link between $H$ and $f\sigma_8$ measurements. In Section \ref{data} we show the datasets used in the analysis and the statistical methodology is reported in Section \ref{methodology}. In Section \ref{results} we report the results of our analysis. | In our work we implemented the conjoined $H(z)-f\sigma_8(z)$ method in order to test an entire family of ten dark energy models; we started with the simplest model, $\Lambda$CDM which is described by three parameters only, and we systematically increased the level of complexity of the model by adding extra parameters, being the non-flat CPL with dark energy perturbation the most complex model (with seven parameters). For each model, we first found the best fit using MCMC analysis by combining the most recent cosmic chronometer and growth data available. Subsequently, we compared the dark energy models with five different statistical criteria, aiming at highlighting the potentiality and the weakness of each criterion. As expected, we found that the evidence is the most accurate statistical test to compare different models as it takes into account the information of the entire likelihood of the parameters and it does not always penalize a model with extra parameters. The 3-FoM better characterizes the sensitivity of the parameters according to the data used. This criterion takes into account simultaneously the errors from both $f\sigma_8(z)$ and $H(z)$; in particular, we showed that the errors of the Hubble parameter increase with redshift and this has an important effect on the constraining power of the test. The FoM instead is limited only to $f\sigma_8(z)$, hence neglecting the information from $H(z)$, which might be crucial if the analysis is extended at high redshift. As a complementary test, we performed the same analysis in the same redshift range as in~\cite{Basilakos:2017rgc} and we found consistent results. For the last two criteria, BIC and AIC$_c$, we showed that they always penalize the addition of extra parameters; in fact, if we consider the two extreme models, i.e. $\Lambda$CDM with only three parameters and non-flat CPL with dark energy perturbations, which has seven parameters, we find that $\Delta \text{BIC} \sim 40$ manifesting a {\em very} strong evidence in favor of the $\Lambda$CDM model. Similarly, but less decisive is $\Delta$AIC$_c$ for which we find a value of $\sim 10$, which still favors strongly $\Lambda$CDM but more moderately than BIC. To demonstrate the power of the 3-FoM, we compute the FoM and 3-FoM at different redshifts starting from $z=0$ up to the $z_{max}$. These results are shown in Fig.~\ref{fig:FoM-3FoM-joint} where we plotted the relative difference of the FoM (top panel) and the 3-FoM (lower panel) for each model with respect to $\Lambda$CDM. It is interesting to notice that at low redshifts the FoM for $w$CDM, $w$CDM-nf, CPL, and CPL-nf is larger than $\Lambda$CDM, meaning that the former is better constrained than the latter. This effect is not manifested in the 3-FoM which is always larger for the $\Lambda$CDM model. \begin{table}[h] \centering \begin{tabular}{ccc} \hline \hline Redshift bin & $H(z)$ [km s$^{-1}$ Mpc$^{-1}$] & $f\sigma_8(z)$ \\ \hline 0 $< z \leq$ 0.4 & 76.8 $\pm$ 5.8 & 0.410 $\pm$ 0.025 \\ 0.4 $< z \leq$ 0.8 & 92.0 $\pm$ 8.6 & 0.456 $\pm$ 0.037 \\ 0.8 $< z \leq$ 0.12 & 121.5 $\pm$ 13.3 & 0.390 $\pm$ 0.104 \\ 0.12 $< z \leq$ 0.16 & 161.2 $\pm$ 11.0 & 0.404 $\pm$ 0.056 \\ 1.16 $< z \leq$ 1.2 & 194.2 $\pm$ 32.2 & 0.364 $\pm$ 0.106 \\ \hline \hline \end{tabular} \caption{Binned measurements of $H(z)$ and $f\sigma_8(z)$ with equispaced redshifts points and its uncertainties. These are the gray points shown in Figs.~\ref{figs:L-w-joint} and \ref{fig:CPL-joint}.} \label{tab:binned} \end{table} | 18 | 8 | 1808.05660 |
1808 | 1808.10655_arXiv.txt | {Some thermonuclear (type I) X-ray bursts at the neutron star surfaces in low-mass X-ray binaries take place during hard persistent states of the systems. Spectral evolution of these bursts is well described by the atmosphere model of a passively cooling neutron star when the burst luminosity is high enough. The observed spectral evolution deviates from the model predictions when the burst luminosity drops below a critical value of 20--70\% of the maximum luminosity. The amplitude of the deviations and the critical luminosity correlate with the persistent luminosity, which leads us to suggest that these deviations are induced by the additional heating of the accreted particles. We present a method for computation of the neutron star atmosphere models heated by accreted particles assuming that their energy is released via Coulomb interactions with electrons. We computed the temperature structures and the emergent spectra of the atmospheres of various chemical compositions and investigate the dependence of the results on the velocity of accreted particles, their temperature and the penetration angle. We show that the heated atmosphere develops two different regions. The upper one is the hot (20--100 keV) corona-like surface layer cooled by Compton scattering, and the deeper, almost isothermal optically thick region with a temperature of a few keV. The emergent spectra correspondingly have two components: a blackbody with the temperature close to that of the isothermal region and a hard Comptonized component (a power law with an exponential decay). Their relative contribution depends on the ratio of the energy dissipation rate of the accreted particles to the intrinsic flux from the neutron star surface. These spectra deviate strongly from those of undisturbed, passively cooling neutron star atmospheres, with the main differences being the presence of a high-energy tail and a strong excess in the low-energy part of the spectrum. They also lack the iron absorption edge, which is visible in the spectra of undisturbed low-luminosity atmospheres with solar chemical composition. Using the computed spectra, we obtained the dependences of the dilution and color-correction factors as functions of relative luminosities for pure helium and solar abundance atmospheres. We show that the { helium model atmosphere heated by accretion corresponding to 5\% of the Eddington luminosity} describes well the late stages of the X-ray bursts in 4U\,1820$-$30. } | Type I thermonuclear X-ray bursts have been used to obtain neutron star (NS) parameters. In particular, bursts happening during the hard persistent states of the systems can be used for such analysis \citep[see discussions in][]{SPRW11,Kajava.etal:14,Poutanen.etal:14,Suleimanov.etal:16}.\footnote{We note here that the bursts occurring in the soft, high-accretion-rate state never show spectral evolution consistent with theoretical prediction \citep{Kajava.etal:14}. Therefore, this kind of burst cannot be used to determine NS parameters at all. } In order to use full information on the variations of the burst spectrum temperature and normalization, the so called cooling tail method or its modification \citep{SPRW11,Poutanen.etal:14,Suleimanov.etal:17} have been used. More accurate results can be obtained by directly fitting the burst spectra at different flux levels with NS atmosphere models \citep{Nattila.etal:17}. However, these hard-state bursts show a spectral evolution that deviates from the theoretically predicted behaviour for passively cooling NSs when the burst luminosity drops below a certain level. Furthermore, the higher the persistent flux, the more significant are the deviations. For instance, deviations start to be visible at burst luminosities below 20\% of the Eddington value $L_{\rm Edd}$ for the bursts taking place at persistent luminosities of about $0.01\,L_{\rm Edd}$ \citep{Nattila.etal:16}. On the other hand, the X-ray bursts of the helium accreting NS in 4U\,1820$-$30 happen at the relatively high persistent luminosity of about $0.07\,L_{\rm Edd}$, and such deviations begin at significantly higher burst luminosity of $\approx 0.7\,L_{\rm Edd}$ \citep{Sul.etal:17}. This leads us to the conclusion that additional heating by the accreted gas is the cause of the deviations. Therefore we need to develop NS atmosphere models with an additional energy dissipation in the surface layers. Such models would be applicable to the hard-state bursts. On the other hand, in the soft state the accreted matter likely spreads over the NS surface { forming} a spreading/boundary layer \citep{IS99,SulP06}. Its interaction with the NS atmosphere cannot be described by a formalism involving interaction of single particles, but a hydrodynamical treatment would be needed. Thus the modelling of bursts happening in the soft state will not be considered here. X-ray spectra of accreting compact objects such as NSs and black holes in their hard persistent states form in a hot rarefied medium. In black holes, the most likely source of this radiation is the inner geometrically thick and optically thin hot accretion flow \citep[see e.g. reviews by][]{DGK07,PV14,YN14}. In accreting NSs this radiation maybe produced also at the NS surface \citep{DDS01}, which even can dominate the total power. In this paper we consider old weakly magnetized NSs in low-mass X-ray binaries (LMXBs). \citet{ZS69} were among the first to discuss the emission from a NS surface heated by accreted matter. They considered two physical processes that lead to deceleration of protons in NS atmospheres. The simplest one is Coulomb interaction with electrons. In this case, the main part of the proton kinetic energy is released deep in the optically thick atmospheric layers resulting in the rather thermal emergent spectra. Another possibility is excitation of collective plasma processes by the protons. Here the proton energy is dissipated within an optically thin atmospheric layer and the emergent spectrum may have a shape far { different} from the blackbody. \citet{AW73} confirmed these conclusions using first numerical models of accretion heated NS atmospheres. They considered hydrogen NS atmospheres heated by protons falling radially with free-fall velocity. The proton deceleration was considered in a self-consistent way. Two input parameters of the model were considered, namely the NS mass and the accretion luminosity. The most important properties of the accretion heated atmospheres, such as an overheated Compton-cooled outer layers and spectral hardening with the increase of the accretion luminosity were found in this work. Later the problem was considered in detail by \citet{BSW92}. They provided useful relations for describing the proton velocity profiles and the vertical distribution of the accretion heating, which were used by other authors for modeling accretion heated atmospheres \citep{Turolla.etal:94, Zampieri.etal:95, Zane.etal:98}. They also considered hydrogen atmospheres heated by accreting protons. They confirmed the basic properties of the heated model atmospheres found by \citet{AW73}, and additionally declared existing of the so called hot solutions, where the temperature can reach almost 10$^{12}$\,K \citep{Turolla.etal:94}. The properties of such "hot" solutions were investigated by \citet{Zane.etal:98}, who, in particular, demonstrated the importance of electron-positron pair creation. \citet{Zampieri.etal:95} concentrated on the low-luminosity accretion heated atmospheres. They assumed the energy release mainly in the optically thick layers of the atmosphere. Therefore, the temperature structures of these models at the spectrum formation depths were close to that of the undisturbed model atmospheres with the same luminosities. As a result, they obtained spectra harder than the blackbody and very similar to those of undisturbed NS atmospheres \citep[see, e.g.][]{Romani:87, Zavlin.etal:96}. The main reason is the dependence of the free-free opacity on the photon energy $k_{\rm ff} \sim E^{-3}$. Photons prefer to escape from an atmosphere in the most transparent energy bands. \citet{Zampieri.etal:95} also found a clear division between the outer heated atmospheric layers cooled by Compton scattering from the almost isothermal inner part cooled by free-free processes. \citet{DDS01} computed the accretion heated NS atmospheres and proton deceleration self-consistently and used the results for interpretation of the hard persistent spectra of some LMXBs. For the first time they showed the importance of the temperature of accreted protons on the atmosphere properties. They also considered the bulk velocity as a free parameter. The results obtained by \citet{Zampieri.etal:95} and \citet{DDS01} were widely used for interpretation of the NS spectra in LMXBs during low mass-accretion-rate states \citep[see e.g.][]{Homan.etal:14, Wijnands.etal:15}. Also deceleration of the protons in magnetized NS atmospheres was considered before \citep[see e.g.][]{NSW93}, and the results were used for the modeling of such atmospheres by \citet{Zane.etal:00}. In this paper we develop a method to compute models of NS atmosphere heated by accreted matter. It is based on the approach developed by \citet{DDS01}, which we slightly modified. In contrast to previous papers devoted to pure hydrogen atmospheres heated by protons, here we consider arbitrary chemical composition of the atmospheres and an arbitrary mix of protons and $\alpha$-particles for the accreted particles. We present the basic properties of the accretion-heated NS atmospheres and their dependence on the input parameters such as the velocity of accreted particles and their direction, the accretion- and the intrinsic NS luminosities. We also compare the developed models to the observed spectral evolution of hard-state X-ray bursts. | In this paper we presented a computational method to construct NS { model} atmospheres heated by accreted particles. Our results confirm previous findings \citep[e.g.][]{AW73, Zampieri.etal:95, DDS01} that fast heavy particles (protons, $\alpha$-particles, or their mix) penetrate down to some depth in the atmosphere and release most of their kinetic energy in a relatively narrow layer due to Coulomb interaction with electrons. The heating by accretion is balanced by the cooling due to mainly free-free emission and Comptonization. The relatively dense optically thick region, where the cooling is dominated by free-free emission, is only slightly hotter than the corresponding layer of an undisturbed atmosphere. It forms a quasi-isothermal transition region between the inner atmosphere unaffected by accretion heating and the hot ($\sim 10^8 - 10^9$\,K) rarefied upper layer. The upper hot layers cool mainly by Compton scattering. The width of the transition region and the temperature of the upper layer depend mainly on the accretion luminosity $L_{\rm a}$ and the intrinsic NS luminosity $L$. The emergent spectrum of an accretion-heated NS atmosphere is wider than the spectrum of the corresponding undisturbed atmosphere and can be roughly represented as a sum of two components: a blackbody-like radiation from the transition region and Comptonized (cutoff power-law-like) spectrum of the hot upper layers. The relative contribution of the components is determined by the ratio of the intrinsic luminosity $L$ to the accretion luminosity $L_{\rm a}$. If $L_{\rm a} \gg L$, the total spectrum is dominated by the Comptonized spectrum of the upper, hot rarefied layers. This kind of spectra are similar to the observed spectra of LMXBs in their low hard spectral states, as mentioned by \citet{DDS01}, and their slope depends on the input parameters of the accreted particles: velocity, temperature, and the penetration angle. We note, however, that all computed spectra have photon indexes $\Gamma > 2$. The observed spectra of LMXBs in their low hard spectral states often have harder spectra with $\Gamma < 2$. It means that the contribution of the accretion flow emission to the total observed spectra may not be negligible. In the opposite case, when $L_{\rm a} \ll L$, the total spectra are very similar to the spectra of the undisturbed models with harder Wien tails and increased flux at low photon energies. In that case, the final spectra depend very little on the properties of accreted particles. The atmospheres heated by material of solar abundance have an important qualitative feature in comparison to pure helium atmospheres heated by $\alpha$-particles only. The spectra of the relatively low-luminosity ($L \ll L_{\rm Edd}$) undisturbed model atmospheres show significant absorption edges due to photoionisation of hydrogen-like iron. These edges disappear in the spectra of the accretion-heated atmospheres because of a significant change in the temperature structure. We also investigated the influence of the accretion heating on the model curves $w - w\fc^4 \ell$ and $\fc - w\fc^4 \ell$, which are used in the (direct) cooling tail method. As it was expected the color-correction factors $\fc$ are larger for the heated atmospheres in comparison with the undisturbed atmospheres, and the dilution factors $w$ are smaller. The model curve computed for heated helium atmospheres (with $\ell_{\rm a}=0.05$) is well fitted to the low-luminosity part of the observed data $K-F_{\rm BB}$ obtained for X-ray bursts taken place during the hard state of the helium-accreting system 4U\,1820$-$30 \citep[see][]{Sul.etal:17}. The model curves $\fc - w\fc^4 \ell$ computed for the heated solar-abundance atmospheres do not have a dip at $\ell \approx 0.1$, which is clearly seen in the model curves computed for the undisturbed solar-abundance atmospheres. This fact may have important implications for interpretation of the X-ray data on bursting NSs that accrete gas of solar abundance, for example, the Clocked Burster GS\,1826$-$24 \citep[see discussion in][]{ZCG12}. We plan to apply the method described here for interpretation of the spectral evolution of that and other bursters in a follow-up paper. | 18 | 8 | 1808.10655 |
1808 | 1808.10149_arXiv.txt | We propose a model for dark matter and dark radiation, based on a strongly-coupled dark ${\rm SU}(5)$ gauge theory with fundamental and decuplet dark-quarks. The model supports light dark-baryons, respecting the chiral symmetry, which are electrically neutral but have electromagnetic form factors, and also a light dark-axion. Since the coupling of dark baryons to the standard model particles is inversely proportional to the square of the confinement scale, dark baryons become either hot dark matter or cold dark matter, depending on when the dark color confines. For the confinement scale $\Lambda\sim 10-10^3~{\rm GeV}$ the dark baryons of mass about $ 1~{\rm GeV}-1~{\rm MeV}$ become cold dark matter with naturally small magnetic moment and give the correct relic abundance. | Dark matter (DM) is widely believed to be one of the major components that make up our universe~\cite{Akrami:2018vks}. Though its origin is uncovered yet, the most popular explanation for DM is that it is made of weakly interacting massive particles (WIMP) that naturally arise with mass of a few hundred GeV in the models beyond the standard model (BSM) of particle physics. Such weakly interacting particles are thermally produced in early universe to give the correct relic density of our present universe. But, it has so far escaped all the searches that have been performed until now. Another popular theoretical candidate for DM is the axion, which is the pseudo Nambu-Goldstone boson of the spontaneously broken ${\rm U}(1)_{\rm PQ}$ symmetry, introduced to solve the strong CP problem in the standard model~\cite{Peccei:1977hh}. If the axion decay constant or the scale of the Peccei-Quinn (PQ) symmetry breaking is $f_{\rm PQ}\sim 10^{9}-10^{12}~{\rm GeV}$, then the axion becomes very light, $m_{\rm axion}\sim 10^{-2}- 10^{-5}~{\rm eV}$, and abundantly produced in early universe to become the significant component of DM at present~\cite{Preskill:1982cy,Abbott:1982af,Dine:1982ah,Kim:2008hd}. Recently however there have been lots of activities in both theory and also in experiments to probe the dark matter much lighter than a typical WIMP~\cite{Alexander:2016aln}. Among them the strongly interacting massive particles (SIMP) have been studied intensively. The SIMP models are not directly related to the BSM models that address the problems of the standard model, but they offer new windows for the dark matter searches to probe vast ranges of the DM mass~\cite{Hochberg:2014dra}. Furthermore, having significant self-interactions, they could explain the problems in the structure formation of our universe that WIMP models fail to explain~\cite{Hochberg:2014kqa}. In this letter we propose a new model of DM that contains very light and stable dark baryons and also very light dark axions. If the dark-baryon mass is vanishingly small, they become dark radiation. The model is based on a strongly coupled gauge theory that is confined but does not break all the chiral symmetries of the theory in the infrared (IR). In the confined phase the spin-$\frac12$ dark-baryons do exist and saturate the flavor anomalies of the ultraviolet (UV) sector of the theory. The dark baryons therefore remain massless in the IR without breaking the chiral symmetry~\footnote{Another consistent solution to the 't Hooft anomaly matching would be the broken chiral symmetry in the IR as in quantum chromodynamics (QCD), where the Nambu-Goldstone bosons saturate the anomaly.} by the 't Hooft anomaly matching condition~\cite{tHooft:1979rat}. For the dark-baryon DM, we break the chiral symmetry explicitly to give the dark baryons small mass and a magnetic moment, keeping them neutral however. The dark baryons therefore couple to the standard model particles electromagnetically and are produced thermally in the early universe to constitute DM of our universe. The model has in addition a very light dark-axion or dark $\eta^{\prime}$ that is potentially a good candidate for the hot DM. One of dark baryons in the model could be dark radiation, contributing the radiation energy density but well within the constraints from the cosmic microwave background (CMB) data, if we take its mass to vanish, as the mass is tunable in our model. Our model is a minimal model that provides naturally multicomponent dark matter and dark radiation with light dark-baryons and very light dark-axions, which might be needed to explain many observational data from sky that a single dark matter fails to explain~\cite{Zurek:2008qg}. | We have proposed a minimal model for multi-component DM and dark radiation. The model is based on the confining ${\rm SU}(5)$ gauge theory with a doublet of fundamental dark-quarks and another doublet of decuplet dark-quarks. In the infrared the model supports a doublet of massless dark-baryons in the chiral limit that saturates the flavor anomalies of the fundamental dark-quarks. The dark-baryons are however a flavor-singlet of the decuplet dark-quarks and therefore the chiral and axial symmetries of decuplet dark-quarks has to be broken. All the four Nambu-Goldstone bosons remain massless in the chiral limit, because the axial charges of dark-quarks are assigned to cancel the axial ${\rm SU}(5)$ anomaly. The model provides dark-baryons and dark-axions as light dark matter that contribute substantially to the relic density of our universe, if the confining scale of dark-colors $1~{\rm GeV}\lesssim \Lambda\lesssim 1~{\rm TeV}$. The dark baryons of our model are electrically neutral but have the electromagnetic form factors, since the dark-quarks carry electric charges. The dark-baryons interact electromagnetically with the standard model particles. If the fundamental dark-quarks have a small current-mass, the dark-baryons will have a small mass, proportional to the current mass, and also the magnetic moments, belonging to the class of dipolar dark matter~\cite{Sigurdson:2004zp}. Depending on the confining scale, they become either cold DM of mass $\sim1~{\rm MeV}-1~{\rm GeV}$, if $\Lambda\sim1-10^{-2}~{\rm TeV}$, or subdominant hot DM of mass $\sim 1~{\rm eV}$ if $\Lambda\sim\Lambda_{\rm QCD}$. When the dark-baryons are very light or massless, they become dark radiation. As long as $\Lambda>3~{\rm GeV}$, their contribution to the radiation energy is within the CMB bound, $\Delta N_{\rm eff}<0.12$. The model has also a very light dark-axion that could be subdominant hot DM of mass $\sim1~{\rm eV}$, if $\Lambda> 10^7~{\rm GeV}$, avoiding the astrophysical and cosmological constraints. Finally we note the dark-axion of our model becomes the QCD axion of the DFSZ model, solving the strong CP problem, if we identify the electroweak-singlet PQ field as the composite of decuplet dark-quarks, provided that the Higgs doublets $\varphi_u$ and $\varphi_d$ of the DFSZ model are also composite of additional dark-quarks and the confining scale $\Lambda$ is as high as the allowed PQ scale, $\Lambda\sim 10^9-10^{12}~{\rm GeV}$. % | 18 | 8 | 1808.10149 |
1808 | 1808.05456_arXiv.txt | {% Magnetic fields in galaxies exist on various spatial scales. Large-scale magnetic fields are thought to be generated by the $\alpha-\Omega$ dynamo. Small-scale galactic magnetic fields (1\,$\kpc$ and below) can be generated by tangling the large-scale field or by the small-scale turbulent dynamo. The analysis of field structures with the help of polarized radio continuum emission is hampered by the effect of Faraday dispersion (due to fluctuations in magnetic field and/or thermal electron density) that shifts signals from large to small scales. At long observation wavelengths large-scale magnetic fields may become invisible, as in the case of spectro-polarimetric data cube of the spiral galaxy NGC~6946 observed with the Westerbork Radio Synthesis Telescope in the wavelength range 17--23\,cm. The application of RM Synthesis alone does not overcome this problem. We propose to decompose the Faraday data cube into data cubes at different spatial scales by a wavelet transform. Signatures of the ``magnetic arms'' observed in NGC~6946 at shorter wavelengths become visible. Our method allows us to search for large-scale field patterns in data cubes at long wavelengths, as provided by new-generation radio telescopes. } | Magnetic fields of nearby galaxies are quite well investigated. The observational results are compatible with a scenario of magnetic field excitation by a galactic $\alpha-\Omega$ dynamo \citep[for a review see e.g.][]{Betal96}. The main bulk of these observations is obtained from analysis of galactic polarized radio continuum radiation observed with the current generation of radio telescopes, such as the Effelsberg 100-m dish and the Very Large Array (VLA). Galactic magnetic fields are important for the evolution of galaxies, the astrophysics of the interstellar medium, cosmic ray propagation \citep[e.g.][]{Pr}, as well as of fundamental interest. In particular, observations demonstrate prominent magnetic structures in the form of magnetic arms that are usually situated between material arms, as in NGC~6946 \citep{2007A&A...470..539B}. The relation between gas and magnetic fields can also be more complicated, like e.g. in IC~342 \citep{Beck2015} and in M\,83 \citep{Frick16}. The verification of various scenarios for the generation of such fine magnetic structures \citep[e.g.][]{Metal15,Anvetal13} is limited by technical abilities (angular resolution and sensitivity) of current radio telescopes. Further progress can be expected with the new generation of radio telescopes, the European Low Frequency Array (LOFAR), the Australia SKA Pathfinder (ASKAP) and the South-African Karoo Array Telescope (MeerKAT), allowing high-resolution, multichannel polarimetric observations. An important effect for the quantification of magnetic fields is Faraday rotation of polarized radio radiation that requires multi-wavelength observations \cite[e.g.][]{RS79}, while contemporary models for galactic magnetic fields are based on observations on few (and quite often of two) wavelengths only. As the Faraday rotation angle increases with the square of wavelength, observing at substantially longer wavelengths compared to those ones for the Effelsberg and VLA telescopes (3--20\,cm) increases the accuracy of the measurements substantially \citep{Beck2012}. On the other hand, increasing the observation wavelength leads to more severe Faraday depolarization effects, which complicates the extraction of physically valuable information on magnetic fields from observational data \citep{B66,Setal98}. Faraday depolarization, in particular Faraday dispersion due to small-scale fluctuations in magnetic field and thermal electron density within the emitting medium or in the Faraday-rotating medium in the foreground \citep{Setal98}, can already be strong at wavelengths of around 20\,cm. Maps of polarized radio emission obtained with the Westerbork Synthesis Radio Telescope (WSRT) for many (21) galaxies \citep{Heald2009} do not reveal (at least not in a straightforward way) many structures that are well known from observations at shorter wavelengths. In particular, the map of Faraday depths (a measure of the integral of the product of thermal electron density $n_\mathrm{e}$ and magnetic field strength $B$ along the line of sight) of NGC~6946 obtained by \citet{Heald2009} in the wavelength range 17--23\,cm (consisting of two bands of 17.0--18.4\,cm and 20.9--23.1\,cm) (Fig.~\ref{fig:data}, bottom) shows small-scale fluctuations superimposed on a large-scale gradient and does not directly show the well-defined ``magnetic arms'' known from a previous study at $\lambda$ 3.6\,cm and 6.3\,cm by \citet{2007A&A...470..539B}. Fig.~\ref{fig:data} was obtained using the sophisticated method of RM Synthesis \citep{Brentjens2005}, but this does not overcome this problem. The aim of this paper is to demonstrate that by combining the ideas of wavelet analysis and RM Synthesis one can recover the magnetic arm configuration from observations at long radio wavelengths. The key idea of the approach is as follows. Strong Faraday depolarization randomizes almost all information concerning the large-scale magnetic field structure. Faraday rotation angles $0.81 \, B \, n_e \, d \, \lambda^2$ (where $B$ is the average strength of the magnetic field along the line of sight, $d$ is the coherence scale of the magnetic field, $n_\mathrm{e}$ is electron density and $\lambda$ the wavelength) are generally larger than $\pi$ at $\lambda\simeq20$\,cm. Fluctuations in $B$ and $n_\mathrm{e}$ lead to strong gradients in the maps of Stokes Q and U and hence to shifting the signals of polarized intensity from large to small angular scales. The power spectrum of polarized intensity becomes flatter \citep[e.g.][]{2006A&A...457....1L}. The imprints of the large-scale field remain recognizable at smaller spatial scales. By using wavelets we can isolate small-scale magnetic fields obtained as the result of decay of the large-scale ones and then recognize the locations of large-scale fields. We will test our method on the observational data for NGC~6946 by \citet{Heald2009}. The magnetic arm configuration in this galaxy is known from observations at short wavelengths \citep{2007A&A...470..539B}, allowing us to verify the results. \begin{figure} \centering \includegraphics[width=0.35\textwidth]{fullmagn} \includegraphics[width=0.35\textwidth]{fulldepth} \caption{Observations of NGC~6946 in the wavelength range 17--23\,cm and results of RM Synthesis: top - distribution of a (normalized) peak intensity $|F^{\rm max}|$ of the Faraday spectra with overlayed red contours at 15\% of the maximum, bottom - distribution of Faraday depths $\phi^{\rm max}$ (in $\FRM$) at which the maximal values of $|F^{\rm max}|$ are obtained, according to \citet{Heald2009}.} \label{fig:data} \end{figure} | The application of RM Synthesis followed by the determination of maxima of the Faraday spectra at each pixel in the sky plane at any Faraday depth does not clearly reveal any elongated arm structures (Fig.~\ref{fig:data}, top). On the other hand, if we first apply a wavelet decomposition which isolates structures at a given scale and then determine the maxima of the Faraday spectra at each pixel, the resulting image reveals pronounced structures that are invisible when applying RM Synthesis only. We clearly see from Fig.~\ref{fig:scales3}, middle that the isolated structures are organized in the form of spiral arms. The elongated arms consist of a set of local maxima along the arms. Comparing the middle plots of Fig.~\ref{fig:scales2} and Fig.~\ref{fig:scales3}, we find that $\wwtmax_a (l, b)$ has a higher contrast between arm and interarm regions than $\wmaxwt_a (l,b)$. We measure the contrast $c$ as follows. We divide the image into the magnetic arm and interarm space as it comes from 6 cm data (boundaries are shown by black contours in Fig.~\ref{fig:depth}, top) and exclude the very central part (distance from center up to 50 $\arcsec$ (equal to $1.67~\kpc$) and the outer part (more than 340~$\arcsec$ or $11.4~\kpc$). Then we calculate the intensities (root-mean-square values) of the wavelet coefficients in the arms and interarm areas and measure the contrast as the ratio between these intensities. The contrasts are given in Table~\ref{tab:rmss}. Our method enlarges this quantity by about 10\% for the northern arm and by about 15\% for the southern arm. In spite of this improvement, the contrast is still smaller than that at 6 cm. A straightforward recommendation would be to perform observations at 6\,cm and even 3\,cm with higher sensitivity, e.g. with the SKA. As this remains unrealistic for the near future, we have to restrict our demands to longer wavelengths, e.g. with the SKA precursors MeerKAT and ASKAP. Then the suggested method is able to improve the contrast up to 15\%, which can be sufficient to isolate the arms in the images. We remind that we used NGC~6946 as an illustrative example because the position of magnetic arms is known from 6\,cm data. Applications of the method are recommended for galaxies where 6\,cm data are absent. \begin{table} \caption{Comparison of contrast for various distributions between arm and interarm regions for wavelet coefficients at scale $a=16\arcsec\approx 535\pc$, versus 6 cm. Errors were estimated by standard deviation at 30\% bootstrapping of the sets} \label{tab:rmss} \begin{tabular}{cccc} \hline & $\wmaxwt_a $ & $\wwtmax_a$ & 6 cm \\ \hline north. arm & $1.398 \pm 0.007$ & $1.517 \pm 0.007$ & $2.002 \pm 0.009$ \\ south. arm & $1.116 \pm 0.005$ & $1.273 \pm 0.005$ & $2.310 \pm 0.010$ \\ both arms & $1.212 \pm 0.005$ & $1.354 \pm 0.004$ & $2.221 \pm 0.008$ \\ \hline \end{tabular} \end{table} Structures of $\wwtmax_a (l, b)$ fit to the large-scale structures that are visible in the polarization map at 6$\cm$ wavelength (see Fig.~\ref{fig:depth}, top), located between the optical arms (see Fig.~\ref{fig:depth}, bottom). We obtain enlargement of contrast for small $a$ only. It means that the method is sensitive to small-scale details in the image. These details can correspond to real small-scale structures of the magnetic field or be the result of Faraday effects on the polarized emission from the large-scale magnetic field, to be figured out with numerical tests which is presented below. Note that the polarized intensity map at 17-23$\cm$ (see Fig.~\ref{fig:data}, top) revels some detail of magnetic arms which are identified in Fig.~\ref{fig:depth}. The point is however that the details in Fig.~\ref{fig:data}, top are embedded in diffuse surrounding and its relation with magnetic arms remains unclear. \begin{figure} \centering { \includegraphics[width=0.38\textwidth]{6cm+16cont} \includegraphics[width=0.38\textwidth]{scale4depth} } \caption{Top: Isolated magnetic arms (red contours at $1.5\times rms$) obtained from the data at 17--23\,cm with the wavelet transform for scale $a = 16\arcsec$ (Fig.~\ref{fig:scales3}, middle), and model arms for methods comparison (black contours) overlayed on the image of polarized intensity at 6\,cm wavelength (grayscale). The maximum intensity is $340\uJyb$. Bottom: Isolated magnetic arms (shown in colour depicting the Faraday depths $\phi^{\rm max}$ ($\FRM$) of $F^{\rm max}$), overlayed on an optical image (grayscale).} \label{fig:depth} \end{figure} To illustrate the idea of the method we constructed an artificial example, producing a data cube for the same set of channels as in data for NGC~6946 analysed before. We simulate the polarized emission (in the computational box 25$\times$25$\times$5$\kpc$) emerging from a large-scale magnetic field in the galactic disk observed face-on that is embedded in a 3D homogeneous isotropic turbulent field in the halo. The large-scale field has only an azimuthal component perpendicular to the line of sight and Gaussian shape \begin{equation} B_\varphi(r,\varphi,z)=B_0 \exp{\left[-\left(\frac{r}{r_0}\right)^2-\left(\frac{z}{z_0}\right)^2\right]} \, , \end{equation} where the radial Gaussian scalelength is $r_0=10$ $\kpc$, the vertical Gaussian scale is $z_0=1$ $\kpc$ and the strength of the large-scale field is $B_0=3$ $\mu$G. The homogeneous turbulent field was simulated by a numerical model that allows to control the spectral law and the characteristic scale of turbulence $l_t$ (providing the maximum of energy spectrum) \citep{Stepanov2014}. We do not consider any possible inhomogeneity of the turbulent field in the midplane. We choose a spectral slope of -5/3 at smaller scales than $l_t$ and +2 at larger scales. We choose three different values of $l_t=100, 200, 400$ $\pc$ in order to search for a possible dependence. The rms strength of the turbulent field is taken as 1\,$\mu$G. The number densities of relativistic and thermal electrons are assumed to be constant, namely $n_c=1$ cm$^{-3}$ and $n_e=0.1$ cm$^{-3}$. The large-scale disk magnetic field corresponds to a Faraday-thin source of the synchrotron emission which does not cause significant Faraday depolarization. The turbulent field acts as a Faraday screen that does not contribute much to the emission but depolarizes it and disperses it to different Faraday depth. The output of the model is two data cubes of Q and U with coordinates and frequency. Next, we perform RM Synthesis on the artificial data cube using the same wavelength range as in the case of NGC~6946 and calculate the wavelet coefficients $\wmaxwt_a (l,b)$ and $\wwtmax_a (l, b)$. The wavelet filter doesn't influence on Faraday spectra directly, however, it suppresses peaks in 2D map whose scales are different from scale $a$. The intensities (spectral power) of $\wmaxwt_a (l,b)$ and $\wwtmax_a (l, b)$ versus scale $a$ are shown in Fig.~\ref{fig:test4spec}. The intensity of $w^{40}_a$ is on a low level on all scales because the polarized intensity at Faraday depth $\phi=40\FRM$ does not contain spectral power at large scales. The intensity of $\wmaxwt_a (l,b)$ substantially increases with scale because only large scales are prominent, while the intensity of $\wwtmax_a (l, b)$ is practically constant, it detects power on all scales. If the large scales are well represented in the data (i.e. if there is no strong Faraday depolarization and most spectral power is on large scales), the standard method of structure recognition ($\wmaxwt_a (l,b)$) works best and reveals the structure of the large-scale magnetic field. If, however, Faraday depolarization is strong and there is no chance to recognize the large scales directly, our new method opens an additional possibility to find imprints of the large scales at small scales; the small scales are much better recognized by $\wwtmax_a (l, b)$ compared to $\wmaxwt_a (l,b)$. Late application of wavelet transform in $\wwtmax_a (l, b)$ suppresses the regions in the extended disk but keeps the regions in the magnetic arms, because the turbulent field is weaker there and hence there is less Faraday dispersion, so that the structures are less randomized. The scale of the crossing point in Fig.~\ref{fig:test4spec} depends of the properties of turbulence and observation wavelength and gives the upper limit in scale below which the suggested technique is suitable. \begin{figure} \centering \includegraphics[width=0.4\textwidth]{test4spectra} \caption{The root-mean-square spectral power as a function of scale $a$ for the simulated data: thick line -- $w^{40}_a$, dashed line -- $\wmaxwt_a$, dotted line -- $\wwtmax_a$. The characteristic scale of turbulence is $l_t=72\arcsec$, corresponding to 200 $\pc$. The noise level of the synthetic signal is zero. } \label{fig:test4spec} \end{figure} \begin{figure} \centering \includegraphics[width=0.25\textwidth]{diagram1} \caption{Scheme of a heuristic approach. Notation $(\cdot)$ means an operator argument taken as a result from the previous step.} \label{fig:diag} \end{figure} We checked the minimal requirement of recovering the Faraday spectrum from the observational frequency band. The FWHM of the Faraday Point Spread Function (FPSF), or equivalently, the resolution in Faraday depth space $\triangle\phi_{\rm FPSF}$ should be comparable at least to the rms dispersion of Faraday rotation caused by the turbulent magnetic field. Following \citet{Brentjens2005} the estimate $$\Delta\phi_{\rm FPSF}=\frac{2 \sqrt{3}}{\lambda^2_{\rm max}-\lambda^2_{\rm min}}$$ gives $\Delta\phi_{\rm FPSF}\approx170 \FRM $ for the observations used here. This is larger than the dispersion of Faraday rotation (about 40 $\FRM$) in Fig.~\ref{fig:data} (bottom). It explains why enlargement of contrast obtained by our method (see Table~\ref{tab:rmss}) remains moderate however sufficient to isolate magnetic arms in Fig.~\ref{fig:scales3}. If $\Delta\phi_{\rm FPSF}$ is very large, then $\phi^{\rm max}$ is about the same for all lines of sight, so that $\wmaxwt_a \approx \wwtmax_a$ and our method give the same results as the traditional one. Our model is admittedly simplistic. Observations at longer wavelengths generally probe magnetic structures that are farther from the galaxy midplane and hence close to the observer \citep[e.g.][]{2010A&A...514A..42B}. Presuming that the magnetic field morphology has some vertical structure, changes in the large-scale morphological features are expected when observing at progressively lower and lower frequencies. If the large-scale fields in the magnetic arms are tied to the star-forming ISM, we may expect to see them vanish at lower frequencies, where mainly the thick disk or halo is observed rather than emission from the disk that is depolarized by star-formation induced turbulence. In the case of NGC~6946, our result shows that the magnetic arms extend sufficiently high into the thick disk or halo, so that their imprints can still be detected at wavelengths around 20\,cm. The overall scheme of our analysis is shown in Fig.~\ref{fig:diag}. In summary, the distribution of small-scale magnetic fields as recognized by wavelet filtering of spatial scales traces the locations of the large-scale field (e.g. in the magnetic arms), if the imprints of large-scale fields are randomized by Faraday depolarization. This method is a powerful tool to analyze spectro-polarimetric data cubes obtained at long radio wavelengths. It should be applied to further galaxies from the survey by \citet{Heald2009} and to galaxies observed with the VLA in L-band (1--2\,GHz) and with new-generation radio telescopes like LOFAR, ASKAP, MeerKAT and SKA. | 18 | 8 | 1808.05456 |
1808 | 1808.03499_arXiv.txt | We derive the distance and structure of the Perseus molecular cloud by combining trigonometric parallaxes from Very Long Baseline Array (VLBA) observations, taken as part of the GOBELINS survey, and Gaia Data Release 2. Based on our VLBA astrometry, we obtain a distance of $321\pm10$~pc for IC~348. This is fully consistent with the mean distance of $320\pm26$ measured by Gaia. The VLBA observations toward NGC~1333 are insufficient to claim a successful distance measurement to this cluster. Gaia parallaxes, on the other hand, yield a mean distance of $293\pm22$~pc. Hence, the distance along the line of sight between the eastern and western edges of the cloud is $\sim$30~pc, which is significantly smaller than previously inferred. We use Gaia proper motions and published radial velocities to derive the spatial velocities of a selected sample of stars. The average velocity vectors with respect to the LSR are $(\overline{u},\overline{v},\overline{w})$ = $(-6.1\pm1.6, 6.8\pm1.1, -0.9\pm1.2)$ and $(-6.4\pm1.0, 2.1\pm1.4, -2.4\pm1.0)$ km~s$^{-1}$ for IC~348 and NGC~1333, respectively. Finally, our analysis of the kinematics of the stars has shown that there is no clear evidence of expansion, contraction, or rotational motions within the clusters. | \label{sec:intro} The Perseus molecular cloud represents an ideal target for studying the fundamental properties of young stars and their environment, since the complex is sufficiently nearby % that spatial scales down to $\sim$50 au are possible to reach with major observing facilities like ALMA and the VLA. Consisting of an elongated chain of dark clouds, Perseus spans over an area of $7^{\rm o}\times3^{\rm o}$ in the plane of the sky. The most prominent sub-structures are Barnard 5 (B5) and IC~348, at the eastern edge, and Barnard 1 (B1), NGC~1333, L1448, L1451 and L1455 at the western edge of the complex \citep[see e.g.][for a comprehensive review]{Bally2008}. Most of the young stars reside in IC~348 and NGC~1333, which contains about 480 and 200 objects, respectively, with ages of $1-3$~Myr \citep{Luhman2016}, mainly identified from optical and near-IR surveys. The protostellar content within Perseus, on the other hand, has been probed with observations at mid-IR (Spitzer; \citealt{Enoch2009}), far-IR (Herschel; \citealt{Sadavoy2014}), sub-mm (JCMT; \citealt{Sadavoy2010}) and radio (VLA; \citealt{Tobin2016,Tychoniec2018}) wavelengths. A total of 94 Class~0/I protostars and flat spectrum/Class~II objects are known to populate the entire cloud \citep{Tobin2016}. Multiple measurements of the distance to the individual clouds in Perseus have been performed in the past. These measurements suggest that there is a distance gradient across the cloud, with values in the range from 212 to 260~pc for the western component of the cloud \citep{Cernis1990,Lombardi2010,Hirota2008,Hirota2011,Schlafly2014} and 260--315~pc for the eastern component \citep{Cernis1993,Lombardi2010,Schlafly2014}. Direct measurement of distances via the trigonometric parallax have been obtained for only a few sources in these regions. Based on Very Long Baseline Interferometry (VLBI) observations of H$_2$O masers associated with two YSOs in NGC~1333 and L1448, \cite{Hirota2008,Hirota2011} found a distance consistent with 235~pc for both clouds. However, whether or not the gradient in the distance across the whole complex is significant remains inconclusive, since the distance uncertainties on individual lines of sight are large (typically $\sim10-20\%$ for photometric distances), and the number of sources with available direct distance measurements is small. In the past few years, we have used the Very Long Baseline Array (VLBA) to measure the trigonometric parallax of several tens of young stars in nearby star-forming regions \citep[][]{Ortiz2017Oph,Kounkel2017,Ortiz2017Ser,Galli2018} as part of the Gould's Belt Distances Survey (GOBELINS) project. Very Long Baseline Interferometry (VLBI) has the advantage of being able to detect highly embedded sources, where the extinction by dust obscures the optical light from the stellar objects. Given the high angular resolution provided by the VLBA and the fact that the interstellar material in these regions is transparent to radio waves, parallaxes with an accuracy of $1\%$ or better are possible for these kind of sources. In addition, parallaxes toward more than four hundred stars in Perseus with a limiting magnitude G=21~mag and parallax uncertainties $<0.7~$mas have become available during the second Gaia data release (DR2). With this highly accurate astrometric data, we can now investigate the depth of the molecular cloud and the three-dimensional motions of the young stars as well as the global properties of the kinematics of IC~348 and NGC~1333. We first describe the VLBA observations in Section \ref{sec:obs} and the fits to our data in Section \ref{sec:astro-vlba}. Section \ref{sec:gaia} presents the extraction of the astrometric solutions from the Gaia DR2 catalog. We then use both VLBA and Gaia data to investigate the structure of the Perseus cloud, which is discussed in Section \ref{sec:structure}. Sections \ref{sec:pm} and \ref{sec:vel} present the kinematics of a selected sample of cluster members in IC~348 and NGC~1333. Finally, our conclusions are given in Section \ref{sec:conclusions}. | \label{sec:conclusions} We have performed multi-epoch VLBA observations of three objects embedded in IC~348 and one object in NGC~1333, which are located near opposite ends within the Perseus molecular cloud. From the astrometric fits of this sample we derived a mean distance of $321\pm10$~pc to IC~348, representing the most reliable distance determination to the eastern of Perseus. This distance is consistent with the mean of Gaia DR2 parallaxes to a selected sample of known confirmed members of the cluster. The uncertainty on the mean of Gaia parallaxes is, however, 2.6 times larger than the VLBA uncertainty. The source detected with the VLBA in NGC~1333 is a close binary system, for which we derive preliminary orbital parameters. Unfortunately, the VLBA data is not enough to provide a reliable distance for this specific source, and consequently, for the NGC~1333 cluster. Gaia parallaxes, on the other hand, yield a mean distance of $293\pm22$~pc. From these measurements, we conclude that the distance between the western and eastern edges of the clouds is about 30~pc in the direction of the line of sight. We use Gaia proper motions and radial velocities from the literature to derive the spatial velocities for a subset of cluster members. We derive the average spatial velocity vectors of IC~348 and NGC~1333, which are similar in magnitude and direction between them, but significantly different to the mean spatial motion of the Perseus OB2 association. We have estimated the expansion (or contraction) and rotation velocities of each cluster and found no clear evidence of such organized motions. | 18 | 8 | 1808.03499 |
1808 | 1808.06796_arXiv.txt | We discuss the spectral and timing properties of the accreting millisecond X-ray pulsar \swiftj{} observed by \xmm{}, \nicer{} and \nustar{} during the X-ray outburst occurred in April 2018. The spectral properties of the source are consistent with a hard state dominated at high energies by a non-thermal power-law component with a cut-off at $\sim70$ keV. No evidence of iron emission lines or reflection humps has been found. From the coherent timing analysis of the pulse profiles, we derived an updated set of orbital ephemerides. Combining the parameters measured from the three outbursts shown by the source in the last $\sim11$ years, we investigated the secular evolution of the spin frequency and the orbital period. We estimated a neutron magnetic field of $3.1\times 10^{8}\,\,\, \textrm{G}<B_{PC}<4.5\times 10^{8}\,\,\, \textrm{G}$ and measured an orbital period derivative of $-4.1\times 10^{-12}$ s s$^{-1}$ $<\dot{P}_{orb}<7.1\times 10^{-12}$ s s$^{-1}$. We also studied the energy dependence of the pulse profile by characterising the behaviour of the pulse fractional amplitude in the energy range 0.3--80 keV. These results are compared with those obtained from the previous outbursts of \swiftj{} and other previously known accreting millisecond X-ray pulsars. | \swiftj{} is a low-mass X-ray binary discovered on 2007 June 7 during an X-ray outburst observed by the \swiftbat{}. Follow-up observations carried out with the \swiftxrt{} and the {\it Rossi X-ray Timing Explorer} (\rxte{}) provided the localisation of the source with an arcsec accuracy and led to the discovery of pulsations at a frequency of $\sim182$~Hz, classifying the source as an accreting millisecond X-ray pulsar \citep[AMXP, see e.g.][for a review]{Patruno12b}, in a 54.7~minutes orbit \citep{Krimm07}. A second outburst was recorded in July 2009 and the result of the observational campaign carried out with \swift{} and \rxte{} was reported in \citet[][hereafter P10]{Patruno10b}. In both occasions, the source displayed a spectral energy distribution compatible with the so-called ``island/extreme island state'' of an atoll source \citep[see e.g.][and reference therein]{Hasinger1989aa} and reasonably well described by a model comprising a power-law with a photon index of $\Gamma$=1.8-2.0 with no high-energy cut-off and a black-body component with a temperature of $kT$=0.4-0.7 keV \citep{Linares08}. Based on the upper limits derived on the spin-down torque, the neutron star magnetic field was constrained in a range compatible with values expected for an AMXP and observed from other sources of this class \citep[0.4$\times$10$^8$~G$<$B$<$9$\times$10$^8$~G;][]{Patruno10a}. The source was discovered to undergo a new outburst by INTEGRAL on 2018 April 1 \citep{mereminskiy18}. The event was confirmed by \swiftbat{}, and follow-up observations provided the detection of pulsations at the known spin period of the source and a preliminary description of its broad-band X-ray spectrum \citep{krimm18,bult18a,bult18b,cha18,mazzola18,kuiper18, Bult2018c}. In this work, we carried out spectral and coherent timing analysis of the 2018 outburst of \swiftj{}, using \inte{}, \xmm{}, \nustar{} and \nicer{} observations of the source. We updated the source ephemerides and investigated the orbital period evolution over a baseline of almost 11 years by combining the current results with those reported from previous outbursts. We also discuss the broad-band spectral properties of \swiftj{}. \section[]{Observations and data reduction} \subsection{XMM-Newton} \label{sec:XMM} \xmm{} observed \swiftj{} on 2018 April 8 (Obs.ID. 0830190401) for a total exposure time of $\sim$~66 ks. During the observation, the EPIC-pn (hereafter PN) camera was operated in {\sc timing} mode and {\sc burst} mode for $\sim$ 49 ks and $\sim$ 10 ks, respectively. The RGS instrument observed in spectroscopy mode during the entire observation, while the EPIC-MOS1 and EPIC-MOS2 were operated in {\sc full frame} and {\sc timing} mode, respectively. To perform spectral and timing analysis of the source we focused on the PN and MOS2 data (the limited statistics and time resolution of the MOS1 data did not provide a significant improvement in any of the results presented here and in the following sections). These were processed using the Science Analysis Software (SAS) v. 16.0.0 with the up-to-date calibration files and the RDPHA calibrations \citep[see e.g.][]{Pintore15a}. We filtered events within the energy range 0.3-10.0 keV, retaining single and double pixel events only (\textsc{pattern$\leq$4}). We extracted the source events for the PN and MOS2 using RAWX=[29:45] and RAWX=[285:325], respectively. We filtered background events for the PN selecting RAWX=[3:5] and we checked that the selected background was not contaminated by the emission from the source. For the MOS2, we extracted the background using an empty circular region of radius 150'' from the MOS1 dataset. The mean PN and MOS2 observed count rates during the observation were $\sim22$ cts/s and $\sim4.5$ cts/s, characterised by a slow decreasing trend. The background mean count rate in the PN selected RAWX range is of the order of $\sim0.5$ cts/s (0.3-10.0 keV). Thermonuclear (Type-I) X-ray burst episodes \citep[see e.g.][for a review]{Strohmayer2010aa} were not detected in the EPIC data. We extracted RGS data with standard procedures. We checked that the RGS1 and RGS2 spectra were consistent and then we merged them with the task {\sc rgscombine}. \noindent Fig.~\ref{fig:lc} shows the monitoring light curve of the 2018 outburst of \swiftj{} as seen by \swiftxrt{} (black points) and obtained from the on-line \swiftxrt{} data products tool \citep{Evans2009a}. The green star represents the beginning of the \xmm{} observation taken few days after the outburst peak. To perform the timing analysis, we reported the PN photon arrival times to the Solar System barycentre by using the \textsc{barycen} tool (DE-405 solar system ephemeris). We applied the best available X-ray position of the source (reported in Tab.~\ref{tab:solution}) estimated performing astrometric analysis to the available \swiftxrt{} observation of the source \citep{Evans2009a}. The new source coordinates are compatible, to within the associated uncertainties, with the position reported by \citet{Krimm07}. \begin{figure} \centering \includegraphics[width=0.48\textwidth]{lc} \caption{\swiftxrt{} light curve (black points) of the 2018 outburst of the accreting millisecond X-ray pulsar \swiftj{}. Data are shown from MJD = 58210 (2018-04-02). Upper limits on the \swiftxrt{} count rate are shown with empty triangles. The green star, red squares, blue diamonds, and purple square represent the starting times of the \xmm{}, \nustar{}, \nicer{} and \inte{} observations, respectively.} \label{fig:lc} \end{figure} \subsection{\nustar{}} \nustar{} observed \swiftj{} twice during its 2018 outburst. The first observation (Obs.ID. 90402313002) started at 08:31 \texttt{UT} on 2018 April 8 for an elapsed time of $\sim 85$ks, resulting in a total effective exposure time $\sim43$ks. The second observation (Obs.ID. 90402313004) started at 02:56 \texttt{UT} on 2018 April 14 for an elapsed time of $\sim 125$ks, corresponding to a total effective exposure time of $\sim68$ks. The epochs at which \nustar{} observed are shown as red squares in Fig.~\ref{fig:lc}. We screened and filtered the events with the \nustar{} data analysis software (\textsc{nustardas}) version 1.5.1. We extracted the source events from the FPMA and FPMB focal planes within a circular region of radius 90$''$ centered on the source position. A similarly extended region shifted to a position not contaminated by the source emission was used for the extraction of the background events. For each of the two observations, we obtained the background-subtracted light curves. These are characterised by an average count rate per FPM of $\sim10$ and $\sim 0.001$ counts/s, respectively. During the second observation the source was not significantly detected, and we thus discard these data for the remaining analysis. We corrected the photon arrival times for the motion of the Earth-spacecraft system with respect to the Solar System barycentre with the {\sc barycorr} tools (using DE-405 solar system ephemeris), in analogy to what was done for the \xmm{} data. \subsection{\nicer{}} \nicer{} observed \swiftj{} seven times during its 2018 outburst (see Tab.~\ref{tab:obs} for more details). We extracted events across the 0.2-12 keV band applying standard screening criteria using the HEASOFT version 6.24 and NICERDAS version 4.0. Observations 105023105/6/7 showed the presence of high-energy background features. To further proceed with the timing analysis we excluded (when available) data 50 s before the raise and 100 s after the decay of the flares. We then barycentered the NICER photon arrival times with the {\sc barycorr} tool using DE-405 Solar system ephemeris and adopting the source coordinates reported in Tab.~\ref{tab:solution}. \begin{table} \begin{center} \begin{tabular}{ | c | c | c | c} \hline Instrument & Obs.ID. & Date & Exp. (s) \\ & (revolution) & & \\ \hline \hline \xmm{}-PN & 0830190401 & 2018-04-08 & 49072 \\ \hline \multirow{2}{*}{\nustar{}} & 90402313002 & 2018-04-08 & 43457 \\ & 90402313004 & 2018-04-14 & 65763 \\ \hline \inte{} & (1939) & 2018-04-07 & 85000 \\ \hline \multirow{7}{*}{\nicer{}} & 1050230101 & 2018-04-03 & 6716 \\ & 1050230102 & 2018-04-04 & 6424 \\ & 1050230103 & 2018-04-07 & 2201 \\ & 1050230104 & 2018-04-08 & 9490 \\ & 1050230105 & 2018-04-09 & 3861 \\ & 1050230106 & 2018-04-10 & 6141 \\ & 1050230107 & 2018-04-11 & 4470 \\ \hline \end{tabular} \caption{Log of the observations of \swiftj{} used to perform the spectral and timing analysis.} \label{tab:obs} \end{center} \end{table} \subsection{INTEGRAL} \label{sec:integral} \swiftj{} was observed with \inte{} \citep{wink} from 2018 April 7 at 18:58 to 2018 April 8 at 19:56 (UTC), during the satellite revolution 1939. We analysed all data by using version 10.2 of the Off-line Scientific Analysis software (OSA) distributed by the ISDC \citep{courvoisier03}. The \inte{} observations are divided into science windows (SCWs), i.e. pointings with typical durations of $\sim$2-3~ks. We analysed a total of 25 SCWs in which the source was located to within an off-axis angle of 3.5~deg from the center of the JEM-X \citep{lund03} field of view (FoV) and within an off-axis angle of 12~deg from the center of the IBIS/ISGRI \citep{ubertini03,lebrun03} FoV. These choices allowed us to minimise the instruments calibration uncertainties\footnote{http://www.isdc.unige.ch/integral/analysis}. We extracted first the IBIS/ISGRI and JEM-X mosaics. \swiftj{} was detected in the IBIS/ISGRI 20-40~keV and 40-80~keV mosaics at a significance of $20\sigma$ and $13\sigma$, respectively. The corresponding fluxes estimated from the mosaics were 15.3$\pm$0.8~mCrab (roughly 1.2$\times$10$^{-10}$~erg~cm$^{-2}$s$^{-1}$) and 9.5$\pm$0.8~mCrab (roughly 7$\times$10$^{-11}$~erg~cm$^{-2}$s$^{-1}$). The source was relatively faint for JEM-X and detected at $11\sigma$ in the 3-10~keV mosaic obtained by combining all JEM-X data. The correspondingly estimated flux was 26$\pm$3~mCrab (roughly 4.0$\times$10$^{-10}$~erg~cm$^{-2}$s$^{-1}$). We extracted the JEM-X light curves of the source with a bin time of 2~s to search for type-I X-ray bursts, but no significant detection was found. | We reported on the spectral and timing properties of the 2018 outburst of the AMXP \swiftj{} observed with \inte{}, \xmm{}, \nustar{} and \nicer{}. From the phase-connected timing analysis of the \nicer{} and \xmm{} observations, we obtained an updated set of the source ephemerides, compatible within the errors with those obtained from the \nustar{} dataset. Owing to the multiple observations performed during the source outburst, we obtained, for the first time since the decommission of \rxte{}, a reliable constraint on the spin frequency derivative (|$\dot{\nu}|<1.4\times 10^{-12}$ Hz/s) of an AMXP during the accretion state. Combing the timing properties from the previous two outbursts, we estimated a secular spin-down frequency derivative $\dot{\nu}_{sd}=4.8(6)\times 10^{-16}$ Hz/s, compatible with a magnetic field (at the polar caps) of $3.1\times 10^{8}\,\,\, \textrm{G}<B_{PC}<4.5\times 10^{8}\,\,\, \textrm{G}$. Furthermore, we obtained a secular orbital period derivative in the range $-4.1\times 10^{-12}$ s/s $<\dot{P}_{orb}<7.1\times 10^{-12}$ s/s (95\% confidence level), suggesting that more outbursts are required to further constrain the orbital evolution of the system. We also investigated the pulsation spectral energy distribution of \swiftj{} in the energy range 0.3--10 keV and 3--80 keV, using the \xmm{} and \nustar{} datasets, respectively. The pulse fractional amplitude trend shown by the fundamental and second harmonic components present similarities with those reported for other AMXPs likely suggesting a Comptonisation origin. Finally, we found that the broad-band (3--90 keV) energy spectrum of \swiftj{} observed during its 2018 outburst is well described by an absorbed cut-off power law plus a soft thermal component. A photon index of $\sim$1.5 combined with a cut-off at $\sim$ 70 keV strongly suggest that the source was observed in a hard state. Contrary to previous outbursts, we detected no significant reflection features, with a constraining upper limit on the iron line equivalent width ($\sim$ 5 eV). | 18 | 8 | 1808.06796 |
1808 | 1808.01349_arXiv.txt | The study of B-mode polarization in the cosmic microwave background (CMB) has now reached a critical and scientifically exciting phase (see~[\citenum{2016Kamionkowski}] for a recent review of the theoretical and experimental status of the field). In the past five years, the gravitational lensing B-mode signal has been independently detected directly by four different instruments [\citenum{2017PB, 2017ACT, 2015SPT,2016BK}]. At the same time, measurements of B-mode polarization at large angular scales are becoming more sensitive and covering a wider frequency range, with the goal of distinguishing a primordial gravitational wave signal from foreground sources. Measuring the very weak B-mode CMB polarization signal is extremely challenging, and requires an experimental methodology that can reach the necessary sensitivities while simultaneously controlling systematic errors from the instrument to the stringent levels required for the measurement. Success on both of these fronts is intimately tied to the quality of the atmosphere above the site. In the millimeter and sub-millimeter atmospheric windows, water vapor is the primary cause of atmospheric opacity. To minimize excess noise caused by an opaque atmosphere, ground-based CMB telescopes observe in discrete spectral windows, which are bracketed by strong oxygen and water vapor emission/absorption lines. Even within these windows, the water vapor continuum absorbs incoming signal and emits thermally, contributing to noise and reducing the sensitivity of CMB observations. Furthermore, unlike the ``dry'' components of the atmosphere, water vapor is in general poorly mixed, with concentrations that vary dramatically in space and time. CMB telescopes typically rely on fast scanning of the sky to map the relevant spatial scales of the CMB signal onto temporal scales at which their detectors are stable. The time- and space-variable precipitable water vapor (PWV) column along the telescope line of sight causes additional brightness fluctuations in single-dish telescopes (and phase noise/decoherence in interferometers) and limits the spatial scales detectable without relying on other modulation techniques. The best sites are therefore high-elevation locations with exceptionally dry air. The most competitive ground-based CMB observations to date have taken place at southern hemisphere sites such as the South Pole, Antarctica (the \bicep/\keckarray~program [\citenum{bkii}], the South Pole Telescope (SPT-3G) [\citenum{spt3g}]) and the Atacama Desert in northern Chile (POLARBEAR[\citenum{polarbear2017}], ACTPol[\citenum{actpol2017}], CLASS[\citenum{class2016}], ABS[\citenum{abs2018}]). At the South Pole, the noise properties in CMB maps are well understood due to years of observations from \bicepone, \biceptwo, \bicepthree, \keckarray, and SPT. However, no instruments of this scale are located in the northern hemisphere, and northern sites remain largely untested. The next-generation CMB-S4 experiment~[\citenum{cmbs4}] plans to deploy receivers at the South Pole and Atacama, and is considering locations in Greenland and possibly Tibet. \begin{figure}[ht] \begin{center} \begin{tabular}{c} \includegraphics[height=8cm]{merra.png} \end{tabular} \end{center} \caption[] {Full map of the 10-year average (2006-2016) precipitable water vapor (PWV) estimates from the NASA MERRA-2 reanalysis [\citenum{merra_dataset}]. The four highlighted sites (South Pole in Antarctica, Chajnantor Plateau in the Atacama desert in northern Chile, Summit Station on the polar plateau in Greenland, and the Ali Observatory in Tibet) are all located in the driest zones on Earth. Note that $1\,kg\,m{^{-2}}$ is equal to 1~mm precipitable water vapor.\label{fig:merra}} \end{figure} Figure~\ref{fig:merra} shows the 10-year average (2006-2016) Earth map PWV column from the NASA Modern-Era Retrospective Analysis Version 2 (MERRA-2). The MERRA-2 reanalysis assimilates a wide range of measurements from satellites, radiosondes, and surface weather stations into a coupled model encompassing atmospheric and surface properties, processes, and dynamics [\citenum{merra_dataset, gelaro2017}]. We have highlighted the locations of interest for our study, all of which have very low PWV. Atmospheric fluctuations can impact observations even in locations with low total PWV. The effect of atmospheric instability is most obvious in the total-power timestreams of a scanning telescope. The sky variations appear as scan-synchronous fluctuations whose amplitudes grow or shrink as the telescope scans across different parcels of water vapor with more or less PWV. The total power depends strongly on the properties of the atmosphere in a given azimuthal direction, which can vary on a timescale of minutes. Although the emission of the atmosphere is mostly unpolarized (so these fluctuations are highly correlated across detectors and within detector pairs), all polarimeters have only a finite degree of common-mode rejection of unpolarized fluctuations. For any given observing strategy, the amplitude and spectrum of the atmospheric fluctuations dictate the performance and the required level of modulation (observing modulation or software filtering) of the experiment. The effects of water vapor on CMB observations can be mitigated by carefully selecting sites with characteristically low total PWV. Total PWV and its average variation across days and seasons have long been used as a metric for assessing site quality [\citenum{radford_peterson_2016, cortes2016, matsushita2017, sarazin2013, tremblin2011, rose2005}]. Ground-based instruments for recording average PWV are commercially available, and MERRA-2 [\citenum{gelaro2017}] and other reanalyses are now capable of providing accurate site-specific climatologies with temporal resolution on the timescale of hours. For CMB measurements (as for interferometry), however, what matters most are the spatial and temporal fluctuations of the water vapor on even shorter timescales. There is no substitute for direct measurement of these fluctuations. The Atacama Large Millimeter/submillimeter Array (ALMA) has led the way in developing sensitive co-pointed water vapor radiometers (WVRs) on each of their 54 12-meter antennas [\citenum{alma_phasecorr_2013}]. The desire to better assess and compare atmospheric fluctuations from site to site [\citenum{layhalverson2000,bussmann2005,sayers2010}] has spawned an interest in developing a small, ultra-precise, scanning-mode WVR that would be appropriate for predicting sky noise for degree-scale CMB experiments. We have developed and built such instruments, aiming to make datasets of long-term measurements of atmospheric noise and stability. Our goal is to directly compare leading CMB sites using the same observing strategy with identical instrumentation and analysis. Coordinating measurements across different sites in this way is crucial to directly comparing possible sites for future CMB telescopes. We will also use the understanding of noise levels in CMB maps at the South Pole to translate measurements of sky stability to a prediction for the eventual levels of noise that will remain in CMB maps produced at other sites. Finally, we will use this comparison of sites to comment on the suitability of potential northern hemisphere locations (Summit Station, Greenland and Ali, Tibet) as a complement to the current leading sites (South Pole, Antarctica and Atacama Desert, Chile) for ground-based CMB observation. | We have developed two stand-alone, weather-hardened, high-precision scanning WVRs that have taken measurements in Greenland and at the South Pole. We plan to extend these measurements to various locations in Atacama and Tibet. These instruments provide a continuous year-round record of water vapor fluctuations to 1~$\mu$m PWV precision, from sub-degree to dipole scales and on timescales $\ge$1~s. In addition, we intend to use these measurements as absolute zenith temperature calibration for co-located CMB experiments. The goal of this program is to collect and publish one-to-one data from different sites that can inform all groups performing or proposing CMB experiments. It will be the first time there are identical long-term sky noise/stability measurements made to directly compare leading sites. The next-generation, ground-based experiment CMB-S4 is planning to deploy telescopes in both Chile and at the South Pole, and possibly in the Northern Hemisphere, where there are no major current CMB experiments and site quality for CMB observation is largely untested. Our campaign will help compare the expected performance of telescopes at different sites, including northern sites. Correlating the distributions of spatial and temporal water vapor fluctuations with the quality of co-measured CMB data at well-characterized sites (such as the South Pole with the \keckarray, \bicepthree, and \biceptng~in the future) gives us a new portable tool to quantify the quality of proposed sites for future large scale CMB observations. Maintaining this common instrumental and analysis framework is essential to the usefulness of this site characterization data. | 18 | 8 | 1808.01349 |
|
1808 | 1808.07335_arXiv.txt | In the past few years, several independent collaborations have presented cosmological constraints from tomographic cosmic shear analyses. These analyses differ in many aspects: the datasets, the shear and photometric redshift estimation algorithms, the theory model assumptions, and the inference pipelines. To assess the robustness of the existing cosmic shear results, we present in this paper a unified analysis of four of the recent cosmic shear surveys: the Deep Lens Survey (DLS), the Canada-France-Hawaii Telescope Lensing Survey (CFHTLenS), the Science Verification data from the Dark Energy Survey (DES-SV), and the 450 deg$^{2}$ release of the Kilo-Degree Survey (KiDS-450). By using a unified pipeline, we show how the cosmological constraints are sensitive to the various details of the pipeline. We identify several analysis choices that can shift the cosmological constraints by a significant fraction of the uncertainties. For our fiducial analysis choice, considering a Gaussian covariance, conservative scale cuts, assuming no baryonic feedback contamination, identical cosmological parameter priors and intrinsic alignment treatments, we find the constraints (mean, 16\% and 84\% confidence intervals) on the parameter $S_{8}\equiv \sigma_{8}(\Omega_{\rm m}/0.3)^{0.5}$ to be $S_{8}=0.94_{-0.045}^{+0.046}$ (DLS), $0.66_{-0.071}^{+0.070}$ (CFHTLenS), $0.84_{-0.061}^{+0.062}$ (DES-SV) and $0.76_{-0.049}^{+0.048}$ (KiDS-450). From the goodness-of-fit and the Bayesian evidence ratio, we determine that amongst the four surveys, the two more recent surveys, DES-SV and KiDS-450, have acceptable goodness-of-fit and are consistent with each other. The combined constraints are $S_{8}=0.79^{+0.042}_{-0.041}$, which is in good agreement with the first year of DES cosmic shear results and recent CMB constraints from the \textit{Planck} satellite. \\ | The large-scale structure of the Universe bends the light rays emitted from distant galaxies according to General Relativity \citep{Einstein1936}. This effect, known as weak (gravitational) lensing, introduces coherent distortions in galaxy shapes, which carry information of the cosmic composition and history. One of the most common statistics used to extract this information is \textit{cosmic shear}, as inferred by the two-point correlation function of galaxy shapes $\xi_{\pm}(\theta)$ \citep{Bartelmann2001}. Assuming the flat-sky approximation, these two-point functions are connected to the lensing power spectrum $C(\ell)$ via \begin{equation} \xi^{ij}_{\pm}(\theta) = \frac{1}{2\pi}\int d\ell \, \ell J_{0/4}(\theta \ell) \, C^{ij}(\ell), \label{eq:xipm} \end{equation} where $J_{0/4}$ is the 0th/4th-order Bessel functions of the first kind. The $i$ and $j$ indices specify the two samples of galaxies (or in the case of $i=j$, the galaxy sample) from which the correlation function is calculated. Usually these samples are defined by a certain redshift selection. Under the Limber approximation \citep{Limber1953,Loverde2008} and in a spatially flat universe\footnote{For a non-flat universe, one would replace $\chi$ by $f_{K}(\chi)$ in the following equations, where $K$ is the universe's curvature, $f_{K}(\chi)=K^{-1/2}\sin(K^{1/2}\chi)$ for $K>0$ and $f_{K}(\chi)=(-K)^{-1/2}\sinh((-K)^{1/2}\chi)$ for $K<0$.}, the lensing power spectrum encodes cosmological information through \begin{equation} C^{ij}(\ell) = \int_{0}^{\chi_{H}} d\chi \frac{q^{i}(\chi)q^{j}(\chi)}{\chi^2} P_{NL}\left( \frac{\ell + 1/2}{\chi}, \chi \right), \end{equation} where $\chi$ is the radial comoving distance, $\chi_{H}$ is the distance to the horizon, $P_{NL}$ is the nonlinear matter power spectrum, and $q(\chi)$ is the lensing efficiency defined via \begin{equation} q^{i}(\chi) = \frac{3}{2} \Omega_{\rm m} \left( \frac{H_{0}}{c}\right)^{2} \frac{\chi}{a(\chi)} \int_{\chi}^{\chi_{H}}d\chi ' n^{i}(\chi') \frac{dz}{d\chi'} \frac{\chi' - \chi}{\chi'}, \label{eq:lensing_efficiency} \end{equation} where $\Omega_{\rm m}$ is the matter density today, $H_{0}$ is the Hubble parameter today, $a$ is the scale factor, and $n^{i}(\chi)$ is the redshift distribution of the galaxy sample $i$. Since the first detection of cosmic shear in \cite{Bacon2000,Kaiser2000,Wittman2000,Schneider2002}, the field has seen a rapid growth. In particular, a number of large surveys have delivered cosmic shear results with competitive cosmological constraints in the past few years \citep{Heymans2013, Becker2015,Jee2016,Joudaki2017,Troxel2017,Hildebrandt2017,DES2017}, while ongoing and future surveys will deliver data in much larger volumes and better quality [e.g. the Dark Energy Survey \citep[DES,][]{Flaugher2005}, the Hyper SuprimeCam Survey \citep[HSC,][]{Aihara2017}, the Kilo-Degree Survey \citep[KiDS,][]{deJong2015} and the Large Synoptic Survey Telescope \citep[LSST,][]{Ivezic2008,Abell2009}]. One of the surprises that has emerged in the past couple of years is that there seems to be a modest level of discordance between different cosmological probes \citep{MacCrann2015,Freedman2017,Raveri2018}. Even though in many of these cases, the level of tension between the different probes still needs to be quantified more rigorously, one consequence has been that the cosmology community has started to more carefully scrutinize how the datasets are analyzed. This is especially important as we expect the statistical power of the datasets to be orders of magnitude better in the near future. If there is indeed a tension between the different probes, it could point to an exciting new direction where the simple $\Lambda$CDM cosmology cannot explain all the observables and new physics is needed. A variety of studies have been carried out to understand systematic effects in weak lensing measurements. This includes systematics from the instrument and the environment, from modeling the point-spread function (PSF) and measuring galaxy shapes, from estimating the redshift of each galaxy, from the theoretical modeling, and many more \citep[see][and references therein for a comprehensive list of studies]{Mandelbaum2017}. In this work, we focus on understanding the steps between the shear catalog and cosmological constraints: measuring the shear two-point correlation function [\Eref{eq:xipm}], estimating the covariance, modeling of the signal, and inferring cosmological parameters. We build a modular and robust pipeline using the \textsc{Pegasus} workflow engine \citep{Deelman2015} to analyze the datasets in a streamlined and transparent fashion -- this pipeline will serve as the first step towards building up cosmological analysis pipelines for the LSST Dark Energy Survey Collaboration (DESC). In this paper we apply the pipeline to four publicly available datasets that are precursors to ongoing and future cosmic shear surveys: the Deep Lens Survey \citep[DLS,][]{Jee2016}, the Canada-France-Hawaii Telescope Lensing Survey \citep[CFHTLenS,][]{Joudaki2017}, the Science Verification data from the DES \citep[DES-SV,][]{Abbott2015}, and the 450 deg$^{2}$ release of the KiDS \citep[KiDS-450,][]{Hildebrandt2017}. All four surveys were carried out fairly recently and have comparable statistical power, so a uniform pipeline is a powerful way to identify any discrepancies and to understand their origin. A detailed look at the consistency between the four datasets can also inform us about potential systematic issues in the processing that produces the catalogs from which our pipeline begins. It is, however, not the scope of this paper to investigate these issues upstream to our pipeline, where a thorough pixel-level study for each survey may be required. The paper is organized as follows. In \Sref{sec:data} we describe the details of the four datasets used in this work. In \Sref{sec:pipeline} we describe the pipeline that is used to process the data. We then outline in \Sref{sec:comparison_framework} the framework in which we compare the datasets and the elements in the pipeline that are allowed to vary. Our results are shown and discussed in \Sref{sec:results} and we conclude in \Sref{sec:conclusion}. \begin{figure*} \includegraphics[width=1.6\columnwidth]{figures/nofz.pdf} \caption{Estimation of the the tomographic redshift distributions used in the four cosmic shear analyses. For DLS, CFHTLenS and DES-SV, stacked photometric redshift probability distribution functions (PDFs) were used; for KiDS-450, the redshift distribution of spectroscopic samples (weighted to match the source galaxies used for the cosmic shear analysis) were used. We see that DLS and CFHTLenS extend to higher redshift compared to DES-SV and KiDS-450.} \label{fig:nofz} \end{figure*} \begin{figure*} \includegraphics[width=1.6\columnwidth]{figures/footprint.pdf} \caption{The location and footprints of the four surveys analyzed in this paper. There is essentially no overlap between the footprints of the four surveys, except for a very small part of CFHTLenS and KiDS-450 at (RA, Dec) $\approx$ (130, -3) deg. We note that the projection in this plot does not reflect the relative area of the four surveys.} \label{fig:footprint} \end{figure*} | 18 | 8 | 1808.07335 |
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1808 | 1808.07103_arXiv.txt | The transverse peculiar velocities caused by the mass distribution of large-scale structure provide a test of the theoretical matter power spectrum and the cosmological parameters that contribute to its shape. Typically, the matter density distribution of the nearby Universe is measured through redshift or line-of-sight peculiar velocity surveys. However, both methods require model-dependent distance measures to place the galaxies or to differentiate peculiar velocity from the Hubble expansion. In this paper, we use the correlated proper motions of galaxy pairs from the VLBA Extragalactic Proper Motion Catalog to place limits on the transverse peculiar velocity of galaxy pairs with comoving separations $<1500$ Mpc without a reliance on precise distance measurements. The relative proper motions of galaxy pairs across the line of sight can be directly translated into relative peculiar velocities because no proper motion will occur in a homogeneous expansion. We place a $3\sigma$ limit on the relative proper motion of pairs with comoving separations $< 100$ Mpc of $-17.4 \ \mu$as yr$^{-1} < \dot{\theta} / \sin \theta < 19.8 \ \mu$as yr$^{-1}$. We also confirm that large-separation objects ($> 200$ Mpc) are consistent with pure Hubble expansion to within $\sim 5.3 \ \mu$as yr$^{-1}$ ($1 \sigma$). Finally, we predict that {\it Gaia} end-of-mission proper motions will be able to significantly detect the mass distribution of large-scale structure on length scales $< 25$ Mpc. This future detection will allow a test of the shape of the theoretical mass power spectrum without a reliance on precise distance measurements. | The modern standard model of cosmology is a spatially flat, expanding Universe dominated by cold dark matter and a cosmological constant. Small matter density perturbations in the primordial Universe drive the initial gravitational collapse of matter and dark matter into overdensities, which give rise to the formation of large-scale structure (LSS), clusters, and galaxies \citep[e.g.,][]{Blumenthaletal1984} through a hierarchical process. The overall process and cosmological parameters that govern the process are, in general, well constrained \citep[e.g.,][]{HuWhite1996a,HuWhite1996b,Huetal1997,Riessetal2001,Peirisetal2003,Moodleyetal2004,PlanckI2016}. However, the exact values of the parameters remain somewhat uncertain, including the baryon and neutrino density \citep[e.g.,][]{EisensteinHu1999, Eisensteinetal2005}, the Hubble constant \citep[e.g.,][]{Eisensteinetal2005,Riessetal2011,PlanckXIII2016,Riessetal2016,Zhangetal2017,Riessetal2018}, the tensor-to-scalar ratio of primordial gravitational waves \citep[e.g.,][]{PlanckXX2016}, and the spatial curvature \citep[e.g.,][]{Eisensteinetal2005,PlanckXIV2016}. The current distribution of LSS provides a means to test many of the less certain cosmological parameters by comparing the observed LSS to that predicted by cosmological simulations. Most commonly, maps of LSS are created using the sky distribution of visible galaxies with redshift as a proxy for distance \citep[e.g.,][]{deLapparentetal1986,Yorketal2000,Gottetal2005}. However, these maps rely on the assumption that light and visible matter trace the overall dark matter distribution. In general, it is reasonable to assume that galaxies form at the peaks of dark matter \citep{Kaiser1984}, but a bias model \citep[e.g.,][]{Bardeenetal1986,Coles1993,Fry1996,TegmarkPeebles1998} is still required to translate the observed LSS to the dark matter distributions generated by cosmological simulations. Galaxy line-of-sight peculiar velocities are an alternate means to track the dark matter distribution that does not require a translation between light and total mass because peculiar velocities of galaxies are directly caused by the matter density distribution. Line-of-sight peculiar velocities are obtained from the difference between the redshift and a redshift-independent distance. Velocity surveys can probe more distant structures than redshift surveys because objects' velocities can be influenced by distant objects that are outside the range of the survey \citep[e.g.,][]{Doumleretal2013,Tullyetal2014}. The drawback of line-of-sight velocity surveys to detect the dark matter distribution is that small uncertainties in a galaxy's distance can translate to large peculiar velocity uncertainties. Therefore, large samples of galaxy redshifts and distances are needed to statistically detect an average matter density distribution \citep{Tullyetal2014}. Both methods of mapping LSS produce similar matter density distributions \citep[e.g.,][]{Straussetal1992,Dekeletal1993,Kitauraetal2012,Courtoisetal2012}, which is a good confirmation of the standard model of LSS evolution through hierarchical growth and gravitational instability. However, both methods require redshifts or other model-dependent distance measures to either spatially place the galaxies or to translate spectroscopic line shifts into peculiar velocities. Therefore, another method that is independent of the ``distance ladder'' and other distance models is needed to provide an independent test of the model of LSS evolution. Extragalactic proper motions can be used to test models of LSS evolution without a reliance on the ``distance ladder.'' Like line-of-sight peculiar velocities, transverse peculiar velocities directly probe the matter density distribution and do not require any assumptions about the relative abundances of visible matter and dark matter (or how light traces mass). As the Universe expands, both line-of-sight and peculiar velocities will contain peculiar motion from gravitational interactions. However, line-of-sight velocities require an independent distance measure to differentiate between Hubble expansion and peculiar velocity. On the other hand, orthogonal velocities across the line of sight (observed as proper motions) are separable from the Hubble expansion because no proper motion will occur in a homogeneous expansion \citep{Nusseretal2012,Darling2013}. In this paper, we use the relative proper motion of pairs of extragalactic objects from our VLBA Extragalactic Proper Motion Catalog \citep{TruebenbachDarling2017} to constrain transverse peculiar velocities induced by the mass distribution of LSS. We describe the expected signal in Section \ref{pairwise_theory} and calculate the signal for pairs with separations $< 1500$ Mpc is Section \ref{pairwise_measurement}. In Section \ref{limit_section}, we compare our measurement to that predicted by a transverse peculiar velocity two-point correlation statistic calculated from a $z=0$ matter power spectrum \citep{Darling2018}. Finally, we predict {\it Gaia}'s ability to measure the relative convergence of galaxies on small scales (Sec. \ref{future_LSS}) and discuss future improvements to our catalog (Sec. \ref{conclusions}). In this paper we assume $H_0 = 70$ km s$^{-1}$ Mpc$^{-1}$ and a flat cosmology with $\Omega_\Lambda = 0.73$ and $\Omega_M = 0.27$. | In this paper, we measured the relative proper motion of all pairs of objects from the VLBA Extragalactic Proper Motion Catalog \citep{TruebenbachDarling2017} with comoving separations less than 1500 Mpc. We found that large-separation pairs with separations $> 200$ Mpc are consistent with the null signal expected for Hubble expansion to within $\sim 5.7 \ \mu$as yr$^{-1}$ ($1\sigma$). This result demonstrates that our method of measuring the relative proper motion of pairs is able to extract small proper motion signals to high precision despite the large scatter of individual pairs due to the intrinsic proper motions from jets. We also found that we have too few close-separation pairs ($<100$ Mpc) to statistically detect the expected convergence predicted by the peculiar transverse velocity two-point correlation function. We found a 3$\sigma$ limit on the convergence rate of close-separation pairs of $-17.4 \ \mu$as yr$^{-1} < \dot{\theta} / \sin \theta < 19.8 \ \mu$as yr$^{-1}$. Additionally, we estimated that we would need $\sim 30$ more uncorrelated pairs with separations $<50$ Mpc to make a 3$\sigma$ detection of this effect. There are several available avenues for expanding our sample of close-separation pairs. The best source for additional extragalactic proper motions is very long baseline interferometry because of its uniquely high angular resolution. From the uncertainties in Figure \ref{xi_plot}, we predict that we would need $\sim 30$ additional pairs with comoving separations $<50$ Mpc to measure the expected signal from the mass distribution of LSS. However, in order to achieve the proper motion precisions in our catalog, astrometry was obtained for all quasars more than three times over at least a 10 year timespan. Therefore, addition of new close-separation pairs would be a long-term project that would require significant VLBI use. The Next Generation Very Large Array (ngVLA) provides another avenue for measuring new extragalactic proper motions. The ngVLA will provide ten times the effective collecting area and ten times longer baselines (300 km) than the current Karl Jansky Very Large Array \citep[JVLA;][]{Carillietal2015}. The larger collecting area will enable faster astrometry of fainter sources than is currently possible with the VLBA. In one hour of observation time, the ngVLA will detect objects with flux densities of a few tenths of a mJy or brighter \citep{Butleretal2018}. If additional ngVLA antennae are installed at VLBA sites, the collecting power of the ngVLA could be combined with the astrometric precision of the VLBA to quickly create a large sample of close-separation pairs. For example, if 10,000 objects were monitored on a yearly basis to obtain $10 \ \mu$as yr$^{-1}$ astrometry per object, this would enable global detection of $0.1 \ \mu$as yr$^{-1}$ correlated signals at 5$\sigma$ significance \citep{Boweretal2015}. At this precision, we could constrain isotropy to 0.1\% of $H_0$ \citep{Boweretal2015}. If objects are selected to be in close-separation pairs, a new proper motion catalog with the ngVLA and VLBA would contain more than enough objects to detect the mass distribution of LSS; we predict that we need only $\sim 30$ pairs with separations $<50$ Mpc for a $3\sigma$ detection of this effect (Section \ref{future_LSS}). However, it is uncertain whether the ngVLA will include long baselines $>1000$ km. For observations at 3.6 cm with a maximum baseline of 1000 km, the ngVLA will achieve a spatial resolution of $\sim 7$ mas. If objects are observed yearly for 10 years, then we expect to detect proper motions $> 700 \ \mu$as yr$^{-1}$. Therefore, unless the ngVLA includes baselines on order of VLBA baselines or can be used in conjunction with existing VLBI networks, the ngVLA will not be an effective tool to expand our sample of close-separation pairs. With the limitations of current VLBA catalogs and ngVLA baselines, {\it Gaia} proper motions remain the best avenue for increasing our sample of close-separation pairs. As discussed in Section \ref{future_LSS}, {\it Gaia} will produce $\sim 10^6$ new extragalactic proper motions \citep{Robinetal2012} with astrometric precisions of $\sim 1.69$ mas \citep{Lindegrenetal2016}. Although {\it Gaia} proper motions are much less precise \citep[uncertainties $\sim 200 \ \mu$as yr$^{-1}$;][]{Paineetal2017} than those produced through VLBI observations \citep[tens of $\mu$as yr$^{-1}$;][]{TitovLambert2013, TruebenbachDarling2017}, the large number of additional proper motions will allow a statistical detection of correlated signals. Additionally, optical proper motions typically have a smaller contribution from intrinsic proper motions than those measured in the radio because the optical source more frequently traces the galaxy core, rather than tracing a jet driven by the active galactic nucleus \citep{Darlingetal2017}. In Section \ref{future_LSS}, we examined a subset of the {\it Gaia} DR1 catalog with 610,841 simulated extragalactic proper motions to estimate how well {\it Gaia} end-of-mission proper motions will be able to detect the relative velocity of close-separation pairs caused by the mass distribution of LSS. We found that the subset is able to significantly detect the predicted correlation of close-separation pairs. Therefore, although our current sample of close-separation pairs is too small to statistically detect the net convergence of gravitationally-interacting pairs, {\it Gaia} proper motions will be able to make a highly significant detection of the mass distribution of LSS on small scales. Comparison of the measured peculiar transverse two-point correlation statistic to that predicted by the theoretical matter power spectrum will allow a means to test the cosmological principles that contribute to the shape and magnitude of the power spectrum without a reliance on precise distance measurements. | 18 | 8 | 1808.07103 |
1808 | 1808.10769_arXiv.txt | We study the vacuum polarization tensor of QED (quantum electrodynamics) at high temperatures up to the two loop levels and its effect on the electromagnetic properties of a medium. One loop corrections to QED coupling vanish at low temperatures (T$\leq 10^{10}K$), but they play an important role at high temperature ( T$\geq 10^{10}$ K) to study the behavior of QED medium at these temperatures. At low temperatures ( $T \leq m_e$)higher order loops give a tiny correction due to the coupling of radiation with matter and an overlap of hot photon loop with cold fermion loop contributes to this effect. These higher loop contributions does not affect the convergence of perturbative series, and renormalizability of QED is guaranteed at temperatures around neutrino decoupling. We use the renormalization scheme of QED at finite temperature in real-time formalism to study the dynamically generated mass of photon indicating the plasmon production in such a medium. Temperature dependence of this QED plasma parameters is discussed. We explicitly show that this behavior of a thermal medium exists upto temperatures of a few MeV only. We compare the first order and second order effects upto the 4MeV temperature and demonstrate that the higher order contributions are smaller than the lower order contributions proving the renormalizability of the theory. The lowest order contributions are sufficiently smaller than the original value as well. | In quantum field theory, thermal background effects are incorporated through the radiative corrections. Renormalization of gauge theories [1-2] at finite temperature [3-11] requires the renormalization of gauge parameters of the corresponding theory. The propagation of particles and the electromagnetic properties of media are also known to modify in the framework of real-time formalism [7-9]. Masses of particles are shown to increase with temperature at the one-loop level [11-22], the two-loop level [23-26] and presumably to all loop levels [27]. At the higher-loop level, the loop integrals have a combination of cold and hot terms which appear due to the overlapping propagator terms in the matrix element. Higher loop calculations are too cumbersome to be performed analytically using the perturbation theory. Order by order cancellations of singularities [12] cannot be shown in the imaginary time formalism. Therefore, the effective potential approach at finite temperature [28] is used to study the overall effects. However, renormalizability can only be tested in detail when order-by-order cancellation of singularities [29] is shown in real-time formalism. \bigskip However, the gauge bosons acquire dynamically generated mass due to plasma screening effect [16-19], at the one-loop level. It helps to determine the changes in electromagnetic properties of a hot medium. In hot gauge theories $m_{e}$ is the electron mass and corresponds to $10^{10}$ K. The vacuum polarization tensor in order $\alpha$ (the QED coupling parameter) does not acquire any hot corrections from hot photons in the heat bath [16] because of the absence of self-interaction of photons in QED. The photon can only interact with the medium in the presence of electrons, which decouple at high temperatures ($\geq 1MeV or 10^{10} K$). Some of the QED parameters have already been calculated for such systems. The effective value of the electric charge increases due to the interaction of charge in the medium and consequently leads to modifications in electromagnetic properties of a medium due to the enhancement in QED coupling. These type of calculations are discussed up to the two loop level. These QED couplings at high density can change an ordinary fluid in a relativistic plasma and give a nonzero value of dynamically generated mass [16-19], which can be treated as a plasma screening mass and causes Debye shielding. We calculate the Debye shielding length of QED plasma as a function of temperature for a QED system which can be identified as QED plasma. \bigskip In this paper, we use the renormalization scheme of QED in real-time formalism to calculate the parameters of QED plasma in terms of thermally corrected renormalized values of electric charge, mass and wavefunction of electron [29]. These thermal contributions to the fundamental parameters of the theory help to determine the effective parameters of the theory under given statistical conditions. In the real-time formalism, all the calculations are done in the rest frame of the heat-bath to re-establish covariance at the cost of Lorentz invariance [7, 20]. The particle propagators include temperature dependent (hot) term in addition to the temperature independent (cold) term [5]. At the higher loop level, the loop integrals involve an overlap of hot and cold terms in particle propagators. This makes the situation cumbersome and getting rid of these singularities needs much more involved calculations. Sometimes $\delta (0)$ type pinch singularities appear in Minkowski space. Cancellation of singularities is not obvious, and this entire scheme works perfectly fine for $T<4m_e$ up to the two loop level [26-28]. \bigskip In the next section, we discuss the vacuum polarization tensor in different ranges of temperature up to the second order in $\alpha$. Section III is comprised of the calculation of the plasma parameters up to the two-loop levels. The plasma screening mass and plasma frequencies up to the second order in alpha are discussed in section IV. We conclude this paper by the discussions of application of these results to certain physical systems in section V. | We have previously noticed [16] that the low temperature (T$\leq m_e$) effects on the vacuum polarization tensor vanish because of the absence of hot electron background. The selfmass of electron is thermally corrected due to the radiation background. However, we have found that thermal corrections to vacuum polarization tensor are nonzero at the higher loop level even at low temperatures because the cold electron loops can overlap with the hot radiation loops and pick up tiny thermal corrections. \bigskip It is worth mentioning that the sizeable thermal corrections at high temperatures (T$\geq m_e$) are not always strong enough to create QED plasma. High temperatures indicate large kinetic energy and the interaction may still be weak enough to treat the system as an ideal gas. However, if the temperature is large and the volume is small, the mutual interaction between electrons is negligibly small as compared to thermal energy of particles. Therefore, a detailed study is required to understand the behavior of such systems which may have a plasma phase for a short time at some particular temperature. However, due to the exponential dependence on temperature, we can use high or low temperature limits for smaller changes due to the quadratic dependence in above equations (see Eqns.10-18). \bigskip Thermal corrections, which lead to the modifications in electromagnetic properties of a medium itself are expected to give larger contribution at higher loop levels. However, most of these terms are finite and the order by order cancellation of singularities [12] is observed through the addition of all the same order diagrams. However, the renormalization of the theory can only be proved if covariant hot integrals are evaluated before the cold divergent integrals on the mass shell. Once the hot loop energies are integrated out, the usual vacuum techniques of Feynman parameterization and dimensional regularization can be applied to get rid of vacuum singularities. It is also worth-mentioning that all the hot corrections give a similar $T^{2}$ dependence. The incorrect order of integration gives an increased order of T due to the overlap with the vacuum divergences. This unusual behaviour of hot integrals induces additional temperature dependent divergences due to the overlap of hot and cold terms. Whereas, the usual regularization techniques of vacuum theory like dimensional regularization would only be valid in a covariant framework. Higher order terms may then violate renormalizability with the inverted order of integration. \bigskip The presence of the statistical contribution of the photon propagator modifies the vacuum polarization and hence the electron charge which leads to changes (though small) in the electromagnetic properties of the hot medium even at low temperatures. Mass, wavefunction, and charge of electron are renormalized [29] in the presence of heat bath. These renormalized values give the dynamically generated mass of photon and its effective charge in such a background modifying dielectric constant and magnetic permeability of a hot medium[29]. The longitudinal component of vacuum polarization tensor $\Pi_L$ vanishes at low temperature as $p^{2}\longrightarrow 0$. However, the transverse component has an extremely small additional thermal contribution and Eq. (7) reduces to, \begin{equation} \Pi _{T}(p)=\frac{\alpha ^{2}T^{2}}{6}, \end{equation} It is worth-mentioning that in the real-time formalism, the propagator has two additive terms, the vacuum term and the temperature dependent hot term. Therefore, in the second order perturbation theory we get a purely hot term ($~\alpha^2 T^4/m^4$), purely cold terms (~$T^0$) and the overlapping hot and cold terms ($~\alpha^2 T^2/m^2$). These terms can only be obtained in this formalism at the two-loop level. So, the overall result comes out to be a combination of all of these terms. For example, Eq.(9) can be written as, \begin{equation} \alpha _{R}=\alpha (T=0)(1+ \frac{\alpha ^{2}T^{2}}{6m^{2}}+O(\frac{T^{4}}{m^{4}})). \end{equation} We restrict ourselves in the temperature range where the last term ($~\alpha^2 T^4/m^4$), is sufficiently smaller than the second term ($~\alpha^2 T^2/m^2$), and Masood's abc functions ($a_i(m\beta)$) can be evaluated for the extreme values. The leading order contribution at low T is based on the perturbative expansion in QED and the first term should be proportional to $T^2/m^2$ in this expansion. The damping factor $exp(-m/T)$ appears in the effective action when all the contributing terms are included simultaneously. We look at all these terms one by one and work for sufficiently small range of temperature (for a few MeV) to get the simple and approximate results to evaluate some physically measurable parameters. We ignore the magnetic field effect in this calculation. It is reasonable to demonstrate the renormalizability of QED at low temperature. High densities and magnetic field, contributing to the potential energy, change the scenario [30-34] and the plasmon effect may be studied at higher temperatures but higher order effects with high magnetic field are still to be investigated. \bigskip A comparison of the one loop and two loops contributions of the transverse component of the photon frequency is shown in Figure 3. It is clear that the second order contribution (broken line) is very small as compared to the first order (solid line) contribution. \bigskip \begin{figure}[!hth] \begin{center} \begin{tabular}{c} \mbox{\includegraphics[width=4in,angle=0] {3TwoLoopPlasma.jpg}} \end{tabular} \end{center} \caption{Comparison between the first order and second order contributions to the square of photon transverse frequency $\omega_T^2$.} \end{figure} A similar behavior is seen between the first order and second order contribution of the longitudinal component of the wavenumber and is shown in Figure 4. It is clear that the second order contribution (broken line) is significantly small as compared to the first order (solid line) contribution. It is evident from Figs, (3) and (4) that the wavenumber of photon gets a much larger thermal contribution than the frequency. However, everything remains smaller than 1 up to the temperatures of 4MeV or so, and this temperature is expected a little before nucleosynthesis in the early universe. \begin{figure}[!hth] \begin{center} \begin{tabular}{c} \mbox{\includegraphics[width=4in,angle=0] {4TwoLoopPlasma.jpg}} \end{tabular} \end{center} \caption{Comparison between the first order and second order contributions to the square of the longitudinal component of the wavenumber $\kappa_L^2$.} \end{figure} A comparison between Figs.3 and 4 shows that the thermal contributions to $\kappa _L^2$ are smaller than the $\omega _T^2$. We compare the two loop contributions of $\kappa _L$ and $\omega _T$ in Fig. 5. \begin{figure}[!hth] \begin{center} \begin{tabular}{c} \mbox{\includegraphics[width=4in,angle=0] {5TwoLoopPlasma.jpg}} \end{tabular} \end{center} \caption{Comparison between the first order and second order contributions to $\kappa _L$ and $\omega _T$.} \end{figure} It has been shown that the QED coupling behaves the same way in Fig.6. The one-loop thermal correction is much more dominant as compared to the second order contribution which shows the validity of these calculations to prove the renormalizability of the theory. Moreover it indicates the validity of the calculational scheme in this range. Fig.6 shows a comparison between two loop contributions with the corresponding second order contributions of the charge renormalization constant. However, the temperature dependent contribution is smaller than other parameters. \bigskip \begin{figure}[!hth] \begin{center} \begin{tabular}{c} \mbox{\includegraphics[width=4in,angle=0] {6TwoLoopPlasma.jpg}} \end{tabular} \end{center} \caption{Comparison between the first order and second order contributions.} \end{figure} Another interesting outcome of these calculations is Debye shielding in QED plasma, which is plotted as a function of temperature in Fig.7. A comparison between the first order and the second order in $\alpha$ terms clearly demonstrates (Fig.7) that the smaller values correspond to the first order term.It indicates the range of temperatures where plasma phase can exist in QED. \bigskip \begin{figure}[!hth] \begin{center} \begin{tabular}{c} \mbox{\includegraphics[width=6in,angle=0] {7TwoLoopPlasma.jpg}} \end{tabular} \end{center} \caption{Comparison between the first order and second order contributions.} \end{figure} In this paper, the first two orders and the second order contributions are compared to validate the renormalizability of QED between $m_e\leq T\leq10m_e$. For this purpose, we plotted the charge renormalization parameter of QED comparing the first two order contributions and the second order contribution only. The quadratic temperature dependence of QED parameters dominates in the proposed temperature range and ensures the validity of renormalizability of QED. The vacuum polarization components, squares of longitudinal and transverse components of photon frequency and momentum are all quadratic functions of temperature. However, the first order contributions indicate the existence of plasmon due to the dominant radiative self-coupling of photon in this region of temperature, which is found a little before nucleosynthesis and existed throughout the nucleosynthesis. Higher temperature behavior breaks this plasma shielding due to high kinetic energy and photons do not have the ability to couple with the medium at lower temperatures. Therefore the plasmon does not exist outside this range of temperature unless the magnetic field or high density contributions [18,33,34] are incorporated. \bigskip | 18 | 8 | 1808.10769 |
1808 | 1808.03121_arXiv.txt | A robust post processing technique is mandatory to analyse the coronagraphic high contrast imaging data. Angular Differential Imaging (ADI) and Principal Component Analysis (PCA) are the most used approaches to suppress the quasi-static structure in the Point Spread Function (PSF) in order to revealing planets at different separations from the host star. The focus of this work is to apply these two data reduction techniques to obtain the best limit detection for each coronagraphic setting that has been simulated for the SHARK-NIR, a coronagraphic camera that will be implemented at the Large Binocular Telescope (LBT). We investigated different seeing conditions ($0.4"-1"$) for stellar magnitude ranging from R=6 to R=14, with particular care in finding the best compromise between quasi-static speckle subtraction and planet detection. | \label{sec:intro} % SHARK is a coronagraphic camera proposed for Large Binocular Telescope in the framework of the “2014 Call for Proposals for Instrument Upgrades and New Instruments”~\cite{2014ebi..confP4.74F, 2015IJAsB..14..365F, 2016SPIE.9909E..31F, 2016SPIE.9911E..27V}. We are building this tool because first of all we hold an excellent adaptive optic (AO) performance, we are in the northern hemisphere with a strong scientific case and the purpose is going on sky in a very short time. In order to take advantage of these points we proposed a simple camera which will allow direct imaging, coronagraphic imaging and coronagraphic low resolution spectroscopy. SHARK-NIR together with the SHARK-VIS channel, are covering a wide wavelength domain, going from 0.6$\mu m$ to 1.7$\mu m$ (Y to H band). SHARK-NIR will offer extreme AO direct imaging capability on a field of view (FoV) of about 18\textquotedblright x 18\textquotedblright, and a simple coronagraphic spectroscopic mode offering spectral resolution ranging from 100 to 700. \\ The main science case of SHARK-NIR is searching for giant planets, to succeed, an high contrast is necessary. We also emphasize that the LBT AO SOUL upgrade will further improve the AO performance, making possible to extend the exoplanet search to target fainter than normally achieved by other 8-m class telescopes, and opening in this way to other very interesting scientific scenarios, such as the characterization of AGN and Quasars, normally too faint to be observed, and increasing considerably the sample of disks and jets to be studied. | \label{sec:conclusion} After having tested other kind of coronagraphs, as the asymmetric SP, the Apodized Phase Plate and the Vortex, we moved for choosing the final configuration of the SHARK-NIR. The asymmetric masks are attractive because of the good performance near the IWA. For these reasons, the final choice is to implement a symmetric SP and a couple of asymmetric masks. Since asymmetric masks cannot access the entire 360$^\circ$ around the star, they may be used for characterization of known objects. However, they are designed to generate two symmetric 110° high-contrast regions (for a total of 220$^\circ$ discovery space each) which are perpendicular to each other. Hence, if used in sequence, we could in principle access the entire 360$^\circ$ region around the star for discovery. \\ Long ADI sequences are still difficult to generate because of computational time. We are focusing on increasing the number of images to compare our simulations to the real amount of data taken during an on sky observation and also because we are confident that the post processing code performance will become more and more successful by using more images. \\ We also analysed some low Strehl cases, the gain of the pipeline in comparison to a simple ADI is confirmed, even if the advantage is lower and focus on the large separation from the host star. This result and the gain of 1 magnitude in high and medium Strehl case in different coronagraph configurations encourage to implement the code on sky images to verify and adjust the pipeline for the analysis of the scientific images of the SHARK-NIR. | 18 | 8 | 1808.03121 |
1808 | 1808.05531_arXiv.txt | During the evolution of proto-planetary disc, photo-evaporations of both central and external stars play important roles. Considering the complicated radiation surroundings in the clusters, where the star formed, the proto-planetary discs survive in different lifetimes due to flyby events. In this paper, we mainly focus on the disc around a T Tauri star, which encounters with another main-sequence star with different temperatures in hyperbolic orbits with different peri-center distances, eccentricities and inclinations. We find the criterion for gap-opening due to photo-evaporation of central star after the flyby event. A gap is opened in the late stage of gas disc, and induce that the gap only influence the planet formation and migration limitedly. If the flyby orbit has a moderate value of peri-center distance, which weakly depends on the eccentricity and inclination, the external photo-evaporation lead to a maximum mass loss during the flyby event. Flyby stars in orbits with smaller eccentricities or larger inclinations induce larger mass loss. Adopting a simple multiple flyby models, we conclude that in open clusters, gas discs usually survive in typical lifetimes between 1 and 10 Myr, except there are many massive stars in dense open clusters. In globular clusters, discs disperse very quickly and hardly produce the gas giant planets. The fast-depleted discs are probably responsible for the null detection of giant planets in globular clusters. \textbf{Key words}: protoplanetary discs -- ultraviolet: stars -- planets and satellites: gaseous planets -- photodissociation region (PDR) -- globular clusters: general -- open clusters and associations: general | The environments of star formation region are usually clustered. Currently, only two pulsars as planet hosts are detected in globular clusters and tens of planet hosts in open clusters. Most recent works by \citet{2017Natur.548..183M}, indicate the number of free-floating or wide-orbit Jupiter-mass planets in the galactic halo, i.e. less than 0.25 Jupiter per main sequence star on average. Only eight giant planets are detected around stars in clusters, while seven of them are in open clusters \citep{2018aJ....155..173C,1999apJ...523..763T}. Because of the limited planet sample in open clusters, there are debates whether the planet occurrence rate around stars in open clusters and field stars are similar or not. However, some transiting surveys in globular clusters with null results hint the occurrence rate of Jupiter-mass planets in clusters may be lower than field stars \citep{2000apJ...545L..47G,2005apJ...620.1043W,2012a&a...541a.144N,2017aJ....153..187M} The instability of planets in clusters is considered to be an important role \citep{2009apJ...697..458S,2013apJ...772..142L}. A large fraction of planets very close to the host star are probably stable in open clusters or even in the outer region of globular clusters \citep{2011MNRaS.411..859M,2001MNRaS.322..859B}. Beside the dynamical instability in clusters, the metallicity is also considered to influence planet formation \citep{2012Natur.486..375B}. Low metallicity in old globular clusters may lead to lower planet formation rate, because the few dust around stars with poor metallicity can hardly form enough planetary embryos. Additionally, the metallicity is crucial for the core due to gas accretion scenarios, but once the core is formed before gas depletion, the gas giant can form. Due to core accretion scenarios, the formation of giant planet is relative to the proto-planetary discs tightly. The mass loss and evolution of proto-planetary gaseous disc play important roles during planet formation, especially for gas giants. In viscous discs, accretion from central star make the density profile decays smoothly, and lead to disc dispersion. Photo-evaporation is also an important mechanism to influence the evolution of gaseous disc. The evolution, especially the lifetime, of the gaseous disc will determine weather the gas giant can form or not, and how large the gas giants can grow. As one of essential mechanisms to lose mass, photo-evaporation leads to dispersal of proto-planetary discs \citep{2011aRa&a..49...67W}. Photo-evaporation can be driven by photons in the energy range of far-UV( 6 eV<$h\nu$<13.6 eV, here after FUV), extreme-UV(13.6 eV<$h\nu$<0.1 keV, here after EUV), and X-ray($h\nu$>0.1 keV). Photons in such energy ranges influence the discs in different ways. FUV can dissociate molecules like $H_{2}$, while EUV and X-ray can ionize the hydrogen atoms. The proto-planetary disc is heated by these photons to temperatures about 1000-10000 K, and consequently gas materials is evaporated via thermal escape. Previous photo-evaporation models of are usually considering EUV, FUV and X-ray in a viscous disc. Most of early models focused on the radiation from central stars \citep{2011aRa&a..49...67W}. \citet{2001MNRAS.328..485C} first introduce a UV-switch model which includes the internal EUV photo-evaporation and disc viscosity. They indicate a photo-evaporative gap in the late time of disc evolution. With a mass loss rate about $10^{-10}$ M$_{\odot}$ yr$^{-1}$ originated by central ionizing flux, at the early stage in the disc evolution, the internal photo-evaporation has a negligible effect compared with the large accretion rate due the viscosity. As the disc disperses, the internal photo-evaporation becomes crucial and clear out the gas material near the gravitational radius while the viscous evolution can hardly fill the gap. Then \citet{2006MNRAS.369..216A,2006MNRAS.369..229A} develop the model and consider the effect of stellar radiation directly incident on the inner disc edge at late disc times which lead to a quick dispersion of the out disc when the gap is opened. Recent photo-evaporation models include internal X-ray \citep{2009ApJ...699.1639E,2010MNRAS.401.1415O,2011MNRAS.412...13O} or FUV irradiation \citep{2009ApJ...690.1539G}, in addition to EUV photons. It's because that X-ray and FUV photons can penetrate the larger columns of neutral gas than EUV photons and consequently lead to larger mass loss rate $\sim 10^{-8}$ M$_{\odot}$ yr$^{-1}$ which is two order magnitude larger than that driven by EUV photons. To speak more specifically, the FUV photons can remove disc mass at large radii, and truncate the disc at $\sim 100$ AU typically. X-ray photons can penetrate the neutral gas to a few tens AU. While EUV photons can only penetrate neutral gas to several AU efficiently. I.e. the mass loss rate due to EUV photons is more concentrated at a scale of a few AU. The external photo-evaporation is also important during disc evolutions. \citet{1998apJ...499..758J} presented a model for the photo-evaporation of circumstellar discs or dense clumps of gas by a massive external source of ultraviolet radiation. Later, in 2003, they developed a model in the proto-planetary discs around T Tauri stars \citep{2003ApJ...582..893M}. They combined viscous evolution, central EUV photo-evaporation and external EUV or FUV photo-evaporation originated by massive nearby stars. External photo-evaporation will shrink the disc outer edge and accelerate the disc evolution. Photo-evaporation will not only influence the lifetime of proto-planetary disc, but also influence the internal structure of the disc. Similar to internal photo-evaporation models, at late times, there will be a photo-evaporative gap. The gap structure will influence both the runaway accretion of gas giant and the migration of small planets. Thus planet formation and orbital architectures may be various during the disc evolution under photo-evaporation. In this paper, we mainly focus on the gaseous disc around a T Tauri star, which undergoes a close encounter with another main-sequence star. The photo-evaporation effect depends on the orbit parameters of the intruder in a hyperbolic orbit. Previous external photo-evaporation models usually assume that the radiation field is perpendicular to the disc surface, while it's not suitable in the condition of flyby events with random orientation. We combine both orbit parameters and efficient receiving area of disc during the flyby, then develop the photo-evaporation model of external stars to a general way. Besides, the duration of a close encounter is usually accounted until the star is too far to influence the mass lose of gaseous disc. It's our aims to figure out the influence on the viscous disc due to flybys with different stars in different orbits. Furthermore, multiple flyby events will influence the disc evolution continuously, especially in the environments of clusters, where the stars formed. We focus on the lifetime of gaseous disc in different environments (e.g. field stars, open clusters and globular clusters), and provide some clues of gas giant formation in these environments according to different disc lifetime. In section 2, We developed an alternate EUV/FUV photo-evaporation model of external star due to different orbital parameters. In section 3, we analysis the relation between the hyperbolic orbit parameters and mass loss, via integration. In section 4, the one-dimension alpha disc simulations are represented and be compared with the analysis results of photo-evaporation. In section 5, our results are extended in young open clusters and globular clusters. Additionally, we compare our results with previous works and observations in clusters. We summarize our conclusions in section 6. | Since planets formed in proto-planetary gaseous disc, the evolution of gaseous disc is crucial for planet formation, especially for gas giants like Jupiter. The lifetime of the gaseous disc can determine whether the gas giant can form or not, and how large the gas giant can grow due to core accretion scenario. The Photo-evaporation is an important mechanism during disc evolution and lead to different mass loss rate of the proto-planetary disc. In this paper, we use a general one dimension disc model (section 2), including EUV and FUV photo-evaporation, to investigate the disc evolution due to different flyby stars in different hyperbolic orbits. In section 3, we analysis the mass loss during single flyby event due to external star in two cases, which depend on the orbital parameters as shown in Equation \ref{20} and \ref{22}. Using one-dimension diffusing equation, we simulate the disc evolution due to both viscosity and photo-evaporation of both external and central stars for single flyby events in section 4. The timescales for gap-opening and depletion of gaseous disc due to different flyby events are obtained. To model the real star birth environments, we simulate disc evolution via a simple multi-flyby model in section 5. We list the main conclusions in this paper as follows: \begin{itemize} \item Using atlas9 model, we calculate the EUV and FUV flux of different stars with effective temperature from 9500 to 38000 K. The FUV photo-evaporation regions for different external stars are estimated as shown in Table 1. \item We derive analytical equations of total mass loss due to external star with different parameters, including external ionization flux, eccentricity $e$, inclination $inc$ and peri-center distance $q$ of hyperbolic orbit. In the case of EUV dominated photo-evaporation when $q > d_{\rm max}$, single flyby lead to very few mass loss (Equation \ref{19}). Considering the FUV photo-evaporation region, the FUV photo-evaporation is much more efficient than EUV, thus the mass loss of disc during flyby mainly depends on how long the star go across the FUV dominated region. In the case of $T_{\rm eff}=19000 $ K, external star with peri-center $\sim 10000 $ AU lead to the maximum mass loss as shown in Figure \ref{Figure 4}. \item Using the alpha-disc model with $\alpha=0.001$, and consider the photo-evaporation of central star, we find that the ionization flux of external star and the peri-center distance $q$ are crucial for the disc evolution after flybys, while the inclination and eccentricity influence the disc evolution limitedly. Only Stars with $T_{\rm eff} > 15000 $ K in an orbit with peri-center $q=10000$ AU can reduce the disc lifetime. For the single flyby cases of $T_{\rm eff}\in [9500,38000] $ K in Figure \ref{Figure 6}, or $q\in[2000,10^6]$ AU in Figure \ref{Figure 10}, the lifetime of gaseous disc is $>7$ Myr, which is typical for proto-planetary disc, and can not significantly restrain the formation of gas giants. \item By changing the dimensionless viscosity alpha and ionization luminosity $\Phi_{\rm c}$(in internal photo-evaporation and viscosity evolution case), we find that, with a larger $\alpha=0.01$ the disc can disperse at less than 2 Myr, which is well consistent with the observations that protoplanetary gaseous disc can disperse in 3 Myr. While increasing the central ionization luminosity to $10^{42}$ s$^{-1}$, the gap opens at an earlier time. We also find that the disc viscosity, instead of the internal ionization luminosity, dominates the outer disc evolution after gap opens in our model. \item During disc evolution, the gap opens when the mass transfer at $r_{\rm g,euv}$ due to viscosity is less than the mass loss due to photo-evaporation. We derived an criterion to estimate the gap-opening timescale $t_{\rm gp}$, which is consistent with simulations(Figure \ref{Figure 7}). According to our simulations, the gap usually opens when the disc is nearly depleted as shown in Table \ref{Table 2}. Since the gaseous disc is less massive, the gap-opening can hardly influence the migration of planets in the disc. \item Flyby events in early stage of disc can influence the disc evolution, e.g. when $t_{\rm b}\le 5$ Myr, the later the flyby begins, the later gap open and the later disc disperses. while flyby events in late stage of disc can influence the disc evolution limitedly, e.g. $t_{\rm b}\in [8,10]$ Myr, $t_{\rm dp}$ is almost the same as the case with no flyby in Figure \ref{Figure 5}, and $t_{\rm dp} \approx t_{\rm b} +1$ Myr. \item Since flyby events occurs frequently in clusters. We investigate the photo-evaporation of gaseous disc during multi-flyby events. In both cases of $q=10^{4}$ and $10^{5}$ AU, there is an area with stellar density $< 100$ pc$^{-3}$ and no massive stars ($T_{\rm eff}< 30000$ K) in the clusters, e.g. open cluster M67, where the gas dispersal timescale $t_{\rm dp}$ is between 1--10 Myr. I.e. the disc lifetime is typical for the formation of gas giants. For globular clusters with stellar density $>10^5 $ pc$^{-3}$, where there usual are some massive stars, the disc will depleted in 0.1 Myr, therefore gas giants can hardly formed in such an environments. \item When the external star($T_{\rm eff}>16000$ K see figure \ref{Figure 6} and \ref{Figure 11}) is coming close to the peri-center, the outer disc edge will shrink, which is similar to the surface density profile shown by \citep{2003ApJ...582..893M}. When the external star is far away from the center star, the external photo-evaporation can be negligible, the disc evolution is dominated by the viscosity and central photo-evaporation which is similar to the UV-switch model \cite{2001MNRAS.328..485C,2006MNRAS.369..216A,2006MNRAS.369..229A}, when the $\dot M = 2*\dot M{\rm c,euv}$, the disc will open a gap at the soften gravitational radius $\beta r_{\rm g,euv}$, and the inner disc will dissipate at the viscous times scale $ \sim 0.1$ Myr based on our parameter setting. \end{itemize} Previous external photo-evaporation models, e.g. \citep{2003ApJ...582..893M,2013apJ...774....9a,2018ApJ...853...22X} do not consider the change of distance between external star and central star, where our model mainly focuses on. They conclude that the disc dispersal timescale should become much less than 10 Myr. Considering a smaller gaseous disc, \citep{2013apJ...774....9a} obtain the gas depletion timescale of 1-3 Myr, when only FUV radiation with $G_0=300-30000$ works, which is $10^2-10^4$ higher than the interstellar value. Our results show a flyby events with external star with $T_{\rm eff}=38000 $ K(with $G_0 \sim 30000$ at $d=10^5$ AU), including both EUV or FUV depends on the distance. The depletion timescale of gaseous disc is put forward $>$ 7 Myr, but larger than the previous works. In this paper, we only consider the flyby with peri-center larger than 2000 AU, to avoid the influence of gravitational effects. Recently, \citet{2014MNRAS.441.2094R} considers both the hydrodynamical evolution of the discs around their natal stars, as well as the dynamics of the stars in clustered environments. As discs evolve due to viscosity, encounters become more and more important, it will truncate the discs and deprive of the outer potions because of the dynamical tidal effect. For close flybys with $q\sim 10^{2}$ AU, the dynamical effect of external star is significant. While for wide flyby events with $q\ge 10^{4}$ AU, the mass loss due to dynamical effect of flyby can be neglected(see figure 3 in their paper). Their results are consistent with our assumption that, when $q\ge 10^{4}$ AU, the gravitational effects can be ignored. They predict that disc sizes are limited by encounters at stellar densities exceeding $\sim 2$ - $3 \times 10^{3}$ pc$^{-2}$. In our work, we do not consider the gravitational effects, but mainly focus on the external photo-evaporation in clustered environments. And our results show that, in globular clusters with high stellar density $>10^{5}$ pc$^{-3}$, disc lifetimes will be probably reduced to less than 1 Myr. In our model, we do not consider the binary cases, although the binary stars are common in universe \citep{2010apJS..190....1R}. Recently, \citet{2018MNRaS.473.5630R} study the evolution of proto-planetary discs in binaries. They focus on the X-ray photo-evaporation of each star in binary and also the tidal effect between companion star and disc. They conclude that the disc around close binary star disperse quickly than single stars due to tidal effect between the companion and the disc. Based on the results of disc in binary, we can deduce that the close flyby with small peri-center will clear the disc more quickly. However, these flybys are rare according to Equation \ref{26}, thus only a few disc may be influenced by the close flyby events accidentally. Note that the massive star is not the only source of EUV and FUV radiation in clusters. Both the accretion disc around Black hole, and accretion disc in compact binary can produce X-ray emission. These additional radiation can strengthen the photo-evaporation and then accelerate the evolution of a proto-planetary disc. Contrarily, in a young embedded cluster with plenty of gas, the strong extinction of EUV and FUV can reduce the mass loss due to photo-evaporation and protect the proto-planetary disc. We adopt a very simple multiple flyby model to investigate photo-evaporation in clusters. As mention before, to get a more reasonable parameters of flyby, it's better to do N-body simulations to obtain the distribution of the mass, peri-center, eccentricities, and inclinations of flyby stars in clusters. However, these distributions are sensitive to the cluster model, especially the initial conditions and IMF of clusters. The results could be more specific, if we can get the distribution of flyby parameters in some clusters. % We can explain the gas giant formation is hard in globular clusters because of short-lived gaseous disc, which is consistent with the null survey results in globular clusters \citep{2017aJ....153..187M}. Our results in this paper predict that the stars in open clusters without very massive stars can sustain a disc similar to the field star, thus the planet occurrence is probably the same as that of field star. It's not constrained well by observation because of few planets samples in open clusters. Further surveys in different clusters with or without massive stars can provide clues to investigate the influence of planet formation rate in clusters. | 18 | 8 | 1808.05531 |
1808 | 1808.07702_arXiv.txt | We report new, $\sim$1000 AU spatial resolution observations of 225 GHz dust continuum emission towards the OB cluster-forming molecular clump G33.92+0.11. On parsec scales, this molecular clump presents a morphology with several arm-like dense gas structures surrounding the two central massive ($\gtrsim$100 $M_{\odot}$) cores. From the new, higher resolution observations, we identified 28 localized, spatially compact dust continuum emission sources, which may be candidates of young stellar objects. Only one of them is not embedded within known arm-like (or elongated) dense gas structures. The spatial separations of these compact sources can be very well explained by Jeans lengths. We found that G33.92+0.11 may be consistently described by a marginally centrifugally supported, Toomre unstable accretion flow which is approximately in a face-on projection. The arm-like overdensities are natural consequence of the Toomre instability, which can fragment to form young stellar objects in shorter time scales than the timescale of the global clump contraction. On our resolved spatial scales, there is not yet evidence that the fragmentation is halted by turbulence, magnetic field, or stellar feedback. | \label{sec:introduction} Molecular clouds may undergo global collapse \citep[e.g.,][and references therein]{Bate2003,Hartmann2012,Enrique2017}, which may lead to a centrally concentrated density distribution of gas and young stellar objects (YSOs) \citep[e.g.,][and references therein]{Liu2012a,Galvan2013,Lin2017}. As a consequence of the accumulated angular momentum, the central flattened rotating gas clump in the collapsing molecular cloud may present spin-up motions even on parsec scales \citep[e.g.,][]{Ho1986,Keto1987,Welch1987,Ho1996,Zhang1997,Galvan2009,Liu2010}. The central cluster-forming clump may be marginally supported by rotational motion, until the collected gas mass is sufficient to trigger the self-gravitational instability, which will then result in spiral arm-like dense gas structures and a distribution of localized gas fragments \citep[e.g.,][and references therein]{Keto1991,Lee2016a,Lee2016b,Sakurai2016,Mapelli2017}. Such structures may have been resolved in some previous observations \citep[e.g.,][]{Liu2012b,Takahashi2012,Beuther2013,Wright2014,Liu2015,Chen2016,Li2017,Beuther2018,Izquierdo2018,Maud2017}. How these gas structures fragment to subsequently form 10$^{3}$-10$^{4}$ AU scales gas cores and YSOs, remain uncertain. To well resolve the gas structures forming out of the self-gravitational fragmentation in the centralized massive molecular gas clumps in OB cluster-forming molecular clouds, we selected to observe the target source G33.92+0.11 ($d\sim$7.1 kpc). Its very small derived virial mass compared to the enclosed molecular gas mass \citep[][]{Watt1999} indicates that it is likely geometrically thin and is in a face-on projection. It encloses a few thousands $M_{\odot}$ of gas mass in the central parsec scale area \citep[][]{Liu2012b}. However, the previous interferometric observations of molecular lines resolved very small relative motions with respect to its systemic velocity 107.6 km\,s$^{-1}$ \citep[e.g., within $\pm$2 km\,s$^{-1}$ for most of the regions. For more details see][]{Liu2015,Minh2016}. This was interpreted as motions predominantly in the plane of the sky, with the dominant motion being rotation, where the axis of rotation is parallel to the line of sight. If this is indeed the case, then the studies of its matter distribution will be minimally affected by line-of-sight confusion. We have resolved G33.92+0.11 using the Atacama Large Millimeter Array (ALMA) and Atacama Compact Array (ACA), with a $\sim$1000 AU spatial resolution. The observations and data reduction are introduced in Section \ref{sec:observation}. The direct observational results are presented in Section \ref{sec:results}. In Section \ref{sub:Toomre} we address the gravitational instability of the resolved system based on the analysis of the Toomre Q parameter. In Section \ref{sub:dendrogram} we present our identification of the localized candidates of YSOs using the {\tt dendrogram} algorithm \citep{Rosolowsky2008}; and in Section \ref{sub:cluster} and \ref{sub:separation} we discuss the clustering of the identified YSO candidates, and the probable physical mechanism to explain their spatial separations. Our conclusion is given in Section \ref{sec:conclusion}. \begin{figure*} \hspace{0.3cm} \begin{tabular}{ p{8cm} p{8cm} } \multicolumn{2}{l}{ \includegraphics[width=14cm]{g33p92_alma224p5_rob0-eps-converted-to.pdf} }\\ \includegraphics[width=7.5cm]{g33p92A_alma224p5_rob0-eps-converted-to.pdf} & \includegraphics[width=7.5cm]{g33p92B_alma224p5_rob0-eps-converted-to.pdf} \\ \end{tabular} \vspace{0cm} \caption{ The 1.3 mm continuum image generated by combining all existing ALMA and ACA data. The full width of half maximum (FWHM) of the primary beam is 28$''$, as indicated with the gray circle. Blue ellipses show the synthesized beam (\beam = 0$''$.16$\times$0$''$.093; P.A.=74$^{\circ}$). Panel (A), (B) and (C) show respectively the full image, and the smaller regions around the A1 and A2 {\it central massive molecular cores}. } \label{fig:almacycle4} \vspace{-0.55cm} \end{figure*} \begin{figure} \includegraphics[width=8.5cm]{g33p92A_h30alphamnt0_rob0-eps-converted-to.pdf} \caption{The velocity integrated intensity map (i.e., moment 0) of the H30$\alpha$ line (gray), overlaid with the 1.3 mm continuum image is presented in Figure \ref{fig:almacycle4}. The two images share the same synthesized beam. Contour levels are 25 $\mu$Jy\,beam$^{-1}$ $\times$ [3, 6, 12, 24, 48, 96, 192, 384, 768]. \label{fig:h30alpha} } \end{figure} | \label{sec:conclusion} We have performed $\sim$1000 AU spatial resolution observations at $\sim$225 GHz with ALMA towards the flattened and rotating OB cluster-forming clump, G33.92+0.11. The target source is likely in a face-on projection, which is ideal for the studies of the morphology of dense gas. We found that dense gas in this region is Toomre unstable, forming the $\sim$0.3 pc scales arm-like or elongated structures, which are hierarchically embedded with smaller internal sub-structures. The spatial separations of the candidates of Class 0/I YSOs we identified from this source can be explained with Jeans lengths, which were derived based on rather simple assumptions. Our present interpretation is in concert with the previously observed fact that the virial parameters on $\sim$1 pc scale and on the spatial scales of individual dense cores appear small, which indicates that the hierarchically forming dense gas structures are supported by specific angular momentum in the direction approximately along our line-of-sight. | 18 | 8 | 1808.07702 |
1808 | 1808.09974_arXiv.txt | We show that a 10 year \textit{Gaia} mission could astrometrically detect the orbital motion of $\sim 1$ sub-parsec separation supermassive black hole binary in the heart of nearby, bright active galactic nuclei (AGN). Candidate AGN lie out to a redshift of $z=0.02$ and in the V-band magnitude range $10 \lesssim m_V \lesssim 13$. The distribution of detectable binary masses peaks at a few times $\sim 10^7 \Msun$ and is truncated above a few times $\sim 10^8 \Msun$. | The \textit{Gaia} satellite is mapping the positions of the stars with unprecedented precision. Its 5 year mission: to survey the 6D phase space coordinates of a billion stars to an astrometric precision of a few $\mu$as \citep{GaiaI:2001, GaiaII:2016, GaiaDR2:2018}. \textit{Gaia} will observe not only stars, but all optical sources brighter than an apparent magnitude of $\sim20$. This includes active galactic nuclei (AGN), namely distant and powerful sources of multi-wavelength emission driven by gas accretion onto supermassive black holes (SBHs) at the centers of galaxies. AGN are used to calibrate \textit{Gaia} astrometric position measurements, both via \textit{Gaia}'s optical astrometry as well as with radio-frequency VLBI \citep{GaiaDR2:astrmsoln:2018}. The AGN are chosen as calibrators because they are distant and hence expected to exhibit very little proper motion or parallax. Despite this expectation, \textit{Gaia} has detected $\gtrsim 1$mas offsets in optical and radio positions of AGN, probing dislodged AGN or radio/optical jet properties \citep{Makarov+2017, PetrovKovalev:2017, KovalevPetrov:2017, PetrovKovalev:2018}. In this \textit{Letter} we show that on $\lesssim 50\mu$as scales, this expectation is also relevant for AGN that harbor sub-parsec (pc) separation SBH binaries (SBHBs). Orbital motion of one or both accreting SBHs in a SBHB can change the position of the optical emitting region of the AGN by an angle greater than the astrometric precision of \textit{Gaia}. SBHB orbital motion would be distinct from the linear motion expected for a jet or ejected AGN. Because binary-induced motions will only occur for a minority of AGN, there will be little impact on \textit{Gaia}'s calibration. This observation does, however, present a path towards definitive detections of sub-pc separation SBHBs. While solid lines of evidence lead us to expect that SBHBs reside in the centers of some galaxies \citep{Begel:Blan:Rees:1980}, their definitive detection at sub-pc separations is yet to be obtained. The existence of sub-pc SBHBs is of special importance as it embodies the `final-parsec problem' \citep{Begel:Blan:Rees:1980, MerrittMilos:2005:LRR}, determining the fate of SBHBs. If interaction with the environments in galactic nuclei can drive SBHBs to sub-pc separations, then they will merge via emission of gravitational waves (GWs), detectable out to redshifts $z \geq 10$ by the future space-based GW observatory LISA \citep{LISA:2017}, and generating a low-frequency stochastic GW background detectable by the Pulsar Timing Arrays \citep[PTAs;][]{PTAs}. To determine which, if any, proposed mechanisms, \cite[\textit{e.g.},][]{GouldRix:2000, ArmNat:2002:ApJL, MacFadyen:2008, Goicovic+2016, Gualandris+2017:TriaxRefill, Yike+2017}, solve the final-parsec problem in nature, one must characterize a population of sub-pc SBHBs. Current detection methods are indirect and require campaigns that last many years \citep[\textit{e.g.},][]{ShenLoeb:2010, Tsalmantza:2011, Bogdanovic+2009, Eracleous+2012, McKFeZoltan:2013, DecarliDott:2013:SpecMBHBcandI, Shen+2013I, Liu+2014II, LiuEracHalp:2016, NguyenBogdan+2018, HKM09, DHM:2013:MNRAS, PG1302MNRAS:2015a, Farris:2014, Graham+2015b, Charisi+2016, LiuGezariPLCs+2016, DZ:2017, DDLens:2017, Gower+1982, Roos:1993, MerrittEker:2002, Zier:2002, Romero:2000, Kun+2014, Kun+2015:PG1302, KulkLoeb:2016, LiuChen:2009, StoneLoeb:2011, Coughlin+2017}. While these techniques provide a way towards identifying and vetting SBHB candidates via a combination of indirect methods, a more direct approach is desired. Recently, we have shown that mm-wavelength VLBI possesses the astrometric resolution and longevity to repeatedly image SBHB orbits out to redshift $z\sim0.5$, providing direct evidence for SBHBs in radio-loud AGN \citep{D'OrazioLoebVLBI:2018}. The technique that we propose here also directly tracks the SBHB orbit with the advantage that target AGN need not be bright in mm-wavelengths and that unlike VLBI, \textit{Gaia} is conducting a survey mission that will map the entire sky, and, as we show, could find evidence for SBHBs within the next $5-10$ years. \begin{table*} \begin{tabular}{l|l|l|l|l} Parameter & Meaning & Fiducial & Optimistic & Pessimistic \\ \hline \hline $f_{\bin}$ & The fraction of AGN harboring SBHBs & $0.1$ & " & " \\ $f_{\Edd}$ & The Eddington fraction of bright AGN & $0.1$ & " & " \\ $BC$ & Bolometric correction from V-band & $10.0$ & " & " \\ $t_{Q}$ & The AGN lifetime & $10^7$~yrs & $5 \times 10^6$~yrs & $10^8$~yrs \\ $V-I_c$ & A mean color for nearby AGN & $0.7$ & $1.1$ & $0.0$ \\ $P_{\max}$ & Mission lifetime & $10$~yrs \ ($5$~yrs) \ ($20$~yrs) & $10$~yrs \ ($5$~yrs) \ ($20$~yrs) & $10$~yrs \ ($5$~yrs) \ ($20$~yrs) \\ $q$ & Binary mass ratio & $0.1$ & $0.05$ & $1.0$ \\ \hline \hline ${ \bf N_{\rm SBHB} }$ & {\bf Number of detectable ($\text{SNR}\geq2$) SBHBs } & ${\bf 1.1}$ \quad ($0.3$) \quad ($3.1$) & ${\bf 1.3}$ \quad ($0.4$) \quad ($3.8$) & ${\bf 0.8}$ \quad ($0.2$) \quad ($2.0$) \\ \end{tabular} \caption{ Model parameters and the resulting number of \textit{Gaia}-detectable SBHBs (note that a 20 year mission lifetime requires a successor to \textit{Gaia}).} \label{Table:params} \end{table*} | We have shown that a 10~yr \textit{Gaia} mission has the capability to astrometrically track the orbital motion of $\mathcal{O}(1)$ SBHBs in bright ($m_V\lesssim13$), nearby ($z\lesssim0.02$) AGN. The discovery of SBHB orbital motion over the next few years of the \textit{Gaia} mission would open a new field of SBHB demography, generating an enormous boon for our understanding of the mutual growth of SBHs and galaxies, evidence towards resolving the final-parsec problem, the prospect of sources of gravitational waves for PTAs, and a new method for calibrating cosmological distances \citep{D'OrazioLoebVLBI:2018}. There is a strong incentive to analyze astrometric data of bright, nearby AGN from \textit{Gaia} DR2 and onwards for signatures of SBHB orbital motion. | 18 | 8 | 1808.09974 |
1808 | 1808.02899_arXiv.txt | We observed the 43 GHz $v$=1, 2 \& 3 and 86 GHz $v$=1 SiO maser transitions quasi-simultaneously for a Mira-variable dominated sample of over 80 sources from the Bulge Asymmetries and Dynamical Evolution (BAaDE) project, using ATCA, and statistically compared the relative line strengths. On average, the 43 GHz $v$=1 line is brighter than the 86 GHz $v$=1 line by a factor of 1.36$^{+0.15}_{-0.14}$. As a result, an 86 GHz $v$=1 observed sample can be observed to 85.9\% of the distance of a 43 GHz $v$=1 observed sample using the same sensitivity. We discuss what impact this may have on the BAaDE Galactic plane survey using the VLA and ALMA. Despite fewer $v$=3 detections, specific trends are discerned or strengthened when the 43 GHz $v$=3 line is detected. In particular the 43 GHz and 86 GHz $v$=1 lines are on average equal for sources with no detectable 43 GHz $v$=3 emission, but the 43 GHz $v$=1 line strength is on average about twice as bright as the 86 GHz $v$=1 line for sources with detectable 43 GHz $v$=3 emission. Some weak correlations are found between line strengths and Midcourse Space Experiment (MSX) flux densities and colors, which are tightened when considering only sources with detectable 43 GHz $v$=3 emission. We discuss these trends in the context of a radiative pumping model to highlight how the 43 GHz $v$=3 line, when coupled with the $v$=1 and $v$=2 lines, can further our understanding of variable conditions like density in the circumstellar envelopes. | The Bulge Asymmetries and Dynamical Evolution (BAaDE) project aims to explore the dynamical structures of the inner Galaxy and Galactic Bulge, using SiO maser lines from red giant stars. By constructing a sample of stellar point-mass probes, models of the gravitational potential can be tested using a final sample that is expected to provide about 20,000 line-of-sight velocities and positions. The SiO maser transitions occur at radio frequencies where extinction is negligible, thus allowing a dense sampling of line-of-sight velocities in the most crowded and optically obscured regions of the Milky Way. For a more detailed description of the BAaDE survey, see \citet{in...prep}. While using both the Karl G.\ Jansky Very Large Array (VLA) and Atacama Large Millimeter/submillimeter Array (ALMA) ensures a full coverage of the Galactic plane, their different receiver availabilities (43 GHz at the VLA and 86 GHz at ALMA) also dictates that one part of the sample be observed in the 43 GHz $J$=1--0 SiO transitions (VLA), and the other part in the 86 GHz $J$=2--1 transitions (ALMA). A fundamental assumption for BAaDE is thus that stars emitting 43 GHz SiO maser emission also harbor 86 GHz masers, and vice versa. This appears to be a commonly accepted fact, supported, for example, by the work of \citet{2004PASJ...56...45S} who noted that out of 39 sources displaying 86 GHz SiO maser emission, 38 also produced 43 GHz masers. What is less clear, however, is whether there is a statistically significant difference in the integrated flux densities of 43 GHz versus 86 GHz masers. Such a difference could have an impact on the analysis of the BAaDE sample, as the VLA and ALMA samples are in principle observed to the same noise levels and two different Malmquist biases would have to be considered when combining the two samples to draw statistical conclusions. This is especially important, as the 86 GHz ALMA sample, which utilizes the possibly fainter transition of the masers (see below), covers Galactic longitudes of $-110\degree<l<0\degree$, and therefore contains the region of the bar that is furthest from us. Observations of both 43 GHz and 86 GHz SiO maser transitions indicate that the flux-density ratio may depend on the stellar type, or circumstellar shell thickness. The shell thickness can be inferred from infrared colors, as the thicker envelopes produce redder colors. In the optically-thick case, the central star may be completely obscured in the optical regime, while it is observable for the optically thin shells. For Asymptotic Giant Branch (AGB) stars, this has been outlined by \citet{1988A&A...194..125V} using IRAS color-color diagrams, where thick-shell OH/IR stars are found in the redder regions IIIb and IV, and thin-shell Mira variables, with typical pulsation periods longer than 100 days and IR amplitude variations greater than one magnitude, in the bluer regions II and IIIa. \citet{1993A&A...280..551N} observed 40 OH/IR stars and noted that the integrated flux densities of the 43 GHz $v$=1 lines were brighter than the 86 GHz $v$=1 lines (which were observed on average two weeks later), and that the difference was greater for stars with thicker envelopes compared to those with thinner envelopes. Similarly, \citet{2016JKAS...49..261K} report on the 43 GHz v=1 line being brighter than the 86 GHz v=1 in two additional OH/IR stars. The BAaDE sample was chosen from a sample of stars likely to be dominated by Miras and hence possess much thinner envelopes than the OH/IR stars. This sample, by an extension of the results of \citet{1993A&A...280..551N}, might imply a smaller difference between the 43 GHz and 86 GHz line flux densities, and perhaps even a reversal in which the 86 GHz line would be brighter. This would be consistent with a weak trend reported by \citet{1998A&A...329..219P}, where 5 out of 9 sources were brighter in the 86 GHz v=1 line compared to the 43 GHz v1 (the majority of which were Miras). Further, numerical simulations investigating collisionally pumped maser time variability in Miras suggests that, averaged over the stellar cycle, the $v$=1 maser emission is expected to be brighter at 86 GHz than at 43 GHz \citep{2002A&A...386..256H}. However, this collisionally pumped model does not include the effects of line overlap at the pumping frequencies, which can produce quite different relative line strengths in radiatively pumped models \citep[e.g.,][]{2014A&A...565A.127D}. For the BAaDE sample, we have no information about the stellar phase for individual targets, thus averaging information over a large sample of stars will be considered to effectively correspond to a full-sample cycle average. To assess the robustness of the assumption that the strengths between the 43 GHz and 86 GHz SiO maser emission are on average equal, we use quasi-simultaneous observations of the 43 and 86 GHz SiO lines in a sample of BAaDE sources. In section 2 we describe the selected sample, the observations and data calibration. In sections 3 and 4 we share our results and discuss the significance between various line strengths starting with a comparison between the 43 and 86 GHz $v$=1 lines. | \renewcommand{\labelenumi}{\alph{enumi})} \begin{enumerate} \item The $v$=3 lines are formed in the high-density regime. In this regime, near $n_{H_2} \sim 10^{10}$ cm$^{-3}$, the 43 GHz $v$=1 line is modeled to be brighter than the 86 GHz $v$=1 line. In the low-density regime, these two lines are on average more equal in strength. This is in full agreement with observational trend \textit{a}. \item In the high-density regime, the 43 GHz $v$=2 and $v$=1 line brightnesses are expected to follow each other closely and to be of near equal strength. In the low-density regime, there is a much larger distribution of the difference between the lines, with $v$=1 being the brightest, in agreement with the observed trend \textit{b}. \item The low-density regime supports the formation of $v$=1 and $v$=2 masers over a large range of brightnesses, including a range of weak masers at densities $< 10^9$cm$^{-3}$, while the high-density regime does not support the range of weakest masers. This is in overall agreement with observed trend \textit{c}. However, the observational trend \textit{c} may partly be due to sensitivity, as the $v$=3 line in general is expected to be weaker than the $v$=1 and 2, so if observed to the same sensitivity levels we would expect to detect the $v$=3 lines in either more nearby, or brighter sources \end{enumerate} In summary, the statistical differences between line ratios for detections versus non-detections in the 43 GHz $v$=3 line can serve as a useful probe of the density conditions in the circumstellar shell. The large spread in the $\log\Big[\frac{ I(\textnormal{43~GHz}~v=1) }{ I(\textnormal{86~GHz}~v=1) }\Big]$ and $\log\Big[ \frac{ I(\textnormal{43~GHz}~v=3) }{ I(\textnormal{43~GHz}~v=1) }\Big]$ distributions are easily explained by this model. Moreover, the model predicts that the 86 GHz $v$=3 line should be brighter than the 43 GHz $v$=3 line, which can be tested observationally. \begin{figure*}[t] \figurenum{8} \includegraphics[scale=0.475]{consolidated_vs_FA.eps} \caption{Comparison between $I(\textnormal{43~GHz~}v=1)$, $I(\textnormal{43~GHz~}v=2)$, and $I(\textnormal{86~GHz~}v=1)$ versus $F_A$ (left, center and right, respectively). Crosses represent sources where 43 GHz $v$=3 emission was detected.} \label{fig_consolidated_correlate_msx_a} \end{figure*} \subsection{Infrared Correlations} The exact nature and the importance of infrared radiation providing at least part of the pumping energy of the SiO masers have long been discussed \citep{1974ApJ...194L..97K, 1976PASJ...28..307D, 1979ApJ...231..124C, 1981A&A...102...65B, 1984ApJ...284..751L, 1987A&A...175..164B, 1994A&A...285..953B, 2007ApJ...669..446N}. Both 4$\mu$m and 8$\mu$m photons are considered, as the 8$\mu$m could excite via the $\Delta v$=1 and the 4$\mu$m the $\Delta v$=2 transitions. Indeed, \citet{1987A&A...175..164B} demonstrated a correlation between the SiO maser intensity and 8$\mu$m intensity, and \citet{1996AJ....111.1987C} found a correlation between the maser intensity and 4$\mu$m intensity. For the ATCA data, weak correlations\footnote{All correlations were tested using the Spearman's rank correlation test and only correlations at the 99\% significance level are considered.} are found between the integrated flux densities $I(\textnormal{43~GHz}~v=1)$, $I(\textnormal{43~GHz}~v=2)$ and $I(\textnormal{86~GHz}~v=1)$ and the MSX flux densities (not corrected for extinction or reddening) $F_A$, $F_C$, $F_D$ and $F_E$ measured in Jy. The strongest correlations are found with $F_A$ (8.28$\mu$m), giving correlation coefficients of 0.38, 0.40 and 0.39 for the 43 GHz $v=1$, $v=2$ and $v=3$ lines, respectively (see Fig.\ \ref{fig_consolidated_correlate_msx_a}). The strength of the correlations decreases toward longer MSX wavelengths, to the point where the correlation between 86 GHz $v=1$ against $F_E$ falls below the 99\% confidence level. That these correlations are stronger with $F_A$ is expected, given that the wavelength of Band A matches the wavelengths of the fundamental vibrational-rotational pumping transition, and that the photon energies of the longer wavelength bands are insufficient to contribute to the pumping \citep[e.g.,][and references above]{1996A&A...314..883B}. The 43 GHz correlations are notably strengthened if only sources with detected 43 GHz $v=3$ emission are considered, resulting in correlation coefficients of 0.60 and 0.62 for the $v=1$ and $v=2$ lines respectively. No statistically significant correlation is found between the 86 GHz $v=1$ line and the MSX fluxes using the 43 GHz $v=3$ detected sample, which may be partially due to the smaller sample size compared to that of the 43 GHz v=1 and v=2 lines. No correlations are found between the integrated flux densities or line ratios and the MSX [A]--[D] color, which is the primary color used to select BAaDE sources \citep{in...prep}. A weak correlation with correlation coefficient of 0.43 (0.38) is found between $\frac{I(\textnormal{43~GHz}~v=1)}{I(\textnormal{86~GHz}~v=1)}$ and [C]--[E] ([A]--[E]), see Fig.\ \ref{fig_consolidated_correlate_msx_e_c}, which are the alternative colors BAaDE sources were chosen from. However, this correlation is markedly diminished if only a few points are removed from the plot. This correlation may be mostly due to an anti-correlation between $I(\textnormal{86~GHz}~v=1)$ and [C]--[E] (correlation coefficient of $-0.33$, see bottom of Fig.\ \ref{fig_consolidated_correlate_msx_e_c}) but is similarly diminished with the removal of a few data points. Unlike \citet{2007ApJ...669..446N}, no statistically significant correlations are found between $\frac{I(\textnormal{43~GHz}~v=2)}{I(\textnormal{43~GHz}~v=1)}$ and [C]--[E]. Their data include redder sources (e.g. $-1<[\textnormal{C}]-[\textnormal{E}]<1$ whereas the BAaDE ATCA sample is restricted to $-1< [\textnormal{C}]-[\textnormal{E}]< 0$) so a stronger relation could reside in the redder region, or by sampling a larger range, the correlation could be more prominent. Finding statistically significant correlations with infrared colors in the MSX band may require larger samples, as the flux densities in the MSX bands vary with stellar phase, and are not observed simultaneously with the SiO maser data. In addition, the presence of the silicate dust feature at 9.7$\mu$m could affect the $F_A$ and possibly the $F_C$ values. Moreover, the BAaDE ATCA sample selection is biased toward the brighter SiO maser sources, thereby limiting the dynamic range when searching for a correlation between the maser flux density and other variables. A more robust statistical study will be performed with observations from the BAaDE survey which will contain thousands of simultaneous observations of the 43 GHz $v$=0, 1, 2 \& 3 using the VLA and thousands of simultaneous observations of the 86 GHz $v$=0, 1 \& 2 lines. \begin{figure} \figurenum{9} \includegraphics[scale=0.5]{consolidated_vs_C_minus_E.eps} \caption{Top: Comparison between $\frac{I(\textnormal{43~GHz~}v=1) }{ I(\textnormal{86~GHz~}v=1) }$ and [C]--[E]. Bottom: Comparison between $I(\textnormal{86~GHz~}v=1)$ and [C]--[E].} \label{fig_consolidated_correlate_msx_e_c} \end{figure} | 18 | 8 | 1808.02899 |
1808 | 1808.09373_arXiv.txt | Blueshifted absorption lines in the X-ray spectra of AGN show that ultra-fast outflows with typical velocities $v \sim 0.1c$ are a common feature of these luminous objects. Such powerful AGN winds offer an explanation of the observed $M-\sigma$ relation linking the mass of the supermassive black hole and the velocity dispersion in the galaxy’s stellar bulge. An extended \xmm\ study of the luminous Seyfert galaxy \pg\ recently revealed a variable multi-velocity wind. Here we report the detection of a short-lived, ultrafast inflow during the same observation. Previous reports of inflows used single absorption lines with uncertain identifications, but this new result identifies an array of resonance absorption lines of highly ionised Fe, Ca, Ar, S and Si, sharing a common redshift when compared with a grid of realistic photoionization spectra. The redshifted absorption arises in a column of highly ionized matter close to the black hole, with a line-of-sight velocity, $v \sim 0.3c$, inconsistent with the standard picture of a plane circular accretion disc. This may represent the first direct evidence for chaotic accretion in AGN, where accretion discs are generally misaligned to the black hole spin. For sufficient inclinations, the Lense-Thirring effect can break the discs into discrete rings, which then precess, collide and shock, causing near free-fall of gas towards the black hole. The observed accretion rate for the reported infall is comparable to the hard X-ray luminosity in \pg\, suggesting that direct infall may be a significant contributor to inner disc accretion. | It is now well-established that a supermassive black hole (SMBH) lies in the centre of most galaxies, and further that it accretes material through a disc. Over the past 15 years, observations with a new generation of X-ray Observatories (Jansen 2001, Mitsuda:2007) have revealed ultra-fast outflows (UFOs), probably launched from regions of the disc accreting at super--Eddington rates (King \& Pounds 2003) where the momentum in the radiation field released by accretion can overcome the inward pull from the black hole's gravity. UFOs appear to be a common component of luminous AGN (Tombesi \et\ 2010, 2011, Gofford \et\ 2013). With typical velocities of $v\sim 0.1c$, these highly ionized winds imply significant feedback onto the surrounding interstellar gas, offering a likely explanation of the $M-\sigma$ relation (Ferrarese \et\ 2000, Gebhardt \et\ 2000), by simultaneously constraining the growth of a supermassive black hole and star formation in the central bulge of its host galaxy (King 2003,2005; King \& Pounds 2015). The archetypal UFO is found in the luminous Seyfert galaxy PG1211+143 (Pounds \et\ 2003). To further explore the properties of this powerful UFO, the {\it XMM-Newton} X-ray Observatory carried out seven full-orbit ($\sim 100$\,ks) observations over 5 weeks in 2014, with a total on-target exposure of $\sim$ 630 ks. Full details of observing times, data reduction procedures and count rates are given in Lobban \et\ (2016). The stacked data revealed a surprisingly complex spectrum, with the hard X-ray pn camera (Strueder \et\ 2001) finding multiple blue-shifted absorption lines, identified with highly ionized Fe between $\sim$ 6.6keV and 8.8 keV and outflow velocities of $\sim 0.06c$, $\sim 0.13c$ and $\sim$ 0.18c (Pounds \et\ 2016a). Independent support for the multiple outflow velocities was found in higher resolution spectra (Pounds \et\ 2016b) using soft-X-ray data from the co-aligned Reflection Grating Spectrometer (RGS :den-Herder \et\ 2001). While all previous UFO detections report a single velocity, although with repeated observations sometimes finding a different value, the 2014 \xmm\ observation of \pg\ was inconsistent with the unique and stable outflow expected from a static axisymmetric accretion disc (Shakura and Sunyaev 1973, Pounds \et\ 2017). Furthermore, a recent orbit-by-orbit study of the RGS data (Reeves \et\ 2018) has shown significant inter-orbit variability in outflow column densities over the 5-week campaign, perhaps providing a further indication of short-term inner disc variability. The highly ionized state of such ultra-fast AGN winds limits strong absorption to the heavier metals, where features of Fe stand out because of its high astrophysical abundance. For that reason, all current UFO discoveries essentially rest on the detection of blue-shifted Lyman-$\alpha$ and/or He-$\alpha$ resonance absorption lines of highly ionized Fe (respective rest energies of 6.70 and 6.96 keV). While spectral modelling, typically over the 2--10 keV band, also includes absorption in lighter metals such as Ca, S, and Si (eg. Pounds and Page 2006), little attention has been given to other spectral features below $\sim 6$\,keV, where red-shifted Fe K absorption lines might be seen. The few historical exceptions (Nandra \et\ 1999, Dadina \et\ 2005, Reeves \et\ 2005, Cappi \et\ 2009, Giustini \et\ 2017) are of isolated absorption lines where the identification -- and hence velocity -- of the absorber remained unclear. In the most detailed report (Dadina \et\ 2005) two (of five) Beppo-Sax observations of the Seyfert 1 galaxy Mkn 509 detected an absorption line at $\sim 5.5$\,keV (rest-frame), suggesting the feature was variable on timescales as short as $\sim 20$\,ks, which the authors argued was more easily reconciled with inflowing matter than with a pure gravitational redshift or failed jet. Of particular relevance to the present study are unidentified absorption lines at $\sim 4.56$\,keV and $\sim 5.33$\,keV in the {\it Chandra} observation of PG1211+143 (Reeves \et\ 2005). While these earlier reports hint at fast-evolving and complicated dynamics in the inner disc, so far none has provided compelling evidence for a high velocity inflow that could represent a direct challenge to the standard picture of a circular, planar disc slowly accreting on to the central black hole. However, recent theoretical arguments suggest that such infall may be common in AGN, where allowing for the possibility that the accretion disc may be misaligned to the central black hole spin reveals new physical effects. Numerical simulations of misaligned discs around a spinning black hole have shown that the Lense-Thirring effect can overwhelm the internal communication of angular momentum within the disc. This causes the disc to break, leading to individual rings of gas which precess effectively independently -- called disc tearing (Nixon \et\ 2012, Nixon 2015, Dogan \et\ 2015, Nealon \et\ 2015, Dogan \et\ 2018). When these rings interact, their opposed velocity fields create shocks which rob the gas of rotational support, allowing it to fall inwards towards the black hole where residual angular momentum causes the gas to re-circularise at a smaller radius. Motivated by these theoretical ideas and the paucity of relevant observations, we have carried out a thorough search for rapidly inflowing matter, on a range of timescales, as part of an on-going orbit-by-orbit analysis of the hard X-ray data from the 2014 {\it XMM-Newton} observation of PG1211+143. \section {Redshifted absorption} The new analysis is based primarily on the high throughput pn camera, while checking for consistency - as necessary - with simultaneous spectral data from the MOS camera (Turner \et\ 2001) and higher resolution spectra from the RGS. As in the published outflow studies, spectral modelling uses the {\tt XSPEC v12} software package (Arnaud 1996), with all spectral fits including absorption due to the line-of-sight Galactic column of N$_{H}=2.85\times10^{20}\rm {cm}^{-2}$ (Murphy \et\ 1996). We assume the reverberation black hole mass estimate of $4\times 10^{7}$\Msun\ (Kaspi \et\ 2000), and a galaxy redshift of 0.0809 (Marziani \et\ 1996). The starting point in modelling each pn data set is the stacked X-ray spectrum (Pounds \et\ 2016a), where the continuum from 1--10 keV is modelled by a double power law, with discrete spectral features imposed by overlying matter being compared with grids of pre-computed photoionized absorption and emission spectra, based on the {\tt XSTAR code} (Kallman \et\ 1996). To facilitate comparison with previous outflow analyses (eg Pounds \et\ 2016a), for the pn and MOS data we use the publically available grid25 (v21), which assumes a power law continuum of $\Gamma$ = 2 and a fixed turbulence velocity 200 km s$^{-1}$, and custom-built grids derived from the observed source spectrum (Pounds \et\ 2016b) for the RGS spectral analysis. Free parameters in the spectral fitting include the ionization, column density and observed redshift (or blueshift) of photoionized matter, with continuum normalisation also free to vary on re-fitting. We find no significant evidence for strongly redshifted absorption \footnote {Potential confusion with strong Fe K emission constrains the resolution of Fe K absorption lines with present data, setting a lower redshift limit $\sim$ 0.15} in the first observing period (spacecraft orbit number 2652, or rev2652), but that outcome changed for rev2659, just two weeks later. \subsection{A strong inflow in rev2659} Modelling the 1--10 keV pn spectrum from rev2659 revealed a strong absorption component at the extreme redshift $\sim$0.483$\pm$0.008. With a high column density (N$_{H} \sim 4.3\times10^{23}\rm cm^{-2}$) and high ionization parameter (log $\xi \sim 3.7$ erg cm s$^{-1}$), the red-shifted absorber is highly significant, its inclusion improving the fit by $\Delta \chi^{2}$ of 19/3. The corresponding probability of a false detection (P=4.4 $\times 10^{-4}$) is a physically robust measure, being based on a comparison of the observed broad band absorption with a set of physically realistic model photoionization spectra. To check that the model was not impaired by some unknown data artefacts, the redshifted component was removed from the {\tt XSTAR} model and the ratio of data-to-model visually examined. Figure 1 (top) reproduces that ratio plot, showing an array of absorption lines identified with resonance transitions in H- and He-like ions of Fe, Ca, Ar, S and Si. The two strongest absorption lines are identified with Lyman-$\alpha$ and He-$\alpha$ of Fe, with individual redshifts closely matching the {\tt XSTAR} model component, while the array as a whole yields a weighted mean observed redshift of 0.476$\pm$0.005. The Ca line at $\sim$2.6 keV is most likely a blend of H- and He-like resonance lines. Table 1 lists each observed line energy, obtained by scanning a narrow negative Gaussian across the ratio plot, with the individual line widths set comparable to the pn detector resolution at each photon energy. We note the detector resolution corresponds to an intrinsic line width of order 1500 km s$^{-1}$ at 5 keV, substantially greater than the preferred {\tt XSTAR} model value, indicating the observed absorption line widths are essentially unresolved in these data. \begin{figure} \centering \includegraphics[width=6.7cm, angle=270]{fig1_redshift_ratio.ps} \centering \includegraphics[width=6.03cm, angle=270]{fig1_redshift_model.ps} \caption{(top)Ratio of pn data-to-model for the rev2659 observation when the absorption component with redshift gz$\sim$0.483 is removed and the model re-fitted. Absorption lines identified with resonance K-shell transitions in Fe, Ca, Ar, S and Si yield a weighted mean redshift of 0.476$\pm$0.005. (lower) The initial spectral model, with a hard power law (dark blue) showing the imprint of the redshifted absorber, the unabsorbed soft power law found in difference spectra (green), and emission from an ionized reflector (red) and the photoionized outflow (light blue).} \end{figure} \begin{table} \centering \caption{Narrow Gaussian lines sequentially fitted to the identified absorption features in the rev2659 pn data shown in Figure 1. Line widths were set comparable to the relevant pn detector resolution and all line energies are in keV. The proposed identification and corresponding redshift for each line are listed, together with the related improvement in $\Delta \chi^{2}$ after re-fitting. The relatively high redshift of the Ca Ly-$\alpha$ line may be explained by a blend with the corresponding He-$\alpha$ line, though its low statistical weight does not affect the array redshift} \begin{tabular}{@{}lcccc@{}} \hline line i.d. & obs energy & lab energy & obs redshift & $\Delta \chi^{2}$ \\ \hline Fe Ly-$\alpha$ & 4.701$\pm$0.010 & 6.96 & 0.481$\pm$0.004 & 8/2 \\ Fe He-$\alpha$ & 4.511$\pm$0.011 & 6.703 & 0.486$\pm$0.004 & 15/2 \\ Fe Ly-$\beta$ & 5.543$\pm$0.020 & 8.25 & 0.488$\pm$0.005 & 4/2 \\ Fe He-$\beta$ & 5.27$\pm$0.08 & 7.88 & 0.495$\pm$0.023 & 1/2 \\ Ca Ly-$\alpha$ & 2.67$\pm$0.025 & 4.09 & 0.532$\pm$0.014 & 2/2 \\ Ca He-$\alpha$ & 2.67$\pm$0.025 & 3.903 & 0.462$\pm$0.014 & 2/2 \\ Ar Ly-$\alpha$ & 2.257$\pm$0.009 & 3.32 & 0.471$\pm$0.006 & 5/2 \\ Ar He-$\alpha$ & 2.09$\pm$0.09 & 3.140 & 0.50$\pm$0.06 & 1/2 \\ S Ly-$\alpha$ & 1.814$\pm$0.010 & 2.62 & 0.452$\pm$0.016 & 4/2 \\ Si Ly-$\alpha$ & 1.382$\pm$0.007 & 2.006 & 0.452$\pm$0.008 & 5/2 \\ Si He-$\alpha$ & 1.283$\pm$0.031 & 1.865 & 0.454$\pm$0.025 & 7/2 \\ \hline \end{tabular} \end{table} Allowing for the cosmological redshift of \pg\ (0.0809) the observed redshift of $0.48\pm0.01$ corresponds to a Doppler-corrected inflow velocity v$\sim$0.30$\pm$0.01c. As this velocity is substantially higher than any of the 3 contemporary outflow velocities, we suggest the matter being a 'failed outflow' relatively unlikely, since a lower velocity outflow -launched from a larger radius - would have too much angular momentum to fall back at the higher velocity. The more interesting alternative is absorption in matter accreting from out-of-the-plane of the disc, and in line of sight to the hard X-ray continuum source (conventionally assumed to be a compact hot corona above the inner disc). In the context of matter falling freely towards the SMBH, equating the infall velocity to the local free-fall velocity places the absorber near 20 $\rg$, where $\rg$ is the gravitational radius. We note that $\sim5\%$ of the observed redshift would be gravitational at that location. Independent support for the above finding is potentially available from the co-aligned MOS camera (Turner \et\ 2000), which was operated throughout the 2014 campaign. An initial check on the rev2659 MOS data did indeed find a highly redshifted absorption component, consistent with that reported from the pn data analysis, although the outflow components are less well constrained due to the lower sensitivity of the MOS above $\sim$6 keV. To obtain a common spectral fit with both data sets, we therefore re-fitted the pn model, described above, with the addition of the MOS data, and allowing the joint parameters to vary. The outcome was positive, with the redshifted absorber increasing in significance ($\Delta\chi^{2}$ of 26/3), while retaining a similar column density (N$_{H}$$\sim$$6.3\times10^{23}\rm cm^{-2}$) and ionization parameter (log $\xi \sim 3.4$ erg cm s$^{-1}$). The false detection probability for the redshifted absorption component in the joint spectral fit is now further reduced with P $\leq10^{-5}$. \subsection{Absorption in the RGS soft x-ray spectrum of rev2659} Finally, higher resolution soft X-ray spectra (albeit with lower count rates) from the RGS instrument, also co-aligned with the pn camera, are examined to check for possible lower ionization matter in the infalling stream. We recall that detection of cooler embedded matter provided important confirmation of multiple velocities in the outflow spectrum of \pg\ (Pounds \et\ 2016b). Again starting with the mean 2014 outflow spectrum, the rev2659 RGS data was re-modelled, now allowing for an additional red-shifted absorber. That search was successful, with a soft X-ray absorption component ($\Delta \chi^{2}$ of 14/3; P=3 $\times 10^{-3}$) being found at the same extreme redshift ($\sim 0.483 \pm 0.003$), though with a much lower column density (N$_{H} \sim 3\times10^{21}\rm cm^{-2}$) and ionization parameter (log $\xi \sim 0.95$ erg cm. s$^{-1}$). Inclusion of the redshifted absorber improved the overall 12--35 \AA\ soft X-ray spectral fit by $\Delta \chi^{2}$ of 14/3, perhaps again indicating higher density matter embedded in the primary inflow. Interestingly, a soft X-ray emission component is also marginally significant ($\Delta \chi^{2}$ of 6/2) with the same redshift. Confidence in the highly red-shifted soft X-ray absorption and emisiion components is supported by the detection of several key features in the spectral residuals, when the high redshift components are removed from the {\tt XSTAR} model fit (Figure 2). While the plot is much 'noisier' than for the pn hard X-ray data, the strong 1s-2p resonance absorption lines of He-like O and Ne are both indicated at a redshift consistent with the {\tt XSTAR} model value, while the broad absorption line at $\sim$24.2 \AA\ corresponds to a similar redshift when identified with the Fe-M UTA, albeit with lower precision as the rest wavelength of the UTA depends on the ionization distribution. Narrow emission lines identified with of O Ly-$\alpha$ and the forbidden line of the OVII triplet also match the same inflow redshift. \begin{figure} \centering \includegraphics[width=6.7cm, angle=270]{rev2659_rgs_red_spectrum.ps} \caption{Ratio of RGS data-to-model for rev2659 when absorption and emission components with redshift z$\sim$0.48 are removed. Absorption lines identified with the 1s-2p resonance lines of He-like O and Ne are detected at a redshift consistent with the {\tt XSTAR} model value. The broad absorption line at $\sim$24.2 \AA\ corresponds to a similar redshift when identified with the Fe-M UTA, although with lower precision as the rest wavelength of the UTA depends on the ionization distribution. Narrow emission lines of OVIII Ly-$\alpha$ and the forbidden line of the OVII triplet also match the model redshift. The line markers are located at wavelengths corresponding to a redshift of 0.48. } \end{figure} Two less strongly red-shifted components are also found in the RGS spectrum, at z$\sim$0.19$\pm$0.01 ($\Delta\chi^{2}$ of 17/3: P=$7\times 10^{-4}$) and z$\sim$0.17$\pm$0.01 ($\Delta\chi^{2}$ of 15/3: P= $1.8\times 10^{-3}$). While the lower implied velocities might simply represent a line of sight at a larger angle to the flow stream, the required large deviation in the flow vector seems unlikely so close to the SMBH. Alternatively, assuming each flow is along the line of sight, the corresponding inflow velocities would $\sim$0.080c and $\sim$0.097c. Both components are highly ionized, with log$\xi$ $\sim$3.4 erg cm s$^{-1}$, similar to that for the pn detection, but with respective column densities of N$_{H} \sim 4.6\times10^{21}\rm cm^{-2}$ and N$_{H} \sim 1.1\times10^{21}\rm cm^{-2}$, much less than in the pn detection. An intriguing speculation is that the lower velocity and column density of these soft X-ray inflows represent the primary, highly ionized flow seen along a different sight line (in this case to the soft X-ray continuum), providing an upstream view through an accelerating flow approaching the SMBH. Support for that conjecture is provided by the absence of a higher energy spectral counterpart to the lower redshift soft X-ray absorbers in the pn data. If a true up-stream measure, the lower velocity of the highly ionized soft X-ray absorber, v$\sim$0.10c, would correspond to a radial distance for the absorber of $\sim$200 $R_g$, conceivably on a sight-line to the larger scale, thermal disc emission. The much lower column density might then be qualitatively consistent with a converging - as well as accelerating inflow. While such deductions are limited by having only two spectral data points, the potential of future observations with higher spectral resolution and photon grasp is clear. | The detection of strongly red-shifted X-ray absorption in data from an extended \xmm\ observation of the luminous Seyfert galaxy \pg\ has provided compelling evidence for the inflow of matter at $v \sim 0.3c$. We suggest a possible origin for the inflow in the form of disc tearing, arising when an accretion disc misaligned to the spin plane of the SMBH precesses due to the Lense-Thirring effect (frame-dragging), with the precession close to the hole being fast enough for individual rings to break off (Nixon \et\ 2012, Dogan \et\ 2018). Collisions between independently precessing rings then cause shocks and a loss of rotational support and the subsequent infall of disc gas. If this picture is confirmed, there are strong implications for the fuelling of AGN, and the evolution of SMBH masses and spins. For example, misaligned or `chaotic' accretion would in general allow the spin to remain low (King \et\ 2008), allowing efficient mass growth for the black hole. This might then remove the need for massive SMBH seeds in the early Universe (King \& Pringle 2006, 2007). The transient nature of the strong inflow detected in the 2014 {\it XMM-Newton} data, with variability on a timescale of hours, helps explain why a compelling detection has not been reported before. It may also be relatively rare to have a sight-line to the hard X-ray source that passes through infalling matter so close to the black hole, and hence with such a high particle density, as we found in rev2659. Looking ahead, X-ray absorption spectra obtained over different sight lines, with high cadence and resolution, together with parallel advances in theory and computation, offer the exciting prospect of mapping the fine structure of ionized flows and accretion close to the SMBH. The high spectral resolution of the microcalorimeter planned for the forthcoming Japanese {\it XRISM} mission will be well matched for studying the highly ionized matter characteristic of fast inflows. Such new data together with new and a broader search of archival \xmm\ observations should provide the best opportunity for further exploration of accretion disc physics and black hole growth in AGN prior to the launch of {\it Athena}, a decade hence. | 18 | 8 | 1808.09373 |
1808 | 1808.04298_arXiv.txt | {The motion of massless particles on the background of a toroidal topological black hole is analyzed in the context of conformal Weyl gravity. Null geodesics, in terms of the Jacobi elliptic functions, are found analytically. In addition, the Sagnac effect in this space-time is characterized, and we find a strong condition in the theory's parameters that is required for its existence. \PACS{02.30.Gp, 04.20.-q, 04.20.Fy, 04.20.Gz, 04.20.Jb, 04.70. Bw} } \tableofcontents | Conformal Weyl gravity (CWG) was born of an attempt to unify gravity and electromagnetism based on the principle of local invariance of a manifold, described by the metric $g_{\mu \nu}(x)$, under the change \begin{equation} \label{f1} g_{\mu \nu}(x)\rightarrow\Omega^2(x)\,g_{\mu \nu}(x), \end{equation}where $\Omega(x)$ is a smooth, strictly positive function \cite{W17,W18A,W18B,B21}. The CWG theory can be obtained from the conformally invariant action \begin{equation} \label{action} I_W=2\,\alpha_w \, \int {\rm d}^4 x\,\sqrt{-g}\, \left[ R_{\mu \nu}\,R^{\mu \nu}-\frac{1}{3} \left(R^{\mu }_{\mu}\right)^2\right], \end{equation} where $\alpha_w$ is a dimensionless parameter chosen to be positive if (\ref{action}) is a positive definite Euclidean action. The vacuum field equations associated with this action are solved by the static, spherically symmetric line element given by \cite{riegert84,MK89,MK91,MK91B,MK94} \begin{equation} {\rm d}\tilde{s}^{2}=-B(\tilde{r})\,{\rm d}\tilde{t}^{2}+\frac{{\rm d}\tilde{r}^{2}}{B(\tilde{r})}+ \tilde{r}^{2}({\rm d}\tilde{\theta}^{2}+\sin^{2}\tilde{\theta}\, {\rm d}\tilde{\phi}^{2}), \label{metrweyl} \end{equation} where the coordinates are defined in the range $-\infty < \tilde{t} < \infty$, $\tilde{r}\geq0$, $0\leq\tilde{\theta}\leq\pi$, $0\leq\tilde{\phi}\leq 2\pi$, and the lapse function $B(\tilde{r})$ is given by \begin{equation} B(\tilde{r})=1-\frac{\tilde{\beta}\,(2-3\tilde{\beta}\,\tilde{\gamma})}{\tilde{r}}-3\tilde{\beta}\,\tilde{\gamma}+\tilde{\gamma} \tilde{r} - \tilde{k} \tilde{r}^{2}. \label{lpasweyl} \end{equation} Here $\tilde{\beta}$, $\tilde{k}$ and $\tilde{\gamma}$ are positive constants associated with the central mass, cosmological constant and the measurements of the departure of the Weyl theory from the Einstein - de Sitter, respectively. Clearly, taking the limit $\tilde{\gamma}=0=\tilde{k}$ recovers the Schwarzschild case so that we can identify $\tilde{\beta} = M$. A study of the basis and properties, together with applications of the motion of massive and massless particles in this geometry can be found, for example, in \cite{edery98,PI04,PII04,sultana10,sultana12,vo13,said13,said,Lu12,Payandeh:2012mj}, and can be obtained using the standard Lagrange procedure \cite{chandra,COV05,shutz,OSLV11,jaklitsch,VSOC13,Halilsoy:zva,LSV11}, which allows a Lagrangian $\mathcal{L}$ to be associated with the metric and then, the equation of motion to be obtained from the Lagrange's equations \begin{equation} \dot{\Pi}_{q} - \frac{\partial \mathcal{L}}{\partial q} = 0, \label{lageq} \end{equation} where $\Pi_{q} = \partial \mathcal{L}/\partial \dot{q}$ are the conjugate momenta to the coordinate $q$, and the dot denotes a derivative with respect to the affine parameter $\tau$ along the geodesic. Thus, in Sec. \ref{TTSEc}, and following the procedure performed by Klemm \cite{klemm98}, we perform analytical continuations to obtain a non-trivial topology associated with toroidal topological black holes coming from CWG. In particular, we focus on the toroidal AdS black hole. Other studies associated with topological black holes can be found, for example, in Refs. \cite{olivera,ligeia}, among others. Then we obtain the conserved quantities together with the equations of motion for massless particles on these manifolds. In Sec. \ref{RMTT} the radial motion is analyzed for photons going to spatial infinite or to the singularity, while Sec. \ref{AMTT} is devoted to obtaining analytically the trajectory for photons with non-zero angular momentum, for which we employ an analysis in terms of Jacobi elliptic functions. In Sec. \ref{SETT} we apply the methods outlined by Sakurai, Tartaglia, Rizzi \& Ruggiero, among others, to obtain an analogy to the Aharanov-Bohm effect to describe the Sagnac effect for this space-time. Finally, in Sec. \ref{STT} we conclude and summarize our results. | \label{STT} In this paper we have studied the null structure of the geodesics for a toroidal topological space-time which surrounds a black hole in the conformal Weyl gravity. First, we obtained an explicit behavior of radial photons to conclude that, while no changes in the motion to the singularity with respect to the Schwarzschild anti-de Sitter counterpart is found, there is a non-trivial coordinate time $t_{\infty}$ for the description of the motion to the spatial infinite (see Eq. (\ref{tinf})). A similar result was obtained by Villanueva \& V\'asquez but in the context of the Lifshitz space-time \cite{Villanueva:2013gra}. Next, following the standard Lagrangian procedure, we have obtained analytically the trajectories of the confined and unbounded angular motion for photons in terms of Jacobi elliptic functions, Eqs. (\ref{trayrt}) and (\ref{radunb}), and then we have shown our results in Fig. \ref{conf} and Fig. \ref{unbfig}, respectively. Obviously, these trajectories depend on the impact parameter $b$ and, due to the topology, always fall to the singularity, which is a characteristic of AdS space-times. Finally, the Sagnac effect has been studied for this topological space-time. Our result is consistent with those obtained previously in other geometries in the sense that no Sagnac effect arises for a non-rotating frame. In addition, we have found a strong condition for its existence, which depends on the theory's parameters $\{\eta, \ell\}$ (in $r_+$ and $\Omega_{\ell}$) as well as on the radius of the circular orbit $R$. This condition is given in Eq. (\ref{condrw}), c.f. \begin{equation} \nonumber \Omega<\Omega_R \equiv \Omega_{\ell}\,\sqrt{1-\left(\frac{r_+}{R}\right)^3}, \end{equation}for which the upper limit for the angular velocity $\Omega_R$ was plotted in Fig. \ref{om} as a function of the radius $R$. Finally, our study provides a simple physical visualization of the null trajectories and their main issues, and complements other studies carried out in the standard and/or trivial topology \cite{vo13,sultana14}. \appendix | 18 | 8 | 1808.04298 |
1808 | 1808.04584_arXiv.txt | Radio emission is observed from many cataclysmic variable stars. It is distinguished by a large variability of both intensity and polarization, up to complete cessation. Duration of quiet periods can be many times longer than the orbital period of the binary star. Among the sources with continuously observed radio emission, one can note AM~Her, AR~UMa, and AE~Aqr \cite{Mason2007ApJ...660..662M}. Radiation is detected at the frequencies of the order of units and tens of GHz (the range of operation of VLA), but only a small number of systems manifest themselves in this range \cite{Coppejans2015MNRAS.451.3801C}. Among cataclysmic variables with magnetic field, it is customary to distinguish polars and intermediate polars \cite{Warner2003cvs..book.....W}. Such systems usually harbor a white dwarf (accretor) and a late-type star (donor), typically a red dwarf. Orbital period of the binary is several hr. The polars and intermediate polars are distinguished by degree of polarization of the observed radiation. This characteristic is associated with the magnitude of the magnetic field of white dwarf: for the polars it exceeds $10^6$~G, while for intermediate polars its range is $10^5\mathdash10^6$~G. The magnitude of the field $10^6$~G also separates two types of accretion flows: in intermediate polars accretion disks can form, while in polars the radius of the magnetosphere of the white dwarf is so large that the disk does not form \cite{Zhilkin2012PhyU...55..115Z}. Since the first radio observations of the polar AM~Her \cite{Chanmugam1982ApJ...255L.107C}, several authors suggested different mechanisms for generation of radio emission of polars and intermediate polars in the frequency range from one to tens of GHz: cyclotron one on non-thermal or relativistic electrons (gy\-ro\-synch\-ro\-tron and synchrotron, respectively) \cite{Chanmugam1982ApJ...255L.107C,Dulk1983ApJ...273..249D,Benz1989A&A...218..137B} and maser one \cite{Dulk1983ApJ...273..249D}. Cyclotron radiation is indeed observed at the wavelengths of $\sim 6000 \mathdash 7000$~{\AA} and is used for the estimates of the magnitude of magnetic field. Maser amplification is involved as a mechanism for generation of radio flares. In this paper, we suggest one more mechanism for generation of radio emission based on cyclotron radiation of electrons in the fluctuations of magnetic field. The presence of fluctuations is associated with the setting of wave Alfv\'{e}n turbulence in the plasma. This paper continues investigation of the effects of Alfv\'{e}n turbulence in cataclysmic binary systems \cite{Kurbatov2017PhyU...60..798K}. The paper is structured as follows. In \S2 we describe accretion flow in polars and estimates of the characteristics of the flow using as an example AM~Her. In \S3 we present examples of observational data and consider various ways of generation of radio emission suggested earlier. In \S4 we suggest the mechanism of cyclotron emission on magnetic field fluctuations. In \S5 the fluxes of emission from the accretion stream in a polar is calculated within the framework of the suggested emission mechanism. Conclusions are presented in \S6. \section {General picture of the matter flow in the polars} \label{sec:general_review} In the polars, accretors are $\sim 1$~$M_\odot$ white dwarfs with magnetic field $\gtrsim 10^6$~G. Donor star is, most often, an M-type dwarf of lower mass. As a result of Roche lobe overflow by the donor, the matter flows to the main component. The spin of components is synchronized with the orbital revolution. Orbital period of AM~Her is $P_\mathrm{orb} = 3.09$\,hr. Assuming that accretor and donor masses are $M_\mathrm{a} = 0.7~M_\odot$ and $M_\mathrm{d} = 0.3~M_\odot$, respectively, separation of components may be estimated as $A = 1.07$~$R_\odot$. In such a case, the distance from accretor to the point L$_1$ is $R_\mathrm{L_1} \approx A / [ 1.0015 + (M_\mathrm{d}/M_\mathrm{a})^{0.4056} ] = 0.58$~$R_\odot$ \cite{Silber1992PhDT.......119S}. After \cite{Gawronski2018MNRAS.475.1399G}, we assume that the distance to AM Her is $D = 88.6$~pc. The rate of accretion in typical cataclysmic binaries is usually estimated as $\dot{M} = 10^{-9} \mathdash 10^{-8}$ $M_\odot/$yr. In general, it is related to the rate of the outflow of the matter from the donor in a nontrivial manner. The reason for this is that the matter leaving the donor is affected not only by the gravitational field of the accretor, but also by the interstellar medium and magnetic field. However, taking into account that the polar AM~Her in this paper acts as a source of typical accretion flow parameters only, we take $10^{-9}$~$M_\odot/$yr as an estimate of the rate of mass exchange. In the cataclysmic variables without magnetic field, the main mass exchange occurs through the vicinity of the Lagrangian point L$_1$ (though, in principle, the possibility of a flow through the point L$_3$ is not excluded \cite{Sytov2007ARep...51..836S}). The matter, leaving L$_1 $ as an accretion stream, acquires an angular momentum with respect to the accretor and forms an accretion disk under the action of viscosity. The presence of the magnetic magnetic field of primary component can strongly change this picture: in the region where the magnetic pressure exceeds the dynamic one, gas flow is controlled by magnetic field. In the case of polars, magnetic field is strong enough to exclude the possibility of formation of accretion disk \cite{Warner2003cvs..book.....W,Bisikalo2013gasdyntdz}. This is confirmed by numerical modeling too: accretion flow has the shape of a stream that starts at the point L$_1$, reaches the boundary of the magnetosphere and then flows into the polar region of the white dwarf along magnetic field lines \cite{Zhilkin2010ARep...54..840Z,Zhilkin2012PhyU...55..115Z}. This, however, is not the only possible scenario of accretion. Some observational data provide evidence in favor of the accretion flow in the shape of a curtain (see \cite{Warner2003cvs..book.....W} for references). Numerical modeling of accretion in polars suggests that the flow can have a complex hierarchical structure \cite{Isakova2018ARep...62..492I}. In Fig.~\ref{fig:num_amher} we present results of 3D numerical modeling of the flow structure in a typical polar. Computations were performed in the framework of model described by \cite{Isakova2018ARep...62..492I}, with a computational domain containing a $384 \times 384 \times 192$ grid. The parameters of the binary system correspond to AM~Her. Computational domain was chosen in such a way that it included also a part of the Roche lobe of the donor star. Magnetic field at the surface of the white dwarf was set to $10^9$~G. Inclination angle of the magnetic axis of the accretor to its rotation axis was set to $30^\circ$, the angle between magnetic axis and the direction to the donor was $90^\circ$. The color scale shows the logarithm of density, white sphere corresponds to the surface of accretor, red line corresponds to the magnetic axis, and the blue line represents the axis of rotation of the white dwarf, green lines with arrows show magnetic force lines. The stream of matter outflowing from the envelope of the donor splits in the magnetic field into two separate streams, moves along magnetic field lines and hits the surface of the white dwarf near its magnetic poles, forming two hot spots. \begin{figure} \centering \includegraphics[width=0.7\textwidth]{amher.eps} \caption{Three-dimensional structure of the flow in a typical polar. The level surfaces of the logarithm of density are shown in color).Also shown are magnetic lines of force (with arrows), magnetic axis (red line) and axis of rotation (blue line). For clarity, the diagrams show configurations for three phases of the orbital period: $0.375$ (a), $0.5$ (b), and $0.625$ (c). The angle of inclination of the plane of rotation of the binary star is $i = 90^\circ$.} \label{fig:num_amher} \end{figure} Let estimate the cross-section of the accretion stream in the vicinity of L$_1$ applying the method described by \cite{Savonije1978A&A....62..317S,Bisikalo2013gasdyntdz}. The flow of gas in the vicinity of L$_1$ is similar to the expansion of the gas into the void from a cavity with a point hole. This means that velocity of the flow through the point L$_1$ is approximately equal to sound velocity $c_\mathrm{s}$. Deviation of the trajectory of the outflowing gas from the Roche lobe boundary in the vicinity of this point is determined by the balance of the potential and kinetic energy: \begin{equation} \pdiff{^2\Phi}{x^2}\bigg|_\mathrm{L_1} \frac{x^2}{2} + \pdiff{^2\Phi}{y^2}\bigg|_\mathrm{L_1} \frac{y^2}{2} = c_\mathrm{s}^2 \;, \end{equation} where $x$ and $y$ are the coordinates in the plane orthogonal to the stream. As can be seen, cross-section of the stream has the shape of an ellipse, while the derivatives of the potential determine dimensions of the semiaxes. Substituting the expression for the Roche potential into previous equation, one can evaluate the area of the ellipse: \begin{equation} \label{eq:stream_section} S_\mathrm{str} \approx \frac{\pi}{4} \left( \frac{c_\mathrm{s}}{\Omega_\mathrm{orb}} \right)^2 \;, \end{equation} where $\Omega_\mathrm{orb} = 2\pi/P_\mathrm{orb}$ is the angular frequency of the orbital motion of the binary system. The flow of the gas in the stream outside the magnetosphere of the white dwarf is determined, mainly, by gravity (more precisely, by the effective Roche potential), rather than by gas pressure. In addition, after the outflow from L$_1$, the gas rather quickly accelerates to the velocity of several tens of Machs (see Fig.~\ref{fig:accr_stream}). Therefore, expression (\ref{eq:stream_section}) is an estimate of the width of the accretion stream along its entire length. \begin{figure} \centering \includegraphics{AMHer_accr_stream.eps} \caption{Distribution of plasma parameters for the radial flow in the polar. \textit{Upper left panel:} number density of electrons. \textit{Upper right panel:} Alfv\'{e}n velocity, free-fall velocity and sound speed. \textit{Lower left panel:} magnetic field of the accretor (solid line), amplitude of the Alfv\'{e}n fluctuations for different values of $ T_\mathrm{w}$ (the rest of lines). \textit{Lower right panel:} the time of free fall from the point $R_\mathrm{L_1}$ (solid line), setting time of the turbulence for different values of $T_\mathrm{w}$ (remaining lines).} \label{fig:accr_stream} \end{figure} After leaving L$_1$, the plasma will move along a wide arc under the action of Coriolis force and Lorentz magnetic force. However, for simplicity, we shall assume that the plasma has a rectilinear trajectory. Let assume that the magnetic field of the white dwarf has a dipole structure, \begin{equation} B = B_\mathrm{a} \left( \frac{r}{R_\mathrm{a}} \right)^{-3} \;, \end{equation} where $B_\mathrm{a} = 10^7$~G is magnetic field at the surface of accretor; $R_\mathrm{a} = 0.013~R_\odot$ is accretor radius. The radius of the magnetosphere $R_\mathrm{m}$ may be found from condition of equality of Alfv\'{e}n and dynamical velocities under assumption that the cross-section of accretion stream does not change while it travels from L$_1$: \begin{gather} \label{eq:magnetosphere_condition} \frac{B(R_\mathrm{m})}{\sqrt{4\pi \rho(R_\mathrm{m})}} = v_\mathrm{ff}(R_\mathrm{m}) \;, \\ \label{eq:continuity} \dot{M} = S_\mathrm{str} \rho v_\mathrm{ff} \;, \\ \label{eq:free_fall_velocity} v_\mathrm{ff}^2 - \frac{2 G M_\mathrm{a}}{r} = c_\mathrm{s}^2 - \frac{2 G M_\mathrm{a}}{R_\mathrm{L_1}} \;. \end{gather} In the last equality, we took into account the fact that the gas outflows from L$_1$ with the sound velocity. For the accepted binary system parameters, $R_\mathrm{m} = 0.13~R_\odot = 0.22~R_\mathrm{L_1}$. Figure~\ref{fig:accr_stream} shows radial distributions of particle concentrations and gas velocity computed by Eqs.~(\ref{eq:continuity}) and (\ref{eq:free_fall_velocity}) under assumption that the gas temperature is $10^4$~K. As it can be seen, the main part of the accretion time a parcel of the gas spends in the outer part of its trajectory, though in the approximation we use, the trajectory of the motion of matter is a straight line. If we would take into account the fact that the trajectory has curvature, accretion time will increase even more. In the paper \cite{Kurbatov2017PhyU...60..798K} we considered the problems of simulation of plasma flows in strong magnetic fields. Under such conditions, Alfv\'{e}n and magnetosonic waves may significantly affect dynamics of the flow. It is known that Alfv\'{e}n waves are less susceptible to damping than magnetosonic waves. Finite amplitude waves polarized across the background (regular) magnetic field interact with each other and this manifests itself as a wave Alfv\'{e}n turbulence \cite{Iroshnikov1963AZh....40..742I,Galtier2000JPlPh..63..447G}. For its setting, the interaction of waves of two families is necessary: one family should have group velocity co-directed with the background magnetic field, while in the other one, it should be headed in the opposite direction. Turbulence of this type is characterized by the magnitude of the background magnetic field, power spectrum, longitudinal and diagonal spatial scales. In the paper \cite{Kurbatov2017PhyU...60..798K} we proposed a model of modified magnetic hydrodynamics in which Alfv{\'e}n turbulence is taken into account via stochastic sources of momentum and magnetic field acting on the average medium flow. The model was confined to the case of the so-called balanced turbulence, for which the energy spectra of Alfv\'{e}n fluctuations of both families coincide. A consequence of this is that the total energy flux of the Alfv\'{e}n waves along the background field is zero. The energy spectrum was calculated by \cite{Galtier2012arXiv1201.1370G}. For the given spectrum $P(\vec{k})$, let denote the energy of the turbulent fluctuations of the magnetic field, per unit mass, as in \cite{Kurbatov2017PhyU...60..798K} \begin{equation} \label{eq:magnetic_energy} \frac{W}{2} \equiv \frac{\langle |\vec{b}|^2 \rangle}{8\pi\rho} = \int d^3k\,P(\vec{k}) \approx 1.17 \sqrt{\epsilon a}\,\frac{L_\perp}{L_\parallel}. \end{equation} Here $W$ is the specific turbulence energy density (that is, the total energy of velocity pulsations and magnetic field per unit mass); $\epsilon$ is the energy flux through the turbulent cascade; $a$ is the Alfv{\'e}n velocity corresponding to the background field; $L_\perp$ and $L_\parallel$ are the transverse and longitudinal scales of turbulence. As a longitudinal scale of turbulence it is natural to choose the longitudinal scale of accretion, for instance, $R_\mathrm{L_1} \approx 4\times10^{10}$~cm. The transverse scale should be associated with the transverse dimension of the accretion stream $\sqrt{S_\mathrm{str}} \approx 0.025~R_\mathrm{L_1} \approx 10^9$~cm. The value $\epsilon$ is a parameter characterizing turbulence energy and is determined by the specific excitation mechanism of Alfv{\`e}n waves. In \cite{Galtier2012arXiv1201.1370G} the energy spectrum of turbulence was presented in the approximation of a ho\-mo\-ge\-neo\-us background field. In a realistic formulation of the problem, it is necessary to take into account that the accretion flow, in general, changes shape and cross-section, following the lines of the magnetic field. The transfer of energy of Alfv{\`e}n waves was correctly considered by \cite{Dewar1970PhFl...13.2710D,Dudorov2009VestCGU}. In the stationary flow, transport equation has the form \begin{equation} \label{eq:energy_conservation_general} \nabla [\rho W_\pm (\vec{v} \pm \vec{a})] = 0 \;, \end{equation} where $W_\pm$ is turbulent energy, corresponding to one family of waves; $\vec{v}$ is local velocity of the gas; $\vec{a}$ is the vector of Alfv{\`e}n velocity. For the balanced turbulence, $W_+ = W_- \equiv W$. Then, summing-up both equations (\ref{eq:energy_conservation_general}), one obtains \begin{equation} \nabla (\rho W \vec{v}) = 0 \,. \end{equation} As it can be seen, the value $W$ does not vary along the stream lines. In the present study we will assign to the turbulence some kind of ``effective'' temperature $T_\mathrm{w}$: \begin{equation} W = \frac{3 k_\mathrm{B} T_\mathrm{w}}{m_\mathrm{p}} \;. \end{equation} Without addressing a specific mechanism of the excitation of Alfv{\'e}n waves, two limiting cases can be distinguished: (a) turbulent energy is equal to the thermal energy of the medium with the temperature of the order of $10^4$~K; (b) the energy is determined by the temperature of the matter at the base of the accretion column, which is $\sim 10^8$~K \cite{Warner2003cvs..book.....W}. According to the definition (\ref{eq:magnetic_energy}), the first case corresponds to the amplitude of the Alfv{\'e}n waves $b\lesssim 10^3$~G, the second case --- to $b\lesssim 10^5$~G. In both cases, the contribution of the fluctuations dominates the contribution of the regular, dipole component of the magnetic field only in the outer part of the accretion flow, as it is seen in Fig.~\ref{fig:accr_stream}. On the other hand, the plasma, after it leaves L$_1$, spends half of the free-fall time at the distance of $0.8$~$R_\mathrm {L_1} $ and larger one, see Fig.~\ref{fig:accr_stream}. Thus, most of the time, the plasma is subject to the influence of a fluctuating magnetic field. It is useful to estimate the setting time of wave turbulence. In the paper \cite{Kurbatov2017PhyU...60..798K} this time was mentioned as the time of energy redistribution in the turbulent cascade: \begin{equation} \tau_\mathrm{w} = \frac{L_\perp^2 a}{L_\parallel W} \;. \end{equation} As it can be seen in Fig.~\ref{fig:accr_stream}, the concepts of Alfv{\'e}n wave turbulence are quite applicable to the description of the fluctuating component of the magnetic field. | In this paper we suggested a mechanism for generation of radio emission observed from polars. The basis of the mechanism is cyclotron radiation of electrons in a fluctuating magnetic field. The source of the fluctuations is Alfv\'{e}n's wave turbulence. It is assumed that thermal electrons of the temperature of order $10^4$~K participate in this mechanism. The power spectra of the emission with two states of polarization and different directions with respect to the orientation of the background magnetic field were calculated. If the amplitude of the fluctuations of magnetic field is comparable or larger than the background field, radiation spectrum will have the width of the order of the cyclotron frequency corresponding to the characteristic amplitude of the fluctuations. Degree of polarization of radiation in this case is in the range $0$ to $1/3$, depending on the direction to the observer. Suggested emission mechanism was applied to a simple model of the accretion stream in the polar AM~Her. It turned out that the spectrum corresponding to observations is formed in the outer part of the accretion stream, outside the magnetosphere of the white dwarf. The spectrum can be characterized by the effective turbulence temperature $T_\mathrm{w} \sim 10^5 \mathdash 10^6$~K. Emitting region has an optical thickness $\tau_\nu \sim 10^4$. Thus, considered emission mechanism can explain the observed flux of radio emission with a margin of four orders of magnitude. There is no fundamental reason that may prohibit the action of suggested mechanism in intermediate polars and even in the systems of nonmagnetic cataclysmic stars (more precisely, in the systems in which the field does not exceed $10^6$~G). In the present study, however, we confined ourselves to the case of polars, since they are, apparently, distinguished by a simpler flow morphology than the systems with weak magnetic field. In the \S~\ref{sec:observations} we considered briefly the mechanisms of radio emission generation suggested earlier: gyrosynchrotron and maser emission. Under not too restrictive conditions (the presence of a region of reconnection of the magnetic field lines as a source of super-thermal electrons), these mechanisms can provide the observed level of the radiation flux. It should be noted, however, that these emission models can act in the regions of sufficiently low electron concentration only. Otherwise, the radiation will decay as a result of the plasma absorption% \footnote {% Plasma transparency condition, in the sense of absorption at frequencies not exceeding Langmuir's one, is: $\nu \gg [2 e^2 N/(\pi m_\mathrm{e})]^{1/2}$ or $N/(10^{10}~\text{cm}^{-3}) \ll 0.5\,(\nu/\text{GHz})$.}. At this, in the case of gyrosynchrotron emission, the concentration can not be too low, otherwise, because of small optical thickness of the radiating layer, in order to enable observed flux, relativistic electrons will be needed. It is significant that the estimates of radiation fluxes for both mechanisms are optimistic. The mechanism of emission suggested in the present study provides a flux of radiation four orders of magnitude larger than the observed one. It is assumed that the radiation is generated by thermal electrons in the accretion column, where typical electron number density is about $10^{16}$~cm$^{-3}$. Isotropic plasma with such a density is, by no means, opaque. However, magnetic field creates in the plasma transparency windows, which allow radiation to be delivered to the region of low electron density. In the present study we did not consider the processes that determine opacity of the plasma in the radio range. Instead, we limited ourselves to the formal introduction of the opacity coefficient, the magnitude of which was selected from the condition for the correspondence of the observed and model fluxes. In the forthcoming paper, we plan to investigate the problem of radiation transfer in the magneto-active plasma with a fluctuating magnetic field in more detail. The authors acknowledge G.~Tovmassian for consultations in the course of this study. E.~P. Kurbatov and A.~G. Zhilkin were supported by the Program of the Presidium of the Russian Academy of Sciences \textnumero 28 ``Cosmos: studies of fundamental processes and their interrelations'' (Subprogram II ``Astrophysical objects as space laboratories''). | 18 | 8 | 1808.04584 |
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1808 | 1808.03701_arXiv.txt | In this paper, we generalize Coleman-Weinberg (CW) inflation in grand unified theories (GUTs) such as $\text{SU}(5)$ and $\text{SO}(10)$ by means of considering a two complex singlet fields with conformal invariance. In this framework, inflation is a result of spontaneously broken conformal symmetry in addition to the GUT symmetry. The latter implies a potential of CW form as a consequence of radiative corrections and the former flattens the CW potential, as a result we obtain $n_{s}\sim0.96-0.967$ and $r\sim0.003-0.005$ for $50-60$ number of e-foldings. We calculate the corresponding estimations of proton lifetime as $\tau_{p}\apprle10^{40}$ years whose decay mediated by the superheavy gauge bosons. We implement type I seesaw mechanism by weakly coupling the complex singlet which carries two units of lepton number to the three generations of singlet right handed neutrinos (RHNs). The spontaneous symmetry breaking of global lepton number amounts to the generation of neutrino masses. We consider non-thermal leptogenesis in which inflaton dominantly decays into heavy RHNs which sources the observed baryon asymmetry. We constrain the couplings of the inflaton field to the RHNs which gives the reheating temperature as $10^{6}\text{ GeV}\lesssim T_{R}<10^{9}$ GeV. | Primordial inflation is a successful paradigm for the description of the early Universe and it is strongly supported by the current observational data \cite{Ade:2015lrj}. Primordial perturbations, when the scales exiting the horizon $\left(k\sim aH\right)$, are eventually responsible for the structure formation in the Universe. From Planck 2015 \cite{Ade:2015lrj,Ade:2015tva}, the key observables of inflation, namely, the scalar tilt and the ratio of tensor to scalar power spectra, are constrained as $n_{s}=0.968\pm0.006$, $r<0.09$ at $95\%$ confidence level. The CMB power spectra is observed to be nearly adiabatic, scale invariant and Gaussian \cite{Ade:2015ava,Ade:2015lrj}. Although, the physical nature of the inflaton is still uncertain \cite{Martin:2013tda,Martin:2015dha}, the models based on $f(R)$ or canonical scalar field with flat potential are favorable with respect to the data. Since the inflationary scale is in general expected to be $\sim10^{16}\,\text{GeV}$ it is natural consider inflaton to be a scalar field associated with grand unified theory (GUT) groups such as $\text{SU}(5)$ and $\text{SO}(10)$. Shafi-Vilenkin (SV) model \cite{Shafi:1983bd} is one of the first realistic model of inflation which was based on $\text{SU}(5)$ GUT \cite{Georgi:1974sy}. In this framework, inflation is a result of spontaneous breaking of $\text{SU}(5)\to\textrm{SU}(3)_{c}\times\textrm{SU}(2)_{L}\times\textrm{U}(1)_{Y}$ by a GUT field ($\textbf{24-\text{plet}}$ adjoint Higgs) and inflaton which is a SU(5) singlet rolls down to a vacuum expectation value (VEV). The success of SV model is that it can lead to a successful baryogenesis after inflation and predicts proton life time above the current lower bound \cite{Shafi:2006cs,Rehman:2008qs}. In this model, the scalar field potential is given by Coleman-Weinberg (CW) form according to which primordial gravitational waves are constrained by $0.02\le r\le0.1$ \cite{Okada:2014lxa}. Although, SV model is well within the current bounds of Planck, several extensions of this model were studied to get smaller values of tensor to scalar ratio. In \cite{Cerioni:2009kn,Panotopoulos:2014hwa,Barenboim:2013wra}, CW inflation was studied in the context of induced gravity, non-minimal coupling and brane-world scenario where the tensor to scalar ratio was obtained to be $r\sim\mathcal{O}\left(10^{-2}\right)-\mathcal{O}\left(10^{-3}\right)$. After all, these modifications necessarily introduce an additional parameter into the theory that is responsible for flatness of potential. Moreover, extensions of SV model within particle physics offers a rich physics beyond the Standard Model (SM). Therefore, SV model embedded in a higher gauge group like $\text{SO}\left(10\right)$ which can be broken to SM via an intermediate group $\text{G}_{422}=\text{SU}(4)_{c}\times\text{SU}\left(2\right)_{L}\times\text{SU}\left(2\right)_{R}$ \cite{Lazarides:1984pq,Lazarides:1991wu}. Obtaining successful inflation in $\text{SO}\left(10\right)$ is more realistic with additional benefits to explain physics beyond SM such as neutrino physics, matter anti-matter asymmetry through non-thermal leptogenesis, monopoles and dark matter (DM) \cite{Rehman:2008qs}. For example, Ref. \cite{Boucenna:2014uma} considered a complex singlet scalar being coupled to right handed neutrinos (RHNs) followed by implementing type I seesaw mechanism. This approach unified inflation with Majorana DM together with the scheme of generating neutrino masses. In \cite{Okada:2013vxa} an additional $\text{U}(1)_{B-L}$ symmetry was considered in the SM i.e., $\textrm{SU}(3)_{c}\times\textrm{SU}(2)_{L}\times\textrm{U}(1)_{Y}\times\text{U}\left(1\right)_{B-L}$, where $B-L$ symmetry can be spontaneously broken when the scalar field takes the VEV. In this setup, one can explain baryon asymmetry of the Universe through non-thermal leptogenesis \cite{Lazarides:1991wu,Asaka:2002zu,Senoguz:2003hc,Senoguz:2005bc}. Recently, CW inflation was studied in an extension with $\text{SO}(10)$ and $\text{E}_{6}$, pointing out the possibilities of observing primordial monopoles \cite{Senoguz:2015lba}. Apart from models based on GUT theories, the Starobinsky model based on the $R^{2}$ gravity modification and the Higgs inflation \cite{Starobinsky:1980te,Starobinsky:1983zz,Bezrukov:2007ep} occupy a privileged position, with practically equal predictions in the $\left(n_{s},\,r\right)$ plane \begin{equation} n_{s}=1-\frac{2}{N}\quad,\quad r=\frac{12}{N^{2}}\,,\label{targetspot} \end{equation} where $N$ is the number of e-foldings before the end of inflation. There has been a growing interest on embedding these models in string theory and supergravity (SUGRA) aiming for a UV completion \cite{Linde:2014nna,Ferrara:2013rsa}. Recently UV completion of Starobinsky model was proposed in the context of non-local gravity inspired from string field theory \cite{Koshelev:2016xqb}. Starobinsky like models were also developed in $\mathcal{N}=1$ SUGRA, namely, no scale \cite{Ellis:2013nxa} and $\alpha-$attractor models \cite{Kallosh:2013yoa} where an additional physical parameter responsible for any prediction of $r<0.1$. In \cite{Kumar:2015mfa} $\alpha-$attractor models were studied in the non-slow-roll context where a new class of potentials were shown to give the same predictions. On the other side, Higgs inflation is particularly interesting due to the fact that Higgs was the only scalar so far found at LHC. But to be a candidate of inflation one requires a large non-minimal coupling $\left(\xi\right)$ to Ricci scalar. It was known that a scalar field with large non-minimal coupling gives rise to $R^{2}$ term considering 1-loop Quantum corrections. Consequently, renormalization group (RG) analysis shows that Higgs inflation is less preferable compared to Starobinsky model \cite{Salvio:2015kka,Calmet:2016fsr}. This result not only applies to Higgs inflation but also to any arbitrary scalar with large non-minimal coupling. Furthermore, in both $R^{2}$ and Higgs inflation the inflaton field rolls down to zero after inflation\footnote{Although, SM Higgs field rolls to its electroweak VEV it is negligible compared to the energy scale of inflation.}. On the contrary, in GUT theories, inflaton field acquires a VEV due to its interactions with GUT fields. The main goal of this paper is to generalize SV model in order to achieve $r\sim\mathcal{O}\left(10^{-3}\right)$ without introducing any additional parameters being responsible for inflaton potential flatness\footnote{Our construction is different from the models with non-minimally coupled scalars where flat potential comes from requiring $\xi\gg1$ \cite{Broy:2016rfg}. }, instead, we consider an additional conformal invariance (or local scale invariance) in our GUT model. It was long ago shown by Wetterich \cite{Wetterich:1987fm} that scale symmetries plays crucial role in the construction of realistic cosmological models based on particle physics. Moreover, scale symmetries successfully explain the hierarchy of different scales such as Planck and Higgs mass \cite{Hooft:2014daa,Quiros:2014hua,Scholz:2012ev,Bars:2013yba}. Therefore, it is natural to consider scale invariance in constructing inflationary scenario, through which we can obtain dynamical generation of Planck mass, inflationary scale and particle physics scales beyond SM. In this regard, we introduce two complex singlet fields $\left(\bar{X},\,\Phi\right)$ of $\text{SU}(5)$ or $\text{SO}(10)$ and couple them to Ricci scalar and adjoint Higgs field $\left(\Sigma\right)$ such that the total action would be conformally invariant. We promote inflation as a result of spontaneous breaking of conformal and GUT symmetries. The former occurs due to gauge fixing of one singlet field to a constant for all spacetime and the latter occurs due to $\Sigma$ field takes its GUT VEV. Here the inflaton is identified with the real part of the second singlet ($\phi=\sqrt{2}\mathfrak{Re}\left[\Phi\right]$). While, the imaginary part is the corresponding Nambu-Goldstone boson, is assumed to pick up a mass due to the presence of small explicit soft lepton number violation terms in the scalar potential \cite{Boucenna:2014uma}. Here, we assume $\Phi$ carries two units of lepton number and it is coupled to the RHNs in such a way that the coupling is highly suppressed during inflation\footnote{This will be explained in detail in the due course of this paper.}. Near the end of inflation, inflaton is supposed to reach its VEV and also the global lepton number is violated. Thereafter, we study the dominant decay of inflaton into heavy RHNs producing non-thermal leptogenesis. We compute the corresponding reheating temperatures and also discuss the issue of producing observed baryon asymmetry. In summary, our study completes with an observationally viable inflationary scenario predicting proton life time, neutrino masses and producing non-thermal leptogenesis from heavy RHNs. The paper is briefly organized as follows. In Sec. \ref{SCmodel-sec}, we describe toy models with conformal and scale invariance. We identify the interesting aspects of spontaneous symmetry breaking leading to viable inflationary scenario. In Sec. \ref{SU5CW-sec}, we briefly present the SV model and the computation of proton life time. In Sec. \ref{twofieldmodel-sec} we propose our generalization of SV model by introducing an additional conformal symmetry. We report the inflationary predictions of the model together with estimates of proton life time. In Sec. \ref{seesawSec} we later explore the nature of inflaton couplings to the SM Higgs, singlet RHNs through type I seesaw mechanism. We constrain the Yukawa couplings of the inflaton field compatible with the generation of light neutrino masses. In Sec. \ref{ReheatSec} we implement non-thermal leptogenesis and compute the reheating temperatures corresponding to the dominant decay of inflaton to heavy RHNs. We additionally comment on the necessary requirements for the production of observed baryon asymmetry through CP violation decays of RHNs. In Sec. \ref{conclusions} we summarize our results with future outlook. In this paper we follow the units $\hbar=1,\,c=1,\,$ $m_{p}^{2}=\frac{1}{8\pi G}$. | \label{conclusions} Coleman-Weinberg inflation \cite{Shafi:1983bd} has been a successful and realistic model based on GUT and is consistent with current Planck data with $r\gtrsim0.02$ \cite{Okada:2014lxa}. In this work, we have further generalized the framework of CW inflation with an additional conformal symmetry. Spontaneous breaking of conformal symmetry essentially useful to create hierarchy of mass scales, therefore it is natural to realize this symmetry in GUT models. In this respect, two complex singlet fields of $\text{SU}(5)$ or $SO(10)$ are considered and are coupled to the GUT fields in a suitable manner. We have showed that this setup upon spontaneous breaking of GUT and conformal symmetry leads to an interesting inflationary scenario driven by the real part of the singlet field. In our model, the above VEV branch of CW potential gets flattened to Starobinsky plateau giving the predictions for $n_{s}\sim0.96-0.967$ and $r\sim0.003-0.005$ for $50-60$ number of e-foldings. We found that these predictions are independent of the VEV of the inflaton field. However, values of inflaton VEV affects the masses of the superheavy gauge bosons that mediate the proton decay. We calculated the corresponding estimates for proton life time above the current lower bound from Super-K data $\tau_{p}\left(p\to\pi^{0}+e^{+}\right)>1.6\times10^{34}$. In the next step, we introduced a coupling between complex singlet field with the generation of three singlet RHNs where the singlet field is assumed to carry two units of lepton number. We implemented type I seesaw mechanism where spontaneous symmetry breaking of global lepton number results in generating neutrino masses. We put an upper bound the inflaton couplings to RHNs assuming the inflation is dominated by inflaton couplings to GUT field. For the non-thermal leptogenesis to happen, we have considered dominant decay of inflaton into some of the RHNs and obtained the corresponding reheating temperatures as $10^{6}\text{ GeV}\lesssim T_{R}<10^{9}$ GeV. In summary, our new development of CW inflation can be verified in the future CMB data \cite{Creminelli:2015oda}. In this work, we mainly restricted the non-SUSY construction of CW with conformal symmetry. It would be interesting to consider this model in SUGRA with superconformal symmetries which we defer for future investigations. | 18 | 8 | 1808.03701 |
1808 | 1808.10438_arXiv.txt | We present the full disk-fit results VANDAM survey of all Class 0 and I protostars in the Perseus molecular cloud. We have 18 new protostellar disk candidates around Class 0 and I sources, which are well described by a simple, parametrized disk model fit to the 8 mm VLA dust-continuum observations. 33\% of Class 0 protostars and just 11\% of Class I protostars have candidate disks, while 78\% of Class 0 and I protostars do not have signs of disks within our 12 AU disk diameter resolution limit, indicating that at 8 mm most disks in the Class 0 and I phases are $<$10 AU in radius. These small radii may be a result of surface brightness sensitivity limits. Modeled 8 mm radii are similar to the radii of known Class 0 disks with detected Keplerian rotation. Since our 8 mm data trace a population of larger dust grains which radially drift towards the protostar and are lower limits on true disk sizes, large disks at early times do not seem to be particularly rare. We find statistical evidence that Class 0 and I disks are likely drawn from the same distribution, meaning disk properties may be defined early in the Class 0 phase and do not undergo large changes through the Class I phase. By combining our candidate disk properties with previous polarization observations, we find a qualitative indication that misalignment between inferred envelope-scale magnetic fields and outflows may indicate disks on smaller scales in Class 0 sources. | Disks of gas and dust around young protostars are fundamental to protostellar mass accretion and act as the mass reservoir from which stars and planetesimals form \citep{Armitage2011,Williams2011}. Circumstellar disks are expected to form around even the youngest Class 0 protostars which are embedded in their dense natal dust and gas envelope. Class I protostars are less embedded, having cleared a portion of their envelopes \citep{Mckee2007}. Until recently, disks around Class 0 and Class I protostars have remained elusive because $\sim$millimeter wavelengths are required to penetrate through the dense envelope \citep{Looney2000}, and sub-arcsecond resolution is required to spatially resolve the disk. Keplerian rotation is a tell-tale sign of true, rotationally supported disks that exist for long enough timescales to form long-lived disk structures and eventually planets; flattened structures without rotation quickly collapse inward \citep[e.g.,][]{Terebey1984}. So few young Keplerian disks are known and as a consequence, questions concerning disk frequency, disk radii, dust populations, disk evolution, and the presence of planetesimals in the youngest protostellar disks are only beginning to be addressed. Keplerian rotation has been detected in disks around only 4 total low-mass Class 0 protostars to date with $R$ $>$30 AU \citep{Ohashi2014,Tobin2012,Murillo2013,Codella2014,Yen2017,Lee2017}; however they are bright sources and may not represent typical disks at this stage of evolution. Class I protostars have longer lifetimes and have cleared enough of their mass reservoir that more low-mass Class I disks have been detected \citep[$\sim$10 total, to date;][]{Harsono2014} than in Class 0 systems though not nearly as many as the $\sim$100 total in more-evolved Class II sources \citep[e.g.,][]{Andrews2009,Andrews2010}. By the Class II stage the envelope has mostly dispersed, clearly revealing the circumstellar disk and allowing geometrical constraints to be found from the spectral energy distribution (SED) of the dust emission. The dense envelopes in Class 0 and I systems prevent disk parameters to be constrained from examining the SED alone. Recent observations of a disk around a Class II protostar have revealed the earliest known evidence of planet formation \citep{ALMA2015}. The Class 0 and I protostellar stages have the largest mass reservoirs available to form disks and planetesimals; therefore understanding the properties of disks at the earliest possible epochs is crucial to determine the formation mechanism behind circumstellar disks and the initial pathway to planet formation. The morphology and strength of the magnetic fields in protostellar systems also play an important role in star and disk formation \citep[e.g.,][]{Crutcher2012}. Magnetic field effects on the small size scale of circumstellar disks ($\sim$0.5$^{\prime\prime}$ or $\sim$100 AU) of young stellar objects have started to be theoretically and observationally quantified in individual systems \citep{Mellon2008,Hennebelle2008,Stephens2014,SeguraCox2015,Cox2015}. Magnetic field morphology can be inferred from dust emission; spinning dust grains align their long axes perpendicular to the magnetic field, polarizing the dust emission \citep[e.g.,][]{Lazarian2007}. When a strong magnetic field and the rotation axis of the circumstellar disk are aligned, magnetic braking can have a significant effect on disk formation by transporting away angular momentum, limiting forming disks to $R$ $<$10 AU \citep[e.g.,][]{Mellon2008,Dapp2010,Machida2011,Li2011,Dapp2012}. In this scenario, disks only reach $R$ $\sim$100 AU at the end of the Class I phase when the envelope is less massive and magnetic braking becomes inefficient \citep[e.g.,][]{Dapp2012,Mellon2009,Machida2011}. Conversely, recent works have highlighted the critical importance of the magnetic field direction relative to the rotation axis: when the field and rotation axis are misaligned, magnetic braking becomes less efficient, and $\sim$100 AU disks can form \citep{Joos2012,Li2013,Krumholz2013,SeguraCox2015,SeguraCox2016}. Disks can form if the coupling of the magnetic field to the disk material can be lessened by non-ideal magnetohydrodynamic effects, allowing material to accrete from the envelope to the disk without dragging in a flux-frozen magnetic field. Because so few young embedded disks are currently known via observations, expanding the number of known Class 0 and I disks is critical to determining the role of magnetic braking in disk growth at early times. To determine the properties of the youngest disks, we used the Karl G. Jansky Very Large Array (VLA) to make continuum observations for the VLA Nascent Disk and Multiplicity (VANDAM) survey toward all known protostars in the Perseus molecular cloud \citep{Tobin2015a}. The continuum observations trace dust emission (at $\lambda\sim8$ and $10$ mm) and free-free emission from jets near the central protostar (at $\lambda\sim4$ and $6.4$ cm). The VANDAM observations form an unbiased survey of young protostellar disks down to $\sim$10 AU size scales, giving us the opportunity to potentially double the number of known disks in Class 0 and I protostars from $\sim$15 total to over 30. We use the term ``candidate disks" because we do not have kinematic data on small scales to determine whether these structures are rotationally supported. The VANDAM sample contains all currently known Class 0 and I protostellar systems in Perseus, with 37 Class 0 systems, 8 Class 0/I systems, and 37 Class I systems. The 21 resolved sources in the VANDAM survey we examine in this paper (Table \ref{diskfulltab1}, Figure \ref{Per44_2x4}, and Figures in Appendix \ref{vandam:extended_images}) are the most complete sample of embedded sources in Perseus to-date \citep[see][for discussion of target selection]{Tobin2016a}. Per-emb-XX designations originate from \citet{Enoch2009}. We define resolved or extended sources as having spatial extents at least 1.1$\times$ the size of the FWHM of the beam, meaning we include marginally resolved sources in this study. We fit disk models to all protostars with relatively axisymmetric resolved emission (17 of 21 sources, see Table \ref{diskfulltab2} and Section \ref{vandam:modeling}) roughly perpendicular to known outflows; however, only sources with either axisymmetric resolved emission and a modeled disk-like profile or non-symmetric emission and other indirect evidence of a disk are considered candidate disks (see Appendix \ref{vandam:gallery}). In this paper, we present the full results toward the protostellar disk candidates around the Class 0 and Class I protostars from the VANDAM survey. This paper expands on the work done in \citet{SeguraCox2016}, which reported a subset of the candidate disks studied here. The observations, VLA set up, and data reduction are described in Section \ref{obssec}, with estimated masses from observed fluxes of extended sources presented in Section \ref{vandam:masses}. Section \ref{vandam:modeling} describes our {\it u,v}-plane disk modeling procedure, and we describe the modeling results in Section \ref{vandam:results}. The results of our study are discussed in Section \ref{vandam:discussion}, and the summary is given in Section \ref{vandam:summary}. Appendix \ref{vandam:extended_images} shows 8 mm images of protostars with extended emission, Appendix \ref{vandam:gallery} lists previously known information on each candidate disk studied here, and Appendix \ref{vandam:bestfits} presents images of candidate-disk modeling results. | \label{vandam:discussion} \subsection{Candidate Disk Properties} \label{vandam:discussion:diskprop} The 18 VANDAM candidate disks have estimated masses of 0.01--3.2 \Ms~ ($M_{d}$; Table \ref{diskfulltab2}). As discussed in Section \ref{vandam:masses}, the estimated mass of L1448 IRS3B is likely under-estimated, and NGC 1333 IRAS4A is an unusual outlier with $M_{d}$ an order of magnitude larger than all other sources. The remaining 17 candidate disks have $M_{d}$ values of 0.03--0.71 \Ms. Our values for the estimated masses of the candidates are all larger than the Minimum Mass Solar Nebula, the 0.01 \Ms~ amount of material expected to be required to form the planets in our own Solar System \citep{Weidenschilling1977}, indicating that these disks have the potential to eventually form planets. The modeled fit fluxes and hence masses calculated from the fit fluxes ($F_{fit}$ and $M_{fit}$ respectively; see Table \ref{diskfulltab3}) for all modeled candidate disks except NGC 1333 IRAS4A are within 0.7 to 1.3 times the measured fluxes and estimated masses from observations ($F_{8mm}$ and $M_{d}$ respectively; see Table \ref{diskfulltab2}), with the largest deviations occurring in sources with smaller modeled radii. The value of $M_{fit}$ of NGC 1333 IRAS4A is a factor of $\sim$0.56 lower compared to the observed $M_{d}$ value, likely because of the remaining envelope emission seen in this source which was not accounted for in the modeling procedure. When scaled to our opacity, the Class 0 protostar L1527's disk mass is 0.013 \Ms~\citep{Tobin2013b} with T$_{d}$ = 30 K. \citet{Harsono2014} revealed four Class I disks to have masses of 0.004-0.033 \Ms, using T$_{d}$ = 30 K and \citet{Ossenkopf1994} opacities. Compared to these embedded disks, our disk masses appear to be significantly higher. One possibility is that our assumption of dust opacity spectral index $\beta$ = 1 is too large. L1527 was found to have a shallower $\beta$ $\sim$ 0 \citep{Tobin2013b} from $\sim$mm wavelengths, which could be attributed to a population of large ($\sim$cm) dust grains centered in the disk midplane with a smaller scale height than observed at shorter wavelengths. Our 8 mm data indeed do trace large grains settled in the midplane, indicating that values of $\beta$ near 0 may be more common than previously thought for deeply embedded young disks, as suggested by \citet{Kwon2015}. If we assume $\beta$ = 0 instead of $\beta$ = 1, our estimated disk masses would change from 0.02--0.71 \Ms to 0.003--0.10 \Ms. Because we model the disk flux, not mass, any uncertainties in $\beta$ do not impact our modeling results. Our best-fit models for 14 candidate disks give -0.58 $<$ $\gamma$ $<$ 1.65 (Table \ref{diskfulltab3}), with the average value of inner-disk surface density power law $\gamma=0.32$ for our Class 0 and I sources. Negative values of $\gamma$ imply increasing disk surface density with radius, incongruent with typical disk profiles. For all our candidate disks at least one value of $q$, the disk temperature structure parameter, produces a positive best-fit value of $\gamma$. The average value of $\gamma=0.32$ is a shallower profile than more evolved disks. Disks around Class II protostars in Ophiuchus yield an a typical value of $\gamma=0.9$ \citep{Andrews2009}. The steeper values of $\gamma$ in Class II sources indicates that evolved disks are generally more centrally concentrated than our Class 0 and I disks. The few Class 0 disks with Keplerian rotation have relatively large radii, though it is unclear if these are typical radii of young disks or if this is simply detection bias towards large and bright sources. At 1.3 mm, VLA 1623 has $R$ $\sim$189 AU \citep{Murillo2013}, Lupus 3 MMS has $R$ $\sim$100 AU \citep{Yen2017}, and L1527 has $R$ $\sim$54 AU \citep{Ohashi2014}. HH212 has $R$ $\sim$60 AU in 850 $\mu$m ALMA data \citep{Codella2014,Lee2017}. Our Class 0 and I VANDAM candidate disks have 9.1 AU $<$ $R_{c}$ $<$ 42.2 AU. The modeled radii (Table \ref{diskfulltab3}) give disk diameters that are a factor of 1 to 1.5 times larger than the deconvolved sizes from the image-plane 2D Gaussian fits (Table \ref{diskfulltab2}). For most sources, the candidate Class 0 and I disks are larger than the expected upper limit of 10 AU from strong magnetic braking models \citep{Dapp2010}. HH211-mms, NGC 1333 IRAS1 A, and NGC 1333 IRAS2A are the three smallest disks with $R_{c}$ $\sim$10 AU. The remaining candidate disks have $R_{c}$ consistent with the radii of Keplerian Class 0 disks L1527 and HH212. \subsection{8 mm Emission as a Lower Limit on Dust Disk Radius} \label{vandam:discussion:lowerlim} As noted in Appendix \ref{vandam:gallery:singles:per14}, Per-emb-14 was resolved with CARMA in continuum dust emission at 1.3 mm \citep{Tobin2015b} with a dust disk a factor of $\sim$3 larger than our modeled 8 mm continuum radius \citep{SeguraCox2016}. ALMA 1.3 mm data (Tobin et al.~2018, submitted) also show evidence for more-extended emission at 1.3 mm compared to 8 mm, with image-plane Gaussian fit major axes of the 1.3 mm data 1.7$\times$ to 4.3$\times$ larger than the 8 mm modeled radii presented here for disk candidates NGC 1333 IRAS1 A, SVS13B, NGC 1333 IRAS4A, and NGC 1333 IRAS2A. A dependance on disk size with wavelength was also found for the more evolved classical T Tauri stars AS 209, CY Tau, and DoAr25 \citep{Perez2012,Perez2015}, with disk size decreasing with longer wavelength observations. Since the wavelength of thermal emission from dust grains roughly traces the sizes of the dust grains, 8 mm emission traces a population of larger sized grains than in 1.3 mm emission. The larger 8 mm grains experience radial drift to a larger extent \citep{Perez2012}, forming a more compact disk closer in to the central protostar than smaller grains which remain further from the protostar for longer periods of time \citep[e.g.,][]{Birnstiel2010}. At shorter wavelengths, near 1 mm, dust emissivity is higher causing the dust emission to be stronger in the outer parts of the disk and more likely to be detected. We consider our VANDAM modeled disk radii at 8 mm to be extreme lower limits on disk size. Shorter wavelength observations may be better tracers of the full extent of circumstellar dust disks due to these large-grain radial-drift effects and surface brightness sensitivity limits at 8 mm. \subsection{Outflow Orientations and Other Indirect Evidence of Disks} \label{vandam:discussion:outflows} Nearly all VANDAM candidate disks, except for Per-emb-63, have clearly associated outflows roughly perpendicular (60-90$^{\circ}$) to the major axis of the candidate disks (see Appendix \ref{vandam:gallery} for details). No outflows are associated with Per-emb-63. The orientations of outflows have been used as proxies for the disk rotation axis \citep[e.g.,][]{Hull2013}, hence outflows nearly perpendicular to extended continuum emission is a strong indicator of a protostellar disk. Along with our continuum emission disk modeling, we use the perpendicular outflows as indirect evidence of rotationally supported disks in embedded sources. HH211-mms, Per-emb-14, Per-emb-15, Per-emb-25, Per-emb-8, SVS13B have bipolar outflows perpendicular to their candidate disk elongation. Per-emb-30 and Per-emb-62 both have single monopolar outflows, possibly because of dense gas interacting with unseen components of their bipolar outflows. Their detected monopolar outflows are both perpendicular to their candidate disk directions. IC348 mms does not appear to be the central driving source of the bipolar outflow in its binary system though the candidate disk remains perpendicular to the outflow. NGC 1333 IRAS4A has a close binary companion, each with bipolar outflows perpendicular to the estimated disk position angle of our candidate disk. The binary system NGC 1333 IRAS2A drives two bipolar outflows, one coming from each close-separation component. The NGC 1333 IRAS2A candidate disk is perpendicular to the outflow associated with its protostar. NGC 1333 IRAS1 A drives an outflow almost perpendicular to its candidate disk and has a binary companion, which may be causing the outflow to have an S-shape via gravitational interactions between the binary protostars. SVS13C drives a bipolar outflow perpendicular to the candidate disk. We detect evidence of the outflow in our 8 mm data. As seen in Figure \ref{SVS13Cdmr}, right panel, the east-west emission component of SVS13C is well modeled with minimal residuals, while the north-south outflow seen in free-free \citep{Tychoniec2018}, remains. Because we already accounted for a point-source component in our model, we did subtract out any small-scale free-free emission coming from the jet-launching regions of the disk \citep{Anglada1998}, leaving minimal residual at target center. The non-axisymmetric candidate disks (Appendix \ref{vandam:gallery:weird}) cannot be fit with our modeling procedure, and for most the orientation of the candidate disk is unclear from 8 mm continuum data alone. IRAS 03292+3039 has a bipolar outflow perpendicular to a velocity gradient across the protostar on 1000 AU scales \citep{Yen2015}. IRAS 03282+3035 is a very close separation binary, with a velocity gradient along the outflow as well as a gradient perpendicular to the outflow on the southeast side of the envelope \citep{Tobin2011}. The 8 mm data of Per-emb-18 is highly elongated along the direction perpendicular to the outflow \citep{Davis2008}. Finally, the triple system L1448 IRS3B has a velocity gradient perpendicular to the outflows that are centered on two of the three triple components \citep{Tobin2016b}. \subsection{The Frequency of Class 0 and I Candidate Disks} \label{vandam:discussion:freq} With the VANDAM survey, we have detected 18 new candidate disks (14 Class 0 and 4 Class I) in the deeply embedded, young protostellar phases. Our survey has more than doubled the number of known possible disks around Class 0 and I protostars, bringing the total count from $\sim$15 to $\sim$33. With so many young disks and candidate disks now known, we can characterize typical young embedded disk frequency and dust properties, determine the relative rarity of large embedded disks, look for evolutionary trends between the protostellar phases, and begin to study the role magnetic fields play during the early stages of disk growth. Of the Class 0 protostars (including Class 0/I sources) in Perseus, 14/43 (33\%) have candidate disks on scales of 12 AU or larger. Only 4/37 (11\%) of Class I protostars in our sample have large, resolved candidate disks. Here we have included both the modeled and unmodeled complicated candidate disks in our counts. We also find that 62/80 (78\%) of Class 0 and I protostars do not have signs of disks within our 8 AU radius modeling limit. Since disk formation in protostars is expected to be a natural consequence of conservation of angular momentum during core collapse, this implies that at 8 mm most disks in the Class 0 and I phases are small ($<$10 AU). The population of unresolved disks may undergo stronger magnetic braking, be subject to other processes limiting disk growth, or the observations may be limited by surface brightness sensitivity. Small disk size at 8 mm does not necessarily imply that the entire disk is small, because disks may be more extended at shorter wavelengths. The lower proportion of Class I candidate disks compared to Class 0 candidate disks is surprising because naively, the disk is expected to grow from the Class 0 to the Class I stage as the envelope is dissipated in part by accreting onto the disk, though no correlations between disk masses or radii have yet been found between the Class 0 and I phases \citep{Williams2011}. The Class I candidate disks in Perseus may suffer from small-number statistics since only 4 Class I candidate disks were detected at 8 mm and may not reflect typical Class I disk frequency in other molecular clouds. This result only applies at 8 mm and is not a universal result since disk sizes vary with observational wavelength. An alternative explanation for the low proportion of VANDAM Class I candidate disks lies in the size of the dust grains seen in our observations. Our data trace large dust grains ($\sim$8 mm), since thermal emission from dust roughly traces the size of the emitting grains. It is possible that the candidate disks around more evolved Class I protostars have been stable long enough for radial drift \citep{Perez2012} to cause the 8 mm dust grain population to be more centrally concentrated around the protostar relative to Class 0 sources (Figure \ref{tbolvsgamma} may also support this scenario; see Section \ref{vandam:discussion:trends}). Observations made at shorter wavelengths which trace smaller dust grains reflect more extended disk sizes since smaller grains have less pronounced radial drift effects. The detectable 8 mm radius of an embedded disk may shrink as the protostar evolves due to 8 mm grain population radially drifts inwards causing the signal-to-noise ratio in the outer disk to decrease below instrument sensitivity thresholds. If disks detected at 8 mm do become more centrally concentrated as they evolve with a smaller detectable radius, Class I candidate disks may have had larger 8 mm disks in the past which may have been detected with 12 AU resolution, but at present the 8 mm disks may have shrunk to below resolution or sensitivity limits. \begin{figure}[!ht] \centering \begin{minipage}{.47\textwidth} \centering \includegraphics[width=0.95\textwidth]{f4.pdf} \caption{A histogram of the radii of all 14 modeled candidate disks, broken down by protostellar Class. } \label{diskhisto} \end{minipage}% \hspace{0.2cm} \begin{minipage}{0.47\textwidth} \centering \includegraphics[width=0.95\textwidth]{f5.pdf} \caption{A histogram of the radii of all 14 modeled candidate disks, broken down by multiplicity. } \label{diskhistomults} \end{minipage} \end{figure} Figure \ref{diskhisto} shows a histogram of all 14 modeled candidate disk radii, including 10 Class 0 and 4 Class I sources. The four non-axisymmetric sources could not be modeled to derive a disk radius and are not included in this histogram. Both the Class 0 and Class I candidate disk distributions peak at 15-30 AU radii, which are well resolved with the 12 AU resolution observations. The two largest disks, with 8 mm radii 30-45 AU both belong to Class 0 sources. With so few Class I protostars sampled, any further differences between the Class 0 and I candidate disk radii are difficult to distinguish. Figure \ref{diskhistomults} shows a histogram of the 14 modeled candidate disk radii, separated by whether the sources are in a single or multiple protostellar system. Seven of the modeled candidate disks are single systems, and 7 belong to multiple systems. It is unclear if there are variations in the distribution of radius that depend on the multiplicity of the systems. \subsection{Trends of Candidate Disk Characteristics} \label{vandam:discussion:trends} As seen in Figures \ref{tbolvsgamma}-\ref{tbolvsradius}, no tight correlations between protostellar age, 8 mm measured flux, modeled candidate disk radius, and modeled inner-disk surface density power law $\gamma$ are found by eye. We use bolometric temperature, T$_{bol}$, as an indicator of protostellar evolution to probe protostellar age \citep[]{Chen1995}. In these plots, Class 0/I sources are counted as Class 0 and we include all 14 modeled VANDAM candidate disks. Bolometric temperatures were taken from \citet{Tobin2016a}. When possible, we include unresolved sources which are not candidate disks if measurements or upper limits on the plotted parameters exist. We see no significant differences between single and multiple candidate characteristics with age, nor any significant differences between Class 0 and Class I candidate disk characteristics with age. \begin{figure}[h] \centering \includegraphics[width=0.75\textwidth]{f6.pdf} \caption{A mild correlation between T$_{bol}$ and modeled inner-disk surface density power law $\gamma$ is seen by eye, with a weak trend of more evolved sources with higher T$_{bol}$ having higher values of $\gamma$. Only modeled candidate disks are included in this figure.} \label{tbolvsgamma} \end{figure} \begin{figure}[!ht] \centering \begin{minipage}{.47\textwidth} \centering \includegraphics[width=0.95\textwidth]{f7.pdf} \caption{No clear correlations are seen between T$_{bol}$ and 8 mm flux. Filled, large symbols are candidate disks. Open, small symbols are unresolved sources.} \label{tbolvsflux} \end{minipage}% \hspace{0.2cm} \begin{minipage}{0.47\textwidth} \centering \includegraphics[width=0.95\textwidth]{f8.pdf} \caption{No clear correlations are seen between T$_{bol}$ and modeled $R_{c}$. Filled , large symbols are candidate disks. Open, small symbols are unresolved sources, and the upper limit radii reflect the 8 AU modeling limit.} \label{tbolvsradius} \end{minipage} \end{figure} A mild correlation between $\gamma$, the power-law of the inner-disk surface density, and T$_{bol}$ (Figure \ref{tbolvsgamma}) may be present with a tentative trend of higher values of $\gamma$ being found in older sources with higher T$_{bol}$. Higher values of $\gamma$ indicates that flux drops off faster with increasing radius, meaning 8 mm flux is more centrally distributed for younger sources than more evolved sources with higher T$_{bol}$. This may be explained as dust grains in more evolved sources having more time to experience radial drift \citep{Perez2012} and concentrate closer to the central protostellar source. As further evidence, in Section \ref{vandam:discussion:freq} we demonstrated that the typical value of $\gamma$ grows larger from the young Class 0/I stage to the more evolved Class II/III phases. If there are truly no correlations between 8 mm disk flux or radius and age (Figures \ref{tbolvsflux} and \ref{tbolvsradius}), a process independent of evolution could be setting the disk radii. Possible processes independent of age that could influence the disk radii include the initial angular momentum imparted onto the disk from the natal protostellar core \citep[e.g.,][]{Terebey1984} or magnetic braking dominating over any evolutionary effects \citep[e.g.,][]{Dapp2010}. Alternatively using T$_{bol}$ as an indicator of protostellar evolution may be imprecise. Certainly it is clear that overall there are no major differences in 8 mm fluxes or radii between the Class 0 and Class I phases. Additionally, in Figures \ref{fluxvsradius}-\ref{radiusvsgamma}, no correlations are seen between measured flux and modeled radius, or modeled radius and modeled $\gamma$. Again no clear differences are seen in the characteristics between the single and multiple populations or the Class 0 or Class I phases. For Figure \ref{fluxvsradius}, we include the unresolved sources which we do not consider to be candidate disks, with the radii represented as an upper limit set by the 8 AU modeling limit (see Section \ref{vandam:modeling}). \begin{figure}[!h] \centering \begin{minipage}{.47\textwidth} \centering \includegraphics[width=0.95\linewidth]{f9.pdf} \caption{No clear correlations are seen between 8 mm flux and modeled $R_{c}$. Filled, large symbols are candidate disks. Open, small symbols are unresolved sources, and the upper limit radii reflect the 8 AU modeling limit.} \label{fluxvsradius} \end{minipage}% \hspace{0.2cm} \begin{minipage}{0.47\textwidth} \centering \includegraphics[width=0.95\linewidth]{f10.pdf} \caption{No clear correlations are seen between modeled $R_{c}$ and modeled $\gamma$, the power-law of the inner-disk surface density. Only modeled candidate disks are included in this figure. } \label{radiusvsgamma} \end{minipage} \end{figure} We generated empirical cumulative distribution functions (CDF) for the modeled Class 0 and Class I candidate disks, as well as for the modeled candidate disks which lie in single and multiple systems (Figures \ref{CDF_classes_radius} to \ref{CDF_mults_gammas}) to examine differences between the sub-samples with disk radius, flux, and $\gamma$. We performed Anderson-Darling (AD) tests \citep{ScholzStephens1987} on each set of empirical CDFs to determine whether the populations are consistent with being drawn from the same distribution. The AD-test is comparable to the Kolmogorov-Smirnoff (KS) test but is more statistically robust---especially if differences between the samples are at the ends of the distribution or if there are small yet significant deviations through the whole distribution---because the KS-test uses maximum deviation to calculate the probability, which the AD-test does not rely on. The results of our AD-tests for the CDFs of disk radius indicate that there is high probability that Class 0 and Class I disk candidates as well as candidate disks from single and multiple systems are drawn from the same distributions (p-values of 0.46 and 0.85, respectively). The AD-tests for the CDFs of flux for the Class 0 and Class I candidate disks have a probability of being drawn from the same distribution of 0.35, and for the candidate disks in single and multiple systems have a p-value of 0.22. Finally, the AD-tests for the CDFs of $\gamma$ for Class 0 and Class I systems have a p-value of 0.62, and there is a p-value of 0.30 for the candidate disks in single and multiple systems. In all cases, the null-hypothesis that the two sub-samples are drawn from the same distribution cannot be ruled out. This may indicate that the disk properties of protostars are usually defined early in the Class 0 phase and do not vary greatly through the Class I phase. \begin{figure}[!h] \centering \begin{minipage}{0.47\textwidth} \centering \includegraphics[width=0.95\linewidth]{f11.pdf} \caption{Empirical cumulative distribution function versus modeled radius for the Class 0 and Class I protostars. p=0.46} \label{CDF_classes_radius} \end{minipage}% \hspace{0.2cm} \begin{minipage}{0.47\textwidth} \centering \includegraphics[width=0.95\linewidth]{f12.pdf} \caption{Empirical cumulative distribution function versus modeled radius for the single and multiple protostars p=0.85} \label{CDF_mults_radius} \end{minipage} \end{figure} \begin{figure}[!h] \centering \begin{minipage}{.47\textwidth} \centering \includegraphics[width=0.95\linewidth]{f13.pdf} \caption{Empirical cumulative distribution function versus observed flux for the Class 0 and Class I protostars. p=0.35} \label{CDF_classes_fluxes} \end{minipage}% \hspace{0.2cm} \begin{minipage}{0.47\textwidth} \centering \includegraphics[width=0.95\linewidth]{f14.pdf} \caption{Empirical cumulative distribution function versus observed flux for the single and multiple protostars. p=0.22} \label{CDF_mults_fluxes} \end{minipage} \end{figure} \begin{figure}[!h] \centering \begin{minipage}{.47\textwidth} \centering \includegraphics[width=0.95\linewidth]{f15.pdf} \caption{Empirical cumulative distribution function versus modeled $\gamma$, the power-law of the inner-disk surface density, for the Class 0 and Class I protostars. p=0.62} \label{CDF_classes_gammas} \end{minipage}% \hspace{0.2cm} \begin{minipage}{0.47\textwidth} \centering \includegraphics[width=0.95\linewidth]{f16.pdf} \caption{Empirical cumulative distribution function versus modeled $\gamma$, the power-law of the inner-disk surface density, for the single and multiple protostars. p=0.30} \label{CDF_mults_gammas} \end{minipage} \end{figure} \subsection{Candidate Disk Sources with Detected Polarized Emission} \label{vandam:discussion:polarized} There are three proposed mechanisms that cause polarized emission in protostellar sources. Polarization may be caused by dust grains interacting with the magnetic field. Non-spherical spinning dust grains are expected to align their short axis parallel to the magnetic field, resulting in polarization orientation perpendicular to the magnetic field \citep[e.g.,][]{Lazarian2007}. In the bulk of the disk where the magnetic energy is expected to be smaller than the rotational energy, the magnetic field is expected to be wrapped into a toroidal configuration by disk rotation. Polarization may also be due to Rayleigh scattering of radiation from the disk and central protostar if grain sizes are not much smaller than the observational wavelength \citep{Kataoka2015}. A new polarization mechanism of dust continuum emission has been posited: radiative alignment of grains with radiation anisotropy \citep{Tazaki2017}. This new radiative alignment scenario causes larger grains ($>$mm size), which are not expected to align with the magnetic field, to be aligned by radiative torques from an anisotropic radiation source with their short axes along the direction of the radiative flux, causing polarized emission. These polarization mechanisms will produce different patterns of polarization morphology and have different efficiencies at varying wavelengths, though determining the definitive contributing mechanisms in any given source is not straightforward and requires detailed modeling \citep{Yang2016}. Even when a magnetic field is present, dust grains may not align with the magnetic field if high gas densities cause sufficient random collisions to prevent grain alignment. If submillimeter and millimeter sized grains are settled in the midplane of the disk, gas density is so high that grain alignment with the magnetic field is difficult to achieve due to gaseous dampening even when unusually strong magnetic fields and grains with many superparamagnetic inclusions are taken into account \citep{Tazaki2017}. While $\sim$mm sized grains do tend to be well settled in the midplane for Class II disks \citep{Pinte2016,Guilloteau2016}, the same is not necessarily true for less evolved Class 0 disks, where grains have grown to millimeter sizes but not yet settled to the midplane \citep{Yang2017}. The Class 0 disk HH212 has resolved, geometrically-thick vertical disk structure in ALMA 850 $\mu$m data \citep{Lee2017}. The vertical structure of L1527 is also resolved by ALMA and geometrically-thick at $\sim$0.8 mm \citep{Sakai2017}. The $\sim$mm sized grains further from the midplane in these Class 0 disks will be in a less dense gaseous environment and alignment of grains perpendicular to a toroidal disk magnetic field may still be possible. Polarization in the envelope, on larger scales than the disk, is likely to be caused by magnetic fields due to the inefficiency of radiative alignment and scattering due to smaller grain sizes and a less dense environment. Polarized emission has been detected towards Class 0 protostars L1527 and IRAS 16293-2422 B at 1.3 mm and 878 $\mu$m wavelengths respectively \citep{SeguraCox2015,Rao2014}, tracing grain sizes similar to the grains not yet settled in the midplanes of Class 0 disks L1527 and HH212. Indeed, if the polarization mechanism in L1527 and IRAS 16293-2422 B is assumed to be purely due to magnetic field alignment, the inferred magnetic field morphologies from polarization in these sources are well-described by disk-wrapped toroidal magnetic fields. In short, while scattering and radiative alignment likely dominate polarized emission in evolved Class II disks (with $\sim$mm sized grains settled close to the midplane), the geometrically-thick $\sim$mm grain populations in the disks of Class 0 sources mean that for the youngest protostars polarization from magnetic fields cannot be ruled out. Five of our 18 candidate disks also have observations of polarized emission towards them: HH211-mms, NGC 1333 IRAS 4A, NGC 1333 IRAS2A, SVS13B, and L1448 IRS3B (see Appendix \ref{vandam:gallery}). These sources are all Class 0 or Class 0/I protostars, young sources which may have thick disks and the potential for emission from $\sim$mm sized grains to become polarized via the magnetic field. We do assume for discussion here that the polarization is purely a signature of the magnetic field in these systems, yet we note that scattering and radiative alignment may also contribute to the polarized emission. In \citet{SeguraCox2015}, we suggested a tentative trend of misaligned inferred magnetic field from polarization and rotation axes in Class 0 systems with disks, with the misaligned orientation helping to reduce the effect of magnetic braking and grow disks with $R$ $>$10 AU at early times \citep[e.g.,][]{Joos2012}. The scenario with aligned inferred magnetic fields and rotation axes would not inhibit magnetic braking effects and a smaller disk would be expected to form. All five of the candidate disks with polarized data have inferred magnetic fields either perpendicular or misaligned with the rotation axes of the systems. The inferred magnetic field morophology of HH211-mms from the SMA data near disk size scales \citep{Lee2014} resembles with a disk-wrapped toroidal field in the northwest with poloidal field lines expected in envelopes and outflows in the southwest and southeast, though scattering may also play a role in polarization on these scales. The true magnetic field morphology of the SMA data on inner envelope and disk size scales is poorly determined due to lack of polarization detected across the entire protostar. The SCUPOL \citep{Matthews2009} and CARMA \citep{Hull2014} inferred magnetic fields are misaligned with both the HH211-mms candidate disk orientation and outflow orientation. The complicated polarization morphology in this source in particular may be partially explained by contributions from either scattering or radiative alignment, though magnetic field contributions cannot be ruled out. NGC 1333 IRAS4A's polarization morphology was long attributed to magnetic fields, which are misaligned with the outflow of the protostar. The large-scale polarization \citep{Girart2006}, with an hourglass-shaped inferred magnetic field is unlikely to be dominated by scattering due to the column densities and small grain sizes in the envelope \citep{Yang2016}. The small-scale polarized emission from NGC 1333 IRAS4A \citep{Cox2015} was found to be consistent with a mix of scattering and a toroidal magnetic field wrapped by the disk \citep{Yang2016}. Nevertheless, a magnetic field causing the polarized emission cannot be ruled out even on small scales. If we assume the extreme limit of the polarization being caused by purely magnetic fields, the fields are consistent with large scale vertical-poloidal hourglass fields misaligned with the rotation axis of the disk, transitioning to small scale frozen-in disk-wrapped toroidal fields in the magnetized disk \citep{Hennebelle2009,Kataoka2012}. As material infalls from the envelope to an embedded disk, frozen-in magnetic field lines will be drawn inwards as well, changing the magnetic field morphology between the envelope and disk \citep{Li2014}. In the case of NGC 1333 IRAS 4A, the inferred magnetic field from polarization on small scales is congruent with the idea that misaligned magnetic fields and rotation axes do not inhibit disk growth by magnetic braking in the same way as the aligned scenario. The inferred magnetic field morphologies from polarization for the five polarized VANDAM candidate disks are consistent with the scenario of perpendicular or misaligned inferred magnetic field orientations compared to the rotation axes of the systems, inhibiting magnetic braking and allowing the disks to grow larger than 10 AU at early times \citep{Hennebelle2009,Joos2012,Li2013,Krumholz2013}. We note that the envelope-scale polarization, which is less likely to be affected by scattering and more likely due to magnetic fields, for all five sources follow this trend. This tentative trend of misalignment between inferred magnetic fields and outflows in Class 0 sources with disks can be strengthened with previous observations of sources outside the Perseus molecular cloud. In Table 2 of \citet{SeguraCox2015}, the Class 0 sources L1527, IRAS 16293-2422 B, and VLA 1623 also have perpendicular inferred magnetic fields from polarization and outflows with strong evidence of embedded disks \citep{SeguraCox2015}. Combined this brings the total of known and candidate Class 0 disks with misaligned outflow and inferred magnetic field orientations to eight sources, double the number of to-date Keplerian-confirmed Class 0 disks. Misalignment between inferred magnetic fields on envelope scales and outflows may indeed be a signpost of disks on smaller scales in Class 0 sources with geometrically-thick disks. With the 12 AU resolution VANDAM survey (with an 8 AU radius modeling limit), we find a total of 18 Class 0 and I candidate disks the Perseus molecular cloud, more than doubling the number of known Class 0 and I disks to-date, bringing the total count from $\sim$15 to $\sim$33. We were able to model 14 of the 18 candidate disks; four are highly asymmetric and not fit well by our disk model. Because we do not have small-scale kinematic data to confirm that these are rotationally supported disks, we refer to the VANDAM sources as disk candidates. We fit the deprojected, azimuthally averaged, and radially binned VLA 8 mm continuum data in {\it u,v}-space to a disk-shaped profile to determine disk candidacy of the extended sources and to begin to model disk properties. We fix the inclination and position angles of the disks using estimates from the image-plane. We take into account a point-source component to account for free-free emission from the jet-launching regions of the disks. NGC 1333 IRAS 4A and NGC 1333 IRAS 4B have obvious envelope contamination in both the image and {\it u,v}-planes. We apply a \textit{u,v}-cut to the data to account for this in the image plane and do not fit the corresponding removed inner baselines to the disk profile in {\it u,v}-space. Other sources have minimal envelope contamination. Except for sources with envelope contamination or small asymmetric features, the residuals from subtracting the model from the data are $<3\sigma$. For all but one candidate disk, the major axis of the disk is roughly perpendicular to the outflow axis (a proxy for the rotation axis), as expected for rotating protostellar disks. 33\% of Class 0 (and Class 0/I) protostars and just 11\% of Class I protostars have candidate disks with radii larger than the 8 mm VLA 12 AU resolution data in Perseus. There is a mild trend that more evolved sources, as gauged by T$_{bol}$, have higher values of $\gamma$ for the 8 mm data; older sources appear to have more centrally concentrated 8 mm dust grain populations. This may be due to the 8 mm grains in older sources having more time to undergo radial drift towards the central protostar than younger sources, resulting in a steeper inner-disk surface density power law ($\gamma$) for more evolved sources. Additionally, the two largest candidate disks belong to Class 0 protostars. 78\% of Class 0 and I protostars do not have signs of disks within our 12 AU resolution limit; at 8 mm most disks in the Class 0 and I phases are small ($<$10 AU). Our estimated masses of the candidate disks are large compared to masses of known Class 0 and Class I disks, indicating that our assumed value of dust opacity spectral index $\beta$ = 1 is too large. Most disks have best-fit models with $q$ $<$0.5, typical for embedded disks. Values of $\gamma$ are lower and more shallow for Class 0 and I candidate disks than in Class II disks, indicating that evolved disks are generally more centrally concentrated than our Class 0 and I disks. Modeled radii of the candidate disks are $>$10 AU at 8 mm and comparable to known Class 0 disk radii determined from kinematics at $\sim$mm wavelengths. Per-emb-14 has a $\sim$3$\times$ larger disk at 1.3 mm with a smaller grain population, evidence that our 8 mm data is a lower limit on true disk radius. Since our 8 mm data trace a population of larger dust grains which radially drift towards the protostar and are lower limits on true disk size, large disks at early times do not seem to be particularly rare. To examine how multiplicity and evolution affect disk radius, $\gamma$, and 8 mm flux, we performed AD tests on CDFs for the modeled Class 0 and Class I candidate disks, as well as CDFs for our modeled disks which are found in single and binary systems. In all cases, we cannot rule out the hypothesis that Class 0 and Class I protostars or single and multiple systems are drawn from the same distribution. Disk properties may be defined early in the Class 0 phase, without much variation through the Class I phase. Five of the 18 candidate disk sources also have polarization detections. In the extreme case of ignoring scattering and radiative alignment, which may contribute to polarized emission but do not preclude magnetic fields contributing to the polarization signal in geometrically-thick Class 0 disks, the five candidate disks with polarized emission have inferred magnetic field morphologies misaligned with the outflows/rotation axes of the systems. Three additional Class 0 disks from other molecular clouds have misaligned inferred magnetic field and outflow directions, bringing the total number to eight. Misalignment between inferred magnetic fields on envelope scales and outflows may be used in the future as indirect evidence of possible geometrically-thick disks on smaller scales in Class 0 protostars. | 18 | 8 | 1808.10438 |
1808 | 1808.00534_arXiv.txt | The 3.2 gigapixel LSST camera, an array of 189 thick fully-depleted CCDs, will repeatedly image the southern sky and accomplish a wide variety of science goals. However, its trove of tens of billions of object images implies stringent requirements on systematic biases imprinted during shift-and-stare CCD observation. In order to correct for these biases which, without correction, violate requirements on weak lensing precision, we investigate CCD systematics using both simulations of charge transport as well as with a unique bench-top optical system matched to the LSST's fast f/1.2 beam. By illuminating single CCDs with realistic scenes of stars and galaxies and then analyzing these images with the LSST data management pipelines, we can characterize the survey's imaging performance well before the camera's first light. We present measurements of several CCD systematics under varying conditions in the laboratory, including the brightness-dependent broadening of star and galaxy images, charge transport anomalies in the silicon bulk as well as the edges, and serial deferred charge. Alongside these measurements, we also present the development and testing of physics-based models which inform corrections or mitigation strategies for these systematics. Optimization of the CCD survey operation under a variety of realistic observational conditions, including systematic effects from the optics, clocking, sky brightness, and image analysis, will be critical to achieve the LSST's goals of precision astronomy and cosmology. | \label{sec:intro} The focal plane of the Large Synoptic Survey Telescope (LSST) will be tiled with an array of 189 charge-coupled devices (CCDs) totaling over 3 billion pixels. This gigantic camera, coupled with the LSST's uniquely wide and fast f/1.2 optics, will enable rapid surveying of the entire visible sky every few nights \cite{Ivezic2008arXiv}. This is due to the unprecedented etendue of the camera-telescope system. The statistical power of this data set will enable a wide variety of goals, from an inventory of our Solar System and Milky Way to investigating the nature of the accelerating Universe. However, great statistical precision requires greater control of systematics. In order to achieve the diverse goals of the LSST, it is necessary to correct for an increasing number of systematic errors that are being discovered in the quest for sub-percent precision CCD observations \cite{Doherty2014SPIE}. Each of the LSST camera's 3.2 billion pixels are towers of silicon, $100 \mu m$ tall by $10 \mu m$ square, defined and read out by grids of potentials and ion implants at the surface and edges of each chip. In the silicon, electrons and holes liberated by the photoconversion process must find their way to the surface where they can be read out. However, on their way to the potential wells which comprise the pixel array, electrons can be diverted by many sources of stray electric fields. These transverse fields induce charge transport anomalies which affect signal electrons in the bulk of the silicon, at its edges, or upon readout of the charge \cite{Stubbs2014JInst}. Each of them undermine the classical assumption of a CCD as a rectilinear grid of independent pixels that most analysis methods rely upon. Left uncorrected, these anomalies develop into systematic errors which bias the measured centroid, shape, and flux of object images, as well as any downstream measurements of the image properties of stars and galaxies which are needed to understand astrophysics and cosmology to sub-percent precision. In this report we will review a few of the most prominent systematic errors that we have characterized in the lab, including the brighter-fatter (BF) effect, astrometric distortions at the edge and in the bulk, and charge transport inefficiency. We present measurements of these effects on realistic stars and galaxies projected onto prototype LSST CCDs, which are then imaged, processed, and analyzed with state-of-the-art methods. Alongside these lab measurements we also present models of the microphysics of charge transport, inside of which we can ``turn'' some of the same ``knobs'' which are accessible to us in the lab, thereby informing the proposed methods of correction and making model predictions \cite{Lage2017JInst}. In Section \ref{sec:setup} of this report we first describe the experimental setup which utilizes a fast f/1.2 beam simulator \cite{Tyson2014SPIE} to re-image realistic PSFs and galaxies onto prototype LSST CCDs (see Figure \ref{fig:beamsim}). We also briefly describe our image processing and analysis tools. In Section \ref{sec:bf} we describe the measurement and modeling of the BF effect, as well as our ongoing attempts at correction using a model in the literature. In Section \ref{sec:astrometry} we demonstrate the measurement and modeling of astrometric shifts due to readout electronics at the edges, and those due to tree rings in the bulk. In Section \ref{sec:cti} we present the measurement and impact of charge transport inefficiency (CTI), and deferred charge more generally, on the shapes of stars. We discuss these measurements and future plans in Section \ref{sec:discuss}, and conclude in Section \ref{sec:conclude}. \begin{figure}\includegraphics[width=.9\columnwidth]{figures/ccd-beamsim.png}\centering \caption[Beam lab]{The LSST beam simulator at UC Davis. A: cryogenic Dewar with CCD mounted a few centimeters behind the input window. B: LSST f/1.2 beam simulator optics, where a spot mask with thousands of pinholes can be placed at the object plane. C: Illuminating sphere which scatters input light to provide uniform coverage of optics. D: Precision X/Y/Z stage upon which the CCD is mounted, enabling micron (1/10th of a pixel) dithering capability.} \label{fig:beamsim} \end{figure} | \label{sec:conclude} To maximize the scientific output of the survey, optimization of observing strategies and mitigation of systematic errors must be carried out before the survey begins. We present ongoing work in the lab to measure and correct for several prominent charge transport systematics using realistic images and modern processing and analysis tools. We quantify the brighter-fatter effect on realistic PSFs and galaxies, as well as test a correction method in development. We demonstrate the strong dependence of the charge transport systematics on operational conditions using the edge astrometric distortion, and successfully capture these variations in simulations. We also present the detection of astrometric distortion due to tree rings, and the effects of deferred charge ($\sim CTI$) are measured and compared to physics-based models. These laboratory measurements provide valuable insight that informs the development of correction algorithms. To properly correct for these systematic errors, new techniques must be developed to sequentially peel back the layers of physical effects which obscure the truth in modern surveys of the sky. | 18 | 8 | 1808.00534 |
1808 | 1808.05916_arXiv.txt | X-ray observations of bright AGNs in or behind galaxy clusters offer unique capabilities to constrain axion-like particles. Existing analysis technique rely on measurements of the global goodness-of-fit. We develop a new analysis methodology that improves the statistical sensitivity to ALP-photon oscillations by isolating the characteristic quasi-sinusoidal modulations induced by ALPs. This involves analysing residuals in wavelength space allowing the Fourier structure to be made manifest as well as a machine learning approach. For telescopes with microcalorimeter resolution, simulations suggest these methods give an additional factor of two in sensitivity to ALPs compared to previous approaches. | \label{introduction} Axion-like particles are one of the simplest and best motivated extensions of the Standard Model. They arise as a generalisation of the QCD axion, which was introduced as an appealing solution to the strong CP problem of QCD (\cite{PecceiQuinn, Weinberg, Wilczek}) - a recent review of axion physics is \cite{Marsh}. While the original QCD axion requires a coupling to the strong force, it is also interesting to consider more general \emph{axion-like particles} (ALPs) that couple only to electromagnetism. Such ALPs arise generally in string compactifications (for example, see \cite{Conlon:2006tq, Svrcek:2006yi, Arvanitaki:2009fg, Cicoli:2012sz}). An ALP $a$ interacts with photons via the coupling: \bea & g_{a\gamma\gamma}\, a\, \vec{E}\cdot\vec{B}, & \eea where $g_{a\gamma\gamma}$ is a dimensionful coupling parametrizing the strength of the interaction and $\vec{E},\vec{B}$ are the electric and magnetic fields, respectively. While we refer in this paper to ALPs, the physics is also relevant for photophilic models of the QCD axion in which the mass is much smaller (or the photon coupling significantly enhanced) compared to naive expectations, such as those in \cite{Farina, Agrawal}. As they attain masses only by non-perturbative effects, ALPs naturally have extremely small masses. The relevant physics is described by the Lagrangian \bea \mathcal{L} & = & \frac{1}{2} \partial_{\mu} a\, \partial^{\mu} a + \frac{1}{2} m_a^2 a^2 + g_{a\gamma\gamma}\,a\,\vec{E}\cdot\vec{B}. \eea The ALP-photon coupling produced by the $a\vec{E}\cdot\vec{B}$ interaction implies that, within background magnetic fields, the ALP state $a$ has a 2-particle interaction with the photon $\gamma$. Under this mixing, the `mass' eigenstates of the Hamiltonian are a mixture of the photon and ALP `flavour' eigenstates, causing oscillation between the modes in a way analogous to neutrino oscillations (\cite{Sikivie:1983ip, RS}). ALP-photon conversion is enhanced by large magnetic field coherence lengths. As it extends over megaparsec scales and contains coherence lengths up to tens of kiloparsecs, this makes the intracluster medium of galaxy clusters particularly efficient (\cite{Burrage:2009mj, Angus:2013sua, Conlon:2013txa, Powell:2014mda, Schlederer:2015jwa, Conlon:2015uwa}). This has allowed the use of X-ray point sources, located in or behind clusters, to produce leading bounds on ALP parameter space (\cite{Wouters:2013hua, 160501043, 1704.05256, Marsh:2017yvc, Conlon:2017ofb, Chen:2017mjf}). While X-rays are particularly good at probing ALPs with $m \lesssim 10^{-11} {\rm eV}$, more massive ALPs can be probed via gamma-ray astronomy - see \cite{Wouters:2012qd, Ajello, Payez, MajumdarHorns} for some related work in different wavebands. The aim of this work is to improve the theoretical tools available for the analysis of ALP-photon conversion, and in particular the statistical methodology used to analyse data. Existing approaches to bounding ALP-photon conversion involve, roughly, a comparison of the overall goodness-of-fit (as measured by the reduced $\chi^2$) of models with ALPs compared to models without ALPs. For sufficiently large ALP-photon coupling, models with ALPs lead to a bad fit. However, this approach makes little use of the actual structure of ALP-photon conversion. While the precise form of ALP-photon conversion depends on the (unknown) magnetic field along the line of sight, there is a quasi-sinusoidal oscillatory structure that is common to all realisations of the magnetic field. The aim of this paper is to develop analysis techniques to isolate this structure, and thereby allow greater sensitivity to constraining ALPs. This development of improved statistical methods is also important in view of the launch over the next decade of microcalorimeter-based X-ray telescopes such as XARM and ATHENA; the use of an overall $\chi^2$ fit would then fail to take full advantage of the energy resolution provided by such instruments. With an eye on the opportunities that will come from future satellites such as XARM and ATHENA, we focus on energies and parameter ranges that are applicable for the case of X-ray photons converting to ALPs within galaxy cluster magnetic fields. | X-ray observations of AGNs offer one of the most sensitive probes of axion-like particles in the low mass regime and it is therefore important to ensure that maximal statistical sensitivity is obtained in searches for them. In this paper we have developed analysis methods aimed at exploiting the characteristic quasi-sinusoidal modulations that are induced by ALPs. Based on simulated data, these methods appear to improve sensitivity to photon-axion conversion by a factor of two (in conversion probability). | 18 | 8 | 1808.05916 |
1808 | 1808.09309_arXiv.txt | Results from regular monitoring of relativistic compact binaries like PSR 1913+16 are consistent with the dominant (quadrupole) order emission of gravitational waves (GWs). We show that observations associated with the binary black hole central engine of blazar OJ~287 demand the inclusion of gravitational radiation reaction effects beyond the quadrupolar order. It turns out that even the effects of certain hereditary contributions to GW emission are required to predict impact flare timings of OJ~287. We develop an approach that incorporates this effect into the binary black hole model for OJ~287. This allows us to demonstrate an excellent agreement between the observed impact flare timings and those predicted from ten orbital cycles of the binary black hole central engine model. The deduced rate of orbital period decay is nine orders of magnitude higher than the observed rate in PSR 1913+16, demonstrating again the relativistic nature of OJ~287's central engine. Finally, we argue that precise timing of the predicted 2019 impact flare should allow a test of the celebrated black hole ``no-hair theorem" at the $10\%$ level. | \label{sec:intro} OJ~287 (RA: 08:54:48.87 \& DEC:+20:06:30.6) is a bright blazar, a class of active galactic nuclei, situated near the ecliptic in the constellation of Cancer. This part of the sky has been frequently photographed for other purposes since late 1800's and therefore it has been possible to construct an exceptionally long and detailed light curve for this blazar using the historical plate material. It is at a redshift ($z$) of 0.306 corresponding to a luminosity distance of 1.6 Gpc in the standard $\Lambda$CDM cosmology which makes it a relatively nearby object as blazars go. The optical light curve, extending over 120 yr \citep{sil88,hud13}, exhibits repeated high-brightness flares (see Figure~\ref{fig:lightcurve}). A visual inspection reveals the presence of two periodic variations with approximate timescales of 12 yr and 60 yr which have been confirmed through a quantitative analysis \citep{val06}. We mark the $\sim 60$ year periodicity by a red curve in the left panel of Figure~\ref{fig:lightcurve} and many observed outbursts/flares are separated by $\sim 12$ years. The regular monitoring of OJ~287, pursued only in the recent past, reveal that these outbursts come in pairs and are separated by a few years. The doubly peaked structure is shown in the right panel of Figure~\ref{fig:lightcurve}. The presence of double periodicity in the optical light curve provided an early evidence for the occurrence of quasi-Keplerian orbital motion in the blazar, where the 12 year periodicity corresponds to the orbital period timescale and the 60 year timescale is related to the orbital precession. The ratio of the two deduced periods gave an early estimate for the total mass of the system to be $\sim 18\times 10^9\, M_{\odot}$, provided we invoke general relativity to explain the orbital precession \citep{pie98}. It is important to note that this estimate is quite independent of the detailed central engine properties of OJ~287. The host galaxy is hard to detect because of the bright nucleus; however, during the recent fading of the nucleus by more than two magnitudes below the high level state it has been possible to get a reliable magnitude of the host galaxy. It turns out to be similar to NGC~4889 in the Coma cluster of galaxies, i.e. among the brightest in the universe. These results will be reported elsewhere (Valtonen et al. 2018). These considerations eventually led to the development of the binary black hole (BBH) central engine model for OJ~287 \citep{leh96,val08a}. \begin{figure} \centering \includegraphics[width=\linewidth]{light_curve.pdf} \caption{ The left panel displays the optical light curve of OJ~287 from 1886 to 2017. We draw a fiducial curve for easy visualization of the inherent long-term variations. The right panel shows the observed double-peaked structure of the high-brightness flares. The positions of the two peaks are indicated by downward arrows from the top of the panel.} \label{fig:lightcurve} \end{figure} According to the BBH model, the central engine of OJ~287 contains a binary black hole system where a super-massive secondary black hole is orbiting an ultra-massive primary black hole in a precessing eccentric orbit with a redshifted orbital period of $\sim 12$ yr (see Figure~\ref{fig:model}). The primary cause of certain observed flares (also called outbursts) in this model is the impact of the secondary black hole on the accretion disk of the primary \citep{leh96,pih16}. The impact forces the release of two hot bubbles of gas on both sides of the accretion disk which radiate strongly after becoming optically transparent, leading to a sharp rise in the apparent brightness of OJ~287. The less massive secondary BH impacts the accretion disk twice every orbit while traveling along a precessing eccentric orbit (Figure~\ref{fig:model}). This results in double-peaked quasi-periodic high-brightness (thermal) flares from OJ~287. Furthermore, large amounts of matter get ejected from the accretion disk during the impact and are subsequently accreted to the disk center. This ensures that part of the unbound accretion-disk material ends up in the twin jets. % The matter accretion leads to non-thermal flares via relativistic shocks in the jets which produce the secondary flares in OJ~287, lasting more than a year after the first thermal flare \citep{val09}. \begin{figure} \centering \includegraphics[width=\linewidth]{OJ287EN-PN4_5scaled1.pdf} \caption{ Artistic illustration of the binary black hole system in OJ~287. The present analysis provides an improved estimate for the spin of the primary black hole. \label{fig:model}} \end{figure} The BBH model of OJ~287 can be used to predict the flare timings \citep{sun97,val08b,val11b} and the latest prediction was successfully verified in 2015 November. The optical brightness of OJ~287 rose above the levels of its normal variations on November 25, and it achieved peak brightness on December 5. On that date, OJ~287 was brighter than at any time since 1984 \protect\citep{val16}. Owing to the coincidence of the start of the flare with the date of completion of general relativity (GR) by Albert Einstein one hundred years earlier, it was termed as the GR centenary flare. Detailed monitoring of the 2015 impact flare allowed us to estimate the spin of the primary BH to be $\sim 1/3$ of the maximum value allowed in GR. This was the fourth instance when multi-wavelength observational campaigns were launched to observed predicted impact flares from the BBH central engine of OJ~287 \citep{val08b,val16}. The latest observational campaigns confirmed the presence of a spinning massive BH binary inspiraling due to the emission of nano-Hertz gravitational waves in OJ~287. These developments influenced the Event Horizon Telescope consortium to launch observational campaigns in 2017 and 2018 to resolve the presence of two BHs in OJ~287 via the millimeter wavelength Very Long Baseline Interferometry. Predictions of impact flare timings are made by solving post-Newtonian (PN) equations of motion to determine the secondary BH orbit around the primary while using the observed outburst times as fixed points of the orbit. The PN approximation provides general relativistic corrections to Newtonian dynamics in powers of $(v/c)^2$, where $v$ and $c$ are the characteristic orbital velocity and the speed of light, respectively. The GR centenary flare was predicted using 3PN-accurate (i.e., third PN order) BBH dynamics that employed GR corrections to Newtonian dynamics accurate to order $(v/c)^6$ \citep{val10a,val10b,val11a}. Additionally, earlier investigations invoked nine fixed points in the BBH orbit, which allowed the unique determination of eight parameters of the OJ~287 BBH central engine model \citep{val10a, val11a}. The GR centenary flare provided the tenth fixed point of BBH orbit, which opens up the possibility of constraining an additional parameter of the central engine. Moreover, the GW emission-induced rate of orbital period decay of the BBH in OJ~287, estimated to be $\sim 10^{-3}$, makes it an interesting candidate for probing the radiative sector of relativistic gravity \citep{Wex14}. These considerations influenced us to explore the observational consequences of incorporating even higher-order PN contributions to the BBH dynamics. Therefore, we introduce the effects of GW emission beyond the quadrupolar order on the dynamics of the BBH in OJ~287 while additionally incorporating next-to-leading-order spin effects \citep{LB_review, BBF_06,wil17}. Moreover, we incorporate the effects of dominant order hereditary contributions to GW emission, detailed in \cite{BS93}, on to the binary BH orbital dynamics. It turns out that these improvements to BBH orbital dynamics cause non-negligible changes to our earlier estimates for the BBH parameters, especially for the dimensionless angular momentum parameter of the primary BH in OJ~287, and the inclusion of present improvements to BBH orbital dynamics should allow the test of the black hole ``no-hair theorem" during the present decade. This is essentially due to our current ability to accurately predict the time of the next impact flare from OJ~287, influenced by the present investigation. This paper is structured as follows. In Section~\ref{sec:pnd}, we discuss briefly the improved BBH orbital dynamics. Details of our approach to obtain the parameters of the BBH system from optical observation of OJ~287 are presented in Section~\ref{sec:bbhorbit}. How we incorporate the effects of dominant-order ``hereditary" contributions to GW emission into BBH dynamics is detailed in Section~\ref{sec:radfac}. Implications of our improved BBH model on historic and future observations are outlined in Section~\ref{sec:lc_comparison}. In Appendix~\ref{app:kdndldetdl}, we display PN-accurate expressions used to incorporate ``hereditary" contributions to BBH dynamics. | \label{sec:dis_con} The present paper provides the most up-to-date and improved description for the binary black hole central engine of OJ~287. This is mainly because of the use of an improved PN prescription to describe the BBH orbit evolution. We incorporate in the BBH dynamics the effects of next-to-next-to-next-to leading (or quadrupolar) order GW emission. This includes effects due to the dominant-order hereditary contribution to the GW-induced inspiral. It turns out to be crucial to incorporate the effect of hereditary contributions to GW emission on the BBH dynamics, and we develop an approach to model the effect into ${\ddot{\vek x}}$ with the help of an unknown parameter $\gamma$. The observationally determined value $\gamma_{obs}$ shows remarkable agreement with its GR-based estimate $\gamma_{GR}$, obtained by adapting GW phasing formalism for eccentric binaries. This formalism is required to construct accurate inspiral templates to model GWs from compact binaries that are inspiraling along PN-accurate eccentric orbits. Furthermore, we incorporate next-to-leading-order spin-orbit contributions to the compact binary dynamics, influenced by \cite{BBF_06} and \cite{wil17}. This leads to a noticeably different estimate for the Kerr parameter of the primary BH, namely $\chi= 0.381 \pm 0.004$, compared to $\chi=0.313$ in \cite{val16}. Additionally, the rate of decay of the binary orbit is slower than in earlier models by about $6.5\%$ \citep{val10b}. The improved description allows us to demonstrate excellent agreement between the observed impact flare timings and those predicted from the BBH central engine model. These improvements should allow us to employ the BBH central engine model to test GR in the strong-field regime that at present is not accessible to any other observatories or systems. The first such test is possible in 2019 July. The next major flare will peak on July 31, around noon GMT in our model. The model without higher-order gravitational radiation reaction terms gives the brightness peak 1.57 days earlier, in the early hours of July 30 GMT. These two models are easily differentiated by observations, provided we are able to monitor OJ~287 during late July 2019. The closeness to the Sun in the sky makes such an effort extremely difficult. However, a successful observational campaign should provide us the unique opportunity to test the black hole no-hair theorem at the $\sim 10 \%$ level during the present decade. Additionally, we demonstrate the possibility of predicting the general shape of the expected optical light curve of OJ~287 during the impact flare season. This should be helpful in analyzing the optical light curve of OJ~287 during the next two accretion impact flares, expected to happen during 2019 and 2022. These observational campaigns will be challenging owing to the apparent closeness of the blazar to the Sun. However, the monitoring of these impact flares should allow us to test general relativity in the strong-field regime that is characterized by $ (v/c) \approx 0.25$ and $ m \approx 18\times 10^9 M_{\odot}$. It will be exciting to extend the preliminary results, displayed in Figure 6 of \cite{val12}, that provided an independent estimate for the mass of the central BH in OJ 287. It turned out that the dynamically estimated total mass in OJ~287 and the measured absolute magnitude of the bulge of the host galaxy is fully consistent with the black hole mass - K-magnitude correlation pointed out in \cite{kormendy11}. | 18 | 8 | 1808.09309 |
1808 | 1808.00472_arXiv.txt | {The fragmentation mode of high-mass molecular clumps and the properties of the central rotating structures surrounding the most luminous objects have yet to be comprehensively characterised.} {We study the fragmentation and kinematics of the high-mass star-forming region \W, as part of the IRAM NOEMA large program CORE.} {Using the IRAM NOrthern Extended Millimeter Array (NOEMA) and the IRAM 30-m telescope, the CORE survey has obtained high-resolution observations of 20 well-known highly luminous star-forming regions in the 1.37~mm wavelength regime in both line and dust continuum emission.} {We present the spectral line setup of the CORE survey and a case study for \W. At $\sim$$0\farcs35$ (700~AU at 2.0~kpc) resolution, the \W\ clump fragments into two cores (West and East), separated by $\sim$2300~AU. Velocity shifts of a few \kms\ are observed in the dense-gas tracer, \mc, across both cores, consistent with rotation and perpendicular to the directions of two bipolar outflows, one emanating from each core. The kinematics of the rotating structure about \W\,W shows signs of differential rotation of material, possibly in a disk-like object. The observed rotational signature around \W\,E may be due to a disk-like object, an unresolved binary (or multiple) system, or a combination of both. We fit the emission of \mckr{4}{6} and derive a gas temperature map with a median temperature of $\sim$165~K across \W. We create a \tq\ map to study the stability of the rotating structures against gravitational instability. The rotating structures appear to be Toomre unstable close to their outer boundaries, with a possibility of further fragmentation in the differentially-rotating core, \W\,W. Rapid cooling in the Toomre-unstable regions supports the fragmentation scenario.} {Combining millimeter dust continuum and spectral line data toward the famous high-mass star-forming region \W, we identify core fragmentation on large scales, and indications for possible disk fragmentation on smaller spatial scales.} | Fundamental questions pertaining to the fragmentation of high-mass clumps and the accretion processes that result in the birth of the most massive stars ($M\gtrsim 8$~\mo) still remain unanswered. This is in part due to the clustered nature of star-formation and the typically large distances involved. For a long time, the existence of high-mass stars had been puzzling as it was thought that the expected intense radiation pressure would prevent the accretion of enough material onto the protostar (\eg\ \citeads{1974A&A....37..149K}; \citeads{1987ApJ...319..850W}). More recently, two- and three-dimensional (magneto)hydrodynamical simulations of collapsing cores have validated the need for accretion disks in the formation of very massive stars, analogous to low-mass star formation (\eg\ \citeads{2002ApJ...569..846Y}; \citeads{2009Sci...323..754K}; \citeads{2010ApJ...711.1017P}; \citeads{2010ApJ...722.1556K}, \citeyear{2011ApJ...732...20K}; \citeads{2013ApJ...772...61K}; \citeads{2016ApJ...823...28K}). Furthermore, different fragmentation processes can contribute to the final stellar mass distribution within a single region, including fragmentation from clouds down to core scales (\eg\ \citeads{2010A&A...524A..18B}; \citeads{2013ApJ...762..120P}, \citeyear{2015MNRAS.453.3785P}; \citeads{2018arXiv180501191B}; see review by \citeads{2017arXiv170600118M}), and disk fragmentation at smaller spatial scales (\eg\ \citeads{2003ApJ...595..913M}; see review by \citeads{2016ARA&A..54..271K}). In the disk-mediated accretion scenario, the non-isotropic treatment of the radiation field reduces the effect of radiation pressure in the radial direction, such that radiation can escape through the poles along the disk rotation axis, while the disk is shielded due to the high densities. Observationally, the existence of such disks is expected due to ubiquitous observations of collimated outflows (\eg\ \citeads{2002A&A...387..931B}; \citeads{2009A&A...504..127F}; \citeads{2011A&A...530A..12L}; \citeads{2014prpl.conf..451F}; \citeads{2015MNRAS.453..645M}), which has also been predicted by theoretical models (\eg\ \citeads{2007prpl.conf..277P}). Although some accretion disks in differential Keplerian-like rotation about B-type (proto)stars have been found in recent years (\eg\ \citeads{2012ApJ...752L..29C}; \citeads{2013A&A...552L..10S}; \citeads{2014A&A...571A..52B}; see reviews by \citeads{2007prpl.conf..197C}, and \citeads{2016A&ARv..24....6B}), the existence of such rotating structures around the most massive, O-type protostars is still elusive, with only a few cases reported so far (\citeads{2015ApJ...813L..19J}; \citeads{2016MNRAS.462.4386I}; \citeads{2017A&A...602A..59C}). As higher resolution observations are becoming more accessible, thus allowing structures to be resolved on scales $<$1000~AU, it is important to determine whether disks around intermediate to high-mass stars (OB-type) are ubiquitous and if so, to characterise their properties. What is the typical extent of these disks? Are they in differential rotation about a centrally-dominating protostar, similar to their low-mass counterparts and if so, over what range of radii? Is there any scale where a core stops fragmenting?\footnote{Here, a core is defined as a gravitationally-bound region that forms a single or multiple stars, following \citeads{2000prpl.conf...97W}.} At what scales do we see the fragmentation of disks? Are close binary/multiple systems an outcome of disk fragmentation as suggested by, for example, \citetads{2018MNRAS.473.3615M}? If so, stability analyses of these high-mass rotating cores and disks are needed to shed light on fragmentation at disk scales. These questions can only be answered with a statistical approach for a large sample of high-mass star-forming regions. We have undertaken a large program at IRAM, called CORE \citepads{2018arXiv180501191B}, making use of the IRAM NOrthern Extended Millimeter Array (NOEMA, formerly Plateau de Bure Interferometer) at 1.37~mm in both line and continuum emission to study the early phases of star formation for a sample of 20 highly luminous ($L>10^4$~\lo) star-forming regions at high angular resolution ($\sim$0.4\arcsec), to analyse their fragmentation and characterise the properties of possible rotating structures. Additionally, observations with the IRAM 30-m telescope are included to complement the interferometric data, allowing us to understand the role of the environment by studying high-mass star formation at scales larger than those covered by the interferometer. Observations in the 1.3~mm wavelength regime of the CORE project began in June 2014 and finished in January 2017, consisting of a total of more than 400 hours of observations with NOEMA. The sample selection criteria and initial results from the observed level of fragmentation in the full sample are presented in \citetads{2018arXiv180501191B}, and details of the 30-m observations and the merging of single-dish with the interferometric observations can be found in Mottram et al. (in prep.). In this work, we describe our spectral setup and present a case study of one of the most promising star-forming cloud in our sample, \W. \W, also known as the ``Turner-Welch object'', resides in the W3 high-mass star-forming region and was initially identified through observations of the dense-gas tracer HCN at 88.6~GHz \citepads{1984ApJ...287L..81T}. It is located $\sim$0.05~pc (5\arcsec) east of the well-known ultra-compact \hii\ region (UCHII) \WOH. The name, \W, stems from the existence of water masers in the vicinity of the source \citepads{1981ApJ...245..857D}, allowing for an accurate distance measurement of 2.0~kpc for this region (\citeads{2006ApJ...645..337H}; \citeads[cf.][]{2006Sci...311...54X}). The relative proper motions of these masers are further explained by an outflow model oriented in the east-west direction \citepads{2006ApJ...645..337H}. A continuum source elongated in the east-west direction and spanning the same extent as the water maser outflow has been observed in sub-arcsecond VLA observations in the radio regime with a spectral index of --0.6, providing evidence for synchrotron emission (\citeads{1995ApJ...443..238R}; \citeads{1999ApJ...513..775W}). This source of synchrotron emission has been characterised by a jet-like model due to its morphology, and the point symmetry of its wiggly bent structure about the center hints at the possibility of jet precession. Moreover, \citetads{2004A&A...418.1045S} attribute this radio emission to a circumstellar jet or wind ionised by the embedded (proto)star at this position. Additional radio continuum sources have been detected in the vicinity of the synchrotron jet, the closest of which is to the west of the elongated structure and has a spectral index of 0.9 (\citeads{1999ApJ...513..775W}; \citeads{2006ApJ...639..975C}), consistent with a circumstellar wind being ionised by another embedded protostellar source. In fact, the high angular resolution ($\sim$0\farcs7) observations of \citetads{1999ApJ...514L..43W} in the 1.36~mm band allowed for the detection of three continuum peaks in thermal dust emission, one of which peaks on the position of the water maser outflow and synchrotron jet, and another on the position of the radio continuum source with positive spectral index, confirming the existence of a second source at this position. The detection of two bipolar molecular (CO) outflows further supports the protobinary scenario, suggesting that \W\ may be harbouring (at least) two rotating structures \citepads{2011ApJ...740L..19Z}. The two cores within \W\ have individual luminosities on the order of $2\times10^4$ \lo, suggesting two 15~\mo\ stars of spectral type B0\footnote{The luminosity and spectral type calculations are described in detail in Section~\ref{ss: mass_estimates}.}. In this paper, we aim to study the fragmentation properties of \W\ and the kinematics of the rotating structures within it. We use this source as a test-bed for what will be expanded in a forthcoming paper which will focus on the kinematic properties of a larger sample within our survey. The structure of the paper is as follows. Section~\ref{s: obs_reduction} presents our spectral line setup within the CORE survey with the details of our observations and data reduction for \W. The observational results are described in Section~\ref{s: obs_results}. The kinematics, temperature, and stability analysis of \W\ is presented in Section~\ref{s: analysis_discussion}. The main findings are summarized in Section~\ref{s: conclusion}. \begin{table} \caption{Observations of \W\ and \WOH.} \centering \label{t:obs_list} \begin{tabular}{lccc} \hline\hline Observation Date & Array & Time On-source & Bandpass \\ & & (h) & Calibrator \\ \hline 2014-Oct-31 & D & 3.9 & 3C454.3 \\ 2015-Mar-18 & A & 2.6 & 3C84\\ 2015-Apr-3 & B & 0.9 & 3C84 \\ 2015-Apr-6 & B & 1.3 & 3C84\\ 2016-Mar-11 & A & 2.2 & 3C84\\ \hline \end{tabular} \tablefoot{The phase and flux calibrators were 0059+581 and MWC349, respectively, for all observations.} \end{table} \begin{table} \caption{Correlator units and frequency ranges observed with NOEMA.} \centering \label{t:correlator_table} \begin{tabular}{lccc} \hline\hline Correlator & Spectral Unit & Pol. & Frequency Range \\ & & & (MHz) \\ \hline Narrow-band & L01 & H & 220\,690.6--220\,769.7\\ & L02 & H & 220\,630.6--220\,709.7\\ & L03 & H & 220\,570.6--220\,649.7\\ & L04 & H & 220\,130.6--220\,209.7\\ & L05 & H & 218\,860.6--218\,939.7\\ & L06 & H & 218\,415.6--218\,494.7\\ & L07 & H & 218\,280.6--218\,359.7\\ & L08 & H & 218\,180.6--218\,259.7\\ \hline WideX & L09 & H & \multirow{2}{*}{218\,878.6--220\,859.5}\\ & L10 & V & \\ & L11 & H & \multirow{2}{*}{217\,078.6--219\,059.4}\\ & L12 & V & \\ \hline \end{tabular} \tablefoot{H and V correspond to horizontal and vertical polarisations.} \end{table} | } Our IRAM large program, CORE, aims to study fragmentation and the kinematics of a sample of 20 high-mass star forming regions. We have performed a case-study for one of the sources in our sample, the prototypical hot core \W. In this paper we present details of the spectral line setup for our NOEMA observations in the 1.37~mm band which covers transitions of important tracers of dense gas and disks (\eg\ \mc, \mf, CH$_3$OH), inner envelopes (\eg\ $\mathrm{H_2CO}$), and outflows (\eg\ $\mathrm{^{13}CO}$, SO). We cover the range of 217\,078.6~$-$~220\,859.5~MHz with a spectral resolution of 2.7~\kms, and eight narrower bands centered on specific transitions to provide higher spectral resolution of 0.5~\kms. With the aim of studying the fragmentation and kinematic properties of \W, the following is a summary of our findings: \begin{itemize} \item At an angular resolution of $\sim$0$\farcs35$ ($\sim$700~AU at 2~kpc), \W\ fragments into two cores, which we refer to as \W~E and \W~W, separated by $\sim$2300~AU, as seen in both line and thermal dust continuum emission. \item Based on the integrated dust continuum emission, \W\ has a total mass of $\sim$26.8~\mo, with 15.4~\mo\ contributed from \W\ E, and 11.4~\mo\ from \W\ W. \item On large scales, there exists a clear velocity gradient in the east-west direction across \W, spanning $\sim$6000~AU, attributed to a combination of circumbinary and circumstellar motions. On smaller scales, velocity gradients of a few \kms\ are observed across each of the cores, perpendicular to the directions of two bipolar molecular outflows, one emanating from each core, as traced by $^{12}$CO and $^{13}$CO. The direction of motion of the gas around the individual cores deviates little from the overall larger-scale rotation of \W, suggesting that these motions, which we interpret as rotation, seen around the cores may be inherited from the large-scale rotation. \item The kinematics of the rotating structure about \W\ W shows differential rotation of material about a (proto)star as deduced from the redshifted part of its PV plot, suggesting that the rotating structure may be a disk-like object. The radio source with a rising spectrum at this position can be attributed to a circumstellar jet or wind ionised by the embedded (proto)star at this position. \item The PV plots of the rotating structure about \W\ E is inconclusive on whether the observed rotation is due to a disk-like rotating object, an unresolved binary (or multiple) system, or a combination of both. \item We fit the emission of \mckr{4}{6} and \mcisokr{0}{3} with \textsc{xclass} and produce temperature, column density, peak velocity, and velocity dispersion maps. On average, the entire structure is hot ($\sim$165~K) with no particular temperature structure. The column density map of \mc\ is doubly peaked, similar to the continuum and line emission maps, with a median column density of $\sim$1.4$\times10^{15}~\mathrm{cm^{-2}}$. Close to the center of the cores, the H$_2$ column density is estimated to be $\sim 5 \times 10^{24}~\mathrm{cm^{-2}}$. \item We investigate the axis-symmetric stability of the two rotating structures using the Toomre criterion. Our \tq\ map shows low values in the outskirts of both rotating structures, suggesting that they are unstable to fragmentation. Some regions with low \tq\ values in the vicinity of \W\ W coincide with unresolved thermal dust continuum peaks (in our highest resolution observations), hinting at the possibility of further fragmentation in this core. \item The Toomre-unstable regions within \W~E and \W~W are able to cool rapidly and any local collapse induced by the gravitational instabilities will lead to further fragmentation. \end{itemize} In this work, we showcased our in-depth analysis for the kinematics and stability of the rotating structures within \W. We showed that high-mass cores can be prone to fragmentation induced by gravitational instabilities at $\sim$1000~AU scales, and core fragmentation at larger scales. Therefore, different modes of fragmentation contribute to the final stellar mass distribution of a given region. The question still remains, how universal are these findings for the high-mass star formation process? To this end, we aim to benchmark our methods using hydrodynamic simulations and extend our analysis to the full CORE sample in future papers. | 18 | 8 | 1808.00472 |
1808 | 1808.09848_arXiv.txt | { CubeSats are tiny satellites with increasing capabilities. They have been used for more than a decade by universities to train students on space technologies, in a hands-on project aiming at building, launching and operating a real satellite. Still today, one shortcoming of CubeSats is their poor ability to transmit large amounts of data to the ground. A possible way to overcome this limitation relies on optical communications. Universit\'e C\^ote d'Azur is studying the feasibility of a student's CubeSat whose main goal is to transmit data with an optical link to the ground at the moderate rate of 1\,kb/s (or better). In this paper, we will present the current state of the project and its future developments. } | \label{sec:intro} CubeSats are small satellites (``SmallSats'' class) made of 10\,cm-side cubes that form a ``unit'' (or a 1U CubeSat). A large number of units can be combined, although the bulk of current developments range from 1U (e.g. Robusta1B), 2U (Spacecube, X-cubesat), up to 3U (e.g. Picsat, Eyesat, NIMPH \cite{fernandez:hal-01274184}). However today, we can start to see 12U CubeSat projects under development by several universities (for example the Grenoble/Toulouse project ATISE \cite{lecoarer:hal-01401693}). These SmallSats are getting more and more attention from the universities because they offer the possibility to teach space-related techniques to students on a hands-on experiment, with a budget that can be reached by a medium-size university. CubeSats also get more and more attention from companies (e.g. Nexeya\footnote{\url{https://www.nexeyaonline.com/small-sats-satellite-platforms}}, Planet Labs\footnote{\url{https://www.planet.com}}) because they offer fast development cycles for new technologies with reduced costs compared to more ``traditional'' satellites. The drawbacks are an extremely small payload volume and mass, a lack of redundancy, and a perfectible reliability, that can be mitigated by payload miniaturization and the use of satellite constellations, or ``flocks'' \cite{boshuizen2014results}. The consequence of this tiny size is the limited data transmission capacities that can be integrated into a CubeSat: most of the radio transmitters (67\% of the 1630 radio emitter-receivers included in CubeSats\footnote{according to \url{http://www.nanosats.eu/index.html\#figures}, consulted in June 2018}) use the amateur radio frequencies to transmit data (UHF -- 437\,MHz \& VHF -- 146\,MHz), a few (6\%) use S-band (2.2-3.4\,GHz), some (25\%) use X-band (10\,GHz), and the remaining 2\% use other frequencies. The use of radio-frequencies to transmit data from the satellite to the ground has some drawbacks, like the crowding of used frequencies (potentially producing interference), or the poor directivity of the radio beam (enabling hacking of the data reception), not to say the poor data rate of UHF and VHF (about 1\,kb/s). In addition, considering the small available volume in the satellite, the UHF/VHF antennas have to be mechanically deployed, which is adding a risk to the success of the mission. Mitigating this risk by finding antenna schemes robust to deployment failure is an interesting track to look for. To alleviate this poor data rate, an optical transmission chain (light source -- beam launcher -- telescope -- photodiode) can be considered instead of a radiofrequency chain (transmitter -- TX antenna -- RX antenna -- receiver). An optical transmission chain has some advantages over a radio chain: it has a high directivity, making it difficult to intercept, there is no need to allocate a frequency, and there is a potential to have a high-speed data link (several hundred of Mb/s \cite{Janson2013}) | \label{sec:conclu} We are about to finish the phase 0 of the Nice$^3$ cubesat mission. During these first 6 months, students worked on the project and brought significant progress to our understanding of the context and difficulties of building a satellite. We have today a set of first boundaries of the satellite mission. This will enable us to progress further in the mission specifications in the coming months. Students working on this project come from many horizons. They may come from university masters, like MAUCA, but also from engineering schools, like Polytech Nice Sophia Antipolis, and of course from other formations, like optics BTS. The acknowledgements below list further the formations that follow closely the project. Making a satellite project with students is an exciting experience and we are preparing for the end of the phase 0 with enthusiasm. \subsection*{Acknowledgements} \emph{This project is supported by the C\^ote d'Azur University (UCA), the C\^ote d'Azur Observatory (OCA) and the Centre National des \'Etudes Spatiales (CNES). } \emph{The Centre Spatial Universitaire (CSU) of the C\^ote d'Azur University (UCA) offers students practical training as part of courses from the Astrophysics master (MAUCA) or the Geophysics master course (master 3G) of UCA, from the Mines Paristech and Polytech Nice-Sophia Antipolis engineering schools, as well as professional experience in research laboratories and institutes Lagrange, Geoazur, LEAT, I3S, CEMEF, Inphyni and INRIA.} \emph{The authors would like to thank the joint laboratory between Université\'e C\^ote d’Azur, CNRS and Orange, ``Centre de Recherche Mutualis\'e pour les Antennes" (CREMANT), for its support in the microstrip patch antenna characterization.} | 18 | 8 | 1808.09848 |
1808 | 1808.05582_arXiv.txt | Galactic globular clusters (GCs) show overwhelming photometric and spectroscopic evidence for the existence of multiple stellar populations. The question of whether or not there exists a GC that represents a true `simple stellar population' remains open. Here we focus on Ruprecht~106 (R106), a halo GC with [Fe/H]=$-1.5$ and [$\alpha$/Fe]$\simeq0$. A previous spectroscopic study found no sign of the Na-O anticorrelation among 9 of its brightest red giants, which led to the conclusion that R106 is a true simple stellar population GC. Here we present new \emph{Hubble Space Telescope} (HST) Wide Field Camera 3 photometry of R106 that, when combined with archival HST images spanning a 6-year baseline, allows us to create proper motion cleaned color-magnitude diagrams spanning the ultraviolet (F336W) to the near-infrared (F814W). These data allow us to construct the pseudo-color $\CUBI$ that is sensitive to the presence of light-element abundance spreads. We find no evidence of a split along the red giant branch (RGB) in the $\CUBI$ diagram but the width of the RGB ($\sigma_{\CUBI}=0.015$) is marginally broader than expected from artificial star tests ($\sigma_{\CUBI}=0.009$). The observed spread in $\CUBI$ is smaller than any other Galactic GC studied to date. Our results raise important questions about the r{\^o}le of formation environment and primordial chemical composition in the formation of multiple stellar populations in GCs. | \label{sec:intro} Observational evidence for the presence of multiple stellar populations (hereafter MPs) in Galactic globular clusters (GCs) is overwhelming \citep{gratton2012,piotto2015,milone2017}. This begs the question: What conditions are required in the formation and early evolution of a GC that allow it form and retain more than one stellar generation to the present day? Insight into this question may be found in searching for GCs that do \emph{not} harbor MPs, if any exist. In this study we focus on Ruprecht~106 (hereafter R106), which resides in the halo at a Galactocentric radius of $\sim18.5$ kpc \citep{harris1996}. R106 has [Fe/H]=$-1.5$ \citep[][hereafter V13]{brown1997,francois1997,villanova2013} and, remarkably, near solar-scaled abundance ratios among the $\alpha$-capture elements with [O/Fe]$\sim 0$ and [$\alpha$/Fe] $\sim 0$ \citep[][ V13]{brown1997}. The low [$\alpha$/Fe] ratio is unique among Galactic GCs with [Fe/H] $< -1$ \citep{pritzl2005}. In addition to atypical chemical abundances, R106 has a fairly low mass for a Galactic GC. Mass estimates range from $\log_{10}(M/M_{\odot})=4.77$ \citep{baumgardt2010} to 4.92 \citep{gnedin1997} placing it toward the low end of the Galactic GC mass spectrum. However, its mass is not anomalously low and \citet{milone2017} have analyzed a handful of GCs with masses lower than R106 that possess strong photometric evidence of MPs. The current mass of any GC is of course only a lower limit on its initial mass, and it is the latter that is likely the most important variable in the formation of MPs \citep[see the discussion of mass budget in][]{renzini2015}. R106 has previously been found to be 1-3 Gyr younger than other Galactic GCs with similar [Fe/H] \citep{buonanno1990,dacosta1992,buonanno1993,dotter2011}. The age, chemical composition, and location lead to the suggestion that R106 was accreted rather than formed \emph{in situ} \citep{forbes2010}. A recent photometric and spectroscopic study of another likely-accreted GC, IC~4499, found compelling evidence for the presence of MPs \citep{dalessandro2018}. Evidence that intermediate- and old-age GCs in the Small Magellanic Cloud (SMC) also host MPs \citep{dalessandro2016,niederhofer2017,hollyhead2018}. These results indicate that the host galaxy mass, at least down to the level of the SMC, is not a limiting factor in the formation of MPs in GCs. In the largest spectroscopic study of R106 to date V13 targeted 9 of its brightest red giants. They found [Na/Fe] $< -0.5$: lower than any other Galactic GC but consistent with Local Group dwarf galaxies (see their Figs.~4 and 6). The driving motivation for this paper is that V13 found no evidence of the canonical Na-O anticorrelation in R106. They wrote ``Ruprecht 106 is the first convincing example of a single-population GC (i.e., a true simple stellar population), although the sample is relatively small.'' These results from V13 raise the possibility that the peculiar abundance pattern found in R106 may have some role in the apparent lack of MPs. Photometry has come to play a major role in the study of MPs, particularly those filters that are sensitive to light-element variations. The medium-band Str{\"o}mgren pseudo-color $c=(u-v)-(v-b)$ is sensitive to NH variations and, therefore, traces light-element (CNO) variations in GC red giants \citep[see, e.g.,][and references therein]{yong2008,sbordone2011}. \citet{monelli2013} showed the efficacy of constructing a broadband version of $c$, $\CUBI$=(U$-$B)$-$(B$-$I), in the study of MPs. A variety of WFC3 filters have been used to study GCs with aims of untangling the MPs within, including combinations of $F275W$, $F336W$, $F438W$, and $F814W$ to construct the so-called `chromosome map' \citep[][and references therein]{milone2017}. \citet{larsen2014} applied the $F343N$ narrowband filter, which captures the NH feature around 3,300 $\AA$, to the study of GCs. $F343N$ has subsequently been used in combination with broadband filters, see $\S$2 of \citet{bastian2017} for a review of filter combinations used to study MPs. Inspired by the above we undertook an HST WFC3 observing program to study R106 in F336W (U) and F438W (B) filters, known to be sensitive to light-element abundance variations that are common in GCs \citep{marino2008}. Here we use an HST version of $\CUBI$ with F336W (U), F438W (B), and F814W (I). For convenience we refer to (F336W$-$F438W)$-$(F438W$-$F814W) simply as $\CUBI$. We use these data to test the assertion that R106 is a simple stellar population by searching for broadened sequence---or multiple sequences---in the HST $\CUBI$ color-magnitude diagram (CMD). | \label{sec:conclusions} Even before the V13 spectroscopic study R106 was an outlier in the Galactic GC population \citep{pritzl2005}. V13 identified a new dimension in which R106 is peculiar: an apparent absence of the Na-O anti-correlation that has become synonymous with GCs. In this study, we find no compelling evidence for a split RGB and only a marginal spread in $\CUBI$ among RGB stars. Taken together, these findings reinforce the claim by V13 that R106 is the best candidate for a simple stellar population GC. However, this can only be confirmed through further study of this enigmatic GC. | 18 | 8 | 1808.05582 |
1808 | 1808.00969_arXiv.txt | Wide-field optical surveys have begun to uncover large samples of fast ($t_{\rm rise} \lesssim 5$d), luminous ($M_{\rm peak} < -18$), blue transients. While commonly attributed to the breakout of a supernova shock into a dense wind, the great distances to the transients of this class found so far have hampered detailed investigation of their properties. We present photometry and spectroscopy from a comprehensive worldwide campaign to observe AT\,2018cow (ATLAS\,18qqn), the first fast-luminous optical transient to be found in real time at low redshift. Our first spectra ($<2$ days after discovery) are entirely featureless. A very broad absorption feature suggestive of near-relativistic velocities develops between $3-8$ days, then disappears. Broad emission features of H and He develop after $>10$ days. The spectrum remains extremely hot throughout its evolution, and the photospheric radius contracts with time (receding below $R<10^{14}$ cm after 1 month). This behaviour does not match that of any known supernova, although a relativistic jet within a fallback supernova could explain some of the observed features. Alternatively, the transient could originate from the disruption of a star by an intermediate-mass black hole, although this would require long-lasting emission of highly super-Eddington thermal radiation. In either case, AT\,2018cow suggests that the population of fast luminous transients represents a new class of astrophysical event. Intensive follow-up of this event in its late phases, and of any future events found at comparable distance, will be essential to better constrain their origins. | The development of sensitive, wide-area digital optical sky surveys has led to the discovery of populations of rare, luminous extragalactic transients that evolve on timescales of just a few days---much faster than typical supernovae, whose light curves are governed by the decay of \Nifs\ within a massive envelope and typically take weeks to months to fade. Many of these have been reasonably well-explained by known phenomena: shock-breakout flashes from supernovae \citep[e.g.,][]{Ofek+2010,Shivvers+2016,Arcavi+2017}, early emission from relativistic supernovae \citep{Whitesides+2017}, or the shockwave afterglows from gamma-ray bursts \citep{Cenko+2013,Cenko+2015,Stalder+2017,Bhalerao+2017}. Other objects are more mysterious, however, and still lack a convincing explanation or firm spectroscopic identification. In particular, populations of optical transients with luminosities comparable to or exceeding those of the most luminous core-collapse supernovae, but rise times of only a few days, have been reported by a variety of different surveys \citep{Arcavi+2016,Drout+2014,Tanaka+2016,Pursiainen+2018,Rest+2018}. Nearly all of these events (dubbed fast-evolving luminous transients by \citealt{Rest+2018}) were found at great distances ($z>0.1$) where they are difficult to study. Furthermore most were not recognized as unusual events in real time, preventing the acquisition of essential follow-up observations. The few spectra that are available tend to show only featureless blue continuua. Because of their origins in star-forming galaxies these transients are widely interpreted as supernovae, but strong constraints are lacking. Fortunately, our ability to find and identify fast transients continues to improve, and several surveys are now monitoring almost the entire sky at cadences of a few days or less. The Asteroid Terrestrial-impact Last Alert System (ATLAS; \citealt{Tonry+2018}) observes most of the visible Northern sky down to 19 mag every $\sim2$ nights. The Zwicky Transient Facility (ZTF; \citealt{ATEL11266}) observes a similar area to 20.5 mag every 3 nights, and a significant fraction of it at much higher cadence. ASAS-SN \citep{Shappee+2014} monitors both hemispheres nightly to $\sim$17 mag. With these capabilities, it is now possible to find and identify transients in (almost) real time over most of the night sky. In this paper, we present a detailed observational study of the first fast high-luminosity transient to be identified in the nearby universe in real time: AT\,2018cow, discovered by the ATLAS survey and independently detected by ZTF and ASAS-SN. We present our extensive, worldwide observational campaign in \S \ref{sec:observations}, focusing on observations at ultraviolet, optical, and near-infrared wavelengths (the multiwavelength view of this transient is presented by \citealt{Ho+2018}). We summarize the key properties of this event in \S \ref{sec:properties}, and illustrate the ways in which AT\,2018cow is distinct from any well-established class of transient in \S \ref{sec:comparisons}. In \S \ref{sec:interpretation} we consider two possible explanations for its origin: a jet-driven supernova erupting into a dense envelope of circumstellar matter, or alternatively the tidal disruption of a star around an intermediate-mass black hole located in a small galaxy's spiral arm. Both models have significant difficulties explaining the full suite of observations, and our observations suggest that the origins of fast luminous transients may be significantly more exotic and complex than previously assumed. We summarize our results and examine future directions in fast-transient research in \S \ref{sec:conclusions}. | \label{sec:conclusions} Prior to AT\,2018cow, fast high-luminosity transients were widely attributed to an extreme variant of the shock-breakout scenario that has already been widely appealed to in order to explain a variety of nearby supernovae. To our surprise, the first real-time detection of a nearby event belonging to this empirical class has only deepened the mystery surrounding these events. While the off-nuclear location within a star-forming region seems to imply the explosion of a star as a supernova, the actual observational properties---including high-velocity absorption in early spectra, a long-lived hot photosphere, a complete lack of narrow lines during the first week, and luminous X-ray through radio emission--- are all difficult to explain under any existing supernova model. If nothing else, any stellar explosion must involve a radically different progenitor structure and/or explosion mechanism compared to known SNe. In contrast, disruption by an intermediate-mass black hole provides an excellent description of the qualitative behaviour of the transient and its later-time spectra. However, the highly super-Eddington luminosity of the transient is a formidable challenge for IMBH TDE models, and it remains to be seen whether alternative explanations for the early heating (e.g. circularization of infalling material) provide an adequate explanation. Studies of fast optical transients are still in their infancy, and there is much more to learn both observationally and theoretically. While an event as close as AT\,2018cow may not be a regular occurrence, its sheer brightness suggests that others of a similar nature are likely to be observed in the near future at somewhat greater distances. Samples of the spatially-resolved galaxy environments, total energetics, and spectroscopic properties of such events are likely to shed light on their nature. | 18 | 8 | 1808.00969 |
1808 | 1808.10089_arXiv.txt | We present high-spatial-resolution ($\sim 0\farcs2$, or $\sim$3\,pc) CO(2--1) observations of the nearest young starburst dwarf galaxy, NGC\,5253, taken with the Atacama Large Millimeter/submillimeter Array. We have identified 118 molecular clouds with average values of 4.3\,pc in radius and 2.2\,\kms\, in velocity dispersion, which comprise the molecular cloud complexes observed previously with $\sim$100\,pc resolution. We derive for the first time in this galaxy the $I{\rm (CO)}$--$N$(H$_2$) conversion factor, $X$ = $4.1^{+5.9}_{-2.4}\times10^{20}$\,cm$^{-2}$(K\,\kms)$^{-1}$, based on the virial method. The line-width and mass-to-size relations of the resolved molecular clouds present an offset on average toward higher line-widths and masses with respect to quiescent regions in other nearby spiral galaxies and our Galaxy. The offset in the scaling relation reaches its maximum in regions close to the central starburst, where velocity dispersions are $\sim$ 0.5 dex higher and gas mass surface densities are as high as $\Sigma_{\rm H_2}$ = 10$^3$\,\Msol\,pc $^{-2}$. These central clouds are gravitationally bound despite the high internal pressure. A spatial comparison with star clusters found in the literature enables us to identify six clouds that are associated with young star clusters. Furthermore, the star formation efficiencies (SFEs) of some of these clouds exceed those found in star-cluster-forming clouds within our Galaxy. We conclude that once a super star cluster is formed, the parent molecular clouds are rapidly dispersed by the destructive stellar feedback, which results in such a high SFE in the central starburst of NGC\,5253. | Our knowledge of the properties of molecular clouds is steadily increasing as more wide-field and high-resolution observations are made available. Large number statistics of the properties of resolved molecular clouds identified using a similar method now exist for different galaxies \citep[e.g.][]{2001ApJ...551..852H,2009ApJ...699.1092H,2011ApJS..197...16W,2008ApJ...686..948B,2013ApJ...772..107D,2015ApJ...801...25L,2014ApJ...784....3C,2012ApJ...761...37M,2008ApJS..178...56F}. These studies showed that while the cloud properties and scaling relations are compatible within the Galaxy and external galaxies across a wide range of environments \citep{2008ApJ...686..948B}, those in starburst (SB) regions or in the Galactic Center substantially differ \citep{2001ApJ...562..348O,2015ApJ...801...25L}. For instance, \citet{2015ApJ...801...25L} showed that the clouds in the prototypical SB galaxy NGC\,253 have large line-widths, high surface gas densities, and short free-fall times, which suggest that its efficient star formation is enabled by the surrounding dense gas reservoir. An important environment where the different conditions of the molecular gas around SBs can be proved is blue compact dwarf galaxies (BCDs). BCDs are defined as faint and compact galaxies, and most of them have compact SB cores with high star formation rates \citep{2004AJ....128.2170H}. These compact SBs comprise a relatively small number of star clusters compared to larger SB galaxies such as NGC\,253, and thus the effect of the SB on the surroundings is less complex. Studying gas properties in smaller-scale SBs in BCDs allows us to better understand larger SB systems. Other dwarf galaxies have in general low gas densities and low star formation efficiencies \citep[SFEs;][]{2010AJ....140.1194B}, which may imply that some mechanism exists in BCDs to change the cloud properties of a typical dwarf galaxy to trigger an SB such as in NGC\,253 (although at smaller scale). Several external triggering mechanisms to explain the bursts in BCDs have been proposed and explored through numerical simulations, including merger or interaction between gas-rich dwarf galaxies \citep{2008MNRAS.388L..10B} and external gas infall \citep{2014MNRAS.442.1830V}. These external triggering events can increase the turbulence of the gas and transfer gas from the outer regions to the center, which is prone to be dense enough for gravitational collapse, and then an SB event follows. These SBs may blow out the gas from the central region, although the outflow gas may fall back and become the seeds of the next SB \citep{2014MNRAS.442.1830V}. Recent Atacama Large Millimeter/submillimeter Array (ALMA) observations toward the BCD galaxy II Zw 40 BCD have shown that the molecular gas is clumpy and is characterized by larger line-widths and more compact sizes as well as higher molecular gas surface densities with respect to other quiescent molecular clouds in other objects \citep{2016ApJ...828...50K}. However, still only a small number of observations toward BCDs exist to study the cloud properties under these peculiar environments. In this paper, we study the nearby\footnote{There are still uncertainties in the distance measurement of $2.8$--$5.2$\,Mpc \citep[e.g.][]{1994ApJ...421L..87B, 2004ApJ...608...42S,1996ApJ...473...88R}. We adopt a distance to NGC\,5253 of 3.15\,Mpc \citep{2001ApJ...553...47F,2007AJ....134.1799D} throughout this paper.} SB galaxy NGC\,5253, which is argued to be a BCD that has experienced \ion{H}{1} gas infall that has triggered its central SB \citep{2012MNRAS.419.1051L}. This galaxy is known to have a relatively low metallicity \citep[$12+\log({\rm O/H})\sim 8.18$--8.30;][]{2007ApJ...656..168L,1997ApJ...477..679K,1999ApJ...514..544K}. NGC\,5253 hosts a young nuclear SB, and because of its proximity it has been the object of many studies. This galaxy has a disk-like morphology extending SE-NW as seen in optical images (see Fig.\,\ref{almafov}, and also Table\,\ref{opt} for a summary of the main galaxy properties) and has several young clusters with ages ranging from 1 to 15\,Myr are distributed within the central 300\,pc \citep{2015ApJ...811...75C}. Older clusters can also be found over the galaxy and their ages are typically 1\,Gyr or more \citep{2013MNRAS.431.2917D}. The H$\alpha$ image shows multiple filamentary and bubble-like structures extending perpendicular to its optical major axis \citep{2004AJ....127.1405C}. The multiage nature of the clusters and the existence of bubbles indicate that NGC\,5253 has experienced a few bursts in the past \citep{1998MNRAS.295...43D,2010ApJ...721..297M}. One of the peculiarities of this galaxy is that its central SB is found to be powered by two very young (1\,Myr) and massive ($\sim3\times10^6\Msol$) compact stellar clusters \citep[so-called super star clusters; hereafter SSC;][]{2015ApJ...811...75C}, separated by 6\,pc \citep{2004ApJ...612..222A,2004A&A...415..509V}. The SSCs are deeply embedded radio compact (1--2 pc size) \ion{H}{2} regions illuminated by an equivalent of 4000 O stars \citep{2000ApJ...532L.109T}. The submillimeter hydrogen recombination line H30$\alpha$ is detected toward the location of the obscured SSCs, and \citet{2017MNRAS.472.1239B} suggest that this region is not only the predominant site of star formation but also potentially accounts for about $\sim$90\% of the total in the nuclear region. As a consequence, an unusually high SFE, 37\,\%--75\,\% on a 100 pc scale, is estimated for the central SB region \citep{2015Natur.519..331T,2015PASJ...67L...1M,2002AJ....124..877M}. Another peculiarity of the galaxy is the extended atomic gas (\ion{H}{1}), which does not follow galactic rotation even though the galaxy seems to have a stellar disk \citep{2008AJ....135..527K,2012MNRAS.419.1051L}. Instead, the distribution of \ion{H}{1} is elongated along the optical minor axis (so-called ``\ion{H}{1} plume''), which is almost at the same direction as the dust lane. Previous \ion{H}{1} observations revealed that there is a velocity gradient along the dust lane, which was interpreted as cold \ion{H}{1} falling into the galaxy, rather than gas rotation along the minor axis like a polar-ring galaxy \citep{2012MNRAS.419.1051L}. The velocity structure of CO emission also show a similar trend and supports this scenario \citep{2002AJ....124..877M,2015PASJ...67L...1M}. The origin of the burst in NGC\,5253 is believed to be due to gas infall, which may have been caused by the past interaction with the spiral galaxy M\,83 or other subgroup members \citep{1999AJ....118..797C,2008AJ....135..527K}, and probably started more than 100\,Myr ago \citep{2012MNRAS.419.1051L}. The main goal of this paper is to study the gas distribution and cloud properties in NGC\,5253 with parsec-scale resolution to be able to resolve clouds close to the SB and to gain insight into the origin of the triggering of an SB event in BCD galaxies. The outline of this paper is as follows. The observations and data reduction are summarized in Section\,\ref{obs}. In Section\,\ref{result} we present the overall molecular gas distribution, the identification of cloud properties, the scaling relations, and the comparison with other galaxies. In Section\,\ref{discussion}.1, we cross-match the molecular clouds with a compilation of young star clusters from the literature, and we compare differences between the properties of the star forming clouds. In Section\,\ref{discussion}.2, the overall molecular gas distribution and kinematics are studied and compared with numerical simulations to shed light on the origin of the SB in NGC\,5253. In a forthcoming paper, we will report the star formation activities around the central SB using H30$\alpha$ emission as well as 230\,GHz continuum data (Paper\,{\sc ii}, Miura, R. E. et al. 2018, in preparation), which were obtained with the same observing setup. A comparison of the star formation rate (SFR) derived from H30$\alpha$ emission with those from other tracers have already appeared in a separate paper \citep{2017MNRAS.472.1239B}. A spatial comparison between the 230\,GHz continuum and the CO molecular gas distribution for each cloud will be presented in a separate paper (Paper\,{\sc iii}, Miura, R. E. et al. 2018, in preparation). | \label{discussion} \subsection{Association with SSCs and Feedback} In this section, we discuss the star formation activities and the properties of clouds that appear associated with young clusters. To do this, we compare the spatial association of the identified CO clouds with previously identified star clusters \citep{2015ApJ...811...75C,2013MNRAS.431.2917D,2004ApJ...603..503H}. \citet{2015ApJ...811...75C} have revised age and mass estimations in \citet{1997AJ....114.1834C} toward 11 star clusters using multiple-wavelength data and a more accurate spectral energy distribution model fitting. We also used the catalog in \citet{2013MNRAS.431.2917D} to cover other star clusters. This includes 181 identified clusters in NGC\,5253 using UV/optical/NIR {\it HST} data, covering an area of about $500\times430$\,pc$^2$. For sources that do not have any measurement in \citet{2013MNRAS.431.2917D} nor \citet{2015ApJ...811...75C}, we use the star cluster catalog by \citet{2004ApJ...603..503H}, obtained from optical broadband {\it HST} data. The uncertainty of absolute astrometry of the {\it HST} image is $0\farcs1-0\farcs3$ and the two bright clusters in the {\it HST} image have a $0\farcs2$ offset from the two 1.3 cm radio components in \citet{2000ApJ...532L.109T} \citep{2015ApJ...811...75C}. It is likely that the two bright clusters correspond to the two radio components and we correct the coordinates of all star clusters in the catalog accordingly. After correcting the coordinates of the clusters, we compare their location with that of our identified clouds. The {\it HST} images also cover most of the observed field in CO (N5253-D, N5253-F and the western half of N5253-C). We search for clusters whose central position spatially overlaps with the extent (at the level of $3\,\sigma$) of any of the molecular clouds, taking into account the uncertainty of astrometry. Among the 118 clouds, we found that six clouds appear spatially associated with 10 star clusters. Figure\,\ref{starcloud1} shows the CO integrated intensity images of the six clouds overlaid on their corresponding {\it HST} composite image. The identifications of these clouds, their associated star clusters, and their properties are listed in Table \ref{sfe}. The majority of the associated clusters are young, less than 10\,Myr, which is the maximum age acceptable to infer that they may have originated from the clouds, taking into account the lifetime of a giant molecular cloud \citep{2009ApJS..184....1K,2012ApJ...761...37M,2011ApJ...729..133M}. Specifically, in the case of cloud\,8, we found that its velocity and location is very close to those of the H30$\alpha$ peak or the central cluster \citep{2017MNRAS.472.1239B,2000ApJ...532L.109T,2004ApJ...610..201S}, which together with the young age of the cluster likely suggest a physical relation between the cloud and the star clusters. Note that cluster G-9, which is paired with cloud\,60, has a large uncertainty in age, which can be up to 100\,Myr depending on the stellar synthesized model that is used. If G-9 is that old, it probably did not originate from that molecular cloud. We estimated the SFEs for these clouds. SFE is defined as $M_{\ast}/(M_{\rm gas}+M_{\ast})$, where $M_{\ast}$ is the mass of a star cluster, and $M_{\rm gas}$ is the mass of the cloud. The derived SFE is also listed in Table \ref{sfe}. The SFEs we obtain for NGC\,5253 range from 5\,\% to 64\,\%. Compared with cluster-forming clouds in our Galaxy, in which the SFE ranges from 10\.\% to 30\,\% \citep{2003ARA&A..41...57L}, two among all, cloud\,5 and cloud\,8, show high SFE. Note that we exclude the case of cloud\,60 here because of the large uncertainty in its age, as explained above. Cloud\,8 has clumpy structure, and if one of its clumps is associated with one of the two clusters (cluster\,11), the SFE would be 80\,\%. In the SSC-forming clouds in an SB system, whether gas collapses or disperses is a competitive process between the inward forces such as gravity and outward forces such as that caused by a jet, expansion of a \ion{H}{2} region, stellar wind, and radiation pressure. The most dominant force in the SSC-forming clouds is radiation force ($F_{\rm rad}$) among the outward forces and gravity ($F_{\rm grav}$) for the inward forces \citep{2010ApJ...709..191M}. Following \citet{2011ApJ...729..133M}, we estimate $F_{\rm rad}$ for five clouds using $F_{\rm rad}=L/c$\,dynes, where $L$ is the luminosity in erg\,s$^{-1}$, and $c$ is the light speed in cm\,s$^{-1}$. We use $L=\xi Q$, where $Q$ is the ionizing photon rate and $\xi=8\times10^{-11}$\,ergs \citep{2010ApJ...709..191M}, and the conversion between $M_{\ast}$ and $Q$ is $M_{\ast}=1.6\times10^4 (Q/10^{51}\,{\rm s}^{-1})\,\Msol$ \citep{2011ApJ...729..133M}. In this conversion between $M_{\ast}$ and $Q$, \citet{2011ApJ...729..133M} assumed a minimum stellar mass of $0.1\,M_{\odot}$, a maximum stellar mass of $120\,M_{\odot}$, a slope of the initial mass function of $-1.35$ and constrained to young star clusters ($\sim4$\,Myr, the averaged main-sequence lifetime of massive stars, which are mainly responsible for emitting ionizing photons). For the force of gravity, we use $F_{\rm grav}=GM_{\rm gas}^2R^{-2}$, where $R$ is the cloud radius \citep{2011ApJ...729..133M}. Table\,\ref{sfe} shows the estimated forces for each cloud and star cluster. Since the conversion from $M_{\ast}$ to $Q$ is limited to young stellar age populations, the $Q$ might be overestimated in the case of an older stellar cluster \citep[beyond $\sim5$\,Myr;][]{2011ApJ...729..133M}, and thus for clouds\,5, 53 and 59, $F_{\rm rad}$ would be an upper limit. Note that we calculate $F_{\rm grav}$ both using $M_{\rm gas}$ only and $M_{\rm gas}+M_{\ast}$ (as shown in the parentheses in the table). In the latter, it is assumed that the star clusters are embedded inside the molecular cloud. In the case of cloud\,8, this scenario is plausible since its associated star clusters are still very young and located at one of the CO emission peaks. Besides, the virial mass of cloud\,8 is much larger when compared to the gas mass only, but it is comparable to the summation of the stellar clusters and the gas masses, unlike in other clouds. Note that for cloud\,5 and cloud\,53, since the ages of the associated clusters G-125, G-127, and G-86 are not available, we used values from only clusters H-2 and C-4, respectively, and thus $M_{\ast}$ and consequently $F_{\rm rad}$ are a lower limit. If only part of cloud\,8 is associated with the star cluster, the radiation force would be more than five times larger than the gravitational force. Clouds\,5 and 8 (if only part of cloud\,8 is associated with the star cluster) are unique clouds in the sense that the radiation force exceeds the inward force of gravity, while the rest of the clouds are governed by gravity. This means that these two clouds will likely dissipate the parental clouds, due to stellar feedback after the SSC is formed, while the other clouds (e.g. clouds\,28, 31, and 53) have still have the possibility to form stars in the gravity-dominant phase. Cloud\,5 is associated with a cluster with an age of 10\,Myr, while cloud\,8 is associated with a younger cluster (1\,Myr), which is the most massive cluster in the galaxy. Taking into account that the clouds in the central SB region have high gas densities, it is reasonable to expect that a more massive cluster will tend to be formed, but then its destructive stellar feedback will remove the gas quickly. Note that for an $X_{\rm CO}$ factor 2.6 times larger than the value applied here, all clouds would be in the gravity-dominant phase. In this case, the parental molecular clouds would have enough content to form more stars rather than being dissipated. \begin{figure*} \includegraphics[width=.45\textwidth, trim={50 30 150 40}, clip]{co_mom0_hst_idx5.pdf} \includegraphics[width=.45\textwidth, trim={50 30 150 40}, clip]{co_mom0_hst_idx8.pdf}\\ \includegraphics[width=.45\textwidth, trim={50 30 150 40}, clip]{co_mom0_hst_idx28.pdf} \includegraphics[width=.45\textwidth, trim={50 30 150 40}, clip]{co_mom0_hst_idx31.pdf}\\ \includegraphics[width=.45\textwidth, trim={50 15 150 40}, clip]{co_mom0_hst_idx54.pdf} \includegraphics[width=.45\textwidth, trim={50 15 150 40}, clip]{co_mom0_hst_idx60.pdf} \caption{{\footnotesize CO integrated intensity map (contours) for individual clouds\,5, 8, 28, 31, 53 and 59 over the {\it HST} composite images. The contour levels are 3, 5, 10, 15, 20, 25, 30, 35 and 40\,$\sigma$, where $\sigma$=0.01\,Jy\,\kms. The RGB {\it HST} composite images are composed of F300W in blue, F658N or H$\alpha$ in green, and F814W in red. The star symbols indicate the star clusters which are associated with the molecular clouds and each color (white, yellow, orange) indicates the different star cluster catalog (\citealt{2013MNRAS.431.2917D}, \citealt{2015ApJ...811...75C}, \citealt{2004ApJ...603..503H}, respectively). The cross symbols are star clusters in our compilation which are not associated with any of the identified clouds. All panels are $2\arcsec$ $\times$ $2\arcsec$. }\label{starcloud1}} \end{figure*} \subsection{Origin of the Starburst: Comparison of the Global Gas Distribution with Simulations} NGC\,5253 is argued to have experienced external gas infall because the kinematics of the atomic and molecular gas cannot be explained by galaxy rotation. The gas infall has likely triggered the powerful SB \citep{2002AJ....124..877M,2015PASJ...67L...1M,2012MNRAS.419.1051L,2014A&A...566A..71L}. \citet{2014MNRAS.442.1830V} showed that external gas infall can trigger a starburst in dwarf galaxies. Those simulations show that the SBs can happen several times within 1\,Gyr depending on the initial conditions. Once external gas encounters the dwarf galaxy, the kinematics/distribution of the gaseous component in the host galaxy is strongly disturbed. The gas loses momentum and is driven inward to form dense clouds and consequently cause the SB. The new stars blow out the surrounding gas to large radii, but it falls back again into the galaxy to be the fuel to form the next SB event. The simulations also predict that the gaseous and stellar components of the galaxy become more compact and result in higher densities toward the center. As the burst settles down, the rotation curve becomes steep, as observed in several BCDs \citep{2001AJ....122..121V,2014MNRAS.441..452K}. From the CO observations, we do not find any clear sign of simple disk rotation, in agreement with previous \ion{H}{1} observations \citep{2012MNRAS.419.1051L}. The distribution of CO emission is very irregular, and the direction of the molecular gas structure does not match that of the H$\alpha$ outflow. This support the scenario that the gas might be mostly infalling to the center and the SB event is most likely recent. The central SB region has high surface gas densities (Section\,\ref{density}), which supports simulation results that show that before and during the burst, gas flows to the center and the density rises significantly compared to larger radii. Taking into account that several bubbles caused by expansion of \ion{H}{2} regions and supernovae are found within the galaxy \citep[][]{1998AJ....116.1212T,1997AJ....114.1834C,1999MNRAS.306...43S} and \ion{H}{1} emission follows (although the resolution is much coarser) the H$\alpha$ bubbles \citep{2012MNRAS.419.1051L}, part of these molecular gas components may also be gas that is blown out to large radii and falling back to the galaxy. First, we calculate the gas inflow rate through the dust lane (i.e. N5253-C). If the gas of N5253-C were infalling, as suggested in previous studies, and using a mass of $M_{\rm gas}$ = $7.6\times10^6\,\Msol$ (equivalent to the gas mass derived from the emission within the ellipse of N5235-C in Figure\,\ref{mom}), the velocity difference of 30\,\kms\ between N5253-C and the central SB as the radial inflow speed, and the distance of N5253-C to the central SB, $\sim$100\,pc, we obtained a gas inflow rate via the dust lane of $\dot{M}_{\rm flow}$ $\simeq$2\,\Msol yr$^{-1}$. This was obtained by using equation $\dot{M}_{\rm flow}$[\Msol\,yr$^{-1}$]= $1.0\times10^{-6} M_{\rm gas} V_{\rm grad, corr}$. Here, $V_{\rm grad,corr}$ is the velocity gradient in units of $\kms\,$pc$^{-1}$, and given by $ V_{\rm grad,corr} =V_{\rm grad}/\tan i$, where $i$ is an inclination of the flow along the line of sight, and we assume $i=45$\arcdeg. A caveat with the infalling scenario is that in the data presented here, we have not found molecular gas features in the PVDs clearly connecting the dust lane (N5253-C) and the central SB (N5253-D). We show in Figure\,\ref{pvcut} the expected P--V curves for a disk-like model assuming a maximum rotation velocity of 30\,\kms\ (Appendix\,\ref{a1}). Although the kinematics of some of the structures may be explained by galaxy rotation, in general this is not true, and the difference is most noticeable for components at PAs 70$^{\arcdeg}$, 100$^{\arcdeg}$ and 175$^{\arcdeg}$. We have found from the PVDs that multiple components connect to the central massive molecular cloud (cloud\,8), showing a smooth velocity gradient (Section\,\ref{filament}). We measure the total gas mass along these structures, $M_{\rm gas}$, and the velocity gradient, $V_{\rm grad}$, in order to estimate the gas flow rate as done before for the dust lane. The obtained gas flow rate is about 0.06--0.07\,\Msol yr$^{-1}$ each (Table\,\ref{outflow}). If these are all inflowing, then the accumulated gas would be $1.9\times10^5$\,\Msol\, in 1\,Myr. To form a molecular cloud like cloud\,8 only from these gas inflows, it would take $\sim10$\,Myr, which is an equivalent timescale for the formation of a molecular cloud of $<20$\,Myr \citep{1994LNP...439...13L}. Note that the calculated inflow rate is lower than the previously estimated value in \citet{2015PASJ...67L...1M}, because it was assumed that the whole molecular cloud complex N5253-D contributed to the inflow. The SFR of NGC\,5253 in the literature is in the range of 0.1--0.4\,\Msol yr$^{-1}$ \citep{2000ApJ...532L.109T,2002AJ....124..877M,2004AJ....127.1405C,2012MNRAS.419.1051L,2015ApJ...811...75C,2017MNRAS.472.1239B}. Even taking into account that other possible mass losses such as mass ejection in the form of hot gas \citep[$\sim$0.2\,\Msol yr$^{-1}$;][]{2004MNRAS.351....1S} and ionized gas outflow \citep[(0.4--2.6)\,$\times 10^{-3}$\,\Msol yr$^{-1}$;][]{2015PASJ...67L...1M}, the gas inflow rate is much larger than the total amount of outflow rate and SFR, which indicates that enough molecular gas is being fueled to maintain star formation. That there is enough amount of gas being fueled to the center may support \citet{2015ApJ...811...75C}, who suggested the continuous star formation at least over 15\,Myr in the region within the central 300\,pc of NGC\,5253. The multiple components connecting to the central massive molecular cloud, which show a smooth velocity gradient and deviate from galaxy rotation, might be due to gas that has been blown out once and rain down on the galaxy. Even if these small components were all outflows rather than inflows, a positive net inflow rate is likely preserved, because in this case the molecular outflow rate would only be 0.2\,\Msol yr$^{-1}$ (i.e. the summation of the flow rates via all three identified structures in the PVDs). | 18 | 8 | 1808.10089 |
1808 | 1808.09895_arXiv.txt | In secular evolution of finite dense star clusters, the statistical acceleration of stars is the most essential non-collective relaxation process whose effect can be mathematically modeled by Kandrup's generalized-Landau (g-Landau) kinetic equation for distribution function of stars. Understanding of the g-Landau equation is of significance in finite system since only the equation can correctly define the total- energy and - number of stars in phase spaces in case the effect of gravitational polarization is neglected. The present paper shows a kinetic formulation of an orbit-averaged generalized Landau equation in action-angle spaces beginning with BBGKY hierarchy and shows the conservation laws, $H$-theorem and anti-normalization condition. Furthermore, the orbit-averaged g-Landau equation is rewritten for anisotropic spherical system. It is shown that the statistical acceleration can be replaced by typical acceleration of star for the relaxation process in secular evolution of any anisotropic spherical systems. The derivation of the equations is made by generalizing the formulation for the inhomogeneous Landau equation done in (Polyachenko 1982) to the g-Landau equation. | 18 | 8 | 1808.09895 |
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1808 | 1808.07136_arXiv.txt | Turbulence, magnetic reconnection, and shocks can be present in explosively unstable plasmas, forming a new electromagnetic environment, which we call here \textit{turbulent reconnection}, and where spontaneous formation of current sheets takes place. We will show that the heating and the acceleration of particles is the result of the synergy of stochastic (second order Fermi) and systematic (first order Fermi) acceleration inside fully developed turbulence. The solar atmosphere is magnetically coupled to a turbulent driver (the convection zone), therefore the appearance of turbulent reconnection in the solar atmosphere is externally driven. Turbulent reconnection, once it is established in the solar corona, drives the coronal heating and particle acceleration. | } In the 80's, the link between the spontaneous formation of current sheets inside fully developed turbulence and the evolution of unstable current sheets to fully developed turbulence has been established with the use of 2D numerical simulations of the MHD equations \cite{Matthaeus86, Biskamp89}. Several recent reviews discuss the way turbulence can become the host of reconnecting current sheets and how reconnecting current sheets can drive turbulence \cite{Matthaeus11, Cargill12, Lazarian12, Karimabadi13a, Karibabadi2013c}. The link between shocks and large amplitude magnetic disturbances and current sheets has also been analyzed \cite{Karimabadi2014}. We will use the term ``turbulent reconnection" to denote an environment where large scale magnetic discontinuities of size $\delta\!B$, with $\delta\! B/B >1$, coexist with randomly distributed Unstable Current Sheets (UCS) \cite{Matthaeus86, Lazarian99}. The importance of turbulent reconnection in many space and astrophysical systems has been discussed in detail in recent reviews \cite{Lazarian15, Matthaeus15}. In the solar atmosphere, turbulence is externally driven by the convection zone, and the spontaneous formation of a turbulent reconnecting environment has been analyzed in several articles \cite{Parker83, Parker88, Galsgaard96, Galsgaard97a, Gasgaard97b, Einaudi94a, Georgoulis96}. This review is divided into three sections. In the first section we pose the question: How the three well known non linear MHD structures appearing in many astrophysical and laboratory plasmas, i.e.\ Turbulence, Current Sheet(s), and Shocks, can lead asymptotically to turbulent reconnection. We will outline briefly the current literature, which addresses this question with the use of MHD, Hybrid and Particle In Cell (PIC) simulations. In section 3 we analyze the question: How the solar convection zone drives fully developed turbulence in the solar atmosphere. Finally, in section 4, we attempt to reply to the question: How the plasma is heated and the high energy particles are accelerated by turbulent reconnection. In section 5, we summarize our main points. | In this review we have stressed the following points: \begin{itemize} \item We present evidence from numerical simulations that supports the fact that all well known nonlinear structures (e.g.\ turbulence, current sheets, and shocks) asymptotically lead to a new nonlinear state, which we call \textit{turbulent reconnection}. \item Turbulent reconnection is a non linear state of the plasma, where large scale magnetic disturbances and UCSs co-exist. \item Based on numerical simulations that are still far from realistic, we suggest that the Solar convection zone may generate and drive the Solar corona into a turbulent reconnection state. The emergence of new magnetic flux, the random stressing of emerged magnetic flux by turbulent photospheric flows, and large scale instabilities of the emerged magnetic field topologies drive a variety of global and/or localized volumes into the state of turbulent reconnection. \item The synergy of the large scale magnetic disturbances and the UCS in turbulent reconnection provides the heating and the acceleration of high energy particles (electrons and ions). We claim that during turbulent reconnection the two well known Fermi mechanisms (stochastic and systematic) co-exist, forming a new very efficient mechanism for the energization of the plasma. \item The stochastic interaction of the particles with the large amplitude magnetic fluctuations is responsible for the heating and the synergy of stochastic and systematic acceleration for the formation of the high energy tail. \item The key elements for the efficient heating and acceleration of particles are (1) the strength of the magnetic field in the energy release volume, (2) the mean free path $\lambda$ the particles travel between scatterers, (3) the size of the energy release volume. \end{itemize} The attempts made so far to analyse the solar corona using simple monolithic magnetic topologies, i.e.\ a single loop, a single current sheet or a shock, fail to grasp the importance of the turbulent state of the solar corona and its consequences, as discussed here. We hope that the Parker Solar Probe will capture the dynamics of the fully developed turbulence in the Solar Corona and let us discover the way it is coupled to the turbulent solar wind. The results reported here can be applied to many astrophysical, space or laboratory plasmas, whenever the state of turbulent reconnection is established. | 18 | 8 | 1808.07136 |
1808 | 1808.02525_arXiv.txt | We show that the stellar specific angular momentum $j_{\star}$, mass $M_{\star}$, and bulge fraction $\beta_{\star}$ of normal galaxies of all morphological types are consistent with a simple model based on a linear superposition of independent disks and bulges. In this model, disks and bulges follow scaling relations of the form $j_{\star{\rm d}} \propto M_{\star{\rm d}}^{\alpha}$ and $j_{\star{\rm b}} \propto M_{\star{\rm b}}^{\alpha}$ with $\alpha = 0.67 \pm 0.07$ but offset from each other by a factor of $8 \pm 2$ over the mass range $8.9 \leq \log (M_{\star}/M_{\odot}) \leq 11.8$. Separate fits for disks and bulges alone give $\alpha = 0.58 \pm 0.10$ and $\alpha = 0.83 \pm 0.16$, respectively. This model correctly predicts that galaxies follow a curved 2D surface in the 3D space of $\log j_{\star}$, $\log M_{\star}$, and $\beta_\star$. We find no statistically significant indication that galaxies with classical and pseudo bulges follow different relations in this space, although some differences are permitted within the observed scatter and the inherent uncertainties in decomposing galaxies into disks and bulges. As a byproduct of this analysis, we show that the $j_{\star}$--$M_{\star}$ scaling relations for disk-dominated galaxies from several previous studies are in excellent agreement with each other. In addition, we resolve some conflicting claims about the $\beta_\star$-dependence of the $j_{\star}$--$M_{\star}$ scaling relations. The results presented here reinforce and extend our earlier suggestion that the distribution of galaxies with different $\beta_{\star}$ in the $j_{\star}$--$M_{\star}$ diagram constitutes an objective, physically motivated alternative to subjective classification schemes such as the Hubble sequence. | Specific angular momentum ($j = J/M$) and mass ($M$) are two of the most basic properties of galaxies. Together, they largely determine another basic property---characteristic size (such as half-mass radius $R_{\rm h}$)---especially for disk-dominated galaxies. Thus, the correlation between $j$ and $M$ constitutes one of the most fundamental scaling relations for galaxies, as important as those between rotation velocity, velocity dispersion, characteristic size, and mass. We have studied the galactic $j$--$M$ relation from both observational and theoretical perspectives (Fall 1983; Romanowsky \& Fall 2012; Fall \& Romanowsky 2013, hereafter Papers 0, 1, and 2). The present paper is a continuation of this series. In the following, when relevant, we distinguish between the stellar, baryonic, and halo parts of galaxies with the subscripts $\star$, bary, and halo, and between their disk and bulge components with the subscripts d and b. We have found that both disk-dominated galaxies and bulge-dominated galaxies (mainly ellipticals) obey power-law scaling relations of the form $j_{\star} \propto M_{\star}^{\alpha}$ with essentially the same exponent, $\alpha = 0.6 \pm 0.1$, and normalizations that differ by a factor of $\sim 5$. Our results are based on a sample in which most galaxies have classical bulges (genuine spheroids) rather than pseudo bulges (disk-like structures). In a plot of $\log j_{\star}$ against $\log M_{\star}$, galaxies of different morphological type and bulge fraction $\beta_{\star} \equiv M_{{\star}{\rm b}} / ( M_{{\star}{\rm d}} + M_{{\star}{\rm b}} )$ follow nearly parallel relations, filling the region between the sequences of disk-dominated and bulge-dominated galaxies. Based on this finding, we have proposed that the distribution of galaxies with different $\beta_\star$ in the $j_{\star}$--$M_{\star}$ diagram constitutes an objective, physically motivated alternative to subjective classification schemes such as the Hubble sequence. The parallel sequences of galaxies of different bulge fraction in the $j_{\star}$--$M_{\star}$ diagram suggest a picture in which galactic disks and spheroids are essentially independent objects, formed by distinct physical processes. Disks likely formed relatively quiescently by diffuse gas settling within dark-matter halos, while spheroids likely formed more violently by colliding streams and clumps of cold gas and by merging of smaller galaxies. Disk-dominated galaxies are those in which major mergers played little or no role, while spheroid-dominated galaxies either never acquired a substantial disk or else acquired one and later lost it by stripping or merging. In this picture, most normal galaxies may be regarded, in a first approximation, as a linear superposition of a flat disk and a round spheroid, each of which lies along the corresponding $j_{\star}$--$M_{\star}$ sequence. The primary purpose of this paper is to make a quantitative test of this picture. The observed $j_{\star}$--$M_{\star}$ scaling relations also link well with galaxy-formation theory. The galactic halos that form by hierarchical clustering in a dark-matter dominated universe (such as $\Lambda$CDM) obey the scaling relation $j_{\rm halo} \propto M_{\rm halo}^{\alpha}$ with $\alpha = 2/3$, an exponent remarkably similar to that for the stellar parts of galaxies. The halo scaling relation follows directly from the fact that the spin parameter $\lambda_{\rm halo}$ and mean internal density $\bar\rho_{\rm halo}$ are independent of $M_{\rm halo}$. By comparing the $j_{\star}$--$M_{\star}$ and $j_{\rm halo}$--$M_{\rm halo}$ relations, mediated by an $M_{\star}$--$M_{\rm halo}$ relation, we found that galactic disks have a fraction of specific angular momentum relative to their surrounding halos of $f_j \equiv j / j_{\rm halo} \approx 0.8$, while galactic spheroids have a fraction $f_j \approx 0.15$. The first of these agrees well with the postulated value $f_j \approx 1$ in simple analytical models of galactic disk formation (Fall \& Efstathiou 1980; Dalcanton et al.\ 1997; Mo et al.\ 1998). The $j$--$M$ scaling relations have been the focus of further observational study, both for low-redshift galaxies (Obreschkow \& Glazebrook 2014; Cortese et al.\ 2016; Butler et al.\ 2017; Chowdhury \& Chengalur 2017; Elson 2017; Kurapati et al.\ 2018; Lapi et al.\ 2018b; Posti et al.\ 2018a; Rizzo et al.\ 2018; Sweet et al.\ 2018) and for high-redshift galaxies (Burkert et al.\ 2016; Contini et al.\ 2016; Harrison et al.\ 2017; Shi et al.\ 2017; Swinbank et al.\ 2017; Tadaki et al.\ 2017; Alcorn et al.\ 2018). Several recent studies have examined the relation between galaxy sizes and halo sizes, a corollary of the $j$--$M$ relation (Kravtsov 2013; Kawamata et al.\ 2015; Shibuya et al.\ 2015; Huang et al.\ 2017; Kawamata et al.\ 2018; Okamura et al.\ 2018). The $j$--$M$ scaling relations have been a benchmark for some recent analytical and semi-analytical models (Stevens et al.\ 2016; Shi et al.\ 2017; Lapi et al.\ 2018a; Posti et al.\ 2018b; Zoldan et al.\ 2018). They have also been the targets of many recent hydrodynamical simulations, some with large volume but relatively low resolution (Genel et al. 2015; Pedrosa \& Tissera 2015; Teklu et al.\ 2015; Zavala et al.\ 2016; DeFelippis et al.\ 2017; Lagos et al.\ 2017; Stevens et al.\ 2017; Lagos et al.\ 2018; Schulze et al.\ 2018) and others with small volume (zoom-in) but higher resolution (Agertz \& Kravtsov 2016; Grand et al.\ 2017; Soko{\l}owska et al.\ 2017; El-Badry et al.\ 2018; Obreja et al.\ 2018). The studies cited above generally confirm our $j_{\star}$--$M_{\star}$ scaling relations, particularly the exponent $\alpha \approx 0.6$ for disk-dominated galaxies. The exceptions to this near-consensus are the works by Obreschkow \& Glazebrook (2014) and Sweet et al.\ (2018), which found $\alpha \approx 1.0$ for galaxies of the same bulge fraction, including $\beta_{\star} = 0$. Obreschkow \& Glazebrook (2014) interpreted this to mean that the angular momenta of galactic disks are influenced in some way by the prominence of galactic bulges, i.e., that these two components are not independent, in contradiction to the picture discussed above. Complicating this comparison, however, is the fact that most of the galaxies in the Obreschkow \& Glazebrook (2014) and Sweet et al.\ (2018) samples have pseudo bulges rather than classical bulges. Thus, a secondary purpose of this paper is to resolve the apparent discrepancy between their work and ours. The remainder of this paper is organized as follows. In Section~2, we compare and contrast four determinations of the $j_\star$--$M_\star$ relation for galaxies of different bulge fraction, revealing some important similarities and differences. In Section~3, we present the corresponding two-dimensional (2D) surfaces defined by these relations in the three-dimensional (3D) space of $j_\star$, $M_\star$, and $\beta_\star$, and show that they are consistent with our picture of independent disks and spheroids. We summarize our results and discuss their implications in Section~4. We make some detailed comparisons between our estimates of $j_\star$, $M_\star$, and $\beta_\star$ and those of others in the Appendix. | The main conclusion of this paper is that the observed values of specific angular momentum $j_{\star}$, mass $M_{\star}$, and bulge fraction $\beta_{\star}$ of the stellar parts of most normal galaxies are consistent, in a first approximation, with a simple model based on a linear superposition of independent disks and bulges. The disks and bulges in this model follow scaling relations of the form $j_{\star{\rm d}} \propto M_{\star{\rm d}}^{\alpha}$ and $j_{\star{\rm b}} \propto M_{\star{\rm b}}^{\alpha}$ with $\alpha = 0.67 \pm 0.07$ but offset from each other by a factor of $8 \pm 2$ over the mass range $8.9 \leq \log (M_{\star}/M_{\odot}) \leq 11.8$. Separate fits for disks and bulges alone give $\alpha = 0.58 \pm 0.10$ and $\alpha = 0.83 \pm 0.16$, respectively. This simple model correctly predicts that galaxies will lie on or near a curved 2D surface specified by Equations (2) and (3) in the 3D space of $\log j_{\star}$, $\log M_{\star}$, and $\beta_\star$. These results reinforce and extend our earlier suggestion that the distribution of galaxies with different $\beta_{\star}$ in the $j_{\star}$--$M_{\star}$ diagram constitutes an objective, physically motivated alternative to subjective classification schemes such as the Hubble sequence. For disk-dominated galaxies in the mass range considered here, the $j_{\star}$--$M_{\star}$ scaling relation is now quite secure, as shown in Figure~1 by the excellent agreement between our determination (from Paper~2) and those of Obreschkow \& Glazebrook (2014) and Posti et al.\ (2018a). Two factors contribute to the robustness of this scaling relation. First, no special efforts are required to obtain photometric and kinematic data that extend to large enough radii (in units of $R_{\rm e}$) to estimate reliably the disk contributions $j_{{\star}{\rm d}}$ and $M_{{\star}{\rm d}}$ to the total values $j_{\star}$ and $M_{\star}$, which usually turn out to be close to those for ideal disks with exponential surface density profiles and flat rotation curves. Second, any uncertainties in the bulge contributions $j_{{\star}{\rm b}}$ and $M_{{\star}{\rm b}}$, even when substantial, have only a minor impact on the total values $j_{\star}$ and $M_{\star}$. For most of the giant galaxies studied here, cold gas (\hi\ and H$_2$) makes a relatively small contribution to their specific angular momentum and mass, and the stellar $j_{\star}$--$M_{\star}$ scaling relation is a good proxy for the baryonic $j_{\rm bary}$--$M_{\rm bary}$ scaling relation (Obreschkow \& Glazebrook 2014).\footnote {The baryonic scaling relations mentioned here include only stars and cold gas, not the warm and hot diffuse gas in galactic halos, which might actually dominate the total $j_{\rm bary}$ and $M_{\rm bary}$ budgets of some galaxies.} This is no longer true, however, for gas-rich dwarf galaxies, which contain more specific angular momentum and mass in cold gas than in stars. Several recent studies have extended the $j_{\rm bary}$--$M_{\rm bary}$ scaling relation down into the mass range $7 \lea \log (M_{\rm bary}/M_{\odot}) \lea 9$, with somewhat confusing claims about whether it lies above or matches onto the extrapolated $j_{\rm bary}$--$M_{\rm bary}$ and $j_{\star}$--$M_{\star}$ scaling relations from higher masses (Butler et al.\ 2017; Chowdhury \& Chengalur 2017; Elson 2017; Kurapati et al.\ 2018). Continuing and refining this work is important, because it has the potential to place constraints on the mass dependence of the retained or sampled fraction of specific angular momentum in galaxies $f_j$ (see below). For bulge-dominated galaxies, the $j_{\star}$--$M_{\star}$ scaling relation is based almost entirely on our work. In this case, the main challenge is obtaining kinematic data that extend to large enough radii (in units of $R_{\rm e}$) that the estimates of $j_{{\star}{\rm b}}$ have converged. This is important because the stellar rotation profiles of bulge-dominated galaxies, unlike the H$\alpha$ and HI rotation curves of disk-dominated galaxies, exhibit a great variety of behaviors; some are flat, while others rise or fall. All of our estimates of $j_{{\star}{\rm b}}$ are based on kinematic data that extend to $\sim 2R_{\rm e}$ and some to much larger radii, thus capturing as much angular momentum with as little extrapolation as possible. Nevertheless, additional studies of the $j_{\star}$--$M_{\star}$ scaling relation for spheroid-dominated galaxies, based on 2D kinematic data that reach even larger radii (for example, $R > 5 R_{\rm e}$), would certainly be desirable. For intermediate-type galaxies, the main challenge to deriving the $j_{\star}$--$M_{\star}$ scaling relation is in disentangling the contributions to $j_{\star}$ and $M_{\star}$ from the superposed disks and bulges. The specific angular momenta of bulges in such galaxies have been approximated in three different ways, by assuming that their rotation velocity is either (1) zero, (2) the same as the rotation velocity of their associated disks, or (3) the same as the mean rotation velocity for bulges of the same velocity dispersion and ellipticity. Method (1) and (2) clearly lead to systematic under- and overestimates of $j_{{\star}{\rm b}}$, respectively, while method (3), the one we have adopted, contributes some scatter but little if any bias to the $j_{\star}$--$M_{\star}$ scaling relation. More accurate results will require careful modeling of extensive 2D photometric and kinematic data to disentangle the velocity fields and hence the specific angular momenta of superposed disks and bulges (as in the recent work of Rizzo et al.\ (2018) on lenticular galaxies). The bulge fraction $\beta_\star$ is inherently uncertain because it depends on the adopted method for decomposing galaxies into disks and bulges, either by pre-specifying their 3D shapes (flat versus round) or by pre-specifying their surface brightness profiles (exponential versus S\'ersic). These two methods generally give similar values of $\beta_\star$ for bulge-dominated galaxies (elliptical, lenticulars, and early-type spirals), but they can give substantially different values of $\beta_\star$ for disk-dominated galaxies (late-type spirals). A related complication is the lack of consensus on the definition of pseudo bulges, including whether they must always be flat (like disks) or may sometimes be round (like spheroids). This ambiguity adds substantially to the uncertainty in estimates of $\beta_\star$ for disk-dominated galaxies, where pseudo bulges are much more common than classical bulges. We find no statistically significant indication that galaxies with pseudo bulges and classical bulges follow different relations in the space of $\log j_{\star}$, $\log M_{\star}$, and $\beta_\star$. This does not mean that both types of galaxies follow exactly the same relation, of course, merely that any differences must be small enough to hide within the scatter. Obreschkow \& Glazebrook (2014) found a different relation (with $\alpha \approx 1$) from a small sample of spiral galaxies with a preponderance of pseudo bulges (13/16) covering a narrow range in $\beta_\star$. This result, however, is based on adopted (partial) errors in $j_\star$, $M_\star$, and $\beta_\star$ that neglect the inherent uncertainties mentioned above and are therefore unrealistically small. As we have shown here, the statistical significance of the Obreschkow \& Glazebrook (2014) relation disappears when we adopt more realistic (total) errors in these quantities. Sweet et al.\ (2018) also found a different relation (again with $\alpha \approx 1$), based on a dataset with large systematic errors in $j_{\star}$. Finally, we offer a few remarks on the astrophysical implications of our results, following the precepts of Paper~1. Comparing the scaling relation for the stellar components of galaxies in the form $j_\star = j_0 (M_\star / M_0)^{\alpha}$ with that for dark-matter halos in the standard $\Lambda$CDM cosmology, we derive the relation $f_j / f_M^{2/3} = 6.8 (j_0/10^3 \,{\rm kpc\,km\,s^{-1}})(M_\star / M_0)^{\alpha -2/3}$ between the fractions of specific angular momentum and mass in stars relative to dark matter, $f_j \equiv j_\star / j_{\rm halo}$ and $f_M \equiv M_\star / M_{\rm halo}$ (with $M_0 = 10^{10.5} M_\odot$ again). With the exponent $\alpha$ and normalizations $j_0$ from our 3D fit to Equations~(2) and~(3), this relation becomes $f_j / f_M^{2/3} = 10.0 \pm 0.6$ for disks and $f_j / f_M^{2/3} = 1.2 \pm 0.4$ for bulges. Then, with the separate relations between $f_M$ and $M_\star$ for late-type and early-type galaxies from Dutton et al.\ (2010), we obtain $f_j \approx 1.0$ for disks and $f_j \approx 0.1$ for bulges at $M_\star \sim 10^{10.5} M_\odot$ and only mild variations over the range $10^{9.5} M_\odot \lea M_\star \lea 10^{11.5} M_\odot$. These estimates of $f_j$ differ slightly from the ones derived in Paper~2 for the same dataset because the new 3D and old 2D fits return slightly different values of $j_0$. The relations $f_j / f_M^{2/3} = {\rm constant}$ derived above imply that $f_j$ and $f_M$ must have qualitatively similar dependences on $M_\star$, namely a broad peak near $M_\star \sim 10^{10.5} M_\odot$, a shallow decline to lower $M_\star$, and a somewhat steeper decline to higher $M_\star$. This is why we find only mild variations in $f_j$. Recent analyses by other authors also indicate $f_j \approx {\rm constant}$ near $M_\star \sim 10^{10.5} M_\odot$ (see Fig.\ 12 of Lapi et al.\ 2018b and Fig.\ 3 of Posti et al.\ 2018a). Over much wider mass ranges, the deviations from a constant $f_j$ may become more pronounced. The model of disk formation preferred by Posti et al.\ (2018a) has $f_j \propto f_M^s \propto M_{\star}^{\gamma}$ with $\gamma = s (2 - 3\alpha)/(2 - 3s)$, which, when fitted to their full dataset ($10^{7.0} M_\odot \leq M_\star \leq 10^{11.3} M_\odot$), gives $\alpha =0.59 \pm 0.02$, $s = 0.4 \pm 0.1$, and thus $\gamma = 0.12 \pm 0.03$. However, even this weak dependence of $f_j$ on mass could be erased if the $j_{\rm bary}$--$M_{\rm bary}$ relation for gas-rich dwarf galaxies turns out to be shallower than the $j_\star$--$M_\star$ relation by only $\Delta\alpha \approx 0.15$ (again, see Fig.\ 3 of Posti et al.\ 2018a). This is why it is important to refine estimates of the baryonic relation at low masses. The fractions $f_j$ and $f_M$ for disks and bulges and the corresponding $j$--$M$ scaling relations (stellar and baryonic) are potentially determined by a large number of astrophysical processes. These include tidal torques, dynamical friction of baryonic structures within dark-matter halos, shocks and radiative cooling in the interstellar and circumgalactic media, star formation and its associated feedback, inflow, outflow, and recycling of gas, merging of gas clumps and dwarf galaxies, and tidal stripping of the outer parts of halos and their circumgalactic media by neighboring halos. We reviewed these processes and their potential impact on the $j$--$M$ scaling relations for disks and bulges at some length in Paper~1. Here, we note only the growing interest in biased-collapse models in which the fractions $f_j$ and $f_M$ are determined by the hypotheses that the baryons and dark matter in protogalaxies start with similar distributions of specific angular momentum and mass and that, at any given time, only the baryons within some critical radius are able to collapse and form the visible parts of galaxies. Analytical models of this type and their implications for the $j$--$M$ scaling relations are explored in several recent papers (Shi et al.\ 2017; Lapi et al.\ 2018a; Posti et al.\ 2018a, 2018b; see also Paper~1 and references therein). In the past few years, hydrodynamical simulations of forming galaxies have succeeded in reproducing, at least approximately, the observed $j$--$M$ scaling relations (Genel et al.\ 2015; Pedrosa \& Tissera 2015; Teklu et al.\ 2015; Agertz \& Kravtsov 2016; Zavala et al.\ 2016; DeFelippis et al.\ 2017; Grand et al.\ 2017; Lagos et al.\ 2017; Soko{\l}owska et al.\ 2017; Stevens et al.\ 2017; El-Badry et al.\ 2018; Lagos et al.\ 2018; Obreja et al.\ 2018). One of the main lessons from these simulations is that feedback in an essential ingredient to match the observed relations for both disk-dominated and spheroid-dominated galaxies. Without feedback, the simulations suffer from the well-known overcooling and angular momentum problems and fail to produce the full range of galactic morphologies. Another important ingredient is merging, which appears to explain, at least partially, the slow rotation of spheroids relative to disks. Despite the success of recent analytical models and hydrodynamical simulations, we do not yet have definitive answers to some important theoretical questions about the $j$--$M$ scaling relations, such as the following. Given the potential complexity of galaxy formation, why are the observed $j$--$M$ relations so simple? In particular, why are the specific angular momentum and mass fractions $f_j$ and $f_M$ so closely linked that they result in power-law $j$--$M$ relations over 3--4 decades in mass (at least for disks)? Why do the disks of massive galaxies have nearly the same specific angular momentum as their dark-matter halos ($f_j \sim 1.0$) and why do their bulges have much less ($f_j \sim 0.1$)? Answering these questions will require a better understanding of how much the specific angular momentum of mass elements inside forming galaxies is redistributed and which physical mechanisms are most responsible for this redistribution. This is a promising direction for future analysis of hydrodynamical simulations (as already begun by DeFelippis et al.\ 2017). | 18 | 8 | 1808.02525 |
1808 | 1808.03523_arXiv.txt | {The precise measurement of the masses and radii of stars in eclipsing binary systems provides a window into uncertain processes in stellar evolution, especially mixing at convective boundaries. Recently, these data have been used to calibrate models of convective overshooting in the cores of main sequence stars. In this study we have used a small representative sample of eclipsing binary stars with $1.25 \leq M/\text{M}_\odot < 4.2$ to test how precisely this method can constrain the overshooting and whether the data support a universal stellar mass--overshooting relation. We do not recover the previously reported stellar mass dependence for the extent of overshooting and in each case we find there is a substantial amount of uncertainty, that is, the same binary pair can be matched by models with different amounts of overshooting. Models with a moderate overshooting parameter $0.013 \leq f_\text{os} \leq 0.014$ (using the scheme from \citealt{1997A&A...324L..81H}) are consistent with all eight systems studied. Generally, a much larger range of $f_\text{os}$ is suitable for individual systems. In the case of main sequence and early post-main sequence stars, large changes in the amount of overshooting have little effect on the radius and effective temperature, and therefore the method is of extremely limited utility.} | \label{sec:introduction} The treatment of mixing at convective boundaries is a fundamental uncertainty for stellar evolution calculations. Basic arguments imply there must be some mixing beyond locally-determined convective boundaries according to, for example, the Schwarzschild criterion. Theoretical estimates for the extent of overshooting vary considerably, ranging from very little to a zone of complete mixing around two pressure scale heights in depth. The amount of overshooting in convective cores affects the main sequence lifetime and therefore the inferred age of stellar clusters and individual post-main sequence stars. Convective core overshooting also increases the luminosity and speed of evolution of post-main sequence stars. Several independent lines of evidence -- colour-magnitude diagrams of star clusters, double-lined eclipsing binary (DLEB) stars, and asteroseismology -- strongly suggest there is mixing beyond the Schwarzschild boundary of convective cores in main sequence stars. By increasing the availability of hydrogen that can be burnt in the convective core, this mixing significantly extends the predicted main sequence lifetime. There is currently no universally accepted theoretical basis to predict the extent of such mixing: it is typically dependent on a parameter (with or without a physical model). In subsequent phases of evolution, the mixing beyond the Schwarzschild boundary of convection zones is equally crucial, but the relative scarcity of observational constraints means that the evolution is even more uncertain. Characterizing and quantifying the processes operating in main sequence convective cores and convective boundaries may also help improve the models of later phases of stellar evolution. Historically, several authors have proposed extensions to mixing length theory in order to quantify the amount of overshooting. \citet{1965MNRAS.130..223R} argued that convective core overshooting region is of the order $10^{-3}$ times the stellar radius, which is up to about ten per cent of the radius of the convective core. \citet{1973ApJ...184..191S} determined an average extent of convective overshoot of $0.01\,\text{M}_\odot$. Adding more sophistication to the approach of \citet{1973ApJ...184..191S}, by accounting for the convective flux carried by overshooting elements and the resultant effect on the temperature gradient, \citet{1975ApJ...201..637C} arrived at 0.23 pressure scale heights of core overshooting for a 3$\,M_\odot$ star, in line with empirical estimates. The applicability of these methods to stellar evolution calculations is limited by our lack of knowledge about the properties of convection in stellar cores and the difficulty of relating these penetration arguments to chemical mixing. The best constraints for core overshooting so far have an empirical basis. The most common approach has been to compare the width of the main sequence and shape of the turnoff observed colour-magnitude diagrams of stellar clusters with theoretical predictions. \citet{1992A&AS...96..269S}, for example, {found that models with initial masses $1.25 \leq M/\text{M}_\odot \leq 25$ with 0.2\,pressure scale heights of overshooting were the best fitting for 65 observed clusters}. Other studies have concluded that a similar magnitude of overshooting is needed. This amount is often used as a default in stellar evolution codes and in published isochrones that are used widely by the astrophysical community. The rationale for using DLEB stars to constrain main sequence overshooting is exactly the same as it is for using stellar clusters: they comprise stars born at the same time and with the same composition but different mass. Although each system offers only a very limited insight compared with an entire stellar cluster, there is compensation from the high measurement precision. | In response to recent findings that the amount of main sequence overshooting required to explain the observations of double-lined eclipsing binary stars is strongly dependent on stellar mass, we have conducted a detailed exploration of a sample of such systems and tested the sensitivity of the results to some important uncertainties. We took a representative selection of eight eclipsing binary systems, covering a wide mass range and including stars in various phases of evolution, from the samples used by \citet{2016A&A...592A..15C,2017ApJ...849...18C,2018ApJ...859..100C}. We modelled overshooting (and any other mechanisms for mixing near the boundary of the convective core) by varying the free parameter $f_\text{os}$ in the scheme from \citet{1997A&A...324L..81H}, where there is an exponential decay in the diffusion coefficient in formally stable regions. We investigated an array of models for each system to establish a range of overshooting parameters that yielded acceptable solutions according to effective temperature, radius, age, and metallicity constraints. These results are presented in Table~\ref{table_results_summary}. In general, our results are indicative of the range of overshooting consistent with the observations but do not necessarily reach the possible extremes. We compare our determinations for the amount of overshooting with \citet{2017ApJ...849...18C,2018ApJ...859..100C} in Figure~\ref{figure_results_summary}. Our findings are usually consistent with the best fit models from \citet{2017ApJ...849...18C,2018ApJ...859..100C} but we find a large range of acceptable $f_\text{os}$ that makes it very difficult to detect any trend with mass. We confirm earlier results that the evidence strongly supports the requirement for overshooting in models of stars with $M \gtrsim 2\,\text{M}_\odot$. We could match all of the eight pairs with $0.013 \leq f_\text{os} \leq 0.014 $ (and seven of the pairs with $0.013 \leq f_\text{os} \leq 0.018 $), which is remarkably consistent with the range of best fit $f_\text{os} \approx 0.016$ found by \citet{2017ApJ...849...18C} for stars with $M > 2\,\text{M}_\odot$. None of the five DLEB pairs of main sequence or subgiant stars were particularly useful for constraining core overshooting. We were, however, able to more tightly constrain the overshooting parameter in models in later phases of evolution. Unfortunately, this presents new challenges because the radius and effective temperature evolution of those models are more strongly dependent on the mixing length parameter and metallicity, and stars can pass through the same place in the HR diagram multiple times, which complicates the search for the most favourable parameters. We have shown that in most cases a valid solution exists with a range of overshooting parameter, even without conducting an exhaustive search of the parameter space, which additionally includes metallicity, helium abundance, possible discrepancies between $f_\text{os}$ or $\alpha_\text{MLT}$ for the two components, and uncertainties in the helium-burning reaction rates for evolved systems. We also caution that in this study we have not formally weighted the solution likelihoods where it may be possible, by considering the duration of the windows of valid solutions with each combination of parameters, for example. In their recent paper, \citet{2018arXiv180307058V} raised the question of whether their conclusions about the difficulty of precisely constraining the overshooting from an eclipsing binary pair apply generally. We have identified that in most cases it is indeed difficult to definitively determine the extent of overshooting from the available measurements of stellar masses, radii, and effective temperatures. In many examples, the allowed range of the overshooting parameters could be reduced with more precise determinations of effective temperature and metallicity. The situation may also be helped by complementary approaches such as asteroseismology and hydrodynamical models which are now being applied to the same problem. We also wish to emphasize the value of the recent approach of \citet{2016A&A...592A..15C} where models for large numbers of systems are assessed together, especially as observations improve in both quantity and quality, which will reduce the uncertainties in each specific case and therefore overall. Overall, we do not find evidence to support a mass dependence for the amount of overshooting, other than that it is necessary for models with mass above about 2\,$\text{M}_\odot$. We find that a constant overshooting parameter provides an adequate fit to the data. | 18 | 8 | 1808.03523 |
1808 | 1808.01183_arXiv.txt | {Despite their activity, low-mass stars are of particular importance for the search of exoplanets by the means of Doppler spectroscopy, as planets with lower masses become detectable. We report on the discovery of a planetary companion around HD\,180617, a bright ($J = 5.58$\,mag), low-mass ($M = 0.45$\,M$_\odot$) star of spectral type M2.5\,V. The star, located at a distance of 5.9\,pc, is the primary of the high proper motion binary system containing vB\,10, a star with one of the lowest masses known in most of the twentieth century. Our analysis is based on new radial velocity (RV) measurements made at red-optical wavelengths provided by the high-precision spectrograph CARMENES, which was designed to carry out a survey for Earth-like planets around M dwarfs. The available CARMENES data are augmented by archival Doppler measurements from HIRES and HARPS. Altogether, the RVs span more than 16 years. The modeling of the RV variations, with a semi-amplitude of $K = 2.85_{-0.25}^{+0.16}\,\metrepersecondnp$, yields a Neptune-like planet with a minimum mass of $12.2_{-1.4}^{+1.0}$~M$_\oplus$ on a $105.90_{-0.10}^{+0.09}$~d circumprimary orbit, which is partly located in the host star's habitable zone. The analysis of time series of common activity indicators does not show any dependence on the detected RV signal. The discovery of HD\,180617\,b not only adds information to a currently hardly filled region of the mass-period diagram of exoplanets around M dwarfs, but the investigated system becomes the third known binary consisting of M dwarfs and hosting an exoplanet in an S-type configuration. Its proximity makes it an attractive candidate for future studies. } | \label{sec:intro} The search for exoplanets by means of Doppler spectroscopy has significantly advanced since the first confirmed discoveries at the end of the last century \citep[e.g.,][]{Queloz_1995}. The development in instrumentation within this field led to detection limits on the order of meters per seconds, as has been demonstrated, for example, by HIRES at the Keck Observatory \citep{Vogt_1994}, HARPS at the La Silla Observatory \citep{Mayor_2003}, HARPS-N at El Roque de Los Muchachos Observatory \citep{Consentino_2014}, or the Automated Planet Finder at the Lick Observatory \citep{Rado_2014}.\\ \indent Still, the choice of the parent stars is crucial. It allows pushing the boundaries with respect to the exoplanet parameter space, and exoplanet programs targeting M dwarfs have become of major interest (\citealp[e.g.,][]{Char_2008}; \citealp{Zechmeister_2009a}; \citealp{Bon_2013}). Because of their low masses, M dwarfs are particularly suitable for Doppler surveys, since the semi-amplitude of a radial velocity (RV) signal induced by a companion increases with decreasing stellar mass. Thus, lower mass planets are more easily detectable around low-mass stars. M dwarfs have luminosities between $10^{-4}$ and $10^{-1} \mathrm{L}_\odot$ and produce significantly less flux than their more massive counterparts. Consequently, their habitable zones (HZs) are located closer to their host stars at distances of about 0.05 to 0.4\,au (\citealp[e.g.,][]{Joshi_1997}; \citealp{Tarter_2007}; \citealp{Kopp_2014}; \citealp{Dress_2015}).\\ \indent However, the intrinsic stellar activity of M dwarfs can hinder the search for exoplanets around these stars. Signatures of stellar activity such as active regions on the stellar surface and their cyclic variability can mimic radial velocity signals that have previously been misinterpreted as arising from a planetary companion. Since the strength of such activity-related signals has been shown to be wavelength dependent \citep[e.g.,][]{Desort_2007, Reiners_2010}, a wide wavelength coverage in RV measurements can help distinguish planetary signals from those due to activity \citep[e.g.,][]{Sarkis_2018}.\\ \indent The high-resolution fiber-fed spectrograph CARMENES\footnote{\tt http://carmenes.caha.es}, which is installed at the 3.5~m telescope at the Calar Alto Observatory in Spain, was specifically designed with a broad wavelength coverage. It consists of two spectrograph channels, which together cover wavelengths from 520 to 1710~nm, with a resolution of 94\,600 in the visual and 80\,500 in the near-infrared channel \citep{Quirrenbach_2014, Quirrenbach_2016}. The instrument has been operating since January 2016 and is performing a search for exoplanets around M dwarfs. More than 300 targets that have been selected from the CARMENES input catalog Carmencita \citep{Caballero_2016b, Reiners_2018b} are regularly monitored for this purpose. The capability of the visual channel of achieving an RV precision of 1--2~$\metrepersecondnp$ has been demonstrated by \citet{Seifert_2016}, \citet{Trifonov_2018}, and the first CARMENES exoplanet discovery by \citet{Reiners_2018a}. \\ \indent In this paper we analyze the RV data of HD\,180617, one of the M dwarfs monitored by CARMENES. The measurements indicate the presence of a planet with a minimum mass comparable to Neptune on an orbit partly located within the habitable zone of the host. In Section~\ref{sec:stars} we characterize the observed star, followed by a short description of the RV data compilation in Section~\ref{sec:data}. The results from the RV analysis are presented in Section~\ref{sec:RV_ana}, and we conclude in Section~\ref{sec:conclusions}. | \label{sec:conclusions} We presented radial velocity time series of the inactive early-M dwarf HD\,180617 using RV measurements from the visual channel of the high-resolution spectrograph CARMENES in addition to HARPS and HIRES archival data. The combined data sets indicate the existence of a planet with a minimum mass of 12.2\,M$_\oplus$ on a 105.9\,d orbit around its stellar host. Our investigation of periodicities in the activity indicators for H$\alpha$ emission, RV chromaticity, and line profile variations shows no obvious indication that the planetary signal is due to activity-induced RV variability. The orbit of the Neptune-like planet, with semi-major axis $a=0.3357$\,~au and eccentricity $e = 0.16$, is partly located in its host star's liquid water habitable zone. Under the condition of an edge-on view onto the system, the semi-amplitude astrometric signature at Earth's distance is estimated to be about 4.6$\,\mu$as. With decreasing orbital inclination, the planetary mass and therefore the astrometric signature increase, therefore this is only a lower limit and the astrometric orbit could potentially be observed by \textit{Gaia}. Because of the proximity of HD\,180617 ($d=5.9$\,pc), the angular separation of the planet from its host star is rather large, comparable to the ${\sim}50$\,mas inner working angle of a high-performance coronograph on a 4~m class space telescope. While the mass of HD\,180617\,b is too high to be considered a close Earth analog, it is thus nevertheless a prime target for future missions - such as the HabEx concept \citep{Menn_2016} - aimed at studying the atmospheres of potentially life-bearing planets, or ESO's ground-based Extremely Large Telescope \citep{Gil_2007}.\\ \indent Only 15 closer stars, 10 of which are M dwarfs, are currently known to have an exoplanet. Within this group, HD\,180617 is the seventh brightest in the $V$ band. Moreover, it is the primary of the wide binary containing the low-mass star vB\,10. We provide new measurements on the binary separation and new determinations of the primary mass, radius, luminosity, and Galactocentric space motion. Of the 88 currently confirmed binary systems that host exoplanets, only a handful possess an M dwarf as the primary \citep{Schwarz_2016}. Except for Gliese\,676\,A/B and Gliese\,15\,A/B, the system investigated in this work is the third of this kind with an exoplanet in an S-type\footnote{In this context, ``S'' stands for satellite.} configuration \citep{Rabl_1988}, which means that the planet is in a close orbit around only one of the systems' stars. Because of the wide separation ($s \ge$ 450\,au) between the system's stellar components (more than 1300 times larger than the semi-major axis of HD\,180617\,b), the perturbation of vB\,10 on the formation and dynamical evolution of the planet is expected to be negligible \citep[e.g.,][]{Whit_1998, Quin_2007, Quin_2008}. In this configuration, the gravitational pull from the host star on the planet is about 10$^7$ times stronger than from the secondary. Furthermore, given the very low luminosity of 425$\pm$4~10$^{-6}$\,L$_\odot$ of the M8.0\,V companion \citep{Cif_2017}, the presence of vB\,10 is not expected to affect the radiative equilibrium of HD\,180617\,b. Consequently, in this particular case, binarity does not affect habitability.\\ \indent In line with HD\,147379\,b, the first exoplanet discovery by CARMENES \citep{Reiners_2018a}, our work demonstrates the potential and capability of the instrument of finding exoplanets within the mass range of Neptunes and mini-Neptunes within the habitable zones of M dwarfs. Both HD\,180617\,b and HD\,147379\,b have orbital periods of about 100\,d. Although dozens of exoplanets with lower masses on shorter orbits around M dwarfs are known, these two discoveries belong to only a few known Neptunes at longer orbits with periods $\gtrsim 100$\,d. While they improve the sampling within this barely filled region of the exoplanet parameter space, even longer periods for exoplanets of similar masses around M dwarfs are becoming accessible as a result of the steadily increasing time baselines and the improved precision of current and future instruments. | 18 | 8 | 1808.01183 |
1808 | 1808.05978_arXiv.txt | Covariance matrix estimation is a persistent challenge for cosmology. We focus on a class of model covariance matrices that can be generated with high accuracy and precision, using a tiny fraction of the computational resources that would be required to achieve comparably precise covariance matrices using mock catalogues. In previous work, the free parameters in these models were determined using sample covariance matrices computed using a large number of mocks, but we demonstrate that those parameters can be estimated consistently and with good precision by applying jackknife methods to a single survey volume. This enables model covariance matrices that are calibrated from data alone, with no reference to mocks. | No matter the scope of a cosmological survey, we have only one sky to observe. This complicates the statistical analysis of cosmological surveys. A common approach is to generate a large number of independent, synthetic skies, then apply standard sample statistics to them. The readily apparent limitations of this approach are that it is challenging to ensure that the synthetic skies reflect the physics of the actual universe, and that the computational cost of generating these ``mock catalogues'' can be substantial. In this paper we take an approach introduced in \cite{OConnell:2015src}, which generates the covariance matrix for a galaxy correlation function with correct long-distance physics and survey geometry, and extend it so that the \emph{short-distance} physics can be calibrated directly against a survey, without reference to mocks. Using mock catalogues to generate a covariance matrix requires a large number of reasonably accurate mocks. The consequences of having an insufficient number of mocks have received significant attention. If $n_{\text{samples}}$ mocks are used to generate a covariance matrix for a correlation function estimated using $n_{\text{bins}}$ different bins the scale for noise in the covariance matrix is set by $n_{\text{bins}}/n_{\text{samples}}$. Noise in the covariance matrix propagates through to become additional noise on cosmological parameter estimates, increasing the parameter covariance by a factor of $n_{\text{bins}}/n_{\text{samples}}$\textbf{ }\citep{Dodelson:2013uaa,Percival:2013sga}. We note that the next generation of surveys, including Euclid \citep{EuclidRedBook}, DESI \citep{Levi:2013gra}, and WFIRST \citep{Dore:2018smn} aim at tomographic analyses and that a simultaneous analysis of multiple redshift bins will dramatically increase the number of correlation function bins used. In \cite{OConnell:2015src} we developed a method built on the observation, due to Bernstein \citep{Bernstein1994}, that the covariance matrix of a 2-point correlation function can itself by written in terms of correlations between four points, integrated over the survey volume. We perform these integrals using a realistic 2-point correlation function and accurately representation of the survey geometry, to produce a covariance matrix that accurately reflects the long-distance physics and structure of the survey. Related work includes \cite{Pearson:2015gca}, which constructed a simple model of the power spectrum covariance matrix, and \cite{Grieb:2015bia}, which took a similar approach to that in\textbf{ }\cite{OConnell:2015src} but on a cubic, uniform survey. The four-point correlations noted above include contributions from the connected 3- and 4-point galaxy correlation functions, which are not as well-understood as the 2-point function. \cite{OConnell:2015src} approximated these contributions by introducing a shot-noise rescaling parameter, $a$, which effectively modelled the short-distance contributions to the covariance matrix from the 3- and 4-point functions as an increase in shot-noise. $a$ was estimated using mock catalogues and the resulting fit covariance matrix was found to be both accurate and precise. Critically, this required a tiny fraction of the computational time required to generate a mock covariance matrix of comparable precision. The procedure is illustrated in figure \ref{fig:OldFlowChart}. This approach was further tested in \cite{Vargas-2016}, where it was found that the resulting covariance matrix performed at least as well as a mock covariance matrix for BAO measurements with BOSS \citep{2013AJ....145...10D}. In this paper, we propose to estimate the shot-noise rescaling $a$ using actual survey data, rather than mock catalogues. The essential observation is that since $a$ is being used to model short-distance physics, we need not use or mimic the entire survey volume in order to estimate it. Instead we propose to use the actual survey to generate a jackknife covariance matrix, then use the jackknife covariance matrix to estimate $a$. The computational cost to do this is quite low, since generating the jackknife covariance matrix requires very little computation beyond counting the pairs in the survey. The estimated value of $a$ can then be used in the original model covariance matrix. This new procedure is illustrated in figure \ref{fig:NewFlowChart}. We find that the level of precision on $a$ that can be achieved with a single survey volume is ample for many applications. It is therefore possible to perform covariance matrix estimation for upcoming surveys without reference to mock catalogues. The observation that relatively small volumes can provide usable information about the covariance matrix has generated recent interest. \cite{Klypin:2017} investigated the power spectrum covariance matrix using small-volume cubic mocks. Small-volume cubic simulations were also used in\textbf{ }\cite{Howlett:2017vwp} to generate a scaled covariance matrix for the 2-point correlation function. The jackknife approach introduced here allows us to utilise small-scale information while accurately reflecting the true survey geometry. In light of the urgency of the covariance matrix problem for upcoming surveys, many approaches to the problem are currently being developed: \begin{itemize} \item New techniques in mock generation aim to increase $n_{\text{samples}}$. For overviews of recent progress see \cite{Chuang:2014toa} and \cite{Lippich:2018wrx}. \item Compression of the correlation function can reduce $n_{\text{bins}}$. This can be particularly helpful in analysing tomographic data. A prominent example is the ``redshift weights'' approach\textbf{ }introduced in \cite{Zhu:2014ica} and most recently applied in \cite{Zhu:2018edv}. \item Several empirical techniques have been developed to smooth sample covariance matrices computed from mocks. These take advantage of resampling methods \citep{Escoffier:2016qnf}, shrinkage \citep{Joachimi:2016xhk}, or the sparse structure of the precision matrix (the inverse of the covariance matrix) \citep{Padmanabhan:2015vlf}. \end{itemize} The result is that practitioners can combine a variety of physical and statistical insights when analysing a cosmological survey. We hope that our contribution will be useful in this regard. This paper is organised as follows. In section \ref{sec:Review} we briefly review the results of \cite{OConnell:2015src}, including the full and Gaussian model covariance matrices and the 1-parameter model for non-Gaussian contributions. In section \ref{sec:Jackknife} we specify how we will compute a jackknife covariance matrix from a single survey volume and how we will compute the corresponding jackknife \emph{model} covariance matrix. In section \ref{sec:Validation-Methods} we use mocks to verify that the values of $a$ estimated from single survey volumes, using a jackknife, are consistent with the values that would be estimated from those mocks using a sample covariance. This establishes the consistency of our method and provides evidence for our claim that $a$ is modelling short-distance physics. We conclude in section \ref{sec:Outlook}. In appendix \ref{sec:Inversion} we present a jackknife-inspired method for accurately inverting a model covariance matrix. That method is used in this paper and may be of interest to researchers working on model covariance matrices in other contexts. \begin{figure} \includegraphics[width=\textwidth]{Figs_v2/Flow_Chart_Old.pdf} \caption{\label{fig:OldFlowChart}The procedure introduced in \citep{OConnell:2015src} for generating and calibrating a model covariance matrix. The model includes one unknown parameter, $a$, the shot-noise rescaling, which is calibrated using mock catalogues.} \end{figure} \begin{figure} \includegraphics[width=\textwidth]{Figs_v2/Flow_Chart_New.pdf} \caption{\label{fig:NewFlowChart}The procedure proposed in this paper for generating and calibrating a model covariance matrix. The unknown parameter $a$ is calibrated using the data directly, rather than mock catalogues. This is accomplished using jackknife methods.} \end{figure} | 18 | 8 | 1808.05978 |
|
1808 | 1808.09151_arXiv.txt | {Although the \emph{Gaia} catalogue on its own is a very powerful tool, it is the combination of this high-accuracy archive with other archives that will truly open up amazing possibilities for astronomical research. The advanced interoperation of archives is based on cross-matching, leaving the user with the feeling of working with one single data archive. The data retrieval should work not only across data archives but also across wavelength domains. The first step for a seamless access to the data is the computation of the cross-match between \emph{Gaia} and external surveys.} {We describe the adopted algorithms and results of the pre-computed cross-match of the \emph{Gaia} Data Release 2 (DR2) catalogue with dense surveys (Pan-STARRS1~DR1, 2MASS, SDSS~DR9, GSC~2.3, URAT-1, allWISE, PPMXL, and APASS~DR9) and sparse catalogues (Hipparcos2, Tycho-2, and RAVE~5).} {A new algorithm is developed specifically for sparse catalogues. Improvements and changes with respect to the algorithm adopted for DR1 are described in detail.} {The outputs of the cross-match are part of the official \emph{Gaia}~DR2 catalogue. The global analysis of the cross-match results is also presented. } {} | The \emph{Gaia} satellite allows determining high-accuracy positions for $\sim$1.7 billion sources and parallaxes and proper motions for $\sim$1.3 billion sources observed all-sky down to magnitude G$\sim$20.7. Compared to the first intermediate \emph{Gaia} Data Release (DR1, see \citealt{Brown2016} for a summary of the astrometric, photometric, and survey properties, and \citealt{Prusti2016} for the scientific goals of the mission), the second intermediate \emph{Gaia} Data Release \citep{Brown2018} provides 48\% additional sources, parallaxes, and proper motions with an unprecedented accuracy for 77\% of all observed sources, which are complemented by a precise and homogeneous multi-band photometry and a large radial velocity survey for more than 7\,000\,000 sources with G magnitude in the $4-13$ range. Astrophysical parameters for $\sim$160 million sources, data on more than 500\,000 variable stars, and $\sim$14\,000 solar system objects are also available in DR2\footnote{A more exhaustive overview of the mission and DR2 details can be found at \url{ https://www.cosmos.esa.int/web/gaia/dr2-papers}}. The main goal of adding a pre-computed cross-match to \emph{Gaia}~DR2 data is complementing \emph{Gaia} with existing astrophysical quantities (that are widely used by the scientific community). This allows the full exploitation of the scientific potential of \emph{Gaia} . The general principles of the adopted cross-match algorithm are given and discussed in \citealt{Marrese2017} (hereafter Paper~I). We here briefly recall that any cross-match algorithm is a trade-off between multiple requisites, and a fraction of mismatched and/or missed objects is always present. Our aim is to define and implement a cross-match algorithm that on one hand should be general enough to be exploited for different scientific cases, and on the other should have complete results that can later be filtered to better fullfil a specific scientific problem. We tried to find a reasonable compromise between the completeness and correctness requirements, which implies that we needed to avoid adding too many spurious matches. In Sections~\ref{sec:gen} and \ref{sec:details} we describe the general principle and the details of the cross-match algorithms defined for \emph{Gaia}~DR2, respectively. Section~\ref{sec:extcat} contains the list of the external catalogues that we matched with \emph{Gaia}~DR2 data and a short description for the newly added catalogues, together with some issues or caveats that are relevant to the cross-match. In Sections~\ref{sec:output} and \ref{sec:results} we describe and discuss the cross-match output content and the results. Finally, Appendix~\ref{sec:app} contains a discussion of the effective angular resolution of external catalogues and its influence on the cross-match. | } \begin{figure} \centering \includegraphics[width=1.0\linewidth]{figure7.pdf} \caption{Issues encountered when cross-matching Hipparcos2 with \emph{Gaia}. For a detailed description and explanation of the results shown in this figure, we refer to the main text (Subsection~\ref{subsec:hipissue}).} \label{Fig:HipIssue} \end{figure} We presented the algorithms we developed for the official cross-match of the high-accuracy \emph{Gaia} DR2 astrometric data with eight large dense surveys and three sparse catalogues. The defined algorithms are positional and are able to fully exploit the enormous number of \emph{Gaia} sources with accurate proper motions and parallax measurements using the full five-parameters astrometric covariance matrix on an object-by-object basis. In addition, we included an improved definition of the surface density of observed objects for each catalogue, which allows a better evaluation of the local environment. The external catalogues and cross-match results were also described. In particular, we analysed the global behaviour of the cross-match results by evaluating their sky distribution, statistical indicators, magnitude, and angular distance distributions. More importantly, we tried to supply scientists, both in the output tables and in the analysis performed in this paper, with all the means to verify the quality of the cross-match results and to understand whether this cross-match is appropriate for their scientific needs. The excellent data provided by the \emph{Gaia}~DR2, and in particular the proper motions, substantially improve the quality of \emph{Gaia} counterparts that are found in external catalogues. The high accuracy of the current \emph{Gaia} data gives a strong drive and powerful tools for understanding and quantifying known complex issues (such as resolution effects, and the presence of astrometric binaries and of duplicated sources) that influence the cross-match results and require non-trivial solutions. The issues will be tackled in the forthcoming \emph{Gaia} data releases. | 18 | 8 | 1808.09151 |
1808 | 1808.09960_arXiv.txt | {In this article, we introduce a de-sitter model in favor of compact stars in low-mass X-ray binary. Here, we merge the presence of the cosmological constant on a small scale to discussion the stellar structure and conclude that this doping is very well suitable with the familiar physical mode of the low-mass X-ray binary compact stars. We calculate the probable radii, compactness (u) and surface red-shift ($Z_{s}$) of six compact stars in low-mass X-ray binaries namely Cyg X-2, V395 Carinae/2S 0921-630, XTE J2123-058, X1822-371 (V691 CrA), 4U 1820-30 and GR Mus (XB 1254-690). We also offer possible equation of state (EOS) of the stellar object. | % As compact stars (neutron stars/strange stars) plays a crucial role to relate astrophysics, nuclear physics $\&$ particle physics, it becomes a great interesting topic to study for long time. Commonly neutron stars are built almost fully of neutrons whereas strange stars are can be composed entirely of strange quark matter (SQM) or the conversion (up-down quarks to strange quarks) may be confined to the core of the neutron star (Haensel et al.~\cite{Haensel1986}; Drago et al.~\cite{Drago2014}). It is well known that neutron stars are bounded by gravitational attraction and on the other hand strange stars are bounded by strong interactions as well as gravitational attractions. Therefore, for lower mass neutron stars gravitational bound becomes much weaker than the strange stars. Hence neutron stars become larger in size in comparison to the strange stars of same mass. All the present EOS of neutron star have zero surface matter density, whereas available EOS of strange star obtained a sharp surface (Farhi \& Jaffe~\cite{Farhi1984}; Haensel et al.~\cite{Haensel1986}; Alcock et al.~\cite{Alcock1986}; Dey et al.~\cite{Dey1998}). Since within very few seconds of life of a neutron star that temperature reduced to less than Fermi energy, hence for a given equation of state the mass and radius of the star depends solely on central density and also it is very hard to find out mass and radius of a neutron star simultaneously. For a detail study we suggest a review work of Lattimer \& Prakash~(\cite{lattimer2007}). Conceptual account of mass and radii of spherically-symmetric non-rotating compact stars are results of analytical or numerical solutions of Tolman-Oppenheimer-Volkoff i.e., TOV equations. From observational point of view, some promising area for surveying of mass and radius of a compact star (neutron stars/strange stars) are thermal emission from cooling stars, pulsar timing, surface explosions and gravity wave emissions. For the experimental scientist face the recent challenges are to use giant dipole resonances, heavy-ion collisions and parity-violating electron scattering techniques to measure the density dependability pressure of nuclear matter. Actually, the most challenging task is to determine the proper EOS to describe the internal formation of a neutron star (Ozel~\cite{Ozel2006}; $\ddot{O}$zel et al.~\cite{Ozel2009a}; $\ddot{O}$zel \& Psaitis~\cite{Ozel2009b}; $\ddot{O}$zel et al.~\cite{Ozel2010}; G$\ddot{u}$ver et al.~\cite{Guver2010a}, \cite{Guver2010b}). Though a several dozen compact star masses have been determined very exactly (to some extend) in binaries (Heap \& Corcoran~\cite{Heap1992}; Van et al.~\cite{Van1995}; Stickland et al.~\cite{Stickland1997}; Orosz \& Kuulkers~\cite{Orosz1999}; Lattimer \& Prakash~\cite{lattimer2005}, \cite{lattimer2007}), no radius information is obtainable for these systems. Therefore, theoretical study of the stellar structure is essential to support the correct direction for the newly observed masses and radii. Here, some of the researcher's work on compact stars (Lobo~\cite{Lobo2006}; Bronnikov \& Fabris~\cite{Bronnikov2006}; Hossein et al.~\cite{Hossein2012}; Rahaman et al.~\cite{Rahaman2012a}, \cite{Rahaman2012b}; Maharaj et al.~\cite{Maharaj2014}; Pant et al.~\cite{Pant2014}; Ngubelanga et al.~\cite{Ngubelanga2015}; Paul et al.~\cite{Paul2015}; Kalam et al.~\cite{Kalam2012}, \cite{Kalam2013a}, \cite{Kalam2013b}, \cite{Kalam2014a}, \cite{Kalam2014b}, \cite{Kalam2016}, \cite{Kalam2017}; Jafry et al.~\cite{Jafry2017}; Maurya et al.~\cite{Maurya2016}; Dayanandan et al.~\cite{Dayanandan2016}; Bhar et al.~\cite{Bhar2017}) are to be mentioned. Casares et al.~(\cite{Casares2010}) surveyed the mass of the compact star in Cyg X-2 by using new high-resolution spectroscopy and it comes out as $1.71 \pm 0.21 M_{\odot}$. In another work, Steeghs \& Jonker~(\cite{Steeghs2007}) measured the mass of the compact star in V395 Carinae/2S 0921-630 with the help of MIKE echelle spectrograph on the Magellan-Clay telescope by using high-resolution optical spectroscopy and it comes out as $1.47 \pm 0.10 M_{\odot}$. On the other hand, Gelino et al.~(\cite{Gelino2003}) surveyed the mass of the compact star in XTE J2123-058 as $1.53^{+0.30}_{-0.42} M_{\odot}$. Mu\~{n}oz-Darius et al.~(\cite{Munoz2005}) surveyed the mass of the neutron star in low-mass X-ray binary (LMXB) X1822-371 (V691 CrA) by perusing the K-correction for the case of ejection lines formed in the X-ray illuminated atmosphere of a Roche lobe filling star and that appear as $1.61 M_{\odot} \leq M_{NS} \leq 2.32 M_{\odot}$. In a recent work, G$\ddot{u}$ver et al.~(\cite{Guver2010b}) surveyed the mass of the compact star in 4U 1820-30 by using time resolved X-ray spectroscopy of the thermonuclear burst of 4U1820-30 and it comes out as $1.58 \pm 0.06 M_{\odot}$. Barnes et. al~(\cite{Barnes2007}) have also determined the mass of the compact object in GR Mus (XB 1254-690) as $1.20 M_{\odot} \leq M_{NS} \leq 2.64 M_{\odot}$. Wilkinson Microwave Anisotropic Probe (WMAP) measurement indicates that in the Universe nearly 73\% of total mass-energy is Dark Energy (Perlmutter et al.~\cite{Perlmutter1998}; Riess et al.~\cite{Riess2004}) and the guidance theory of dark energy is risen on the cosmological constant, $\Lambda$ characterized by expulsive pressure which was initiated by Einstein in 1917 to achieve a static cosmological model. Later Zel'dovich (\cite{Zel'dovich1967}, \cite{Zel'dovich1968}) turned it as a vacuum energy of quantum fluctuation. However, for viability of the present-day accelerated Universe the earlier cosmological constant $\Lambda$, commonly, accepted it time-dependent in the cosmological domain (Perlmutter et al.~\cite{Perlmutter1998}; Riess et al.~\cite{Riess2004}). At the same time, space-dependent $\Lambda$ has an desired outcome in the astrophysical point of view as argued by other researchers (Chen \& Wu~\cite{Chen1990}; Narlikar et al.~\cite{Narlikar1991}; Ray \& Ray~\cite{Ray1993}) in respect to the behaviour of local massive objects kind of galaxies and elsewhere. In the present motto of compact stars, however, we take cosmological constant, $\Lambda$ as a absolutely constant quantity. This constancy of $\Lambda$ unable to ruled out for the scheme of very small dimension like as compact star systems or elsewhere with various physical needs (MaK~\cite{MaK2000}; Dymnikova~\cite{Dymnikova2002}; Dymnikova~\cite{Dymnikova2003}; B$\ddot{O}$hmer \& Harko~\cite{Bohmer2005}). To estimate mass and radii regarding neutron star Egeland (\cite{Egeland2007}) incorporated the presence of cosmological constant proportionality trust on the density of vacuum. Egeland have done it by application the Fermi equation of state along with the Tolman-Oppenheimer-Volkoff (TOV) equation. Voluntary by the above knowledge, we organize the presence of cosmological constant in a small scale to exercise the construction of compact stars in low-mass X-ray binaries namely Cyg X-2, V395 Carinae/2S 0921-630, XTE J2123-058, X1822-371 (V691 CrA), 4U 1820-30 and GR Mus (XB 1254-690) and attained to a finality that incorporation of $\Lambda$ tells the compact stars in good manners. | \label{sec:3} \begin{figure}[htbp] \centering \includegraphics[scale=.3]{Graph74.eps} \caption{Possible pressure ($p$)-density ($\rho$) relation (EOS) at the stellar interior taking a=0.0016 $km^{-2}$, C=1.133, where $\alpha, \beta, \eta, \delta, \xi$ are constants and all are in units of $km^{-2}$.} \label{fig:11} \end{figure} It is to be noted here that the model described by Heint IIa~(\cite{Heintzmann1969}) is useful to study both neutron and strange stars depending upon the choice of the metric parameter $a$, $C$ (Kalam et al.~\cite{Kalam2016}; Kalam et al.~\cite{Kalam2017}). In this article, we have investigated that whether the same Heint IIa metric is capable to explain the compact stars within low-mass X-ray binaries or not. For which, another demonstrated the physical behavior of the six compact stars within the low-mass X-ray binary (LMXB) namely Cyg X-2, V395 Carinae/2S 0921-630, XTE J2123-058, X1822-371 (V691 CrA), 4U 1820-30 and GR Mus (XB 1254-690) by considering isotropic pressure in nature. Here we have also merged the previously cosmological constant $\Lambda$ in the Einstein's field equation in favour of study the stellar construction. Effectively, we obtained an analytical solution fot the fluid sphere another really interesting attach to diverse physical property, which are as follows: \begin{enumerate} \item[(i)] In our model at the interior of the compact stars density and pressure well funtion (positive definite at the centre) (Fig.~1). It is to be mentioned here that pressure and density are both maximum at the origin and interestingly pressure fall to zero (monotonically decreasing) towards the boundary while density does not. Therefore it is justified to designated these compact stars as strange stars therein the surface density does not vanishes in place of the neutron stars dissimilar the surface density vanishes at the boundary. Here, we assume the values of constants ($a$, $C$) in the metric and $\Lambda$ in suchlike that pressure must dissolve at the boundary. By assuming of the constant's values $a$, $C$ and $\Lambda$, we calculate the central density, $\rho_{0}$ as $567 \times 10^{-6} km^{-2} (7.651 \times 10^{14} gm/cm^3)$ and central pressure, $p_{0}$ as $2224.51 \times 10^{-7} km^{-2} (5.557 \times 10^{35} dyne/cm^2)$ (Table~1). It satisfies energy conditions, TOV equation and Herrera's stability condition. It is also stable regard to infinitesimal radial perturbations. From mass function (equation~8), all desired inside properties of a compact star be possile to evaluated which satisfies Buchdahl mass-radius relation ($\frac{ 2M}{R} < \frac{8}{9}$) (Figs.~4, 5(left)). The surface redshift in respect of compact stars are found under the standard measure ($Z_{s}\leq 0.85$) that is favourable (Fig.~5(right)) (Haensel et al.~\cite{Haensel2000}). We estimated the EOS and that would be like $ p = \alpha e^{(-\rho/\beta )} + \eta e^{(-\rho/\delta )} + \xi $ whereinto $\alpha, \beta,~ \eta,~ \delta, ~\xi$ are constants and theirs unit of $km^{-2}$. Fig.~7 indicates that a stiff equation of state ($\ddot{O}$zel~(\cite{Ozel2006}); Lai \& Xu~(\cite{Lai2009}) and Guo et al.~(\cite{GUO2014})) rather be a soft equation of state. \item[(ii)] From our mass function graph Fig.~6, equation (9) and equation (10), we obtain the radii, compactness and surface red-shift of six compact stars within the low-mass X-ray binary (LMXB) as like Cyg X-2, V395 Carinae/2S 0921-630, XTE J2123-058, X1822-371 (V691 CrA), 4U 1820-30 and GR Mus (XB 1254-690). The detail comparison chart are shown in Table~2. \end{enumerate} It is to be mentioned here that we actually considering Heint IIa metric with de-Sitter spacetime to describe the compact stars within low-mass X-ray binaries where inlaid metric parameters $a$, $C$ are assess by computing all modes of necessary situations. When metric parameters values are known, the EOS additionally the central density are settled. In general, the mass-radius curve are considered under a conferred equation of state for different values of central density; with a definite value of the central density, the mass and radius of a compact star are settled. In spite of our model is diverse and theoretically attractive. According to our model, six compact stars within the low-mass X-ray binary (LMXB) namely Cyg X-2, V395 Carinae/2S 0921-630, XTE J2123-058, X1822-371 (V691 CrA), 4U 1820-30 and GR Mus (XB 1254-690) derive the identical values of $a$, $C$ and therefore the identical central density and the identical equation of state. Further interestingly in our stellar model, if we begin out of the center by a particular central density, the construction of a compact star be possible determined with preventing on any radius whereinto pressure arrive to zero. Therefore, our conclusion is that we may find useful relativistic model in the sake of compact stars within low-mass X-ray binaries by suitable choice of the values of the metric parameters $a$, $C$ in the metric given by Heint IIa~(\cite{Heintzmann1969}). | 18 | 8 | 1808.09960 |
1808 | 1808.10456_arXiv.txt | We present and apply a method to infer the mass of the Milky Way (MW) by comparing the dynamics of MW satellites to those of model satellites in the \eagle{} cosmological hydrodynamics simulations. A distribution function (DF) for galactic satellites is constructed from \eagle{} using specific angular momentum and specific energy, which are scaled so as to be independent of host halo mass. In this 2-dimensional space, the orbital properties of satellite galaxies vary according to the host halo mass. The halo mass can be inferred by calculating the likelihood that the observed satellite population is drawn from this DF. Our method is robustly calibrated on mock \eagle{} systems. We validate it by applying it to the completely independent suite of 30 \auriga{} high-resolution simulations of MW-like galaxies: the method accurately recovers their true mass and associated uncertainties. We then apply it to ten classical satellites of the MW with 6D phase-space measurements, including updated proper motions from the \textit{Gaia} satellite. The mass of the MW is estimated to be $M_{200}^{\textnormal{MW}}=1.04_{-0.14}^{+0.23}\times10^{12}M_{\odot}$ (68\% confidence limits). We combine our total mass estimate with recent mass estimates in the inner regions of the Galaxy to infer a halo concentration of $c_{200}^{\textnormal{MW}}=12.0^{+3.5}_{-2.3}$, which is higher than the concentration of typical $\sim 10^{12}M_\odot$ cold dark matter haloes. | The mass of the Milky Way (MW) is a fundamental astrophysical parameter. It is not only important for placing the MW in context within the general galaxy population, but it also plays a major role when trying to address some of the biggest mysteries of modern astrophysics and cosmology. The intricacies of galaxy formation are highly dependant on feedback and star formation processes, which undergo a crucial physical transition around the MW mass \citep[e.g.][]{Bower2017}. Apparent discrepancies with the standard $\Lambda$CDM model, such as the missing satellites \citep{Klypin1999,Moore1999} and the too-big-to-fail problems \citep{BK11_TBtF_MNRAS.415L..40B} depend strongly on the MW halo mass \citep[e.g.][]{Purcell2012,Wang12_MissingSat_MNRAS.424.2715W,VeraCiro2013, Cautun2014b}. In addition, tests of alternative warm dark matter models \citep{Kennedy2014:WDMParticlefromMWSat_MNRAS.442.2487K,Lovell2014:WarmDM_MNRAS.439..300L} are also subject to the total halo mass. Thus, by considering the cosmological context of the MW and its population of dwarf galaxy satellites, important inferences about large-scale cosmology can be made. With the recent \textit{Gaia} DR2 release \citep{GaiaDR22018arXiv180409365G}, we now have significantly more information than ever before about our galaxy, and are better placed to make progress on these problems. There have been many attempts to infer directly the MW mass through a variety of methods. The total MW mass is dominated by its dark matter (DM) halo, which cannot be observed directly. Instead, its properties must be inferred from the properties of luminous populations, such as the luminosity function of MW satellites \citep[mostly the Large and Small Magellanic Clouds, e.g.][]{Busha2011b,Gonzalez2013,Cautun2014MNRAS.445.2049C} and the kinematics of various dynamical tracers of the Galactic halo. The dynamics of halo tracers are mostly determined by the gravitational potential of the MW halo, and provide a key indirect probe of the total halo mass. Examples of halo tracers used for this purpose are satellite galaxies \citep[e.g.][]{WilkinsonEvans1999, Watkins2010}, globular clusters \citep[e.g.][]{Eadie2016, BinneyWong2017, Sohn2018, WatkinsMWGC2018arXiv180411348W}, halo stars \citep[e.g.][]{Xue2008,Deason12_BrokenDegenMNRAS.424L..44D, Kafle2012,Kafle2014}, high velocity stars \citep[e.g.][]{Smith2007, Piffl2014,Fragione17_HVSMWMassNewA...55...32F,Rossi2017,Monari2018} and stellar streams \citep[e.g.][]{Koposov2010, Newberg10_StreamMWMass_ApJ...711...32N, Gibbons2014, Kupper2015, Bowden2015}. There are a variety of methods for inferring the Galactic halo mass using dynamical tracers. A common approach is to model the tracers as distributions in equilibrium whose parameters are determined by fitting the model to observational data \citep[e.g.][]{Evans2003,Han2016a}. Advances in the calculation of action-angle coordinates \citep[e.g.][]{Vasiliev18_AGAMAarXiv180208239V} have led to a new generation of analytical galaxy modelling, centred around distribution functions (DFs) in action-angle space. Examples include modelling the MW population of globular clusters \citep[e.g.][]{PostiHelmi2018arXiv180501408P} or individual DFs of components such as the thick and thin disc, bulge, stellar halo and DM halo \citep{ColeBinneyCore2017MNRAS.465..798C}. The recent availability of large cosmological simulation has enabled a new class of methods based on comparing the observed properties of MW satellites to those of substructures in cosmological simulations \citep[e.g.][]{Busha2011ApJ...743...40B,Patel2017Orbits2_MNRAS.468.3428P}. Although over the past decades a large amount of effort has been dedicated to inferring the Galactic halo mass, its value remains uncertain to within a factor of two, with most mass estimates ranging from $0.5$ to $2.5\times10^{12}M_{\odot}$ \citep[e.g.][and our Fig. \ref{fig:LitComparison}]{Wang15_DMHaloDynTracers_MNRAS.453..377W}. While many studies claim uncertainties smaller than this range, the analytical models upon which they rely require several assumptions such as dynamical equilibrium and a given shape of the density or the velocity anisotropy profiles. These assumptions can lead to additional systematic errors, which are difficult to quantify but can be the dominant source of error \citep[e.g. see][]{Yencho2006, Wang15_DMHaloDynTracers_MNRAS.453..377W,Wang2018}. Furthermore, most methods typically estimate the mass within the inner tens of kiloparsecs, since this is the region where most tracers (such as halo stars and globular clusters) reside, necessitating an extrapolation to the virial radius. This extrapolation requires additional assumptions about the radial density profile of the MW and can lead to further systematic uncertainties. Large-volume high-resolution cosmological simulations offer a unique test-bed for analytical mass determination methods \citep[e.g.][]{Han2016b,Penarrubia2017,Wang2017} and, importantly, enable new methods for inferring the Galactic halo mass with a minimal set of assumptions. The simulations have the advantage of self-consistently capturing the complexities of halo and galaxy formation, as well as the effects of halo-to-halo variation. However, with a few exceptions, the limited mass resolution of current simulations means that they can resolve satellite galaxies but not halo stars or globular clusters \citep[although see e.g.][]{Pfeffer2018,Grand2018}. This is not a major limitation since satellite galaxies, due to their radially extended spatial distribution, are one of the best probes of the outer MW halo. This is especially true now that the \emph{Gaia} DR2 release has provided a large sample of MW satellites with full 6D phase space information \citep[][]{GaiaDR2_MotionsDwarfGC_2018arXiv180409381G, Fritz2018,Simon2018}. Galactic halo mass estimates that rely on cosmological simulations are relatively recent. \citet{Busha2011ApJ...743...40B} pioneered the approach of inferring halo properties by finding the best match between the MW satellites and satellites of simulated haloes. The MW mass is then determined by weighting the host haloes according to the quality of the satellite match, a technique known as importance sampling. \citeauthor{Busha2011ApJ...743...40B} used the distance, velocity and size of the Large and Small Magellanic Clouds to constrain the MW mass. The distance and velocity of satellites can vary rapidly, especially when close to the pericentre of their orbit, so very large simulations are needed in order to find enough counterparts to the MW system. \citet{Patel2017Orbits2_MNRAS.468.3428P} pointed out that approximately conserved quantities, such as angular momentum, are better for identifying satellite analogues in simulations. This makes it easier to find MW counterparts; applying the criterion to a larger number of satellites results in a more precise mass determination \citep{PatelMWEnsemble2018ApJ...857...78P}. A further advance was achieved by \citet{Li2017ApJ...850..116L} who showed that, when scaled appropriately, the DF of satellite energy and angular momentum becomes independent of halo mass. This scaling allows for a more efficient use of simulation data, since any halo can be rescaled to a different mass, and thus a better sampling of halo formation histories and halo-to-halo variation can be achieved. This approach represents a major improvement over importance sampling methods, in which the statistically relevant systems are those in a small mass range. In this paper we improve and extend the \citet{Li2017ApJ...850..116L} mass determination method. We start by constructing the phase-space distribution of satellite galaxies using a very large sample of host haloes taken from the \eagle{} (Evolution and Assembly of GaLaxies and their Environments) galaxy formation simulation \citep{EAGLE2015MNRAS.446..521S,Crain2015}. We then describe and calibrate three mass inference methods based on the satellite distributions of: i) angular momentum only, ii) energy only, and iii) a combination of both angular momentum and energy. We test these methods by applying them to an independent set of simulations, taken from the \auriga{} project \citep{Auriga_2017MNRAS.467..179G}; this is a very stringent test because of the much higher resolution and rather different galaxy formation model implemented in \auriga{} compared to \eagle{}. Finally, we apply our methods to the latest observations of the classical satellites to determine the MW halo mass; we are able to estimate this mass with an uncertainty of only $20\%$. The structure of the paper is as follows. Section \ref{sec:satellite_phase_space} describes the construction of the phase-space DFs using the \eagle{} data. Section~\ref{sec:method} describes our mass inference methods, their calibration and validation with tests on mock systems. In Section~\ref{sec:MW}, we apply this method to the observed MW system and discuss our results. Finally, Section~\ref{sec:conclusion} summarises and concludes the paper. | \label{sec:conclusion} We have developed a method to determine the total mass of the Milky Way (MW) dark matter (DM) halo by comparing the energy and angular momentum of MW satellites with the respective distributions predicted in the \eagle{} galaxy formation cosmological simulations. When scaled appropriately by host halo mass, the energy and angular momentum of the satellites become independent of the host halo mass (see Fig.~\ref{fig:appendix_scaled_E_L}). Thus, we can use a large sample of \eagle{} haloes, and associated satellites, in our estimate of the MW halo mass. For this, we constructed the satellite distribution function in $(E,L)$ space from the simulations and carried out a maximum likelihood analysis to infer the halo mass from the phase-space properties of the ten brightest satellite galaxies (excluding the disrupting Sagittarius galaxy). Using mock samples from \eagle{} we analysed the performance of the method and quantified its statistical and systematic uncertainties. A key test of our method was to apply it to estimate the masses of the DM haloes of 30 MW analogues simulated in the \auriga{} project. These simulations have much higher resolution and employ different baryonic physics models than \eagle. They produce realistic MW-like galaxies \citep{Auriga_2017MNRAS.467..179G,Grand2018} and thus provide a rigorous and completely independent external test of our method. We find that our method provides an unbiased estimate of the total halo masses of the \auriga{} galaxies, with a precision of $\sim 20\%$, in very good agreement with the expectations from the \eagle{} simulations. Our main conclusions are: \begin{itemize} \item Applying our method to ten classical MW satellites gives an estimate for the total mass of the MW halo of $M_{200}^{\textnormal{MW}}=1.04_{-0.14}^{+0.23}\times10^{12}M_{\odot}$. This result agrees well with most previous estimates in the literature but with a rigorously tested accuracy (${\sim}20\%$) which is better than most other estimates. \item Combining our total DM halo mass estimate with recent estimates of the halo mass within $20\kpc{}$ we infer a concentration for the MW halo of $C =12.0^{+3.5}_{-2.3}$. This is higher than predicted by the \eagle{} simulations, which include the contraction of the halo due to the formation of the galaxy at the centre. \eagle{} haloes with masses of $10^{12}M_\odot$ have a median concentration of 8.2, with only ${\sim}3\%$ of them having concentrations of 12 or higher. However, while this suggests that the MW is an outlier, the inferred MW halo concentration has large uncertainties, and the true value could potentially be lower and in better agreement with the \eagle{} haloes. \item Our halo mass estimate can be improved by increasing the number of halo tracers and/or reducing the observational uncertainties. We found that the observed proper motions of the ten classical satellites are already so precise that further improvement will make little difference to the halo mass estimate. Increasing the number of satellites, on the other hand, for example by including the ${\sim}50$ currently known satellites in the MW, would reduce the mass errors to ${\sim}14\%$. Further improvements would be possible by analysing all ${\sim}125$ satellites that are predicted to reside in the MW \citep{Newton18_MWSatPop_MNRAS.479.2853N}, which would result in a ${\sim}10\%$ mass uncertainty, a factor of two improvement over our current estimate. \end{itemize} \noindent In summary, our MW halo mass estimate is precise and accurate and has been thoroughly tested on realistic model galaxies and their satellite populations. Mass estimates that rely on cosmological simulations are relatively new but the use of simulations enables a robust and testable methodology. Indeed, the accuracy we are now able to achieve ($\sim 20\%$; see also \citealt{PatelMWEnsemble2018ApJ...857...78P}) is a significant step forward from the factor of two uncertainty that has plagued MW mass estimates for years. This theoretical boost, coupled with the exquisite 6 dimensional data that \textit{Gaia} and complementary facilities are now providing, brings us closer to what may be called the era of ``precision'' near-field cosmology --- when we can go beyond rough estimates of the MW halo mass and, instead, remove this important degree of freedom when making use of the properties of the MW to inform cosmological models and dark matter theories. | 18 | 8 | 1808.10456 |
1808 | 1808.07716_arXiv.txt | Transient ultraluminous X-ray sources (ULXs) provide an important link bridging transient low-mass X-ray binaries and ULXs. Here we report the first discovery of both a canonical sub-Eddington outburst and an ultraluminous super-Eddington outburst from an unusual transient ULX, NGC~7793~P9 with a variability factor higher than $10^3$. Its X-ray spectrum switches between the typical high/soft state and the steep power-law state during the canonical outburst. The inner radius of the accretion disk and the disk temperature--luminosity correlation suggest that P9 harbors a stellar-mass black hole (BH). At the beginning of the ultraluminous outburst, we observe a cool outer disk with a hard Comptonized spectrum, implying a transition to the ULX regime. When the luminosity increases to $L\gtrsim3\times10^{39}$\,\lumcgs, P9 shows a significantly curved spectrum that can be described by either a slim disk or a strongly curved Comptonized corona. The phenomenological model suggests that the hot disk observed near the peak of the ultraluminous outburst is coincidentally consistent with the extension of the thermal track. Utilizing more physical Comptonized disk models, we suggest that the corona cools down and the apparent disk-like spectrum is dominated by soft Comptonization. The significant variability above 1\,keV supports this two-component scenario. The spectral evolution can also be interpreted with the supercritical accretion model. All these indicate that a canonical black hole X-ray binary can show properties of a ULX. | Accreting black holes (BHs) are bright X-ray sources powered by the gravitational potential energy of accreting materials. The luminosities of Galactic BHs are strongly limited to the Eddington luminosity ($L_{\rm{Edd}}$), at which the radiation pressure balances the inward gravitational force of the accreting materials. Thanks to the revolution in astronomical instrumentation over the past several decades, the discovery of ultraluminous X-ray sources (ULXs) challenged the Eddington limit \citep[see reviews by][]{FrankKR2002, FengS2011, KaaretFR2017}. ULXs are extragalactic off-nuclear X-ray point sources with luminosities exceeding the Eddington limit of a stellar-mass BH. Most ULXs are believed to be powered by stellar-mass BHs with mild beaming and/or moderate super-Eddington accretion rates. They show ultraluminous states distinct from the canonical spectral states of BH X-ray binaries \citep[BHXBs, ][]{FengS2011, GladstoneRD2009, KaaretFR2017}. Unlike the spectral behavior of a BHXB in the steep power-law (SPL) state that shows a hot disk and an optically thin corona, a genuine ULX usually shows a much cooler disk with a larger inner radius plus an optically thick corona and has a luminosity of $L\gtrsim3\times10^{39}$\,\lumcgs. ULX states can be further divided into the hard ultraluminous (HUL), soft ultraluminous (SUL), and supersoft ultraluminous spectral states \citep{SuttonRM2013, KaaretFR2017}. In addition, those low-luminosity ULXs with $L\lesssim3\times10^{39}$\,\lumcgs usually show a curved spectrum that can be described with a broadened disk (BD) model, implying that the disk component is no longer a geometrically thin accretion disk but the luminosity remains sub-Eddington \citep{SuttonRM2013}. The best studied ULXs are persistent sources with low variability factors \citep{FengK2006,KaaretF2009}. A stable high mass accretion rate is necessary to interpret the long lifetime of the ultraluminous state. As long as the companion star of a ULX is a massive star with $M\gtrsim10M_\odot$, super-Eddington mass transfer occurs when the companion evolves to the Hertzsprung gap without a common envelope \citep{KingDW2001, RappaportPP2005, PavlovskiiIB2017}. Most of the currently known ULXs, including the ultraluminous pulsar P13 in NGC 7793, have early-type companions with high optical luminosities \citep{BachettiHW2014, MotchPS2014, LauKB2016}. It has also been proposed that some ULXs may exhibit a low-mass companion and be powered by a long-lasting super-Eddington outburst \citep{King2002}. In contrast, BHXBs with low-mass companions are usually transient X-ray sources. Therefore, transient ULXs is key to testing this scenario and showing that a low-mass X-ray binary (LMXB) can evolve to a ULX with typical UL spectral behaviors. The Galactic LMXBs GRS 1915+105 and XTE J1550$-$564 showed optically thick Comptonized spectra during the peak of the outburst with a luminosity of $\sim10^{39}$\,\lumcgs. They could fill in the link bridging BHXBs and ULXs \citep{KubotaD2004, Soria2007, Vierdayanti2010}. Several transient ULXs were reported as evolved from canonical BHXBs although they showed either canonical BHXB spectral states or ultraluminous outburst only \citep[see, e.g., ][]{MiddletonSR2012, KaaretF2013}. Therefore, it remains unclear whether a canonical BHXB showing sub-Eddington outbursts can evolve into a ULX with a much higher luminosity and typical UL spectral behaviors. In this paper, we report the discovery of a transient ULX, NGC~7793~P9 (hereafter P9), that showed both a canonical BHXB outburst and an ultraluminous outburst that fill in the gap. NGC~7793 is an SA(s)d type galaxy in the Sculptor Group with a distance of 3.6--3.9\,Mpc \citep{KarachentsevGS2003, Crowther2013, Radburn-SmithJS2011, TullyCS2016}. It has a bulge with filamentary spiral structures and several H{\sc ii} regions \citep{ReadP1999}. P9 is located at the northeast edge of one filament extending from the host galaxy and has no obvious association with any H{\sc ii} regions or supernova remnants \citep{ReadP1999}. P9 was first characterized as a highly variable point source in a \rosat\ survey \citep{ReadP1999}. The X-ray spectrum of P9 was well fit with either a PL or a thermal bremsstrahlung model. The absorption column density ($N_{\textrm{H}}$) was rather large. A deep \chandra\ observation (ObsID \dataset[3954]{http://cda.harvard.edu/chaser/viewerContents.do?obsid=3954}) revealed a faint source, CXOU J235808.7$-$323403, at the position of P9, but its spectrum was not well characterized \citep{PannutiSF2011}. Therefore, the luminosity was not well constrained. By simply assuming a PL with a photon index $\Gamma=1.5$ and $N_{\textrm{H}}=1.15\times10^{20}$\,cm$^{-2}$, the luminosity of P9 was estimated to be $4.5 \times10^{37}$\,\lumcgs. No optical or radio counterparts have been found \citep{ReadP1999, PannutiSF2011}. We report a detailed analysis of the transient ULX P9 based on the \emph{Neil Gehrels Swift} (hereafter \swift) monitoring data and several dedicated observations made with \chandra, \xmm, and \nustar. We introduce the observation and data analysis in Section \ref{observation}. The detailed spectral analysis, including the variability of $\Gamma$, the stacked \swift\ spectroscopy, and tests of more physical models during the ultraluminous outburst are presented in Section \ref{result}. We further discuss the implications of the spectral variability, especially the evolution of the disk component and the accretion models, in the super-Eddington accretion regime in Section \ref{discussion}. Finally, we summarize this work in Section \ref{summary}. | \label{discussion} \begin{deluxetable*}{ccccccccc} \tabletypesize{\footnotesize} \tablecaption{Best-fit Spectral Parameters of Stacked \swift\ Spectra During the Ultraluminous Outburst. } \label{swift_parameter} \tablehead{\colhead{Bins} & \colhead{Luminosity Range} & \colhead{Model} & \colhead{$L_{\textrm{0.3--10 keV}}^a$} & \colhead{$L_{\textrm{MCD}}$} &\colhead{$R_{\rm{in}}\sqrt{\cos\theta}$} &\colhead{$kT_{\rm{in}}$} & \colhead{$\Gamma$ or $p$} & \colhead{$C^2/\rm{dof}$} \\ \colhead{} & \colhead{($10^{39}$\,\lumcgs)} & \colhead{} & \colhead{($10^{39}$\,\lumcgs)} & \colhead{($10^{39}$\,\lumcgs)} & \colhead{(km)} & \colhead{(keV)} & \colhead{} & \colhead{} } \startdata S1 & $L<1$, $\Gamma<1.8$ & MCD & $0.8\pm0.1$ & $0.8\pm0.1$ & $33^{+11}_{-9}$ & $1.3^{+0.3}_{-0.2}$ & \nodata & $143.9/167$ \\ S2 & $L<1$, $\Gamma>1.8$ & MCD+PL & $0.86\pm0.8$ & $0.3\pm0.1$ & $150_{-100}^{+120}$ & $0.5_{-0.1}^{+0.3}$ & $1.9(5)$ & $224.8/254$ \\ & & DiskPBB & & \nodata & $<48$ & $1.1_{-0.2}^{+0.4}$ & $0.53_{-0.03}^{+0.04\,b}$ & $223.2/255$ \\ S3 & $L=1$--1.3 & MCD+PL & $1.2\pm0.8$ & $0.4\pm0.2$ & $300_{-150}^{+180}$ & $0.39_{-0.08}^{+0.12}$ & $1.4_{-0.9}^{+0.5}$ & $149.1/184$ \\ & & DiskPBB & & \nodata & $<25$ & $1.3_{-0.4}^{+1.5}$ & $0.53_{-0.04}^{+0.05\,b}$ & $154.3/184$ \\ S4 & $L=1.3$--2 & MCD+PL & $1.5\pm0.2$ & $0.6_{-0.3}^{+0.5}$ & $150_{-80}^{+90}$ & $0.6_{-0.1}^{+0.2}$ & $1.5_{-0.9}^{+0.4}$ & $258.3/267$ \\ & & DiskPBB & & \nodata & $23_{-21}^{+18}$ & $1.4_{-0.3}^{+0.8}$ & $0.56_{-0.03}^{+0.04\,b}$ & $264.8/267$ \\ S5 & $L=2$--3 & MCD+PL & $2.2\pm0.2$ & $1.7_{-0.6}^{+0.3}$ & $46^{+14}_{-13}$ & $1.4^{+0.2}_{-0.3}$ & $2.3^{+6.2}_{-1.2}$ & $280.0/324$ \\ & & DiskPBB & & \nodata & $34^{+20}_{-17}$ & $1.5^{+0.4}_{-0.2}$ & $0.65^{+0.07\,b}_{-0.05}$ & $279.8/325$ \\ S6 & $L>3$ & MCD & $3.6_{-0.3}^{+0.4}$ & $3.6_{-0.3}^{+0.4}$ & $56^{+11}_{-10}$ & $1.5^{+0.2}_{-0.1}$ & \nodata & $253.5/312$ \enddata \tablenotetext{a}{Derived from the best-fit model.} \tablenotetext{b}{This is the $p$-value of the $p$-free disk.} \vspace{-0.9cm} \end{deluxetable*} \begin{figure} \includegraphics[width=0.47\textwidth]{swift_eeufspec.eps} \caption{\swift\ stacked spectra and the best-fit MCD+PL model P9. } \label{swift_spec_state} \end{figure} \subsection{Implications from the MCD+PL Model} \subsubsection{Spectral Evolution of P9} During the canonical outburst, the X-ray spectrum of P9 switched between an MCD and a PL. The MCD spectra were observed in mid 2011 and mid 2012. The disk had an apparent inner radius of $R_{\rm{in}}\sqrt{\cos\theta}\approx50$\,km and an inner-disk temperature $kT_{\rm{in}}=0.7$--0.9\,keV. This radius corresponds to a true inner radius of 90\,km after applying corrections for the color–temperature and inner boundary condition of an MCD and assuming an inclination angle of 60$^{\circ}$ \citep{ShimuraT1995, KubotaTM1998}. Assuming a non-spinning BH, the mass can be estimated to be $\sim10$\,$M_\odot$. For an extremely spinning BH, the apparent value corresponds to a BH mass of $\sim60$\,$M_\odot$. In addition, the disk could be weakly Comptonized and result in a smaller $R_{\rm{in}}$ \citep{KubotaM2004}. The above mass estimation represents the lower limit of the BH mass in P9. This indicates that P9 has a regular or massive stellar-mass BH. Its X-ray spectrum switched between the SPL and high/soft states of a canonical outburst before 2013 \citep{TetarenkoSH2016}. When the luminosity exceeds $10^{39}$\,\lumcgs, $\Gamma$ shows a clear decreasing trend with the luminosity (see Figure \ref{pl_index_all}), and the high-quality \xmm\ spectra in this ultraluminous state cannot be well fit by single-component models. They are significantly different from those in the canonical outburst. When the luminosity is slightly higher than $L\approx10^{39}$\,\lumcgs\ (X3, S3, and S4), P9 has a cool disk and a hard PL tail. Moreover, the spectral curvature increases with flux when $L\gtrsim2\times10^{39}$\,\lumcgs\ (X4, S5, S6). These data sets can be described by disk-like spectral models. As described in Section \ref{introduction}, the BD usually dominates the spectra of ULXs with $L\lesssim3\times10^{39}$\,\lumcgs, and those of ULXs with $L\gtrsim3\times10^{39}$\,\lumcgs\ showed two-component spectra like HUL and SUL \citep{GladstoneRD2009, SuttonRM2013}. However, some ULXs like NGC 1313 X-2 and Holmberg IX X-1 showed an inverse transition, i.e., an HUL spectrum in the low-luminosity state, and a disk-like BD spectrum near their peak luminosities \citep{PintoreZ2012, MiddletonHP2015, LuangtipRD2016}. The BD feature of these sources is likely to be misclassified due to the strong spectral curvature caused by the highly optically thick Comptonized spectrum \citep{SuttonRM2013}. Similar to these two sources, P9 has an apparent inverse transition between these two states. However, the X4 light curve shows a significant variability dominated by the emission above 1~keV ($F_{\rm{var}}=0.15(2)$ for 1--10\,keV and $F_{\rm{var}}=0.05(2)$ for 0.3--1\,keV). It is difficult to interpret with a single slim disk model since the variability should be independent of the energy for a single emission component. Hence, we suggest that the apparent BD spectral shape observed in X4 consists of two components like NGC 1313 X-2 and Holmberg IX X-1. \subsubsection{Variability of the Disk Component} P9 evolved from the canonical BHXB regime to the ULX regime when its luminosity crossed the Eddington limit of a $\sim10$\,\ms\ BH. To obtain the disk evolution, we calculated the 0.3--10\,keV luminosity of the disk to check its correlation with the inner-disk temperature. Assuming a fixed inner radius, the luminosity of a standard disk shows a positive correlation with the disk temperature: \begin{equation} L_{\rm{MCD}}=\frac{\pi\sigma G^2M^2T_{\rm{in}}^4}{6c_0^4f_{\rm{col}}^4c^4}, \end{equation} where $\sigma$ is the Stefan--Boltzmann constant, $M$ is the mass of the central compact object, $c_0\approx0.1067$, and $f_{\rm{col}}^4$ is the color--temperature correction, which is usually larger than unity \citep{GierlinskiZP1999, GierlinskiD2004}. Many ULXs (especially supersoft ones) show an anticorrelation between the disk luminosity and the temperature \citep[see, e.g.,][]{KajavaP2009, UrquhartS2016}. On the other hand, some ULXs show a positive correlation close to $L_{\rm{disk}}\propto T^4$ \citep{MillerFM2003}. Since P9 has both canonical and ultraluminous outbursts, it provides a unique chance to test the accretion disk model. We collected the 0.3--10\,keV luminosity of the MCD component and disk temperature from both the high/soft state (C2, X1, and S1), HUL state (X3, S3, S4, and possibly S2), and disk-like state in the ultraluminous outburst (X4, S5, and S6). We show the disk luminosity--temperature relation of P9 together with that of a BHXB, XTE J1550$-$564, in Figure \ref{kt_lt_1550}. The disk configuration of P9 during the high/soft state in the canonical outburst (X1 and C2) was fully consistent with the thermal track of XTE J1550$-$564. We attempted to add an NTHCOMP component to these two data sets but the Componized component contributed negligible emission. The disk component in the disk-like spectral states of P9 during the ultraluminous outburst well lies on the extension of the thermal track. This is likely a coincidence since they can also be well described by Comptonized models. The MCD component in the HUL state clustering at the lower-left corner of this figure is similar to the ultraluminous branch of XTE J1550$-$564. However, P9 shows a much harder ($\Gamma\sim$1.6--1.8) PL tail compared to XTE J1550$-$564 ($\Gamma\sim$2.5--2.8). The presence of this state and the apparent disk-like state makes P9 as a unique source among canonical BHXBs. \begin{figure} \begin{center} \includegraphics[width=0.5\textwidth]{kT_LT_1550.eps} \end{center} \caption{Comparison of the $kT_{\rm{in}}$--$L_{\rm{MCD}}$ relation between P9 and XTE J1550$-$564 \citep{Soria2007}. The red dashed line is the best-fit $L_{\rm{MCD}}\propto T_{\rm{in}}^4$ relation for XTE J1550$-$564. The red squares are obtained from the high/soft state of XTE J1550$-$564 and the green ones are from the SPL state and ultraluminous branch. The black data points denote the thermal and BD states of P9, and the blue ones denote the HUL spectral state. The black circle with the dashed line is the best-fit thermal component of X4 spectra with a two-component model containing an MCD and a PL with a high-energy cutoff. } \label{kt_lt_1550} \end{figure} \subsection{Implications from the Comptonized Disk} The simple MCD+PL model is helpful for classifying the spectral states. The result suggests that the innermost region of the accretion disk was covered by the Comptonized corona. We further used physical models to deal with the coupling of both the disk and the corona. We noticed that all of the spectra with luminosities above $L\gtrsim10^{39}$\,\lumcgs, including X3, X4, and S3--S6, show clear spectral difference above $\sim$1\,keV (see Figures \ref{chandra_xmm_spec} and \ref{swift_spec_state}). This implies that the mass accretion rate during the ultraluminous outburst has no significant difference, and the observed variability over (1--5)$\times10^{39}$\,\lumcgs\ is dominated by the change of the Comptonized component. If we use NTHCOMP to describe the Comptonized corona, then a cool inner-disk temperature, a large re-estimated radius, a hard PL with $\Gamma\sim1.8$, and a large optical depth are necessary for both X3 and X4 spectra. In contrast, X4 shows a higher normalization and a clear cutoff at 4--6\,keV. This implies that the plasma temperature of the corona is hot at the beginning of the ultraluminous outburst, and a huge amount of disk photons are scattered to hard X-ray. As the corona temperature cools down, fewer photons can be scattered to hard X-rays. The spectrum shows a cutoff in the \xmm\ energy range and an apparently higher luminosity in 0.3--10\,keV. We also described the UL spectra with the OPTXAGNF model. The BH mass can be estimated as $\sim20$\,$M_\odot$ for a non-spinning BH or $\sim120$\,$M_\odot$ for an extremely fast spinning BH. This result is a bit higher than that obtained from the simple estimation based on the apparent inner-disk radius in the high/soft state. In the case of $M_{\textrm{BH}}\sim20M_\odot$, most of the disk emission below $R_{\rm{cor}}$ is scattered into the hard Comptonized component ($f_{\rm{pl}}\sim1$) and produces a hard PL with $\Gamma=1.6$ in X3 spectrum. When the X-ray luminosity increases to near the peak of the ultraluminous outburst, a significant power is emitted in the low-temperature soft Comptonization component. This result is consistent with the cooling of the corona, which is implied from the NTHCOMP model. The luminosity is close to or slightly higher than the Eddington luminosity, which suggests that P9 belongs to the major ULX population: the high-luminosity tail of BHXBs. Of course, we could not exclude the explanation that X4 shows a slim disk since the spectrum can be well fit by both the DiskPBB model and the slim disk model. If this is true, the corona just fades away when the luminosity is close to the peak. However, the short-term variability is only observed above 1\,keV in the X4 observation. Therefore, the apparent BD spectrum likely contains two components. A BH with a mass of $\sim120M_\odot$ can also produce the observed spectra with a smaller coronal radius (in units of $r_g$). The luminosity is roughly 0.2\,$L_{\rm{Edd}}$, consistent with that of canonical outbursts in BHXBs. If this is true, P9 is consistent with a massive stellar-mass BH that was recently found via gravitational wave observations \citep{AbbottAA2016, AbbottAA2017}. However, the hard PL tail observed in X3 is not expected in canonical outbursts. Hence, this scenario is less favored. \subsection{Supercritical Accretion} The UL spectra can also be physically interpreted with the supercritical accretion model. The inner disk is geometrically thick and an optically thick outflowing wind is blown. A funnel-shaped cavity forms along the rotational axis of the BH \citep{KingDW2001, OhsugaM2007, MiddletonHP2015}. In this case, the soft emission component is contributed by the wind, and the hard emission component is from the inner accretion disk or the accretion flow, although detailed modeling is still under development \citep{PoutanenLF2007, KajavaPF2012}. The beaming factor was proportional to the mass accretion rate. Hence, the inclination angle is expected to be low since the hard component dominates the observed flux during the HUL state. The mass accretion rate is maximum during the HUL state (X3), where the beaming factor is large and we observed the innermost accretion flow. When the mass accretion rate decreased, the inner-disk temperature slightly cooled down and the beaming factor decreased. The resulting spectrum is best described by a BD if the emission from the wind becomes less dominant. As the disk luminosity drops to sub-Eddington levels, we observe a thin disk in the tail of the outburst. We attempted to fit the hard component with DiskPBB in the X3 data set, but the disk temperature is extremely high ($kT\sim4$\,keV) and the disk radius is not well constrained. Therefore, this scenario is less favored although not entirely excluded. Of course, a corona can possibly form around the supercritical accretion disk and cause a much more complex X-ray spectrum. Finally, P9 shows the HUL spectrum when $L\approx10^{39}$\,\lumcgs, much lower than several genuine ULXs showing the same spectral behavior \citep[e.g., 1--1.2$\times10^{40}$\lumcgs\ for Holmberg IX X-1, ][]{SuttonRM2013}. This could imply that those genuine ULXs have much higher BH masses if the spectral state classification only depends on the Eddington ratio. However, the luminosity of a ULX can be strongly affected by the beaming factor and viewing angle. An important counterexample is NGC 5907 ULX-1, which shows an HUL spectral behavior at $L\gtrsim10^{40}$\,\lumcgs. Recently, it was found to be powered by a neutron star \citep{IsraelBS2017}. Therefore, the spectral state and the mass of the central compact object cannot be directly associated. The empirical luminosity range of the spectral states may stop being valid as the ULX sample grows. Comprehensively analyzing the timing and spectral behaviors of ULXs with next-generation X-ray telescopes and improving the supercritical accretion spectral model will be helpful for describing the accretion disk structure of ULXs. | 18 | 8 | 1808.07716 |
1808 | 1808.10510_arXiv.txt | We present a combined, homogenized analysis of archival Submillimeter Array (SMA) and Atacama Large Millimeter/submillimeter Array (ALMA) observations of the spatially resolved 340\,GHz (870\,$\mu$m) continuum emission from 105 nearby protoplanetary disks. Building on the previous SMA survey, we infer surface brightness profiles using a simple model of the observed visibilities to derive the luminosities ($L_{\rm mm}$) and effective sizes ($R_{\rm eff}$) of the continuum emission. With this sample, we confirm the shapes, normalizations, and dispersions for the strong correlations between $L_{\rm mm}$, $M_\ast$ (or $L_\ast$), and $\dot{M}_\ast$ found in previous studies. Moreover, we verify the continuum size--luminosity relation determined from the SMA survey alone (extending to an order of magnitude lower $L_{\rm mm}$), demonstrating that the amount of emission scales linearly with the emitting surface area. Moreover, we identify new, although weaker, relationships between $R_{\rm eff}$ and the host and accretion properties, such that disks are larger around more massive hosts with higher accretion rates. We explore these inter-related demographic properties with some highly simplified approximations. These multi-dimensional relationships can be explained if the emission is optically thick with a filling factor of $\sim$0.3, or if the emission is optically thin and disks have roughly the same optical depth profile shapes and normalizations independent of host properties. In both scenarios, we require the dust disk sizes to have a slightly sub-linear relationship with the host mass and a non-negligible dispersion ($\sim$0.2\,dex at a given $M_\ast$). | } The physical processes most relevant to the formation and early evolution of planetary systems in the disks orbiting young stars are remarkably complex and, at present, only weakly constrained by the data. Nevertheless, the prospect that the most dominant of these processes will imprint deterministic signatures in the ``bulk" characteristics of the disk population serves as a strong motivation for demographics studies \citep[e.g.,][]{ida04,ida08b,mordasini09,mordasini12}. Early successes in such population-oriented research were manifested in basic evolutionary trends. Surveys found that the infrared excess fractions \citep{haisch01,hernandez07,richert18} and accretion rates ($\dot{M_\ast}$; \citealt{muzerolle00,sicilia-aguilar05}) in different clusters decreased with stellar host age ($\tau_\ast$), indicating that gas and dust in the innermost regions of disks are (exponentially) depleted over a (half-life) timescale of a few Myr \citep{mamajek09,fedele10}. Those same metrics also depend on the stellar host mass ($M_\ast$), such that (at a given age) warm dust is more depleted and gas accretes faster in the disks around more massive stars \citep[e.g.,][]{carpenter06,muzerolle03,manara17}. Such behavior is broadly consistent with predictions for the viscous evolution of disk structures \citep{hartmann98}, but these specific diagnostics are not obviously associated with the more {\it global} disk properties that are most relevant for testing planet formation models. The global property of most interest is the disk mass ($M_d$). By far, the most accessible estimates of $M_d$ for large-scale demographics studies come from the (presumed) optically thin millimeter-wavelength continuum emitted by dust grains in the disk \citep{beckwith90}. Millimeter continuum luminosities ($L_{\rm mm}$) are expected to scale roughly linearly with $M_d$, modulo assumptions about the dust opacities ($\kappa_\nu$), dust-to-gas ratio, and temperature ($T_d$). In population studies of several young ($\sim$1--3\,Myr) clusters, $L_{\rm mm}$ is found to have a steep dependence on $M_\ast$ that has been interpreted as a linear or super-linear scaling between $M_d$ and $M_\ast$ \citep{andrews13,ansdell16,pascucci16}. Similar surveys in older clusters demonstrate that this relationship evolves (steepens) on few Myr timescales \citep{barenfeld16,ansdell17,ruiz-rodriguez18}. Building on these results, \citet{manara16} also identified a tentative $M_d$ -- $\dot{M_\ast}$ relationship. Taken together, these studies affirm the promise of demographics to reveal fundamental links between the bulk properties of disks that are relevant to planet formation \citep[e.g.,][]{mulders17}. They also confirm that the general disk population is rife with high-dimensionality connections: so far, there is evidence that dynamics and/or heating (via $M_\ast$), evolution ($\tau_\ast$), accretion ($\dot{M_\ast}$), and mass ($M_d$) are all related. That said, the substantial amount of scatter around these confirmed relationships hints at other parameters of interest. Some of that scatter, and perhaps some of those relationships, might be associated with the ways in which mass is spatially distributed in these disks \citep[e.g.,][]{hartmann98,rafikov17}. The relevant observable that is a complement to $L_{\rm mm}$ ($\sim$$M_d$) and suitable for demographics studies in that case is the spatial extent, or {\it size}, of the mm continuum emission. \citet{andrews10b} first tentatively identified a positive trend between disk sizes and masses inferred from mm continuum data \citep[see also][]{pietu14}, which was robustly verified for a much larger sample by \citet{tripathi17}. Recent work with other samples has broadly validated and extended this size--luminosity correlation \citep[albeit with somewhat different definitions of ``size";][]{tazzari17,barenfeld17}. One challenge in linking the sizes with the other relevant parameters has been that the \citet{tripathi17} sample has a comparatively heterogeneous construction: it was primarily based on availability in the Submillimeter Array (SMA) archive, and thereby comprises smaller collections of disks in different cluster environments. In this article, we aim to alleviate some potential sample biases and re-examine the mm continuum size--luminosity relationship for a large and well-defined collection of disks in the Lupus star-forming region \citep{ansdell16}, with the goal of facilitating a more holistic look at its links to the larger disk demographics landscape. Section~\ref{sec:data} describes the associated data, collected from the Atacama Large Millimeter/submillimeter Array (ALMA) archive. Section~\ref{sec:sample} describes the sample. Section~\ref{sec:analysis} explains the size measurements, and Section~\ref{sec:results} quantifies some relationships with other parameters. Finally, Section~\ref{sec:discussion} considers potential origins for these connections in the context of the general disk population. | } We have used archival ALMA and SMA observations to measure the sizes ($R_{\rm eff}$) and luminosities ($L_{\rm mm}$) of the 340\,GHz continuum emission from 105 nearby protoplanetary disks. After collating measurements of the effective temperatures, luminosities, and accretion luminosities of the corresponding stellar hosts from the literature, we computed a homogenized set of host masses, ages, and accretion rates and then examined the multi-dimensional relationships among these basic demographic parameters. With this analysis, we confirm the $L_\ast$--$L_{\rm mm}$ (or, equivalently, $M_\ast$--$L_{\rm mm}$) scaling relations inferred previously in various $\sim$Myr-old young clusters \citep{andrews13,ansdell16,pascucci16}, the recent claims of a $L_{\rm mm}$--$\dot{M}_\ast$ correlation \citep{manara16,mulders17}, and the classical $M_\ast$--$\dot{M}_\ast$ scaling \citep[e.g.,][]{muzerolle03,manara15}. Moreover, we verify the $L_{\rm mm}$--$R_{\rm eff}$ scaling inferred by \citet{tripathi17} and present corresponding, albeit weaker, relationships between $R_{\rm eff}$ and the host and accretion parameters. Although we have now identified and quantified the higher dimensional connections between these properties (Table~\ref{table:regressions}), the origins of these relationships, and particularly the lingering scatter around them, are not yet clear. Such ambiguity will likely remain until larger samples with considerably better angular resolution and sensitivity are available. There are simply too many forking paths to generalize with the limited information at hand, since the detailed distribution of the continuum emission may indeed provide the hidden links to the demographic connections studied here. Nevertheless, we hope to illuminate some broadly relevant issues by exploring these parameter relationships in highly simplified scenarios. With that more modest goal in mind, we will explore two limiting approximations. First, we need to design a population model for the host properties in this sample to enable a more general interpretation of their connections to the disk parameters. Figure~\ref{fig:stars_covar} summarizes that model characterization, showing probability densities in grayscale with the mean behaviors in red, and compares it to the individual measurements. This host population model was manually designed and tuned as described in Appendix~\ref{appendixB}. Briefly, it employs a three-part linear model of (visual, though not necessarily physical) sub-populations in \{$\log{T_{\rm eff}}$, $\log{\tau_\ast}$\}-space (bottom left corner of Figure~\ref{fig:stars_covar}). And second, we make two assumptions that will enable a more intuitive (both physically and mathematically) exploration. The first is that the continuum emission distribution has a distinct outer boundary, $R_o$. This rigidity compared to the surface brightness models used in Section~\ref{sec:analysis} is solely designed to simplify the interpretations and focus on the basic results. In the parlance of the ``Nuker" profile models, it is worth noting that the vast majority of the disks in this sample have $\lesssim$20\%\ of their total emission originating outside $R_t$ (i.e., crudely approximated here as $R_o$), due to their relatively large $\beta$ indices. We will briefly touch on the effects of ignoring this component of the emission profile below. The second assumption is that the radial dust temperature profile has a simple parameterization, \begin{equation} T_d(r) = T_0 \, \left(\frac{L_\ast}{L_\odot}\right)^{0.25} \, \left(\frac{r}{r_0}\right)^{-q} \, , \label{eq:tdust} \end{equation} with a fiducial normalization of $T_0 = 30$\,K at $r_0 = 10$\,au for $L_\ast = L_\odot$. This behavior is motivated by the (admittedly coarse) radiative transfer grid calculated by \citet{andrews13}, and is roughly consistent with measurements that locate the CO condensation fronts at $\sim$20\,K in the disks around two benchmark sample members, TW Hya \citep{qi13} and HD 163296 \citep{qi15} for reasonable values of $q$. \subsection{Optically Thick Scenario} \label{subsec:thick} The first scenario we will consider is the case where all of the emission is optically thick ($\tau_d \gg 1$). We define a corresponding brightness profile, \begin{equation} I_\nu(r) \approx \begin{cases} \mathcal{F} \, B_\nu(T_d) \, & \, \text{if $r \le R_o$} \, .\\ 0 \, & \, \text{otherwise} \, , \end{cases} \end{equation} where $\mathcal{F}$ is a constant, tunable, {\it intensity-weighted} ``filling factor" ($0 < \mathcal{F} \le 1$). To simplify the discussion, let us temporarily assume the disks in this sample can be described in the Rayleigh-Jeans approximation, so that \begin{equation} I_\nu(r) \approx \frac{2 \, \nu^2 \, k}{c^2} \, \mathcal{F} \, T_0 \, r_0^q \, \left(\frac{L_\ast}{L_\odot}\right)^{0.25} \, r^{-q} \end{equation} inside $R_o$. Then, following Equation~\ref{eq:fcum}, the cumulative intensity profile is \begin{equation} f_\nu(r) = \frac{4 \pi \, \nu^2 \, k}{(2-q) \, c^2} \, \mathcal{F} \, T_0 \, r_0^q \, \left(\frac{L_\ast}{L_\odot}\right)^{0.25} \, r^{2-q} \, , \end{equation} and, recalling the definitions of the key observables $L_{\rm mm} = f_\nu(R_o)$ and $x L_{\rm mm} = f_\nu(R_{\rm eff})$, we can relate the continuum luminosity and effective size to the model, \begin{subequations} \begin{equation} L_{\rm mm} = \frac{4 \pi \, \nu^2 \, k}{(2-q) \, c^2} \, \mathcal{F} \, T_0 \, r_0^q \, \left(\frac{L_\ast}{L_\odot}\right)^{0.25} \, R_o^{2-q} \, , \label{eq:Lmm_Rt_thick} \end{equation} \begin{equation} R_{\rm eff} = x^{1/(2-q)} \, R_o \, . \label{eq:Reff_Rt_thick} \end{equation} \label{eq:base_thick} \end{subequations} These are the ``base" relations that connect the disk observables \{$L_{\rm mm}$, $R_{\rm eff}$\} to the stellar population (through $L_\ast$, via the population model shown in Figure~\ref{fig:stars_covar}) and the model parameters \{$q$, $\mathcal{F}$, $R_o$\}. \begin{figure*}[t!] \includegraphics[width=\linewidth]{multid_thick.pdf} \figcaption{The relationships between disk continuum and stellar host parameters for the joint sample (points, symbols as in Figure~\ref{fig:hrd}. The red curves describe the mean behavior of a population model that assumes all of the emission is optically thick, with the assumptions described in Equation~\ref{eq:thickmodel}. The grayscale is a probability density map for that mean population model which folds in some appropriate scatter in $T_0$, $q$, $\mathcal{F}$, and the prescribed $R_o$--$M_\ast$ scaling, along with the underlying stellar population model. \label{fig:thick_covar} } \end{figure*} Combining Equation~\ref{eq:base_thick} and re-arranging terms, we find \begin{equation} R_{\rm eff} \propto L_\ast^{-1/4(2-q)} \, L_{\rm mm}^{1/(2-q)} \, . \label{eq:thick_link} \end{equation} Therefore, in this scenario the {\it shapes} of both the $R_{\rm eff}$--$L_{\rm mm}$ and $L_\ast$--$L_{\rm mm}$ scaling relations derived in Section~\ref{subsec:scalings} (see Table~\ref{table:regressions}) are reproduced when $q = 0.57\pm0.08$, a reasonable temperature gradient for irradiated accretion disks \citep[e.g.,][]{dalessio98}. Going back to the base relations in Equation~\ref{eq:base_thick}, we can re-cast the requirement to reproduce these scaling relations in terms of the model parameters and find that $R_o \propto L_\ast^{0.4}$, or, for the mean $M_\ast$--$L_\ast$ relationship, $R_o \propto M_\ast^{0.7}$. The {\it normalization} of that $R_o$--$M_\ast$ relation needs to be chosen to reproduce the observed $R_{\rm eff}$--$M_\ast$ behavior (or its equivalents; see Figure~\ref{fig:lmmreff_hosts}). However, for the fiducial $T_d$ normalization and a continuous distribution of emitting dust ($\mathcal{F} = 1$), this model scenario generates too much emission (i.e., $L_{\rm mm}$ is too high for any given $L_\ast$, $M_\ast$, or $R_{\rm eff}$). The magnitude of the discrepancy is about a factor of six in $L_{\rm mm}$, although roughly half of that is a result of presuming the Rayleigh-Jeans approximation is valid; a more accurate discrepancy is a factor of three. The only ways to reconcile this offset are to make the disks colder (e.g., decrease $T_0$ from $\sim$30 to 10\,K for solar luminosity hosts), to selectively remove some emitting material (i.e., set $\mathcal{F} \approx 0.3$), or some combination of those options. The first option alone is difficult to achieve for realistic radiation transfer calculations, so we prefer the second. There are many ways to distribute the emitting material to fulfill that criterion, but from the demographics perspective considered here the arrangement itself is not important. Figure~\ref{fig:thick_covar} compares the observed relationships between \{$L_{\rm mm}$, $R_{\rm eff}$, $L_\ast$, $M_\ast$\} for the joint sample with a population model of the optically thick scenario described above (using the full Planck function instead of the Rayleigh-Jeans approximation). The mean behavior in this multi-dimensional parameter-space can be explained well with three key assumptions: \begin{equation} \begin{cases} \,\, q \approx 0.57 \, , \\ \,\, R_o \approx 90 \, (M_\ast/M_\odot)^{0.7} \, \text{au} \, , \\ \,\, \mathcal{F} \approx 0.3 \, . \end{cases} \label{eq:thickmodel} \end{equation} The scatter around these mean relationships has modest contribution ($\sim$30--50\%) from the dispersion in the stellar host parameters. Neither those variations, nor the additional scatter imposed by reasonable uncertainty in $q$ ($\pm$0.08) or even in $T_0$ ($\pm$5\,K), nor any variations in $\mathcal{F}$, provide an explanation for the broad distribution observed in the $M_\ast$--$R_{\rm eff}$ plane. That scatter must be explained in the assumed $R_o$--$M_\ast$ behavior: we find that {\it all} of the observed scatter can be explained if its normalization has a log-normal variation around the mean with a standard deviation of 0.2\,dex and its index has a normal distribution with a standard deviation of 0.1. The permissible range of $\mathcal{F}$ is $\sim$0.1--0.5 before the model scatter becomes considerably broader than is observed. \subsection{Optically Thin Scenario} \label{subsec:thin} \begin{figure*}[t!] \includegraphics[width=\linewidth]{multid_thin.pdf} \figcaption{The same as Figure~\ref{fig:thick_covar}, but for an optically thin disk population model. \label{fig:thin_covar} } \end{figure*} We also consider the opposite limiting scenario, where the emission is optically thin ($\tau_d \ll 1$) and \begin{equation} I_\nu(r) \approx \begin{cases} B_\nu(T_d) \, \tau_d & \, \text{if $r \le R_o$} \, .\\ 0 \, & \, \text{otherwise} \, . \end{cases} \end{equation} This scenario is mathematically similar to the optically thick case, but with the added complexity of permissible variations in $\tau_d$. We assume the generic behavior \begin{equation} \tau_d(r) = \tau_0 \, \left(\frac{L_\ast}{L_\odot}\right)^\eta \, \left(\frac{r}{r_0}\right)^{-p} \, . \label{eq:tau} \end{equation} To again analytically illustrate the scaling behaviors in this model, we (temporarily) assume that the Rayleigh-Jeans approximation is applicable, so \begin{equation} I_\nu(r) \approx \frac{2 \, \nu^2 \, k}{c^2} \, T_0 \, \tau_0 \, r_0^{p+q} \, \left(\frac{L_\ast}{L_\odot}\right)^{0.25+\eta} \, r^{-(p+q)} \end{equation} inside $R_o$. Integrating that brightness profile and relating it to the definitions of the continuum luminosity and effective size, we arrive at the ``base" relations, \begin{subequations} \begin{equation} L_{\rm mm} = \frac{4 \pi \nu^2 k T_0 \, \tau_0 \, r_0^{p+q}}{(2-p-q) \, c^2} \, \left(\frac{L_\ast}{L_\odot}\right)^{0.25+\eta} R_o^{2-p-q} \, , \label{eq:Lmm_Rt_thin} \end{equation} \begin{equation} R_{\rm eff} = x^{1/(2-p-q)} \, R_o \label{eq:Reff_Rt_thin} \end{equation} \label{eq:base_thin} \end{subequations} (cf., Equation~\ref{eq:base_thick} for the $\tau_d \gg 1$ case). Combining the parts of Equation~\ref{eq:base_thin}, we find the somewhat unwieldy scaling behavior \begin{equation} R_{\rm eff} \propto L_\ast^{-(\eta+0.25)/(2-p-q)} \, L_{\rm mm}^{1/(2-p-q)} \, . \end{equation} In this case, the criterion required to simultaneously reproduce the observed shapes of the $R_{\rm eff}$--$L_{\rm mm}$ and $L_\ast$--$L_{\rm mm}$ scaling relations is \begin{equation} \eta \approx 0.4(p+q) - 0.2 \, . \end{equation} For reasonable $p+q$ (based on the $\gamma$ values inferred in Section~\ref{sec:results}), that again implies that $R_o \propto L_\ast^{0.4}$, or equivalently $R_o \propto M_\ast^{0.7}$ in this case (as it must to explain the observed $R_{\rm eff}$--$L_\ast$ or --$M_\ast$ behaviors). Figure~\ref{fig:thin_covar} shows the parameter relationships of interest overlaid with optically thin population models that were manually tuned to match the data. In those model calculations, we assumed the more general surface brightness prescription $I_\nu \approx B_\nu(T_d) (1 - e^{-\tau_d})$, although the simplification described above still provides a reasonable approximation. As was the case for the optically thick models, we need to make three assumptions to reproduce the trends, relating to a gradient, a size-scaling, and a normalization, \begin{equation} \begin{cases} \,\, \eta \approx 0.4(p+q) - 0.2 \, , \\ \,\, R_o \approx 90 \, (M_\ast/M_\odot)^{0.7} \, \text{au} \, , \\ \,\, \tau_0 \approx 0.4 \, (M_\ast/M_\odot)^{1.7 \eta} \, , \end{cases} \label{eq:thinmodel} \end{equation} where in the last criterion we substituted the mass scaling for the original luminosity scaling using the mean $L_\ast$--$M_\ast$ relation appropriate for this sample. In Figure~\ref{fig:thin_covar}, we have assumed $p+q \approx 0.75\pm0.25$ ($\eta \approx 0.10\pm0.11$) to be consistent with the measured surface brightness slopes (note that $p+q \approx \gamma$ in this model). \begin{figure*}[t!] \includegraphics[width=\linewidth]{multid_mdot.pdf} \figcaption{An analogous comparison as Figure~\ref{fig:thick_covar}, but in this case for the added dimension of accretion rate ($\dot{M}_\ast$). These behaviors presume the same stellar population model, the optically thick prescription from Section~\ref{subsec:thick}, and an intrinsic $\dot{M}_\ast$--$M_\ast$ scaling (see Table~\ref{table:regressions}). \label{fig:mdot_covar} } \end{figure*} We want to explicitly point out two consequences of this formulation. First, this behavior, where $\eta$ is roughly zero, is consistent with a scenario where all disks have the {\it same optical depth profile}, regardless of their host properties \citep[cf., Eq.~\ref{eq:tau}; see also][]{pietu14}: the observed scaling relations can all be explained with variations in the disk sizes. And second, if we make the standard assumption that the dust opacity $\kappa_\nu$ does not vary with radius, we can derive a scaling relation between the disk mass and host mass. To do that, under these assumptions ($\eta \sim 0$ and $\kappa_\nu(r) \approx {\rm constant}$) we note that the dust surface density profile would vary like \begin{equation} \Sigma \propto M_\ast^{\xi} \, r^{-p}, \end{equation} where we have permitted an $M_\ast$-dependence in the normalization. The dust mass is the integral of the surface density profile over the disk area, and so \begin{equation} M_d = \int_0^{R_o} \Sigma(r) \, 2\pi \, r \, dr \propto M_\ast^\xi \, R_o^{2-p}. \end{equation} Since $R_o \propto M_\ast^{0.7}$ (Eq.~\ref{eq:thinmodel}), and $p = 0.25$ for the fiducial case where $q = 0.5$, we find that $M_d \propto M_\ast^{1.2+\xi}$. The salient point is that a consideration of disk sizes in the demographic landscape shows that the roughly linear $M_d$--$M_\ast$ scaling relation inferred for the typical set of assumptions \citep{andrews13,ansdell16,pascucci16} implies that disks have similar surface density profiles regardless of their host parameters (i.e., $\xi$ is consistent with zero). It is worthwhile to consider these implications in studies that employ the scaling relations that are inferred with such assumptions. In terms of the scatter in these parameter relationships, we are again forced to assume the same dispersion around the mean $R_o$--$M_\ast$ relation as in the optically thick case. Some additional scatter in the optical depth normalization ($\tau_0$) is also plausible, so long as the standard deviation in a log-normal distribution is $\lesssim$0.4\,dex (and $\sim$0.2\,dex is more appropriate). \subsection{Links to Accretion Rates} \label{subsec:mdot} An example of the extensions of these population models to include the mass accretion rate is shown in Figure~\ref{fig:mdot_covar}, in the case for the optically thick scenario (although the optically thin scenario works just as well). The added assumption when expanding the dimensionality this way is that there is an intrinsic $\dot{M}_\ast$--$M_\ast$ relation like the one observed (including the scatter; see Table~\ref{table:regressions}). We could just as easily have presumed an intrinsic $L_{\rm mm}$--$\dot{M}_\ast$ relationship as observed \citep[e.g., see][]{mulders17}, and thereby would predict an appropriate $\dot{M}_\ast$--$M_\ast$ scaling and amount of scatter. Figure~\ref{fig:mgas} makes the comparison between $L_{\rm mm}$ and a crude diagnostic of the gas mass of the disk, $M_{\rm gas} \approx \dot{M}_\ast \, \tau_\ast$ \citep[cf.,][]{hartmann98}. \dt{The observed correlation is to be expected, since $\dot{M}_\ast$ and $L_{\rm mm}$ are correlated (Figure~\ref{fig:lmmreff_acc}) and neither of those variables depend significantly on $\tau_\ast$. For the assumptions outlined above, the same behavior would be produced for the optically thick or thin population models. That said, the latter models may seem more compelling in this context} because $L_{\rm mm}$ is roughly proportional to the dust mass ($M_{\rm dust}$). The red curve in Figure~\ref{fig:mgas} is not a fit: it represents the \dt{expected behavior for equivalent masses if the disks are optically thin,} have a dust-to-gas mass ratio of 1\%\, a mean dust temperature of 20\,K, and a 340\,GHz dust opacity of 3.5\,cm$^2$\,g$^{-1}$. It is tempting to fold the added dimensionality of {\it size} into an interpretation in the context of the viscous evolution of simple accretion disks \citep[e.g.,][]{hartmann98,andrews09,isella09,pascucci16,rafikov17,tazzari17,lodato17,mulders17}. The effective continuum sizes and luminosities as measured here are not easily associated with the viscous model parameters of interest (the gas mass and characteristic radius) because they are both strongly (and non-linearly) modulated by the coupled evolution of the gas and solids in the disk \citep[e.g.,][]{birnstiel12,birnstiel15}. Moreover, models that perform that coupling are known to have a significant efficiency problem \citep[e.g.,][]{takeuchi02,takeuchi05,brauer07}, presumably requiring disk structures that deviate significantly from the standard gas model configurations \citep[e.g.,][]{pinilla12a}. For those reasons, our preference is to keep the focus on more empirical relationships. Nevertheless, it would be interesting in future work to consider full population models that link the evolution of solids in viscous accretion disks in an effort to more accurately assess their ability (or not) to reproduce the observed scalings on the right timescales. \begin{figure}[t!] \includegraphics[width=\linewidth]{Mgas_Mdust.pdf} \figcaption{A comparison of the continuum luminosities ($\propto M_{\rm dust}$) and a simple diagnostic of the gas mass in standard accretion disk models ($M_{\rm gas} \approx \dot{M}_\ast \, \tau_\ast$; cf., \citealt{hartmann98}). The red curve marks the expected behavior for standard opacity and temperature assumptions if the dust-to-gas mass ratio is 1\%. \label{fig:mgas} } \end{figure} \subsection{Comments on Transition Disks} \label{subsec:trans} To this point, we have discussed the multidimensional relationships among all of the sample targets together, regardless of their continuum emission morphology. The sub-sample of transition disks (see Section~\ref{subsec:results} for our definition) have bulk demographic properties that, broadly speaking, match well with the overall sample. But aside from that general agreement, there are some subtle deviations that merit brief attention. First, and perhaps most obviously, the transition disks in this sample tend to be hosted by more luminous (massive) stars. We presume that this is a selection effect, in that there could well be other transition disks among the less luminous (massive) members of the sample that have not yet been identified \citep[e.g., see][]{ercolano09}. However, this could instead be a real physical effect; very high angular resolution imaging would be required to know for certain. Figure~\ref{fig:stars_covar} suggests that the transition disk hosts are slightly less luminous than their ``normal" counterparts at a fixed effective temperature, and therefore appear older. Their inferred age distribution does indeed peak at 2.5\,Myr, compared to 1.5\,Myr for the normal disks, but a shift of that amount is not statistically significant given the large uncertainties. \citet{pinilla18} recently argued that the $M_d$--$M_\ast$ (or equivalently $L_{\rm mm}$--$M_\ast$) scaling for transition disks is considerably flatter than for the normal disk population. A close examination of the top right panel of Figure~\ref{fig:lmmreff_hosts} indeed shows a similar result for this sample.\footnote{There is substantial sample overlap between this article and the \citet{pinilla18} study, particularly at the low-$M_\ast$ end.} The normal disk relation is the same as quoted in Table~\ref{table:regressions}, but the transition disks alone have a much lower scaling index ($\mathcal{B} = 0.86\pm0.18$; a difference at the 3\,$\sigma$ level). We should caution, however, that both this analysis and the \citet{pinilla18} result are strongly influenced by the 2 or 3 targets at the low-$M_\ast$ end. If those are just the brightest transition disks with low-mass hosts, and there are indeed more such systems with fainter disks that we have not yet identified, this apparent discrepancy could easily be biased by selection effects. We also find marginal support for the claim by \citet{najita07,najita15} that the transition disks have preferentially lower accretion rates for a given disk mass (see Figures~\ref{fig:lmmreff_acc} or \ref{fig:mdot_covar}). A comparison of the regression results in the $L_{\rm mm}$--$M_\ast$ plane finds that the normalization (intercept) is $\sim$0.6\,dex lower for transition disks compared to normal disks (at $\sim$90\%\ confidence); both sub-samples have the same slopes. Finally, we note that the surface brightnesses measured in Section~\ref{sec:analysis} for the transition disks are high. Assuming the standard dust temperature prescription (Equation~\ref{eq:tdust}), we would infer that all of these disks are optically thick around their peaks. One could reasonably argue that the local $T_d$ at these peaks should be enhanced by direct stellar irradiation \citep[e.g.,][]{dullemond01,dalessio05}, but not to a level that would substantially reduce the optical depth estimates. In reality, the ``ring" features for these systems are not resolved well, so if they are more narrow than we can infer then the peak brightness temperatures (and thereby optical depths) could actually be higher. This is all an interesting manifestation of the first scenario to explain the multi-dimensional scaling relations, described in Section~\ref{subsec:thick}: the emission is optically thick and has a large effective size, but the depleted central cavities reduce the continuum luminosity to a level consistent with a ``filling factor" ($\mathcal{F}$) of $\sim$0.1--0.5. Since the transition disks generally have similar demographic behaviors to the general population, it is natural to wonder if they are just a more obvious manifestation of this kind of behavior. Perhaps this is a clue that the multidimensional relationships probed here are produced by optically thick substructures on spatial scales considerably smaller than the available resolution. \subsection{Further Caveats and Future Work} \label{subsec:hybrid} Of course, the limiting approximations in Sections~\ref{subsec:thick} and \ref{subsec:thin} are vast over-simplifications. The reality is probably a messier hybrid of those scenarios, perhaps where there are higher-level dependencies between the model parameters and the host and/or disk properties. In a sense, this could still be treated in the optically thick substructures paradigm if we permit a complicated functional form for the ``filling factor" $\mathcal{F}$. That parameter can include optically thin contributions, depleted zones, or both, in myriad morphological configurations. \begin{figure}[t!] \includegraphics[width=\linewidth]{twhya.pdf} \figcaption{The 340\,GHz continuum surface brightness distribution model inferred from the 0\farcs3 SMA observations of the TW Hya disk (black), overlaid with the observed (deprojected) ALMA emission profile at 20\,mas resolution \citep{andrews16}. The vertical dashed line marks the $\varrho_{\rm eff}$ for this target listed in Table~\ref{table:SBpars2}. The shaded region marks the radial separations where the emission transitions from optically thick to thin, based on the brightness temperatures and spatially resolved spectral index measurements of \citet{tsukagoshi16} and \citet{huang18}. \label{fig:twhya} } \end{figure} This may seem like a good explanation of demographic trends will rely on some knowledge of complicated minutiae. In reality, that is probably the case: it will be difficult to distinguish how to properly formulate population models without observations at much higher resolution. However, we can start to check if even the simple pictures laid out above make sense for specific examples. Perhaps the most obvious case to illustrate the point is the TW Hya disk, which has been observed at very high angular resolution (20--30\,mas; $\sim$2\,au) with ALMA (Andrews et al.~2016). Figure~\ref{fig:twhya} shows the model brightness profile inferred from the more modest-resolution SMA observations \citep{andrews16,tripathi17}, overlaid with the ALMA data. The spectral index of this emission is roughly constant at $\approx$2 inside $\sim$0\farcs5 (and perhaps out to 0\farcs9); coupled with the high brightness temperatures (comparable to the expected $T_d$), it is clear that the inner disk emission is optically thick \citep{tsukagoshi16,huang18}. In this case, the optically thick filling factor is $\mathcal{F} \approx 0.4$--0.6; if we further account for the small-scale radial depletions (``gaps"), $\mathcal{F}$ decreases by another $\sim$0.1. So, in this case, a high-resolution knowledge of where the emission is optically thick lends some credence to an explanation similar to the scenario described in Section~\ref{subsec:thick}. Ultimately, similar evidence as for TW Hya would be needed to definitively connect physical models to the demographic properties discussed here (or perhaps more elaborate ones). That said, some tests of the optically thick scenario might be available by folding in spectral index information (perhaps even unresolved) to the population analyses. Some guidance from linking models of viscous evolution and the coupling between gas and solids would also be welcome. We have combined the archival SMA measurements from \citet{tripathi17} with the comparable archival ALMA survey in Lupus \citep{ansdell16} to conduct a homogeneous analysis of the resolved 340\,GHz continuum emission from 105 nearby protoplanetary disks, with a focus on the multidimensional demographics related to emission {\it sizes}. Our key findings are: \begin{itemize} \item We confirm (in form and normalization) and quantify the previous measurements of strong correlations between the continuum luminosities ($L_{\rm mm}$) and the stellar host masses ($M_\ast$) or luminosities ($L_\ast$) -- $L_{\rm mm} \propto M_\ast^{1.5}$ or $L_{\rm mm} \propto L_\ast^{0.8}$ -- as well as the accretion rates -- $L_{\rm mm} \propto \dot{M}_\ast^{1.0}$. \item We verify the relationship between the continuum emission size ($R_{\rm eff}$) and $L_{\rm mm}$ measured by \citet{tripathi17}: the amount of emission scales linearly with its surface area: $L_{\rm mm} \propto R_{\rm eff}^{2.0}$. \item We identify new, albeit weaker, scaling relations connecting the emission sizes and the host properties -- $R_{\rm eff} \propto M_\ast^{0.6}$ or $\propto L_\ast^{0.3}$ -- and a marginal connection with the accretion rates -- $R_{\rm eff} \propto \dot{M}_\ast^{0.9}$. \item With some simplified approximations, demographic models explain these scaling relations and their associated dispersions for either optically thick or thin emission. The thick case requires an intensity-weighted filling factor of $\sim$0.3. The thin case suggests that disks have a relatively uniform optical depth profile that is independent of the host properties: the distribution of continuum luminosities are explained with variations in the disk size. In both cases, we require a slightly sub-linear scaling (and considerable dispersion) between the dust disk size and the host mass. \item The transition disks appear superficially to exhibit the same behavior as the general population, but may show some subtle differences in their connections to the stellar host and accretion parameters. Selection effects are still an issue in making firm conclusions in this regard. \item The key takeaway point is that the standard demographic analyses show clear connections among the disk masses ($\sim$$L_{\rm mm}$), disk sizes, and host properties. Unambiguously disentangling those connections may require large mm continuum surveys at a few times better angular resolution, preferably with spectral index information, to assess the role of optical depth in shaping the observables. \end{itemize} | 18 | 8 | 1808.10510 |
1808 | 1808.02902_arXiv.txt | We report the results of {\sl Hubble Space Telescope (HST)} and {\sl Chandra X-ray Observatory (CXO)} observations of the { GLIMPSE-C01 (hereafter GC01) star cluster}. Color-magnitude and color-color diagrams suggest a cluster age of { $\gtrsim2$ Gyrs up to $\sim10$ Gyrs (dependent on GC01's metallicity), a distance of 3.3-3.5 kpc, and strong differential reddening with $A_V=14-22$.} After performing astrometric corrections, we find that nine of the 15 X-ray sources have at least one NIR counterpart within the { 2$\sigma$} {\sl CXO} positional error circles. { However, given the very high density of NIR sources in the cluster, most of these counterparts are likely due to chance coincidence. We jointly analyze the X-ray and NIR properties to assess the likelihood of true associations. Based primarily on their X-ray properties, we identify} an LMXB candidate (source X2), a CV candidate (source X1), and { an AB candidate} (source X9). Source X11 is detected during an X-ray flaring episode with a flare luminosity ($L_{X}=2.1\times10^{33}$ erg s$^{-1}$) and has a quiescent luminosity $L_{X}<8.0\times10^{30}$ erg s$^{-1}$, in 0.5--8 keV at the distance of GC01, suggesting that the source is { either} an { AB or CV}. We also discuss the limits on an intermediate mass black hole at the center of GC01 and the challenges of X-ray source classification imposed by the limitations of the existing data and instrumentation along with future prospects in the {\sl James Webb Space Telescope} era. | Thousands of star clusters exist within the Milky Way Galaxy and galactic halo \citep{2015A&A...581A..39S}. Typically these clusters are classified either as old globular clusters (GCs), which have high concentrations of stars, or open clusters, { which are less dense}. Open clusters, typically found in the galactic disk, have masses on the order of 10$^{3}$ M$_{\odot}$, { near-solar metallicities}, and ages ranging between 1 Myr and 1 Gyr, although several old open clusters have also been discovered (e.g., NGC 6791, NGC 188; \citeauthor{2005A&A...438.1163K} \citeyear{2005A&A...438.1163K}). Among these are the Young Massive Clusters (YMCs), which are sometimes considered to be a class on their own, with masses $\gtrsim 10^4 M_{\odot}$ and ages up to a few hundred Myrs \citep{2010ARA&A..48..431P}. Low-metallicity GCs are mostly found in the Galactic halo (although higher metallicity GCs exist in the disk) with masses $\sim10^{4}$--$10^{6}$ M$_{\odot}$ and ages of $8-14$ Gyrs (see e.g., Fig~9 in \citealt{2010ApJ...708..698D}). The GLIMPSE$-$C01 cluster (GC01, hereafter), which was discovered with \emph{Spitzer} during the Galactic Legacy Infrared Mid Plane Survey \citep{2005AJ....129..239K}, is interesting because it does not fit well into any of these categories. GC01 is located at $l=31.3^{\circ}$, $b=-0.1^{\circ}$, and lies within 10~pc of the Galactic mid-plane. It was originally estimated to have a mass of $\sim 10^{5}$ M$_{\odot}$, a half-light radius of 36$''$, and a distance of 3.1-5.2 kpc \citep{2005AJ....129..239K}. This cluster is also highly reddened with $A_V\sim15\pm3$, which likely varies across the cluster (\citeauthor{2005AJ....129..239K} \citeyear{2005AJ....129..239K}; \citeauthor{2005A&A...442..195I} \citeyear{2005A&A...442..195I}; \citeauthor{2016AJ....152..173D} \citeyear{2016AJ....152..173D}). Since the discovery of GC01, widely varying distance and age estimates have been reported. Both \cite{2005A&A...442..195I} and \cite{2016AJ....152..173D} used red clump stars, observed in slightly different filters, to estimate a distance to GC01 of 3.7 and 5.2 kpc, respectively. \cite{2011MNRAS.411.1386D} used near infrared spectroscopy of 50 stars in the cluster to calculate their velocities and, by assuming that GC01 is moving with the disk, derive a kinematic distance of 5.0$\pm0.9$ kpc. The measured velocity dispersion of the stars suggests a virial mass of (8$\pm3)\times$10$^{4}$ M$_{\odot}$ \citep{2011MNRAS.411.1386D}. Due to its centrally-concentrated appearance in the NIR images, GC01 strongly resembles GCs. Therefore, it was initially suggested that GC01 is a GC with an age of at least a few Gyr. The diffuse infrared emission coincident with GC01 in both {\sl Spitzer} IRAC and MIPS images was interpreted in support of a GC passing through the Galactic disk and interacting with the gas and dust in the disk \citep{2005AJ....129..239K}. The advanced age and classification as a globular cluster would be consistent with the lack of radio emission (typically seen in younger open clusters), its high central stellar density, and the large number of giant branch stars with no luminous supergiants \citep{2005AJ....129..239K}. However, \cite{2011MNRAS.411.1386D} found that GC01 is more compact than typical GCs of a similar mass and that the mass density is more similar to that of YMCs. Using the K-band mass-to-light ratio \cite{2011MNRAS.411.1386D} infer an age between 0.3 and 2 Gyr. Models of the Red Giant Branch (RGB) tip brightness in the color-magnitude diagram (CMD) for stars in GC01 also suggests an age between 1-2.5 Gyr \citep{2016AJ....152..173D}. Therefore, GC01 could be a rare intermediate age massive cluster \citep{2011MNRAS.411.1386D}. The origin of massive GCs is a matter of ongoing debate \citep{2017MNRAS.465.3622R}. It has been suggested that at least some massive GCs in the disk could be the outcome of YMCs evolution \citep{2010ARA&A..48..431P}. However, GCs could have also coalesced out of a primordial gas cloud that later collapsed into the Galactic disc (see e.g. \citealt{2003Sci...299...65K}). Furthermore, recent simulations show that low-metallicity GCs may represent the cores of satellite galaxies that merged with the Milky Way \citep{2017MNRAS.465.3622R}. If the age of GC01 is indeed substantially less than the age of the Galaxy, its metallicity is closer to solar \citep{2013MNRAS.436..122L}, and its X-ray binary population is different from those in GCs, then it may represent a missing evolutionary link between the YMCs and massive GCs. { On the other hand, if GC01 has a low metallicity and old age (as suggested by \citealt{2005A&A...442..195I} and \citealt{2005AJ....129..239K}, respectively), then it could be similar to the rare GCs that reside in the Milky Way disk outside of the bulge, such as NGC 6544 \citep{2014AJ....148...18C} and Glimpse-C02 \citep{2008A&A...489..583K}.} Studies of X-ray sources located in open and globular clusters are crucial for understanding the evolution, dynamics, and stellar populations of these objects (\citeauthor{2013ASPC..470..251V} \citeyear{2013ASPC..470..251V}; \citeauthor{2010PNAS..107.7164P} \citeyear{2010PNAS..107.7164P}, \citeauthor{2010AIPC.1314..135H} \citeyear{2010AIPC.1314..135H}). Typically, the X-ray populations of globular and aged open clusters consist of cataclysmic variables (CVs), non-accreting\footnote{These MPSs can be solitary, in wide binaries, or in binaries with a very low mass companion which is ablated by the pulsar's wind (i.e., redbacks or blackwidows; for a recent review see \citealt{2017JApA...38...42M})} millisecond radio pulsars (MSPs), neutron star (NS) or black hole (BH) low mass X-ray binaries which can be in quiescence (LMXBs or qLMXBs, respectively), and active binaries (ABs, such as RS CVn and W Uma type systems; \citeauthor{2013ASPC..470..251V} \citeyear{2013ASPC..470..251V}; \citeauthor{2010PNAS..107.7164P} \citeyear{2010PNAS..107.7164P}, \citeauthor{2010AIPC.1314..135H} \citeyear{2010AIPC.1314..135H}). In dense GCs the relative numbers of these objects are expected to be related to the number of dynamical encounters (see \citeauthor{2006ApJ...646L.143P} \citeyear{2006ApJ...646L.143P}, \citeauthor{2006ApJ...651.1098H} \citeyear{2006ApJ...651.1098H}, and references therein). In less-dense clusters the X-ray source population is more likely to be primordial in nature (see e.g., \citeauthor{2006ApJ...647.1065K} \citeyear{2006ApJ...647.1065K}, \citeauthor{2010ApJ...712..380L} \citeyear{2010ApJ...712..380L}). As a result, GCs tend to be rich with recycled MSPs, many of which are in binaries (e.g., \citealt{2006ApJ...646.1104B}), while in old ($>7$ Gyr) open clusters the observed dominating X-ray source population is ABs (see e.g., \citealt{2017ApJ...837..130V}). If GC01 is indeed an intermediate age massive cluster, the population of X-ray sources would be essentially unknown. This prompted us to carry out multi-band \emph{Hubble Space Telescope} ({\sl HST}) observations to look for counterparts of the X-ray sources seen in the archival \emph{Chandra X-ray Observatory} ({\sl CXO}) image of GC01\footnote{The initial analysis of the {\sl CXO} data was reported by \cite{2007arXiv0708.3365P}. }. Below we report the results of these {\sl HST} observations together with a re-analysis of the archival {\sl CXO} data and a discussion of the existing limitations and challenges imposed by the current instrumentation. We present {\sl HST} and {\sl CXO} observations, and the data reduction methods in Section \ref{obs}. In Section \ref{analysis} we describe the analysis of the {\sl CXO} and {\sl HST} data, including the properties of the X-ray sources, NIR/optical photometry, and the cross-correlation of optical/NIR counterparts with the X-ray source positions. We describe the properties of the cluster in Section \ref{clusprop}. The multi-wavelength classification of X-ray sources and the limit on the mass of a possible intermediate mass BH (IMBH) are discussed in Section \ref{xray}. We discuss the current limitations of this study in Section \ref{outlook} and summarize the results of our findings in Section \ref{summ}. \begin{figure*} \begin{center} \includegraphics[scale=0.5]{Figure1.pdf} \caption{Comparison of the {\sl HST} WFC3 IR and UVIS filters used in this study with Johnson filters.} \label{filter} \end{center} \end{figure*} \begin{figure*} \includegraphics[scale=0.5,{trim=-0 0 0 0}]{Figure2.pdf} \caption{ \emph{CXO} and {\sl HST} images of the GC01 field showing the same area on the sky (North is up, East is to the left). { The solid white (left) and green (right) circles represent the $36''$ half-light radius, while the green cross (right) shows the cluster center \citep{2005AJ....129..239K}.} {\sl Left:} Binned (by a factor of 2) and smoothed (with a Gaussian kernel with a radius of 2$''$) \emph{CXO} false color image ($3-8$ keV - blue, $1.5-3$ keV -- green, and $0.5-1.5$ keV -- red). X-ray sources detected with a significance $>$6 net counts are numbered in correspondence with Table \ref{xsrc}. {\sl Right:} False color {\sl HST} image made from the F127M (blue), F139M (green), and F153M (red) WFC3/IR images. The magenta cross marks the cluster center determined by Kobulnicky et al.\ (2005). % }\label{ROI} \end{figure*} \vspace{2mm} | { The distance estimates using RC stars in the cluster's center implies a distance of $\approx$ 3.0--3.7 kpc depending on the absorption ($A_V=17-19$). This distance estimate is slightly more constrained { for the photometry from} the region outside of the cluster's core where the reddening is more uniform, giving $d=3.3-3.5$ kpc for $A_V=17-16$, respectively. This estimate is at the lower end of previous distance estimates of 3.1--5.2 kpc obtained from the ${}^{13}$CO feature emission and extinction map \citep{2005AJ....129..239K}, but is consistent with the 3.8$\pm0.7$ kpc distance obtained using RC stars \citep{2005A&A...442..195I}. However, these smaller distances still leave open the possibility that the cluster is embedded in the ${}^{13}$CO cloud at 3.1 kpc.} One of the most controversial properties of GC01 is its age. The isochrone matching to the cluster CMD suggests an age of $\gtrsim2$ Gyr, assuming a solar metallicity, $A_V=17$, and distance of 3.3 kpc. However, the 10 Gyr low metallicity isochrones are also consistent with the data. If GC01 does have a low-metallicity, it would then be more likely to be a a GC passing through the disk of the Galaxy, because low-metallicity GCs are typically found in the Galactic halo and not the Galactic disk (see e.g., \citealt{2013MNRAS.436..122L}). We find that ages $\lesssim2$ Gyr are inconsistent with the observed NIR CMD. The WFC3/UVIS CMD also disfavors ages $<1$ Gyr. Recently, \cite{2016AJ....152..173D} have fitted {\sl SPITZER} IRAC CMDs and found the age to be between 1 and 2.5 Gyr, compatible with our estimate for solar metallicity. Thus, it is possible that GC01 is a $\sim$2 Gyr old massive cluster born in the Galactic disk with an age similar to those of the Galactic clusters IC 4651, NGC 752, and M67 (\citealt{2002A&A...386..187M,1972MNRAS.157..147B,2015MNRAS.452.3394M}) but with a much larger mass. Due to its large mass, it could be an aged YMC, such as Westerlund 1, RSGC 03, or Arches \citep{2010ARA&A..48..431P}. { However, with the data at hand, we cannot rule out that GC01 is an old ($\sim$ 10 Gyr) low-metallicity globular cluster plunging into the disc (c.f., NGC 6544; \citealt{2017A&A...608A.140C})}. Future spectroscopic observations to determine the metallicity of GC01 would allow us to differentiate between these two scenarios. Further, the higher angular resolution { and larger field of view} of {\sl JWST/NIRCam} can also help to better constrain the properties of the CMD of GC01. \subsection{Loop-like Structure} \label{loopstruc} \cite{2005AJ....129..239K} reported a loop-like structure seen in the {\sl Spitzer} IRAC images of GC01 from the {\sl GLIMPSE} survey. They rule out both dust shells ejected by stars and a supernova remnant due to the size of the feature and lack of radio emission, respectively. They suggest that the structure could either be an old nova shell or planetary nebula \citep{2005AJ....129..239K}. However, in the WFC3 images the structure is resolved into several stars of a similar brightness arranged in a peculiar loop-like pattern (see Figure \ref{loop}), which can be described as an ellipse with a semi-major axis $a=2\farcs7$ and semi-minor axis $b=1\farcs7$, corresponding to physical sizes of 0.04 pc and 0.03 pc, respectively (at a distance of 3.3 kpc). If the structure is a ring seen in projection onto the sky, the inclination angle would be $\simeq30^{\circ}$ { north through east}. Only seven out of $\sim13$ stars in the loop-like structure have photometry that satisfies the quality criteria of the photometric catalog. These seven stars are plotted as red points on top of our CMD. They all lie towards the top of the CMD (see Figure \ref{loop}) and all but one have a $m_{\rm F127M}$-$m_{\rm F153M}$ color of $\sim$1.7. Given the advanced age of the cluster it is difficult to imagine that the ring-like arrangement of the stars could be maintained since its formation. Most likely the structure is just an accidental arrangement in the projection onto the plane of the sky. Accurate 3D velocity measurements for these stars would provide further information. In addition to spectroscopic radial velocity measurements, future {\sl JWST} observations would allow one to measure tangential velocity components of the stars if the cluster is as close as 3.3 kpc. \begin{figure*} \begin{center} \includegraphics[scale=0.455]{Figure11.pdf} \caption{Left: HST WFC3/IR false color image in F127M (blue), F139M (green), and F153M (red) filters featuring the loop-like structure discussed in Section \ref{loopstruc}. Right: $m_{\rm F153M}$ vs. $m_{\rm F127M}$-$m_{\rm F153M}$ CMDs showing the locations (with the red filled circles) of the stars from the loop-like structure. The stars are bright, and lie on the red side of red giant branch of the isochrones. } \end{center} \label{loop} \end{figure*} \label{summ} Using both {\sl HST} and {\sl CXO} we were able to probe the parameters of GC01 and its X-ray source population. { In the cluster center}, we have detected 1,964 sources in the WFC3/IR F127M, F139M, and F153M images and { 777} sources in the WFC3/UVIS F814W image. A color-color diagram suggests a variable extinction $A_V=18\pm4$. { We have also selected a region with less differential reddening ($A_V=16-17$) and crowding just outside of the cluster's core, and have analyzed 1,354 NIR sources in this region.} The peak of the red clump star distribution { in this alternative cluster region} was used to estimate the distance to the cluster, giving { $d=$3.5--3.3 kpc for $A_V=16-17$, respectively}. The CMDs and stellar isochrones corrected for this extinction and placed at a distance of 3.3 kpc imply an age of { $\gtrsim2$} Gyrs { up to $\sim$10 Gyrs for a lower metallicity.} Therefore, the {\sl HST} photometry by itself does not confidently discriminate between an intermediate age massive galactic cluster or an old globular cluster plunging into the disk. The lack of any radio MSPs could suggest that GC01 is not an old GC (such as Terzan 5 or 47 Tuc), although there remains a possibility that their detections are hampered by the large absorption towards GC01. We have analyzed the 15 brightest X-ray sources located within the central part of GC01. Nine of the X-ray sources have at least one coincident NIR source seen by {\sl HST} and three have multiple coincident NIR sources. { Additionally, seven of the X-ray sources are coincident with at least one optical/NIR source. This suggests that several of the optical/NIR sources are true counterparts of the X-ray sources. However, we cannot confidently determine which sources in particular are true counterparts to the X-ray sources.} Source X1 is a likely CV candidate given its X-ray colors and luminosity. One of the sources (X2) is a likely a qLMXB as it has a soft X-ray spectrum ($\Gamma=5.6$) and relatively large { unabsorbed} X-ray luminosity ($L_{\rm 0.5-8}=3.4\times10^{33}$ erg s$^{-1}$). X9 is likely an AB type system with a { potential} NIR counterpart that has well measured photometry that suggests it { is a red giant}. X11 showed an X-ray flare that lasted about 200 s and reached a peak X-ray luminosity of $L_{0.5-8}=2.1\times10^{33}$ erg s$^{-1}$ and then decayed back to below the detection limit over 8 ks. { This source is likely to be either an AB type system or CV}. X11 does not have a NIR/optical counterpart inside of its X-ray positional error circle. However, there is one NIR/optical counterpart, with well measured photometry that lies in the { red giant} region of the CMD, just on the edge of the X-ray positional error circle and may still be associated with X11. \noindent{ Acknowledgements:} Support for Program number HST-GO-14183.003-A was provided by NASA through a grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555. We are very grateful to Leisa Townsley and Patrick Broos for their help with running {\sl ACIS Extract}, as well as for insightful discussions on X-ray image analysis. { We would like to thank the anonymous referee for their helpful and constructive comments, which improved the paper.} | 18 | 8 | 1808.02902 |
1808 | 1808.04838_arXiv.txt | The primary goals of the STRong lensing Insights into the Dark Energy Survey (STRIDES) collaboration are to measure the dark energy equation of state parameter and the free streaming length of dark matter. To this aim, STRIDES is discovering strongly lensed quasars in the imaging data of the Dark Energy Survey and following them up to measure time delays, high resolution imaging, and spectroscopy sufficient to construct accurate lens models. In this paper, we first present forecasts for STRIDES. Then, we describe the STRIDES classification scheme, and give an overview of the Fall 2016 follow-up campaign. We continue by detailing the results of two selection methods, the Outlier Selection Technique and a morphological algorithm, and presenting lens models of a system which could possibly be a lensed quasar in an unusual configuration. We conclude with the summary statistics of the Fall 2016 campaign. Including searches presented in companion papers (Anguita et al.; Ostrovski et al.), STRIDES followed up \NOBSTOT\ targets identifying \NlensTOT\ new strongly lensed systems, and \NNIQTOT\ nearly identical quasars (NIQs), which could be confirmed as lenses by the detection of the lens galaxy. 76 candidates were rejected and \NIncTOT\ remain otherwise inconclusive, for a success rate in the range 6-35\%. This rate is comparable to that of previous searches like SQLS even though the parent dataset of STRIDES is purely photometric and our selection of candidates cannot rely on spectroscopic information. | \label{sect:intro} In the four decades since the discovery of the first strongly lensed quasars \citep{WCW79,Wey80}, they have morphed from an intellectual curiosity to a powerful and in some sense unique astrophysical tool \citep{CSS02}. Three classes of applications make strongly lensed quasars especially valuable. First, by modeling how the light of the background quasar and its host galaxy is distorted one can reconstruct the distribution of luminous and dark matter in the deflector, and thus address fundamental astrophysical problems like the normalization of the stellar initial mass function \citep{Poo++09,Sch++14} and the abundance of dark matter subhalos \citep{M+S98,M+M01,D+K02,Nie++14,Nie++17,BAR17}. Second, by exploiting magnification, one can study in great detail the distant quasars, the properties of their accretion disks and host galaxies \citep{Peng:2006p236,Din++17}. Third, by measuring time delays between the variable images and stellar kinematics of the deflector one can measure cosmic distances and thus cosmological parameters, especially the Hubble Constant \citep{Ref64,Sch++97,T+K02b,S+S13,Suy++13,T+M16,Suy++14,BAR16,Bon++17,T+K18,STA18}. Unfortunately, most applications to date have been limited by the small number of known suitable systems. Lensed quasars, patricularly the ones with four images which provide the most information, are rare on the sky \citep[of order 0.1 per square degree at present day typical survey depth and resolution;][]{O+M10}. Therefore successful searches for lensed quasars require searches over large solid angles \citep[e.g.,][]{Bro++03,Ina++12,Mor++16} and substantial follow-up to weed out false positives. Furthermore, not every lensed quasar system is suitable for every application: depending on the specifics of the lensing configuration and on the brightness of deflector and source, some systems contain significantly more information than others. Thus, in practice, every application of strongly lensed quasars to date has been limited to samples of one or two dozens at best. The present generation of wide field imaging surveys provides an opportunity to make transformative measurements by increasing dramatically the sample of known lens quasars. Hundreds of strongly lensed quasars are hiding in the thousands of square degrees currently being imaged by the Dark Energy Survey (and similarly, e.g., the Hyper-Suprime-Cam SSP Survey, the VST-ATLAS Survey), waiting to be discovered and followed up. In order to exploit the bounty of data provided by DES, we have formed the STRIDES collaboration (STRong lensing Insights into Dark Energy Survey \footnote{STRIDES is a Dark Energy Survey Broad External Collaboration; PI: Treu. \url{http://strides.astro.ucla.edu}}). The immediate goal of STRIDES is to identify and follow-up large numbers of multiply imaged quasar motivated by two main science drivers: i) study dark energy using gravitational time delays; ii) study dark matter using flux ratio and astrometric anomalies. Additional science goals include the determination of the normalization of stellar mass-to-light ratios of massive early-type galaxies and the properties of accretion disks through the study of quasar microlensing \citep[e.g.,][]{Mot++17}. As we will show in this paper, STRIDES can in principle discover enough strongly lensed quasars to make substantial headway on both its two main science drivers. Strongly lensed quasars' main contribution to dark energy measurements is through the determination of absolute distances in the relatively low redshift universe, and thus of the Hubble Constant $H_0$ \citep{T+M16}. In turn, knowledge of $H_0$ is currently a limiting factor in the interpretation of cosmic microwave background data \citep{Wei++13,Bon++17}. Current measurements of $H_0$ based on the local distance ladder method reach $\sim 2.4$\% precision \citep{Rie++16,Rie++18a,Rie++18b}. The most recent time delay based measurements reach $\sim3.8$\% with just 3 systems \citep{Bon++17}. Reaching 1\% equivalent precision on H$_0$ is extremely important \citep{Suy++12,Wei++13,T+M16} and it will require $\sim40$ lensed quasars \citep{JeeEtal2016,STA18} with data and models of quality equivalent to those presented by \citet{Suy++17,Rus++17,Slu++17,Won++17,Bon++17}. Similarly, current limits on dark matter substructure are based on $\sim10$ lenses \citep{D+K02,Veg++14,Nie++17}. Quadrupling the sample of viable quads will be a major step forward in constraining the properties of dark matter \citep{Gil++18}. Finding lensed quasars in purely imaging datasets of the size of DES is an unprecedented task. It requires the development of new algorithms to identify candidates from the imaging data, and substantial investment of telescope time to follow-up and confirm the candidates. In order to maximize completeness and purity, the collaboration is pursuing multiple independent approaches to identify candidate lenses. The lack of $u$-band imaging data in DES makes it particularly hard to identify QSOs, therefore many of the selection techniques combine DES imaging with WISE photometry. The candidates are then followed up with spectroscopy and higher resolution imaging. Both are necessary to confirm the lensing nature of the systems and obtain the redshift and astrometry necessary for modeling and scientific exploitation. First results from the STRIDES program have been presented by \citet{Agn++15,Lin++17}. Once the candidates are confirmed, the best ones are selected for monitoring either with the 1.2m Euler Telescope or with the MPIA 2.2m Telescope at La Silla as part of the COSMOGRAIL network \citep{Cou++18}. This paper has multiple aims. First, it provides an overview of the STRIDES program and forecasts the number of expected lensed quasars to be found in the complete Dark Energy Survey (DES; \S~\ref{sec:forecasts}). The forecasts show that the DES area depth and resolution should be sufficient to more than double the current sample of known lensed quasars, providing new systems especially in the South hemisphere outside the area covered by previously largest search based on the Sloan Digital Sky Survey \citep{Ina++12}. Second, this paper defines a candidate classification system, and various subclasses of inconclusive and contaminant sources (\S~\ref{sec:class}). The system will be applied throughout the collaboration with the goal to standardize the lens confirmation process and hopefully adopted by other investigators. Third, this paper gives an overview of the Fall 2016 Follow-up campaign (\S~\ref{sec:overview}), listing the candidates selected by two techniques (\S~\ref{sec:outliers} and \S~\ref{sec:morphological}) that did not yield any confirmed lens, except for a possible unusual quadruply imaged quasar. Companion papers in this series present the follow-up of candidate lensed QSOs selected using other techniques \citep[][and Ostrovski et al. 2018, in preparation]{Ang++18}, showing spectra and images for all confirmed lenses and otherwise promising inconclusive systems\footnote{During the follow-up campaign, non-DES targets selected from other surveys were also targeted. Those are described by papers outside of this series \citep[e.g.,][]{Sch++17,Wil++18,Ost++18,Agn++18b}.}. The fourth goal of this paper is to present the summary statistics of the 2016 follow-up campaign, combining the results from every search method, as discussed in \S~\ref{sec:ss}. Target selection for the 2016 campaign was based on early DES datasets, which did not cover the full depth and footprint of the survey. Thus, the Fall 2016 campaign statistics are not sufficient for a detailed comparison with the forecast for STRIDES. However, the follow-up statistics are sufficient for an assessment of the success rate and the completeness of the searches so far. Remarkably, the success rate is comparable to those of previous searches, even though no spectroscopic preselection or u-band imaging was available. A short summary concludes the paper in \S~\ref{sec:summary}. All magnitudes are given in the AB system, and a standard concordance cosmology with $\Omega_m=0.3$, $\Omega_\Lambda=0.7$, and $h=0.7$ is assumed when necessary. | \label{sec:ss} In addition to the Outlier Selection Technique introduced by \citet{Agn17}, and the morphological technique described in Section~\ref{sec:morphological}, other techniques were developed by members of the STRIDES collaboration. Their selection techniques and results of follow-up are described in other papers of this series \citep[][and Ostrovski et al. 2018, in preparation]{Ang++18}. Overall, taking into account all selection methods, \NOBSTOT\ DES-selected candidates were observed. \NlensTOT\ were confirmed as lensed quasars, including \NquadTOT\ quads, \NNIQTOT\ were classified as NIQs. For \NIncTOT\ the observations were inconclusive, and the rest were rejected as contaminants. The scale of the follow-up is sufficient to get a first assessment of the success rate of our candidate selection techniques, and compare it with previous searches. The overall success rate across all techniques is in the range 6-35\%. This is a good success rate considering that the selection is purely photometric and no spectroscopic pre-selection is applied. For comparison, the most recently completed large scale search for lensed quasars is the Sloan Digital Sky Survey Quasar Lens Search \citep[SQLS;][]{Ina++12}. Starting from a sample of 50,836 spectroscopically confirmed quasars, they identified 520 candidates based on color and morphology. Thirty (including 26 in the so-called ``statistical sample'') of those were confirmed as lensed quasars. One important class of contaminants were 81/520 QSO pairs, i.e. approximately 16\%. Another important class of contaminants were QSO+star (at least 100), to which one should probably add most of the objects classified as ``no lens'' based on imaging data (158; spectroscopic classification is not available for this class; these are most likely to be QSO+star, Oguri 2017, private communication), A few objects could not be confirmed as lenses due to small separation (9), although some of them could very well be lenses. Thus, the overall success rate is at least 6\% but possibly a little higher. QSO+star class comprises at least 19\% of the spurious candidates, and perhaps as high as 50\%. We refer to the individual papers of this series for a breakdown in the various class of contaminants for the STRIDES searches. A more recent search is that carried out by the SDSS-III BOSS quasar lens survey \citep[BQLS;][]{Mor++16}. Similarly to SQLS, they start from spectroscopically confirmed quasars and look for evidence for lensing. In their initial study, they confirmed as lenses 13 of their 55 best candidates, i.e. a success rate of 20\%. Of the top 55 candidates 11 are confirmed quasar pairs, some of which might be unrecognized lenses. In addition, we can compare the number of forecasted lenses with the number of confirmed lenses to roughly estimate the completeness of our search so far, keeping in mind that the searches were conducted on partial and different DES data releases. The two search algorithms presented in this paper were applied to the Y1A1 DES data release, which covers approximately 1800 square degrees, i.e. 36\% of the DES footprint, shallower than full depth. The algorithm presented in paper II \citep{Ang++18} was applied to the Y1+Y2 footprint, corresponding to approximately the entire DES footprint, shallower than full depth. The algorithm presented in paper III (Ostrovski et al. 2018) was applied to the part of the Y3 data release that overlaps with the VISTA-VHS survey (approximately half the entire DES footprint, shallower than full depth). Considering only the brighter systems ($i\sim20.2$ or brighter) that should have been detectable at reduced depth, we expected (\S~\ref{sec:forecasts} roughly 60 lensed quasars, including 15 quads. We confirmed \NlensTOT\ lenses, including 2 quads (possibly 8/3 if one wishes to include DESJ2346-5203). It is unlikely that more quads are hiding amongst the 33 inconclusive systems (including NIQs), as those generally tend to be easily to confirm due to their peculiar morphology. Thus, we conclude that a large fraction of quads (and possibly doubles) brighter than $i\sim20.2$ remains to be found in the DES footprint, motivating additional searches in subsquent DES data releases and follow-up campaigns. This conclusion is consistent with the discovery of doubles and quads in the DES footprint, before \citep{Agn++15,Ost++17,Lin++17} and after \citep{Agn++18b} the conclusion of the Fall 2016 campaign. At the moment of this writing, considering all known lensed quasars within the DES footprint including those discovered before and after the STRIDES Fall 2016 campaign, there is a good agreement between the forecasts and the observations for $i\lesssim18.5$ (see Figure~\ref{fig:complot0}). Beyond this limit the number of known lensed quasars increases much more slowly than forecasted, suggesting that many lenses remain to be found. \begin{figure} \centering \includegraphics[angle=0, width=\columnwidth]{complot0.pdf} \caption{Comparison between known lenses (including those discovered before and after the Fall 2016 STRIDES campaign) within the DES footprint (solid lines) and OM10 forecasts (dotted lines). The thin blue lines indicate doubles (excluding NIQs), and the thick red lines indicate quads. The vertical axis shows the cumulative number of lenses, while the horizontal axis shows the total $i$-band magnitude measured within a $5''$-diameter aperture in DES images.} \label{fig:complot0} \end{figure} The public data releases \citep{GaiaDR2} of the Gaia Satellite \citep{Gaia16} have provided another powerful tool in the arsenal of the lens quasar finding community. Gaia's high resolution positions and proper motions have been shown to be extremely powerful by themselves \citep{Kro++18} and especially in combination with optical and mid-IR images for identifying lensed quasars and reject contaminants \citep{Agn++18b,Lem++18,A+S18}. The fast turnaround of these discoveries after the data releases is very encouraging for STRIDES both in terms of the prospects of completeness and success rate of targeted follow-up. Finally, we can make a further comparison between the forecast and the properties of entire sample, by looking at the quasar redshift distribution. Given the small number statistics we combine both confirmed lenses and NIQs, assuming that they are drawn from the same distribution, even though this of course will need to be revisited at the end of the STRIDES multi-year effort. The distribution is shown in Figure~\ref{fig:plottazs}. As forecasted, the distribution peaks at $z_s\sim2$, and drops off below 1 and above 3. Whereas the numbers are still too small for a quantitative comparison between forecast and detections, the qualitative agreement is encouraging, especially because contrary to the SDSS searches we did not rely on u-band imaging or spectroscopic information for selection of candidates. \begin{figure} \centering \includegraphics[angle=0, width=\columnwidth]{plottazs.pdf} \caption{Distribution of quasar redshifts for confirmed lenses (shaded histogram) and NIQs (open histogram).} \label{fig:plottazs} \end{figure} We have presented an overview of the STRIDES program, an external collaboration of the Dark Energy Survey aimed at finding and studying strongly lensed quasars, and outlined some of the results of the first comprehensive follow-up campaign. The main results of this paper can be summarized as follows: \begin{itemize} \item Our detailed forecasts indicate that about 50 quads and 200 doubles should be detectable in DES data. Of those, approximately 60 should be bright enough for time delay measurements with 1-2m class telescopes, while the rest will require a 4m class telescope for monitoring. All the systems will be bright enough to measure stellar velocity dispersion with 8-10m class telescopes. \item The STRIDES lens classification scheme is presented. In addition to confirmed lenses, and inconclusive systems, we adopt the class of Nearly Identical Quasars (NIQ) to identify inconclusive targets which are particularly promising for additional follow-up. \item We detail the results of the follow-up of 42 targets selected by two of the search techniques (Outlier Selection and Morphological; OST and MT respectively). One of those is a candidate quadruply imaged quasar (DESJ2346-5203; see the next bullet item), 11 are inconclusive, and and 30 are contaminants. The contaminants are dominated by QSO+star pairs for the OST and by star pairs for MT. \item Based on the analysis of $0\farcs4$ seeing Gemini-S images of the candidate quad DESJ2346-5203, and we conclude that this is not a quadruply imaged quasar in a classic configuration. If it is a multiply imaged quasar the morphology requires a complex deflector or extreme flux ratio anomalies. High resolution imaging or spectroscopy are required to definitely rule out (or confirm) this system as a lens. \item We summarize the results of our Fall 2016 observing campaign with the Keck, SOAR, NTT, Shane Telescopes. In total we followed up \NOBSTOT\ targets, confirming \NlensTOT\ lenses including \NquadTOT\ quads, and found \NNIQTOT\ NIQs. The observations were inconclusive for \NIncTOT\ targets, yielding a success rate in the range 6-35\%. This success rate is comparable with those of other large searches for lensed quasars even though no spectroscopic information nor $u$-band imaging was available to help in the selection. \end{itemize} In conclusion, the results of the first extensive STRIDES follow-up campaign demonstrate that multiply imaged quasars can be found efficiently from wide field imaging survey even in the absence of $u$-band or spectroscopic preselection. At the conclusion of our multi-year campaign, with the investment of telescope time to carry out imaging and spectroscopic follow-up of DES-selected targets, STRIDES should more then double the current a sample of known lensed quasars, and thus enable significant progress in our understanding of the nature of dark matter and dark energy. | 18 | 8 | 1808.04838 |
1808 | 1808.04715_arXiv.txt | I conduct simple analytical estimates and conclude that mixing by vortices is a more efficient process to transfer the cosmic ray energy of jet-inflated bubbles to the intracluster medium (ICM) than streaming of cosmic rays along magnetic field lines. Jets and the bubbles they inflate transfer heat to the ambient gas in cooling flows in cluster of galaxies and in galaxies. The internal energy of the jet-inflated bubbles is dominate by very hot thermal gas and/or cosmic rays. Cosmic rays that stream along magnetic field lines that connect the bubbles with the ICM heat the ICM as their energy is dissipated there. I find that about half of the cosmic ray energy is dissipated in the bubbles themselves. I also find that the ICM volume that the cosmic ray streaming process heats is only about five times as large as the volume of the bubbles. The outcome of heating by streaming only is that the cosmic rays form a larger bubble filled with very hot thermal gas. Therefore, there is a requirement for a more efficient process to transfer the internal energy of the bubbles to the ICM. I suggest that this process is heating by mixing that operates very well for both cosmic rays and the very hot thermal gas inside the bubbles. This leaves mixing-heating to be the dominant heating process of cooling flows. \newline \textbf{Keywords:} galaxies: clusters: intracluster medium; X-rays: galaxies: clusters; galaxies: jets; (ISM:) cosmic rays | \label{sec:intro} Although the old ``cooling-flow problem'' in clusters and groups of galaxies, and in galaxies has been solved, there is still no consensus on the heating mechanism of the intracluster medium (ICM; for a recent review see \citealt{Soker2016}). The ICM in many clusters of galaxies has short radiative cooling times. If no heating processes of the ICM exist large amounts of ICM should cool to low temperatures (by ICM I refer also to the interstellar medium in galactic cooling flows). The cooling-flow problem refers to the contradiction between the large mass cooling rates of the ICM that the no-heating assumption implies and the observations of much lower mass cooling rates. The solution is simply to relax the no-heating assumption and to consider the heating of the ICM in all cooling flows. The vast majority of heating models involve jets that the active galactic nucleus (AGN) at the center of the cooling flow launches. The jets and the hot low-density bubbles they inflate heat the ICM via a negative feedback mechanism that prevents a run-away heating or cooling (e.g., \citealt{Fabian2012, Farage2012, McNamaraNulsen2012, Gasparietal2013, Pfrommer2013, Baraietal2016, Soker2016, Birzanetal2017, Iqbaletal2017}). The two parts that compose the negative feedback cycle are the feeding of the AGN by the ICM and the heating of the ICM by the AGN. There are tens of observational and theoretical studies in recent years that support models in which the ICM feed the AGN with cold clumps (e.g., limiting the list to the last 3 years, \citealt{ChoudhurySharma2016, Hameretal2016, Loubseretal2016, McNamaraetal2016, Russelletal2016, Tremblayetal2016, YangReynolds2016b, Davidetal2017, Donahueetal2017, Gasparietal2017, Hoganetal2017, Meeceetal2017, Prasadetal2017, Russelletal2017a, Voitetal2017, Babyketal2018, Gasparietal2018, Jietal2018, Prasadetal2018, Pulidoetal2018, Vantyghemetal2018, Voit2018, YangLetal2018}). These models can be grouped under the \emph{cold feedback mechanism} that \cite{PizzolatoSoker2005} developed to replace the then popular Bondi accretion. In the cold feedback mechanism gas flows inward to form cold clouds and to feed the central AGN. Cold clumps are observed now down to the AGN, e.g., in Perseus \citep{FujitaNagai2017}. This implies that a cooling flow does take place, like the extreme cooling flow in the Phoenix cluster (e.g., \citealt{Pintoetal2018}), although in most cases the cooling flow is a moderate one. For that, in this study I use the term cooling flow (rather than other names that were invented later and caused some confusions; \citealt{Soker2010}). On the other hand, there is no agreement on the processes that contribute the most to the heating of the ICM by the jets that the AGN launches and the bubbles that they inflate. I find it useful to distinguish between heating processes where the jets and jet-inflated bubbles do work on the ICM and between heating by energy transport. In the first class of processes the jets and the jet-inflated bubbles do work on the ICM by exciting sound waves (e.g., \citealt{Fabianetal2006, Fabianetal2017, TangChurazov2018}), by driving shocks (e.g., \citealt{Formanetal2007, Randalletal2015, Guoetal2018}), by powering turbulence (e.g., \citealt{DeYoung2010, Gasparietal2014, Zhuravlevaetal2014, Zhuravlevaetal2017}), and/or by uplifting gas from inner regions (e.g., \citealt{GendronMarsolaisetal2017}). However, there are studies that show that these processes do not heat the ICM efficiently enough, despite the fact that these processes themselves take place to some degree. Although turbulence is observed in the ICM (e.g., \citealt{Zhuravlevaetal2014, Zhuravlevaetal2015, AndersonSunyaev2016, Arevalo2016, Hitomi2016, Hofmannetal2016,Hitomi2017}), there are studies that question turbulent heating (e.g., \citealt{Falcetaetal2010, Reynoldsetal2015, Hitomi2016, HillelSoker2017a, Bambicetal2018, MohapatraSharma2018}). Sound waves also are unable to supply the entire heating (e.g. \citealt{FujitaSuzuki2005}). Heating by shocks has some problems as well (e.g., \citealt{Sokeretal2016}). Observations show that bubbles can uplift cooler gas (e.g,. \citealt{Russelletal2017b, Suetal2017, GendronMarsolaisetal2017}), and some studies simulate the uplifting process (e.g., \citealt{Guoetal2018, Churazovetal2001}), but \cite{HillelSoker2018} claim that this cannot be the main heating process of the ICM. In the second class of heating processes the energy from the hot jet-inflated bubbles is carried into the ICM. One possibility, that is the main subject of the present study, is that cosmic rays that are accelerated within the jet-inflated bubbles, stream into the ICM and heat it (e.g. \citealt{Fujitaetal2013, FujitaOhira2013, Pfrommer2013}). The second possibility is heating by mixing that works as follows. As the jets propagate through the ICM and inflate bubbles they form many vortices inside the bubble and on the boundary between the ICM and the bubbles. These vortices mix hot bubble gas into the ICM (e.g., \citealt{BruggenKaiser2002, Bruggenetal2009, GilkisSoker2012, HillelSoker2014, YangReynolds2016b}). Finally, some studies suggest that two or more processes out of the different process can work together, e.g., thermal conduction and cosmic rays (e.g., \citealt{GuoOh2008}), mixing of cosmic rays that were accelerated inside jet-inflated bubbles with the the ICM, \citep{Pfrommer2013}, and heating by turbulence together with turbulent-mixing (e.g. \citealt{BanerjeeSharma2014}). In a series of papers (e.g., \citealt{HillelSoker2017a, HillelSoker2018} for most recent papers) we argued that although the inflation of bubbles also excite sound waves, shocks, and ICM turbulence, heating by mixing is much more efficient than the heating mechanisms where the jets and bubbles do work on the ICM. The heating by mixing can work for bubbles that are filled with very hot thermal gas or with cosmic rays, or a combination of the two. I consider these simulations to describe the inflation of bubbles by jets and the bubble evolution even when weak magnetic fields are presence. Although magnetic fields can change the details of the interaction of the bubbles with the ICM (e.g., \citealt{DursiPfrommer2008, Weinbergeretal2017}), I take the view that the main effect is that the bubbles amplify the ICM magnetic fields (e.g., \citealt{DursiPfrommer2008, Soker2017}). For example, magnetic fields can suppress the Kelvin-Helmholtz instability along the field lines (e.g., \citealt{DursiPfrommer2008, Weinbergeretal2017}), but not perpendicular to the field lines. Therefore, magnetic fields are expected to change the geometry of the instabilities and to some extend the geometry of vortices that mix the ICM and bubbles, but magnetic fields are not expected to change the existence and global roles of the vortices, that include mixing and the amplification of the magnetic fields. We did not compare the heating by mixing with the heating process where cosmic rays from the bubbles stream, instead of being mixed by vortices, into the ICM and heat it. Observations show that in older bubbles the contribution from sources other than the cosmic ray electrons to the pressure tends to be larger than in younger bubbles (e.g., \citealt{DunnFabian2004, Birzanetal2008, Crostonetal2008}). According to the results of the present study this additional pressure component could result from thermal hot gas heated by the mixing-heating process. The mixing-heating proceeds and increases the amount of thermal gas on the expense of the cosmic ray energy as the bubbles age. In light of some recent interest in cosmic ray heating by streaming (e.g., \citealt{Ruszkowskietal2017, Ehlertetal2018, JiangOh2018, Ruszkowskietal2018}), and the possibility that in many cases cosmic ray energy is the main energy content of bubbles (e,g., \cite{Abdullaetal2018} for MS~0735.6+7421), in this study I examine some of the properties of this process. In section \ref{sec:heating} I consider only cosmic ray streaming, without mixing by vortices, and study the partition of cosmic ray energy between heating the bubbles themselves and heating the ICM, and in section \ref{sec:ICMvolums} I estimate the volume of the ICM that can be heated by streaming cosmic rays. I discuss and summarize my results in section \ref{sec:summary} where I consider mixing of cosmic rays by vortices to be much more efficient than streaming. | \label{sec:summary} By means of simple analytical estimates I examined the efficiency by which streaming of cosmic rays along magnetic field lines can transfer energy from the jet-inflated bubbles to the ICM. In section \ref{sec:heating} I found that about half of the cosmic ray energy is dissipated in the bubbles themselves. Even before the dissipation of cosmic ray energy starts, some of the kinetic energy of the jets has been converted to thermal energy of very hot gas inside the bubbles. Over all, more than half of the kinetic energy of the jets ends up as very hot thermal gas inside the bubbles. In section \ref{sec:ICMvolums} I found that by streaming alone the cosmic rays heat a total ICM volume of about five times the volume of the bubbles, and hence this mechanism does not cover a large enough volume of the ICM for an efficient heating to work. Taking these two conclusions together, I argue that the effect of cosmic ray heating by streaming alone turns a bubble filled with cosmic rays to a larger bubble filled with very hot thermal gas. Therefore, although heating by cosmic ray streaming does take place, I think there is a need for an additional and a more efficient process to carry the energy from the bubbles to the ICM. I take this more efficient process to be heating by mixing (for references see section \ref{sec:intro}). The propagation of jets through the ICM and the inflation of the bubbles form many vortices in the ICM and in the bubbles, and these mix the content of the bubbles, i.e., the very hot thermal gas, the cosmic rays, and the magnetic fields, with the ICM. When the bubble content and the ICM are well mixed we expect that magnetic fields from the bubble content and from the ICM reconnect on small scales. Then streaming of cosmic rays on small scale can make the final cosmic ray energy transfer, as much as heat conduction can locally transfer energy from the very hot thermal gas, which is now well mixed with the ICM, to the ICM. The heating by mixing process that I refer to results from the vortices that the inflation of bubbles form (e.g., \citealt{GilkisSoker2012, HillelSoker2016}), and it is indiscriminate to the content of the bubble, and hence works for cosmic rays as well as for the thermal gas. Indeed, \cite{Pfrommer2013} already mentioned the mixing of cosmic rays with the ICM and referred to it as an important process. Here I argue that it is more significant than heating by cosmic ray streaming, and hence I strengthen our earlier claim that heating by mixing is the main process by which jets-inflated bubbles heat the ICM (see section \ref{sec:intro}). I emphasize that the non-relativistic thermal gas inside bubbles must be very hot, $T \ga 10^9 \K$. There are strong observational limits on the mass inside the bubbles (e.g., \citealt{SandersFabian2007, Abdullaetal2018}), and pressure balance between the bubbles and the ICM requires therefore the gas inside the bubbles to be very hot. Both mixing of the ICM with cosmic rays and inflation of the bubbles with jets at velocities of $\ga 10^4 \km \s^{-1}$ can account for this very hot gas. Another observational constrain comes from observations of cold rims around some bubbles (e.g., \citealt{Blantonetal2003}). Our 3D hydrodynamical simulations show that a cold dense gas around rising bubbles exists alongside the operation of the mixing-heating process \citep{HillelSoker2018}. More studies are required to better compare heating-mixing with these and other observations, e.g., radio observations. These studies require new sets of 3D magnetohydrodynamical simulations. The simulation of the feedback process in cooling flows is a very complicated task, e.g., as \cite{Martizzietal2018} show in a very recent study. There are many physical and numerical issues to consider. One of the key ingredients that simulations must include is the launching of kinetic jets. \cite{SternbergSoker2008} show that injecting energy off-center instead of launching jets might miss key processes, such as the formation of many vortices. It might be that by inserting their jets in a sphere off-center \cite{Weinbergeretal2017} and \cite{Ehlertetal2018} underestimated the role of heating by mixing \citep{HillelSoker2017b}. The heating by mixing also has some implications to the mass feeding part of the feedback cycle. \cite{PizzolatoSoker2005} showed that non-linear perturbations in the ICM are required to form the cooling clumps that feed the AGN, as was also confirmed in following studies (e.g., \citealt{Gasparietal2018}). \cite{PizzolatoSoker2005} also suggested that the AGN activity form these non-linear perturbations. This is most likely a result of the vortices and the mixing process. Therefore, in an indirect way, the results of the present study that further strengthen the mixing-heating mechanism, add also to the cold feedback mechanism. I thank Karen Yang, Yutaka Fujita, Prateek Sharma, Christoph Pfrommer and an anonymous referee for useful comments. This research was supported by the Israel Science Foundation, and by the E. and J. Bishop Research Fund at the Technion. | 18 | 8 | 1808.04715 |
1808 | 1808.01306_arXiv.txt | We present \xmm\ imaging spectroscopy of ten weak emission-line quasars (WLQs) at \hbox{$0.928\leq z \leq 3.767$}, six of which are radio quiet and four which are radio intermediate. The new \xray\ data enabled us to measure the power-law photon index, at rest-frame energies $>2$ keV, in each source with relatively high accuracy. These measurements allowed us to confirm previous reports that WLQs have steeper \xray\ spectra, suggesting higher accretion rates with respect to ``typical" quasars. A comparison between the photon indices of our radio-quiet WLQs and those of a control sample of 85 sources shows that the first are significantly higher, at the $\gtsim3 \sigma$ level. Collectively, the four radio-intermediate WLQs have lower photon indices with respect to the six radio-quiet WLQs, as may be expected if the spectra of the first group are contaminated by \xray\ emission from a jet. Therefore, in the absence of significant jet emission along our line of sight, these results are in agreement with the idea that WLQs constitute the extreme high end of the accretion rate distribution in quasars. We detect soft excess emission in our lowest-redshift radio-quiet WLQ, in agreement with previous findings suggesting that the prominence of this feature is associated with a high accretion rate. We have not detected signatures of Compton reflection, \Ka\ lines, or strong variability between two \xray\ epochs in any of our WLQs, which can be attributed to their relatively high luminosity. | \label{sec:introduction} It is common to classify weak emission-line quasars (WLQs) as luminous active galactic nuclei (AGN) having rest frame equivalent widths (EWs) of either $<$15.4\AA\ or $<$10.0\AA\ for the \lya$+$N~{\sc v}~$\lambda1240$ emission complex or C~{\sc iv}~$\lambda1549$ emission line, respectively (Diamond-Stanic \et 2009). These thresholds mark the $3\sigma$ limit at the low-EW tail of the respective EW distributions in Sloan Digital Sky Survey (SDSS; York \et 2000) quasars. Based on this classification, $\approx 10^{3}$ WLQs are known, to date, discovered mainly by the SDSS (e.g., Fan \et 1999; Anderson \et 2001; Collinge \et 2005; Plotkin \et 2010; Meusinger \& Balafkan 2014), but also by other surveys (e.g., McDowell \et 1995; Londish \et 2004). Interestingly, the fraction of WLQs among quasars appears to increase sharply from $\sim0.1\%$ at \hbox{$2 \lesssim z \lesssim 5$} to $\gtrsim15\%$ at $z \gtrsim 6$ (e.g., Fan \et 2006; Diamond-Stanic \et 2009; Ba$\tilde{\rm n}$ados \et 2016). Identifying the cause(s) for their line weakness is therefore important for understanding the physical conditions in the early universe. Multiwavelength and multi-epoch observations of several sub-samples of WLQs have shown that they are unlikely to be high-redshift galaxies with apparent quasar-like luminosities due to gravitational-lensing amplification, dust-obscured quasars, or broad absorption line (BAL) quasars (e.g., Shemmer \et 2006; Diamond-Stanic \et 2009). Additionally, the radio and \xray\ properties of WLQs indicate that they are unlikely to be identified as high-redshift BL Lacertae objects (Shemmer \et 2009; Plotkin \et 2010; Lane \et 2011). Therefore, the emission lines in WLQs are considered to be intrinsically weak. Several proposals have been put forward that attempted to explain the intrinsic emission-line weakness in WLQs. One of these suggested that the broad emission line regions (BELRs) in WLQs have either abnormal physical properties (e.g., lack of line-emitting gas, or a low covering factor), or are in the early stages of formation (e.g., Hryniewicz \et 2010; Liu \& Zhang 2011). Although such ideas may appear promising in their attempt to explain the increasing fraction of WLQs as a function of redshift, they face several difficulties, mainly on physical grounds. A different model, suggesting a relatively cold accretion disk, as a result of an unusually high supermassive black-hole mass and low accretion rate (Laor \& Davis 2011), faces its own challenges. In particular, a predicted sharp cutoff in the spectral energy distribution (SED) at $\lambda_{\rm rest}$\ltsim 1000~\AA\ has not yet been detected following observations of several WLQ sub-samples. The most promising path to identifying the underlying reason for intrinsic BELR line weakness in quasars originates from what is known as the Baldwin effect, which is an anti-correlation between BELR line EW and quasar luminosity (Baldwin 1977). In its modified form, this effect involves an anti-correlation between BELR line EW and the Eddington fraction (i.e., $L$/$L_{\rm Edd}$, where $L$ and $L_{\rm Edd}$ are the bolometric and Eddington luminosity, respectively; e.g., Baskin \& Laor 2004; Dong \et 2009; Shemmer \& Lieber 2015). The idea that a high Eddington fraction, corresponding to a high normalized accretion rate, is responsible for intrinsic line weakness has been proposed in various studies. For example, Leighly \et (2007a,b) have suggested that an extremely high accretion rate would result in a modified, UV-peaked SED lacking high-energy ionizing photons (see also Vasudevan \& Fabian 2007). However, such a model necessarily predicts unusual \xray\ weakness, with respect to the optical emission, which is not observed in all WLQs (e.g., Wu \et 2011, 2012; Luo \et 2015). Quantifying this \xray\ weakness is based on the optical-\xray\ spectral slope, defined as \aox$=\log(f_{\rm 2\,keV}/f_{2500\mbox{\rm\,\scriptsize\AA}})/ \log(\nu_{\rm 2\,keV}/\nu_{2500\mbox{\rm\,\scriptsize\AA}})$, where $f_{\rm 2\,keV}$ and $f_{2500\mbox{\rm~\scriptsize\AA}}$ are the flux densities at 2\,keV and 2500\,\AA, respectively. This parameter is strongly correlated with the luminosity density at 2500~\AA, $L_{\nu}(2500\,{\rm \AA})$ (e.g., Just \et 2007; Lusso \& Risaliti 2016). According to Luo \et (2015), a quasar is considered to be \xray\ weak if it has an observed \aox\ value that is lower by at least 0.2 from the expected value based on the \hbox{\aox-$L_{\nu}(2500\,{\rm \AA})$} correlation, i.e., \hbox{$\Delta$\aox\ $< -0.2$}; otherwise, it is considered \xray\ `normal'. In order to accommodate the wide range of optical-to-\xray\ flux ratios in WLQs, as well as their other properties, Wu \et (2011) and Luo \et (2015) have proposed an alternative model which also predicts extremely high accretion rates as a primary ingredient to explain quasar emission-line weakness. Unlike the modified SED scenario, this model predicts that the highly ionizing photons are absorbed by a shielding-gas component, growing vertically from the inner accretion disk, perhaps as the accretion rate rises above a certain threshold. This shielding-gas component may be physically identified with the thick inner accretion disk. The range in relative \xray\ weakness is thus explained by a range of viewing angles to the central \xray\ source. When viewed at large inclination angles (i.e., closer to a `pole on' view), a WLQ will appear to have `normal' \xray\ emission with respect to its optical emission; when viewed at smaller inclination angles, a portion of the \xray\ emission is absorbed by the shielding gas, resulting in an \xray\ weak WLQ. This model is supported by observations of \xray\ weak WLQs that show considerably harder \xray\ spectra with respect to typical quasars, indicating heavy intrinsic absorption in such sources (Wu \et 2011, 2012; Luo \et 2015). However, when compared to typical quasars over wide ranges of redshift and luminosity, WLQs do not appear to follow the strong EW-$L/L_{\rm Edd}$ anti-correlation, where $L/L_{\rm Edd}$ estimates are based on the H$\beta$ line (Shemmer \& Lieber 2015). The EWs of their \civ\ emission lines predict $L/L_{\rm Edd}$ values that are a factor of $\sim5$ larger than estimated. This discrepancy may imply that either (i) there are other factors that regulate emission-line strength in quasars, or (ii) the H$\beta$ line cannot be used to obtain reliable $L/L_{\rm Edd}$ estimates for all quasars. The \xray\ power-law photon index ($\Gamma$), particularly when measured above $\sim2$\,keV in the rest frame, has been identified as a more robust proxy for estimating $L/L_{\rm Edd}$ in quasars (e.g., Shemmer \et 2006, 2008; Constantin \et 2009; Brightman \et 2013; Fanali \et 2013). In particular, Risaliti \et (2009) found a strong correlation (\hbox{$r = 0.56$} and \hbox{$p < 10^{-8}$}) between $\Gamma$ and \hb-based \lledd\ in a sample of 82 SDSS quasars having \xmm\ observations. Accurate measurements of $\Gamma$ in a sizable sample of WLQs can therefore provide an independent indicator of their accretion rates. The first steps in this direction were taken by Shemmer \et (2009) and Luo \et (2015) who jointly fitted \xray\ data of seven and 18 WLQs, respectively, obtained from shallow \chandra\ X-ray Observatory (hereafter, \chandra; Weisskopf \et 2000) observations (see also Wu \et 2012). The first of these studies measured a \hbox{$\left < \Gamma \right > = 1.81^{+0.45}_{-0.43}$} in the observed-frame 0.5--8 keV range, concluding that this value is consistent with the values measured in typical type~1 quasars. However, their small sample included a mixture of \xray\ weak and \xray\ normal WLQs with extremely limited photon statistics (hence the large uncertainty on $\left < \Gamma \right >$). The second study measured \hbox{$\left < \Gamma \right > = 2.18\pm{0.09}$} in the rest-frame $>$2 keV band for a well-selected sample of \xray\ normal WLQs with considerably better photon statistics, thereby reducing the uncertainty on $\left < \Gamma \right >$ by a factor of $\approx5$. This recent result shows that \xray\ normal WLQs have, on average, a higher than normal photon index that indicates a high \lledd\ value. It also demonstrates the power of sample averaging in the presence of non-negligible intrinsic scatter that is inherent in the \hbox{$\Gamma$-\lledd} correlation. In this work we aim to obtain accurate measurements of $\Gamma$ values in a sample of {\em individual} \xray\ normal WLQs in order to determine the extent that these values deviate from the distribution of $\Gamma$ values in typical type~1 quasars. For this purpose, we obtained \xmm\ (Jansen \et 2001) observations of nine high-redshift WLQs discovered by the SDSS that were detected by \chandra. Prior to this investigation, only two such sources were observed by \xmm; one was targeted, and the other was observed serendipitously. We include these two sources in our analysis. We describe our sample selection, observations, and data reduction in Section 2; in Section 3 we present the results from our \xray\ imaging spectroscopy of WLQs and compare them with similar data for a carefully-selected sample of typical quasars. A summary is given in Section 4. Throughout this work we compute luminosity distances using the standard cosmological model (\hbox{$H_{0}$ = 70 km ${\rm s}^{-1}\,\ {\rm Mpc}^{-1}$}, \hbox{$\Omega_{\Lambda}$ = 0.7}, and \hbox{$\Omega_{\rm M}$ = 0.3}; Spergel \et 2007). Complete source names appear in the Tables, and abbreviated names appear in Figures and throughout the text. Unless noted otherwise, hard \xray\ refers to the $>2$ keV energy range in the rest frame, and \nh\ ($N_{\rm H,Gal}$) refers to the intrinsic (Galactic) neutral absorption column density. | \label{sec:results} \subsection{How Extreme are the Hard-\xray\ Spectral Slopes of WLQs?} \label{sec:L16} Type~1 quasars are known to exhibit a hard-\xray\ spectral slope of $\Gamma \sim 1.8-2.0$ across the Universe (e.g., Reeves \& Turner 2000; Page \et 2005; Piconcelli \et 2005; Shemmer \et 2005; Vignali \et 2005; Just \et 2007; Young \et 2009) that appears to be regulated by \lledd\ (Shemmer \et 2008). Based upon the well-known $\Gamma$-\lledd\ correlation, the small fraction of quasars with measured $\Gamma$ values of $\gtsim2.2$ are interpreted as sources that accrete close to or even above the Eddington limit (e.g., Risaliti \et 2009). A natural explanation for the mean $\Gamma$ value of $2.18\pm0.09$, measured for 18 \xray\ normal WLQs by Luo \et (2015) is that these sources lie at the extreme high end of the \lledd\ distribution in quasars. Quantifying, or constraining, their deviations from that distribution can be done by obtaining a deeper \xray\ observation for each individual source. Our data provide almost an order of magnitude increase in the number of \xray\ counts for radio-quiet and \xray\ normal WLQs, and they confirm the basic Luo \et (2015) finding. Table~\ref{tab:WLQfitresults} shows that most of our radio-quiet sources have extremely high $\Gamma$ values, and the mean $\Gamma$ value of these sources, $\left < \Gamma \right > = 2.30\pm0.06$, based on jointly fitting their spectra (Table~\ref{tab:joint_fitting}) is larger than, yet consistent within the errors with, the Luo \et (2015) value. In order to quantify the extremity of the $\Gamma$ values of these WLQs, we searched the literature and \xray\ archives to identify the most suitable comparison sample of quasars. The sample of Liu \et (2016; hereafter, L16) includes 1786 type~1 quasars observed with \xmm\ as part of the XMM-XXL-North survey; this is one of the largest samples of \xray\ detected quasars to date. Of these, 1731 sources are part of the SDSS Data Release 12 quasar catalog (P\^{a}ris \et 2017) which are covered in the Faint Images of the Radio Sky at Twenty cm (FIRST; Becker \et 1995) footprint. We further limited this sample by requiring each source to meet all of the following criteria: \begin{itemize} \item{pn counts $> 100$} \item{BAL flag equals zero for sources at $z>1.57$, according to the P\^{a}ris \et (2017) catalog} \item{radio-quiet sources having $R < 10$} \end{itemize} The first of these criteria ensures that sources have \xray\ data with comparable quality to our WLQ observations (a higher pn counts threshold would have limited the sample considerably; see Figure~\ref{fig:test1}). The second criterion is set to minimize the effects of \xray\ absorption (e.g., Gallagher \et 2006), and the third criterion is required for minimizing the potential contribution of a jet to the \xray\ emission (e.g., Miller \et 2011). There are 167 L16 sources that meet the first criterion, 48 of which are at $z>1.57$. One of these 48 sources is flagged as a BAL quasar by P\^{a}ris \et (2017) and is therefore removed from our sample. Based on the BAL fraction at $z>1.57$, we expect that $\sim 2-3$ BAL quasars may be present among the remaining 119 sources at $z<1.57$. Cross-matching with the FIRST catalog, using a 2$''$ search radius around the SDSS coordinates of each source (see, e.g., P\^{a}ris \et 2017), yielded 15 radio counterparts to the remaining 166 L16 sources. Two additional sources, out of 166, have FIRST detections with angular offsets of 2.2$''$ and 2.7$''$; we consider these to be physically related to the respective L16 sources, thus raising the total number of radio counterparts to 17. The $R$ values for the radio counterparts were derived by taking the FIRST flux densities at an observed-frame frequency of 1.4 GHz and extrapolating to a rest-frame frequency of 5 GHz, assuming a radio power-law continuum of the form $f_{\nu} \propto \nu^{-0.5}$ (Rector \et 2000). These flux densities were then divided by the flux densities from the $i$-band ($AB_{7672}$) magnitudes taken from the P\^{a}ris \et (2017) catalog and extrapolated to a rest-frame wavelength of $4400\,{\rm \AA}$, assuming an optical power-law continuum of the form $f_{\nu} \propto \nu^{-0.5}$ (Vanden Berk \et 2001). Only two of the 17 radio counterparts were found to be radio quiet; the other 15 were therefore culled from the comparison sample. For the 149 L16 sources that do not have radio counterparts, we computed upper limits on their $R$ values as described above, except that the radio fluxes were derived by multiplying the RMS radio flux at the SDSS position by a factor of 3. In order to meet our third criterion above and to ensure that only radio-quiet sources are considered for the comparison, we further excluded all sources that have upper limits on $R$ that are greater than 10. The final sample, hereafter the L16 comparison sample, includes 85 sources, 62 of which are at $z<1.57$ (if, instead, we used a more conservative constraint on the radio-undetected sources, by multiplying their RMS fluxes by a factor of 5, this would have reduced the sample size to 63 sources). Assuming the BAL quasar contamination fraction above, we can expect the L16 comparison sample to include not more than $\sim1$ BAL quasar at $z<1.57$. Figure~\ref{fig:test1} presents distributions of the redshift, luminosity, and number of pn counts for the L16 comparison sample. \begin{figure*}[t] \epsscale{0.9} \plotone{f6.pdf} \caption{Distributions of the redshifts (left), luminosities (middle), and number of pn counts (right) of the L16 comparison sample. Dashed lines represent the upper and lower limits of our WLQ sample in each panel.} \label{fig:test1} \end{figure*} L16 took a Bayesian approach to the \xray\ spectral analysis, using the {\sc bntorus} model.\footnote{This model includes an intrinsic power-law component and Compton scattering from absorbing material. When fitting unobscured quasars like those in our sample, the results of this model are in excellent agreement with those of the {\sc pexrav} model in XSPEC (Magdziarz \& Zdziarski 1995; see Brightman \et 2015 for more details).} Unlike our analysis in the $>2$ keV rest-frame energy range, L16 fit their spectra in the \textit{observed-frame} $0.5-8$ keV range of each source. By adopting this band, the L16 fitting procedure results in non-uniform sampling of their \xray\ spectra, given the wide range of source redshifts. One effect that may stem from employing this procedure is the measurement of an unrealistically high $\Gamma$ value. As we discuss below, this is most likely a consequence of including a soft-excess component in the spectral fitting. The distribution of the $\Gamma_{0.5-8~{\rm keV}}$ values, as measured by L16, for the L16 comparison sample appears in Figure~\ref{fig:test2}. In order to obtain a meaningful comparison between the L16 $\Gamma_{0.5-8~{\rm keV}}$ values and the $\Gamma$ values of our WLQs, we re-fitted each of our objects in the observed-frame $0.5-8$ keV range, this time with the Galactic-absorbed power-law and Compton reflection ({\sc phabs$*$pexrav}) model (a similar model with an additional intrinsic absorption component, {\sc zphabs}, was used for SDSS J0928$+$1848; see Section~\ref{sec:obs}). We ran {\sc pexrav} while fixing only the redshifts and the Galactic absorptions; all the other model parameters were free to vary. The best-fit $\Gamma_{0.5-8~{\rm keV}}$ values resulting from these fits\footnote{The corresponding best-fit relative Compton-reflection parameters have not yielded meaningful constraints on the Compton-reflection component in any of our WLQs, similar to results we present in Section~\ref{sec:compt_reflect} below.} appear in Column (11) of Table~\ref{tab:WLQfitresults}. The $\Gamma_{0.5-8~{\rm keV}}=4.51$ value for SDSS J1429$+$3859 appears to be unphysical, but can be explained by the indication of excess soft-\xray\ emission at \hbox{$<$ 2 keV} in the rest-frame (see Figure~\ref{fig:radio_quiet_spec}). Furthermore, we searched for a potential systematic offset between the method L16 used to measure their $\Gamma_{0.5-8~{\rm keV}}$ values and the analysis method we used to obtain the $\Gamma_{0.5-8~{\rm keV}}$ values of our WLQs. Therefore, we have reanalyzed the \xmm\ spectra of seven sources from the L16 comparison sample that had greater than 600 total counts per source. Such a threshold on the number of counts ensures that we compare L16 data sets of roughly matched quality to those of our radio-quiet WLQs (see Table~\ref{tab:obs_log}). We employed the same data reduction and analysis as described above on single L16 data sets of those seven sources; since each L16 source typically has 1--10 exposures per pointing with a range of angular offsets from the aimpoint, we used the data set with the longest exposure time in each case. As done in L16, we restricted the fitting range to 0.5--8 keV in the observed frame of each source, then fitted each data set once with a Galactic-absorbed power-law and Compton reflection ({\sc phabs$*$pexrav}) model and a second time with an added intrinsic absorption component {\sc (phabs$*$zphabs$*$pexrav}); based on $F$-tests, none of the spectra warranted a neutral intrinsic absorption component. The results of this analysis are given in Table~\ref{tab:L16_seven}. \textit{Column (1)} gives the \xray\ source identification string used by L16; \textit{Columns (2) and (3)} give the SDSS quasar name and corresponding \xmm\ observation ID number, respectively; \textit{Columns (4) and (5)} give the redshifts and number of pn counts taken from L16, respectively; \textit{Columns (6) and (7)} give the $\Gamma_{0.5-8~{\rm keV}}$ values from L16 and our {\sc pexrav} analysis, respectively. We found no significant systematic difference between the L16 $\Gamma_{0.5-8~{\rm keV}}$ values and the $\Gamma_{0.5-8~{\rm keV}}$ values we obtained for these seven sources.\footnote{We also analyzed the data set of the source from the L16 comparison sample with the largest value of $\Gamma_{0.5-8~{\rm keV}}$ (SDSS J022039.48$-$030820.3 with \hbox{$\Gamma_{0.5-8~{\rm keV}}= 2.91^{+0.08}_{-0.20}$}) in the same way and found a $\Gamma_{0.5-8~{\rm keV}}$ value of $2.66^{+0.18}_{-0.17}$. However, we note that this data set contains only 106 pn counts, and the source is detected close to the edge of the detector.} Similar to the analysis described in Section~\ref{sec:obs} for our WLQs, we also fitted those seven L16 sources with a Galactic-absorbed power-law model and found a systematic offset of $\approx +0.2$ between $\Gamma_{0.5-8~{\rm keV}}$({\sc pexrav}) and $\Gamma_{0.5-8~{\rm keV}}$(power-law), as may be expected, given that the {\sc pexrav} model attempts to include an additional component due to reflection from neutral material. Additionally, we note that the seven L16 sources with the largest number of counts are not identical to those with the highest $\Gamma_{0.5-8~{\rm keV}}$ values. Therefore, our results are not biased by sources with the steepest spectral slopes. Figure~\ref{fig:test2} shows how the $\Gamma_{0.5-8~{\rm keV}}$ values of our WLQs compare with those of the L16 comparison sample. \begin{deluxetable*}{lcccccc}[t] \tablecolumns{7} \tabletypesize{\scriptsize} \tablewidth{0pc} \tablecaption{Best-Fit Parameters of L16 Sources With Largest Number of Counts \label{tab:L16_seven}} \tablehead { \colhead{} & \colhead{} & \colhead{Observation} & \colhead{} & \colhead{} & \colhead{{\sc $\Gamma_{0.5-8~{\rm keV}}$}} & \colhead{{\sc $\Gamma_{0.5-8~{\rm keV}}$}} \\ \colhead{UXID}& \colhead{SDSS Name} & \colhead{ID} & \colhead{$z$} & \colhead{pn Counts} & \colhead{(from L16)} & \colhead{(from {\sc pexrav})} \\ \colhead{(1)} & \colhead{(2)} & \colhead{(3)} & \colhead{(4)} & \colhead{(5)} & \colhead{(6)} & \colhead{(7)} } \startdata N\_38\_68 & \object{SDSS~J021808.24$-$045845.2} & 0112371001 & 0.714 & 3308 & $2.27\pm{0.02}$ & $2.53\pm{0.09}$ \\ N\_38\_117 & \object{SDSS~J021817.45$-$045112.5} & 0112371001 & 1.083 & 1985 & $1.95\pm{0.03}$ & $1.88^{+0.09}_{-0.17}$ \\ N\_20\_50 & \object{SDSS~J022105.64$-$044101.5} & 0037982001 & 0.199 & 974 & $2.11\pm{0.03}$ & $2.16^{+0.09}_{-0.15}$ \\ N\_0\_30 & \object{SDSS~J022224.20$-$034757.3} & 0604280101 & 1.687 & 938 & $1.87^{+0.07}_{-0.06}$ & $1.97^{+0.19}_{-0.27}$ \\ N\_27\_19 & \object{SDSS~J022244.40$-$043347.0} & 0109520601 & 0.761 & 849 & $2.15\pm{0.05}$ & $2.30^{+0.16}_{-0.14}$ \\ N\_113\_13 & \object{SDSS~J022851.50$-$051223.1} & 0677590132 & 0.316 & 619 & $2.13\pm{0.04}$ & $2.17^{+0.12}_{-0.08}$ \\ N\_38\_79 & \object{SDSS~J021830.59$-$045622.9} & 0112371001 & 1.397 & 611 & $2.28\pm{0.05}$ & $2.63^{+0.49}_{-0.35}$ \enddata \end{deluxetable*} The $\Gamma_{0.5-8~{\rm keV}}$ values of four of our WLQs, all of which are radio-quiet, lie above the $3\sigma$ threshold at the high end of the $\Gamma_{0.5-8~{\rm keV}}$ distribution of the L16 comparison sample; similarly, the $\Gamma_{0.5-8~{\rm keV}}$ value of another radio-quiet WLQ lies at the $\sim2 \sigma$ threshold. Importantly, the average $\Gamma_{0.5-8~{\rm keV}}$ value of our six radio-quiet WLQs ($\Gamma_{0.5-8~{\rm keV}} = 2.89^{+0.45}_{-0.30}$) also lies above the $3\sigma$ threshold of the $\Gamma_{0.5-8~{\rm keV}}$ distribution of the L16 comparison sample (this average drops to \hbox{$\Gamma_{0.5-8~{\rm keV}} = 2.28^{+0.51}_{-0.57}$} when SDSS J1429$+$3859 is excluded). \begin{figure*}[t] \epsscale{0.8} \plotone{f7.pdf} \caption{Comparison of power-law photon indices measured in the observed-frame $0.5-8$ keV between our WLQs and sources from the L16 comparison sample. The unshaded histogram represents the L16 comparison sample, the dashed curve is the best-fit Gaussian distribution for this histogram, and the hatched and solid black bars represent our radio-intermediate and radio-quiet WLQs, respectively.} \label{fig:test2} \end{figure*} In order to check whether the $\Gamma_{0.5-8~{\rm keV}}$ values of our WLQs, as a group, are significantly higher than those of typical quasars, we ran a Mann-Whitney nonparametric rank test between the $\Gamma_{0.5-8~{\rm keV}}$ values of our six radio-quiet WLQs and those of the L16 comparison sample. We found that the two distributions are significantly different, with $>99.8\%$ confidence ($> 3\sigma$), one-tailed (another test with the exclusion of SDSS J1429$+$3859 resulted in the two distributions being significantly different with $>99.5\%$ confidence). We also ran a similar Mann-Whitney test between the $\Gamma_{0.5-8~{\rm keV}}$ distributions of our six radio-quiet and four radio-intermediate WLQs, and similarly found that the two distributions are significantly different at the $95\%$ confidence level. This result is also reflected in Table~\ref{tab:joint_fitting}. The lower $\Gamma$ values of the radio-intermediate WLQs, with respect to their radio-quiet counterparts, may be a manifestation of jet contributions to their \xray\ emissions (e.g., Miller \et 2011). Our results therefore indicate that, in the absence of potential \xray\ emission from a jet along our line of sight, WLQs have significantly higher hard \xray\ power-law photon indices than typical quasars. This result reinforces the idea that weak emission lines in quasars may be a direct consequence of a high Eddington fraction. In this respect, our results are in agreement with the Luo \et (2015) model, which suggests that the scale height of the inner accretion disk grows as a function of the accretion rate and acts as a filter that prevents highly ionizing photons from reaching the BELR. However, in order to establish a relationship between BELR line strength and the Eddington fraction across wide ranges of these parameters, the hard \xray\ power-law photon indices of a statistically meaningful sample of quasars should be measured accurately (see, e.g., Shemmer \& Lieber 2015). \subsection{A Soft Excess - Accretion Rate Connection?} \label{sec:soft_excess} The discrepancies in the $\Gamma$ values of our WLQs between those fitted in the $>2$ keV rest-frame band and those in the observed-frame $0.5-8$ keV band (see Table~\ref{tab:WLQfitresults}) could be attributed to the existence of soft \xray\ excess emission, at least in our lowest-redshift sources. The physical nature of this component is uncertain (Porquet \et 2004; Gierlinski \& Done 2004; Vasudevan \et 2014), yet it is present in many AGN spectra which makes it of interest to search for the existence of this component in our sources. In order to check whether any of our sources shows evidence for a soft excess, we extrapolated the best-fit Galactic absorbed power-law model (with added intrinsic neutral absorption for SDSS J0928$+$1848) obtained for rest-frame energies $>2$ keV (see Section~\ref{sec:obs} and Table~\ref{tab:WLQfitresults}) to the \hbox{$>0.3$ keV} observed-frame energies. All but one of our sources show $\chi$ residuals no greater than the $3\sigma$ level and, therefore, no indication of excess soft-\xray\ emission. Only one of our WLQs, SDSS J1429$+$3859, has an indication of soft excess emission with $\chi$ residuals up to $6\sigma$ (see Figure~\ref{fig:radio_quiet_spec}). This is our lowest-redshift WLQ. The non-detection of this feature in the other WLQs is not unexpected given their considerably higher redshifts and the $\sim0.2$ keV energy threshold of \xmm\ (see Shemmer \et 2008). In order to assess the effect of the putative soft excess on the photon index of SDSS J1429$+$3859, we performed an additional spectral fitting on this source in which a thermal component (the {\sc nlapec} model in {\sc XSPEC}) was added to the model employed in Section~\ref{sec:L16} (i.e., {\sc phabs$*$pexrav$+$nlapec}). This fitting resulted in a photon index value of \hbox{$\Gamma_{0.5-8~{\rm keV}} = 2.57^{+0.39}_{-0.68}$}. An $F$-test shows that the addition of the {\sc nlapec} component provides a significantly better fit, with $>90$\% confidence, and that the $\Gamma_{0.5-8~{\rm keV}}$ value is reduced considerably with respect to the one from Section~\ref{sec:L16} ($\Delta \Gamma_{0.5-8~{\rm keV}} \simeq 2$). A soft excess feature is expected to be more pronounced in sources with higher accretion rates (e.g., Done \et 2012). The fact that we detect a feature of this kind in one of our radio-quiet sources is in agreement with the idea that WLQs have extremely high accretion rates. However, \xray\ imaging spectroscopy of additional WLQs is required to establish such a connection. \subsection{Searching for Signatures of Compton Reflection and Iron-Line Emission} \label{sec:compt_reflect} We conducted a search for the existence of a Compton-reflection continuum as well as signatures of a neutral narrow \Ka\ emission line at rest-frame 6.4 keV in our WLQs in order to assess their potential effects on our photon index measurements. This was performed by fitting all the \xmm\ spectra for each source in the $>2$ keV rest-frame energy range with XSPEC, employing a Galactic absorbed power-law with a Compton-reflection continuum model (i.e., the {\sc pexrav} model in XSPEC, using a similar spectral fitting approach as the one performed in Section~\ref{sec:L16}), and a Galactic absorbed power-law with a redshifted Gaussian emission line model ({\sc zgauss} in XSPEC) for the \Ka\ emission line. The Gaussian rest-frame energy and width were fixed at \hbox{$E = 6.4$ keV} and \hbox{$\sigma = 0.1$ keV}, respectively. Table~\ref{tab:EW} lists the best-fit parameters from these fits. \textit{Column (1)} gives the SDSS quasar name; \textit{Column (2)} gives the rest-frame EW of the \Ka\ emission line; \textit{Column (3)} gives the relative-reflection component ($R_{\rm rel}$) of the Compton-reflection continuum expressed as \hbox{$R_{\rm rel} = \Omega/2\pi$}, where $\Omega$ is the solid angle subtended by the continuum source. Due to the relatively low quality of the SDSS J1012$+$5313 observation, this object was not included in this portion of the analysis. \begin{deluxetable*}{lcc}[t] \tablecolumns{3} \tabletypesize{\scriptsize} \tablewidth{0pt} \tablecaption{Compton Reflection and Iron Emission \label{tab:EW}} \tablehead { \colhead{} & \colhead{EW(\Ka)\tablenotemark{a}} & \colhead{} \\ \colhead{WLQ} & \colhead{(eV)} & \colhead{$R_{\rm rel}$\tablenotemark{b}} \\ \colhead{(1)} & \colhead{(2)} & \colhead{(3)} \\ \noalign{\smallskip}\hline\noalign{\smallskip} \multicolumn{3}{c}{Radio Intermediate} } \startdata \object{SDSS~J092832.87$+$184824.3}\tablenotemark{c} & $\le319$ & $\le5.8$ \\ \object{SDSS~J101204.04$+$531331.8} & \nodata & \nodata \\ \object{SDSS~J114153.34$+$021924.3} & $\le531$ & $\le4.6$ \\ \object{SDSS~J123132.37$+$013814.0} & $\le62$ & $\le6.1$ \\ \noalign{\smallskip}\hline\noalign{\smallskip} \multicolumn{3}{c}{Radio Quiet} \\ \noalign{\smallskip}\hline\noalign{\smallskip} \object{SDSS~J141141.96$+$140233.9} & $\le121$ & $\le96.7$ \\ \object{SDSS~J141730.92$+$073320.7} & $\le326$ & $\le50.1$ \\ \object{SDSS~J142943.64$+$385932.2} & $\le90$ & $\le163.5$ \\ \object{SDSS~J144741.76$-$020339.1} & $\le42$ & $\le12.3$ \\ \object{SDSS~J161245.70$+$511816.9} & $\le170$ & $\le9.6$ \\ \object{SDSS~J164302.03$+$441422.1} & $\le160$ & $\le3.4$ \enddata \tablecomments{Best-fit parameters of fitting each spectrum at the $>2$ keV rest-frame energy range with a model consisting of a Galactic absorbed power-law, a Compton-reflection component, and a neutral \Ka\ emission line.} \tablenotetext{a}{Rest-frame equivalent width of a neutral \Ka\ emission line at rest-frame $E = 6.4$ keV and a fixed width of $\sigma = 0.1$ keV.} \tablenotetext{b}{Relative Compton-reflection parameter.} \tablenotetext{c}{For this object we also included, in addition to the model above, an intrinsic neutral-absorption component.} \end{deluxetable*} Previous studies have shown trends where the EW of the narrow \Ka\ line decreases with \xray\ luminosity, i.e. the ``\xray\ Baldwin effect" (e.g., Iwasawa \& Taniguchi 1993; Ricci \et 2013), the origin of which is still unclear. Table~\ref{tab:EW} shows that we have not detected any statistically significant Compton-reflection continua nor any neutral \Ka\ emission in any of our sources. Due to the relatively high luminosities of our sources, these results are in agreement with the \xray\ Baldwin effect. However, our results can be referred to as not being sensitive enough to either confirm or rule out an \xray\ Baldwin Effect. Detection of (or placement of meaningful constraints on) narrow iron lines in such luminous quasars require considerably longer exposures with \xmm. \subsection{X-ray Variability} \label{sec:variability} Table~\ref{tab:WLQfitresults} shows that no unusual \xray\ variations are observed for any of our WLQs between their \chandra\ and \xmm\ epochs, separated by $\approx 1$ yr in the rest frame of each source. The differences between each pair of \aox\ values indicates \xray\ variations of up to a factor of $\sim3.5$ (in SDSS J1447$-$0203), which is consistent with the \xray\ variability of typical quasars having similar luminosities (e.g., Vagnetti \et 2013; Lanzuisi \et 2014). Therefore, we do not expect that our main results are affected by \xray\ variability. We present \xray\ spectroscopy of ten SDSS, \xray\ normal WLQs at \hbox{$0.928\leq z \leq 3.767$} that have sufficient \xray\ counts to allow basic measurements of their \xray\ spectra with \xmm\ observations. Six of these are radio-quiet and four are radio-intermediate. Our analysis provides measurements of the hard \xray\ photon index in these sources. We have compared these data with similar data for a carefully-selected sample of 85 radio-quiet type 1 quasars in order to quantify the extremity of the hard-\xray\ spectral slopes of WLQs with respect to typical quasars. The results of this comparison show that the photon indices of radio-quiet WLQs, as a group, constitute the $>3\sigma$ tail of the photon index distribution in quasars. The radio-intermediate WLQs have considerably lower photon indices which are comparable to those of the bulk of the quasar population; we suggest that \xray\ emission from a jet contributes to the harder \xray\ spectra in these sources. Considering the hard \xray\ power-law photon index as an Eddington-fraction indicator, our results imply that radio-quiet WLQs occupy the extreme high end of the accretion-rate distribution in quasars. Our lowest-redshift radio-quiet WLQ, SDSS J1429$+$3859, is our only source that exhibits soft excess emission which may be another manifestation of its high accretion rate. None of our sources shows signatures of Compton reflection, or the presence of a narrow iron line, and none shows unusual \xray\ variability. These results are in line with the high luminosity of our sources. In the near future, a more rigorous comparison with the $\Gamma$ values of our WLQs could be made using recent and deeper \xmm\ observations of a subset of the L16 comparison sample (Chen \et 2018). The new spectra of these sources will be fitted also at the $>2$ keV rest-frame range to conform with the spectral fitting of our WLQs. Sensitive \xray\ imaging spectroscopy of a large sample of quasars across a wide range of BELR line strength, redshift, and luminosity, is required for establishing connections between BELR line strength, Eddington fraction, and the prominence of a soft excess component in all quasars. | 18 | 8 | 1808.01306 |
1808 | 1808.06183_arXiv.txt | {Endogenous or exogenous, dry or wet, various scenarios have been so far depicted for the origin of water on our Solar System's rocky bodies. Hydrated silicates found in meteorites and in interplanetary dust particles together with observations of abundant water reservoirs in the habitable zone of protoplanetary disks are evidences that support aqueous alteration of silicate dust grains by water vapor condensation in a nebular setting.} {We investigate the thermodynamics (temperature and pressure dependencies) and kinetics (adsorption rates and energies, surface diffusion and cluster formation) of water adsorption on surfaces of forsterite grains, constraining the location in the solar nebula where aqueous alteration of silicates by water vapor adsorption could occur efficiently and lead to the formation of phyllosilicates. We analyze the astrophysical conditions favorable for such hydration mechanism and the implications for water on solid bodies.} {The protoplanetary disk model (ProDiMo) code is tuned to simulate the thermochemical disk structure of the early solar nebula at three evolutionary stages. Pressure, temperature and water vapor abundance within $1$~au from the protosun were extracted and used as input for a Monte Carlo code to model water associative adsorption using adsorption energies that resemble the forsterite [100] crystal lattice.} {Hydration of forsterite surfaces by water vapor adsorption could have occurred within the nebula lifetime already at a density of $10^8$~cm$^{-3}$, with increasing surface coverage for higher water vapor densities. Full surface coverage is attained for temperatures lower than $500$~K, while for hotter grain surfaces water cluster formation plays a crucial role. Between $0.5$ and $10$ number of Earth's oceans can arise from the agglomeration of hydrated $0.1$~$\mu$m grains into an Earth-sized planet. However, if grain growth occurs dry and water vapor processes the grains afterwards, this value can decrease by two orders of magnitude.} {This work shows that water cluster formation enhances the water surface coverage and enables a stable water layer to form at high temperature and low water vapor density conditions. Finally, surface diffusion of physisorbed water molecules shortens the timescale for reaching steady state, enabling phyllosilicate formation within the solar nebula timescale.} | After almost $40$ years of study, the origin of Earth's water is still strongly debated \citep{Drake2005}. One hypothesis is that Earth accreted from a mixture of dry and wet primary building blocks in which water was in the form of hydrous silicates (the \emph{wet-endogenous} scenario); another view supports dry accretion, with water delivered at a later stage during impact of hydrous asteroidal or cometary bodies (the \emph{exogenous} scenario). A direct evidence of aqueous alteration processes in the solar nebula is contained in Carbonaceous Chondrites (CCs). These "undifferentiated" meteorites are considered primitive Solar System objects along with Interplanetary Dust Particles (IDPs) and cometary grains because they show solar composition~\citep{Barrat2012}. Depending on the chemical and mineralogical composition and size of their parent body, several classes of CCs are defined \citep{Weisberg2006}. Among them, CI (Ivuna-like) group, CM (Mighei-like) group and CR (Renazzo-like) carbonaceous chondrites are the most hydrous varieties, with $3$ to $14$~wt\% of water content in CM and CR, and up to $15$~wt\% in CI~\citep{Alexander2010}. Most of the water in these chondrites is structurally bound in phyllosilicates\footnote{Layered silicate platelets with swelling properties~\citep{Schuttlefield2007}.} that formed during aqueous alteration of anhydrous minerals (e.g. olivine and pyroxene) very likely on the meteorites parent bodies~\citep{Brearley2006}. Recent mid-IR spectroscopy measurements revealed that the most aqueous altered samples are ($-$OH)-rich and almost depleted in olivine, the "dry" precursor mineral~\citep{Beck2014}. The high variability in the abundance and in the nature of these hydro-silicates in CCs indicates many levels of aqueous alteration and suggests different possible origins and evolution. Most of the models that have been developed over the last $30$ years are based on the fluid flow and liquid water-rock interaction on their parent bodies: previously accreted water ice melts, the fluid flows through various mineral matrices of different permeability and reacts with the anhydrous precursor mineral finally forming the hydrated products~\citep[see review by][]{Brearley2006}. The chemical composition of phyllosilicates forming by the flow of fluids may be controlled by the composition of the anhydrous precursor mineral and/or the composition of the aqueous solutions~\citep{Howard2011,Velbel2012}. On the other hand, hydrated silicates could have formed by direct condensation of water vapor within the terrestrial planets forming region. First models indicated that silicate hydration would be kinetically inhibited in a nebular setting~\citep{Fegleyprinn1989}. Using a Simple Collision Theory (SCT) model and the activation energy of $8420$~K ($70$~kJ mol$^{-1}$) as the amount of energy required to convert MgO into Mg$\mathrm{(OH)}_2$ (brucite) at $1$~atmosphere, they estimated the formation rates of serpentine and brucite and concluded that formation of hydrous silicates takes too long to occur by nebular condensation ($\mathrm{10}^5$ times the nebular life time of~$10^{13}$~s). However,~\citet{Ganguly1995} used the same SCT approach and estimated a shorter time scale for the hydration of olivine if a lower activation energy (about $3909$~K, that is $32.5$~kJ mol$^{-1}$) is assumed in the calculation. In the attempt to explain the presence of phyllosilicates fine-grained rims (FGRs) in the Murray CM chondrite,~\citet{Ciesla2003} draw a scenario in which, holding the $8420$~K of hydration energy, shock waves pass through an icy region of the nebula, the water vapor partial pressure is locally enhanced, thus increasing the collision rates of water molecules with the bare grains. Therefore, hydrated silicates can form much faster than the solar nebula life time, allowing a nebular origin of the chondrules and the phyllosilicates components as well.~\citet{Woitke2017} consider phyllosilicates in thermo-chemical equilibrium, and found that below $345$~K and at one bar, the dominant phyllosilicate is $\mathrm{Mg}_3\mathrm{Si}_2\mathrm{O}_9\mathrm{H}_4$ (lizardite) which replaces $\mathrm{Mg}_2\mathrm{SiO}_4$ in chemical + phase equilibrium. Phyllosilicates can retain water when heated up to a few hundreds of degrees centigrade~\citep{Beck2014, Davies1996} being able to preserve structural water also in the inner and warmer regions of a protoplanetary disk. Once agglomerated into planetesimals, phyllosilicates could be a potential source of water for terrestrial planets, in line with the wet-endogenous scenario. More recent computer simulations have studied water adsorption energy, binding sites and mechanisms (associative and/or dissociative) on forsterite surfaces and demonstrated that many Earth oceans could efficiently form \emph{in situ} under accretion disk conditions~\citep{Stimpfl2006, Muralidharan2008, King2010, Asaduzzaman2013, Asaduzzaman2015, Prigiobbe2013}. However, these modeling attempts possess some major uncertainties, namely a detailed temperature-pressure structure of the young solar nebula. In the exploratory rate-based warm surface chemistry model of~\citet[][submitted]{Thi2018}, water from the gas-phase can chemisorb on dust grain surfaces and subsequently diffuse into the silicate bulk. The phyllosilicate formation model was applied to a zero-dimensional chemical model and to a 2D protoplanetary disk model (ProDiMo) to investigate the formation of phyllosilicates in protoplanetary disks. In this work we test the possibility of water vapor condensation on bare forsterite grains in the region of the terrestrial planets prior to their accretion into planetesimals. In the endogenous scenario, we want to quantify how much water could have been delivered to planetesimal precursors of Venus, Earth and Mars $4.5$~Gyr ago. We used the astrophysical model for protoplanetary disks, ProDiMo~\citep{Woitke2009}, and the Monte Carlo (MC) simulation optimized for studying accretion of ice mantles on grains~\citep{Cazaux2010, Cazaux2015}, both described in sections~\ref{prodimoinputs} and~\ref{MCinput}. Using T Tauri disks observed in the Orion Nebula as templates, with ProDiMo we carefully build up our early solar nebula model at three time steps in the Sun's evolution. Temperature and pressure radial profiles and water vapor abundance are then extracted specifically for the midplane region close to the protosun. We use the MC simulations to calculate water adsorption rates. Surface coverages at different physical conditions were then estimated and used to quantify possible scenarios on the origin of water on terrestrial planets and meteorites. \begin{table} \centering \caption[]{List of the stellar and disk parameters} \label{parameter} \begin{tabular}{l l l} \hline\hline \noalign{\smallskip} Parameters & Symbol & Value\\ \noalign{\smallskip} \cline{1-3} \noalign{\smallskip} Stellar Mass & $M_{*}$ & $1$~M$_{\odot}$ \\ Stellar Luminosity & $ L_{*}$ & $11.02$,\,$2.17$,\,$0.46$\,L$_{\odot}$ \\ Effective temperature & $T_{\rm eff}$ & $4147$, $4282$, $4290$~K \\ UV luminosity\textsuperscript{a} &$L_{\rm UV}$ &$0.01$~L$_{\odot}$ \\ X-ray luminosity\textsuperscript{a}&$L_{\rm X}$ & $10^{30}$~erg s$^{-1}$ \\ \scalebox{0.95}[1]{Cosmic ray ionization rate} & CRI & $1.7~\times~10^{-17}$~s$^{-1}$\\ \noalign{\smallskip} \cline{1-3} \noalign{\smallskip} Disk mass & $M_{\rm d}$& $0.003$, $0.03$~M$_{\odot}$\\ Disk inner radius & $R_{\rm in}$ & $0.07$~au \\ Disk outer radius & $R_{\rm out}$ & $100$~au\\ Tapering-off radius & $R_{\rm tap}$ & $50$~au\\ Reference scale height & $H_ 0$ & $0.4$~au\\ Reference radius & $R_{\rm ref}$ & $10$~au\\ \noalign{\smallskip} \cline{1-3} \noalign{\smallskip} Dust settling turbulence & $\alpha$ & $0.1$ \\ Column density index\textsuperscript{a} & $\epsilon$& $1.0$\\ Dust-to-gas mass ratio\textsuperscript{a} & $\rho_{\rm d}/\rho_{\rm g}$ & $0.01$ \\ \noalign{\smallskip} \cline{1-3} \noalign{\smallskip} Min. size dust grain\textsuperscript{a} & a$_{\rm min}$ & $0.05~\mu$m\\ Max. size dust grain\textsuperscript{a} & a$_{\rm max}$ & $3000~\mu$m\\ Dust size distr. index\textsuperscript{a} & $p$ & $3.5$\\ Dust composition\textsuperscript{a}: & &\\ $\mathrm{Mg}_{0.7}\mathrm{Fe}_{0.3}$Si$\mathrm{O}_3$ && 60\%\\ amorph. carbon & &$15$\%\\ porosity & &$25$\%\\ Dust material density& $\rho_{\rm gr}$ & $2.076$~g \mbox{cm$^{-3}$}\\ \noalign{\smallskip} \cline{1-3} \end{tabular} \tablefoot{ProDiMo input parameters used to model six early solar nebula, for two disk mass values ($0.003$~M$_{\odot}$ and $0.03$~M$_{\odot}$) at three nebular ages ($0.2$~Myr, $1$~Myr and $10$~Myr), to which the following stellar parameters ($L_{*}$ and $T_{\rm eff}$) correspond respectively ($11.02$~L$_{\odot}$,~$4147~{\rm K}$), ($2.17$~L$_{\odot}$,~$4282~{\rm K}$) and ($0.46$~L$_{\odot}$,~$4290~{\rm K}$).\\ \tablefoottext{a}{Standard values from~\citet{Helling2014} and~\citet{Woitke2009}.} } \end{table} \begin{figure*}[h!tp] \centering {\includegraphics [width=\columnwidth]{TthreeagesMCFost.png}} {\includegraphics [width=\columnwidth]{PthreeagesMCFost.png}} \caption{{Midplane ($z/r=0.00$) dust temperature (left) and pressure (right) radial profiles extracted from ProDiMo models of the early solar nebula at three ages ($0.2$~Myr, $1$~Myr and $10$~Myr) and $0.03$~M$_{\odot}$ disk mass. }} \label{fig:PTprofiles} \end{figure*} | In this work we have investigated the efficiency of water vapor adsorption onto forsterite grains surfaces as one of the mechanisms that contributed to the water on Earth and in asteroids. The astrophysical disk model ProDiMo tailored to the solar nebula properties was combined with Monte Carlo simulations of water adsorption on a [100] forsterite crystal lattice. Water vapor abundances, temperature and pressure radial profiles identify the region in the warm disk midplane, between $0.07 - 0.3$~au from the protosun, where hydration of dust grains could have occurred. Several MC simulations were run to assess the dependency of the adsorption rate and the surface coverage on the parameter space identified by the pairs ($T, n_{\rm H_2O}$). Our MC models show that complete surface water coverage is reached for temperatures between $300$ and $500$~K. For hotter environments ($600$, $700$ and $800$~K), less than $30$\% of the surface is hydrated. At low water vapor density and high temperature, water cluster formation plays a crucial role in enhancing the coverage (see also Appendix~\ref{appendixIII}). The binding energy of adsorbed water molecules increases with the number of occupied neighboring sites, enabling a more temperature-stable water layer to form. Lateral diffusion of water molecules lowers the timescale for surface hydration by water vapor condensation by three order of magnitude with respect to an SCT model, ruling out any doubts on the efficiency of such process in a nebular setting. Finally, the amount of water potentially delivered on Earth drastically varies if we rely on a grain size distribution instead of single sized grains. Grain agglomeration and dust settling to the midplane, the initial steps for planetesimal formation, should clearly lead to a wide grain size distribution as the nebula evolves. In order to improve our initial estimates, detailed dust evolution models should be combined with the water adsorption efficiencies found here. In addition, dynamical simulations of grain growth are required to understand how agglomeration and collision processes affect the amount of water retained on the grain surfaces and how this competes with the diffusion timescale of water molecules into the bulk of the grains. | 18 | 8 | 1808.06183 |
1808 | 1808.01847_arXiv.txt | { We present the results of CANDELSz7, an ESO large program aimed at confirming spectroscopically a homogeneous sample of z$\simeq$6 and z$\simeq$7 star forming galaxies. The candidates were selected in the GOODS-South, UDS and COSMOS fields using the official CANDELS catalogs based on $H_{160}$-band detections. Standard color criteria, which were tailored depending on the ancillary multi-wavelength data available for each field, were applied to select more than 160 candidate galaxies at z$\simeq 6$ and z$\simeq$ 7. Deep medium resolution FORS2 spectroscopic observations were then conducted with integration times ranging from 12 to 20 hours, to reach a Ly$\alpha$ flux limit of approximately 1-3$\times 10^{-18}$ erg s$^{-1}$ cm$^{-2}$ at 3$\sigma$. For about 40\%\ of the galaxies we could determine a spectroscopic redshift, mainly through the detection of a single emission line that we interpret as Ly$\alpha$ emission, or for some of the brightest objects ($H_{160}\leq 25.5$) from the presence of faint continuum and sharp drop that we interpret as a Lyman break. In this paper we present the redshifts and main properties of 65 newly confirmed high redshift galaxies. Adding previous proprietary and archival data we assemble a sample of $\simeq 260$ galaxies that we use to explore the evolution of the Ly$\alpha$ fraction in Lyman break galaxies and the change in the shape of the emission line between $z\sim6$ and $z\sim7$. We also discuss the accuracy of the CANDELS photometric redshifts in this redshift range.} \authorrunning{L. Pentericci et al.} \titlerunning{CANDELSz7} \date{Received ; accepted} | The exploration of the reionization era is surely one of the most challenging and fascinating tasks of present-day extra-galactic astronomy. For the first time we can compare precise results from cosmic microwave background data from Planck \citep{adam16} to observations of primeval galaxies, when the Universe was still largely neutral. These observations help us to understand the exact time-line of the reionization process, how it proceeded spatially, and which were the sources that produced all or most of the ionizing photon budget. The general consensus seems to be that galaxies, and in particular the faintest systems, were those providing most of the ionizing radiation \citep[e.g.,][]{bouwens+16b,fink15}, although faint AGN might also have played a role \cite[e.g.,][]{giallongo+15}. \\ To understand the evolution of the reionization process, one of the key quantities we would like to measure is the fraction of neutral hydrogen present in the Universe and its evolution with cosmic time. In the future SKA (and maybe its precursors) will be able to detect directly the neutral hydrogen content in the early Universe by mapping the 21 cm emission. In the meantime we must rely on alternative observational probes, which allow us to set indirect constrains on the amount of neutral hydrogen. These include deep optical spectra of high redshift QSO where we can analyse the Gunn-Peterson optical depth (e.g. \citealt{fan2006}), the distribution of dark gaps (Chardin et al. 2018, McGreer et al. 2015), the analysis of damping absorption wings as in Schroeder et al. (2013) and the analysis of GRB spectra \citep{totani+14}. For Lyman break galaxies (LBGs) and Ly$\alpha$ emitters (LAEs) the most promising tools are studying the prevalence of Ly$\alpha$ emission in star-forming galaxies (\citealt[][hereafter LP14]{pentericci+14},\citealt[][hereafter LP11]{pentericci+11},\citealt{ono+12,treu+13,schenker+14,caruana+14,tilvi+14}), the evolution of the clustering and luminosity function of LAEs \citep{ouchi+10,tilvi+10,sobacchi+15}. \\ In particular, several groups have focused their attention on the presence of the Ly$\alpha$ line in samples of LBGs. While from redshift $\sim$2 to $\sim$6 the fraction of galaxies showing a bright \lya\ emission line seems to be increasing steadily \citep{stark+10,cassata+15}, there is a strong deficit of such lines in galaxies as we approach z$\sim$7 (LP14; \citealt{tilvi+14,treu+13,ono+12,schenker+12};LP11;\citealt{fontana+10}). This is so far one of the strongest and perhaps the most solid evidence that at z$\sim$7 the Universe is partially neutral, since neutral hydrogen can easily suppress the visibility of the line. It would be much harder to explain the observed drop in the Ly$\alpha$ fraction with a very rapid change in the physical properties of galaxies, such as the intrinsic dust content. The only alternative viable explanation is a sudden increase of the escape fraction of Lyman continuum photons \citep{mesinger15}; however despite the recent progress in the discovery of real Lyman continuum emitters (\citealt{shapley16,izotov+16b,vanzella+16}), this quantity remains elusive and assessing its evolution in the early Universe is extremely difficult. \\ The dramatic decrease of Ly$\alpha$ fraction at high redshift implies that the number of spectroscopically confirmed galaxies above z=6.5 is still very low. This emission line is at present one of the few possible redshift indicators in the reionization epoch, although the Carbon line emission is becoming a viable alternative, from transitions visible in the sub-mm (the [CII] $158 \mu m$ line e.g., \citealt{pentericci+16,bradac+17,smit17}), or in the UV domain (the CIII]1909\AA\ emission line, \citealt{stark+17,lefevre17,maseda17,stark15}). All the above results are based on small data-sets, and statistical fluctuations can be very large. This is particularly true for the faintest LBGs, since most previous observations focused on the brighter candidates ($M_{UV}<-20.5$). The results also came from very heterogeneous observational efforts, in terms of wavelength coverage, sample detection and selection, integration times etc: it is therefore hard to combine them and to assess their global statistical significance. To overcome these problems, in 2013 we started an ESO Large Program with FORS2 (Program ID 190.A-0685) to assemble observations of a much larger and homogeneously selected sample of galaxies at z$\sim$6 and $\sim$7, to place much firmer constraints on the decrease in the visibility of Ly$\alpha$ emission between these two epochs, and to assess if and how this decrease depends on galaxies' brightness. This is of particular relevance since the visibility of Ly$\alpha$-emitting galaxies during the Epoch of Reionization is controlled by both diffuse HI patches in large-scale bubble morphology and small-scale absorbers. Improved constraints on the relevant importance of these two regimes can be obtained by analyzing the full UV-luminosity dependent redshift evolution of the Ly$\alpha$ fraction of Lyman break galaxies \citep{kaki16}. In this paper we will present the observations and the results of our large program, while deferring a full analysis of the properties of the galaxies and on the constraints to reionization models to other papers. In Section 2 we describe the target selection; in Section 3 the observations and data reduction procedure; in Section 4 we present the results and in Section 5 we discuss the observational properties of the detected galaxies. In a companion paper (De Barros et al. 2017), we have discussed the physical properties of the redshift 6 sample, while in \citet{cast17} we have investigated the nature of the GOODS-South and UDS z$\sim$7 target, to probe possible physical differences between \lya\ emitting and non-emitting sources. Finally, in an upcoming work we will discuss more extensively the implication of our results for the reionization epoch. We adopt a $\Lambda$-CDM cosmological model with $\mathrm{H}_\mathrm{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_\mathrm{m} = 0.3$ and $\Omega_\Lambda= 0.7$. All magnitudes are expressed in the AB system \citep{okegunn83}. \begin{figure} \includegraphics[width = 9cm,clip=]{goods.png} \label{goodsfig} \caption{The layout of the two FORS2 masks in the GOODS south field overplotted on the CANDELS $H_{160}$ image} \end{figure} | We have presented the results of an ESO spectroscopic large program aimed at exploring the reionization epoch by observing a large and homogeneous sample of star forming galaxies at redshift between 5.5 and 7.2 to set firm constraints on the evolution of the \lya\ emission fraction at this epoch. Galaxies were selected from the $H_{160}$-band CANDELS catalogs in the GOODS-South, UDS and COSMOS fields, using standard color criteria and/or the official CANDELS photometric redshifts. Spectroscopic observations of 167 high redshift galaxies were carried out with FORS2, using a medium resolution red grating. In 67 objects we could determine a redshift, mostly from the presence of a single \lya\ emission line or, in few cases, from the detection of continuum flux with a sharp drop that we interpret as the Lyman break. Two galaxies are low redshift interlopers. Overall, the success rate for the identification of the high redshift targets for which data could be reduced in a satisfactory way, is 40\%. Our sample increases substantially the number of sources with secure spectroscopic redshifts in the CANDELS fields, especially at z>6.5, including 3 new galaxies at z>7. With the newly confirmed galaxies, as well as previous spectroscopic redshifts available in the same fields, we evaluated the accuracy of the CANDELS photometric redshifts at $z \geq 6$. We found that the fraction of catastrophic outliers is 14\%, i.e. more than 3 times higher than for the lower redshift galaxies in the rest of the CANDELS catalog, where it is only $\sim$3\%. After removing the outliers, the rms uncertainty is 0.036. We also found that photometric redshifts are in general underestimated for galaxies with $H_{160}>27$ and $z>6.8$, probably due to the presence of a strong \lya\ emission line that influences the broad band photometry, and which is not taken into account in the photo-z models employed. Using our medium resolution spectra we have analysed the average shape of the \lya\ line by creating spectral stacks in two redshift bins. We found that at z>6.5 the blue side of the \lya\ emission line is completely erased and it is consistent with the spectral resolution, while at lower redshift a fraction of the blue wing is still transmitted. The \lya\ emission has a smaller FWHM and is slightly more asymmetric at z$\sim$7 compared to z$\sim$6. Finally we have evaluated the distribution of the \lya\ rest-frame EW using the new detections as well as the accurate upper limits determined through extensive simulations, for all the objects where no emission line is observed. The fraction of \lya\ emitters that we measure at $z=6$ is consistent with previous determinations only for REW$\leq 20\AA$, and it is below for larger REW, with a difference that reaches a factor larger than 2 at the highest REWs probed. The fraction of \lya\ emitters at $z\sim6$ is actually consistent with the one determined at $z\sim5$ (e.g., \citealt{stark+10}) indicating a possible flattening in the evolution with redshift between z$\sim$5 and z$\sim$6, instead of a steady increase up to z$\sim$6. The frequency of \lya\ drops by an average factor of 2 between z$\sim$6 and z$\sim$7 for galaxies with $M_{UV}>-20.25$. Overall our results might indicate a possibly slower and more extended reionization process, and in a future paper we will use our new data to set more stringent constraints on the models. In particular, it was shown by \cite{kaki16} that improved constraints can be derived by analyzing the full $M_{UV}$-dependent redshift evolution of the \lya\ fraction of Lyman break galaxies, such that it would be possible to distinguish between the effect of a `bubble' model, where only diffuse H I outside ionized bubbles is present, and the `web' model, where H I exists only in over-dense self-shielded gas. | 18 | 8 | 1808.01847 |
1808 | 1808.08966_arXiv.txt | Radio images of the Galactic Center supermassive black hole, Sagittarius~A$^\ast$ (\sgra), are dominated by interstellar scattering. Previous studies of \sgra\ have adopted an anisotropic Gaussian model for both the intrinsic source and the scattering, and they have extrapolated the scattering using a purely $\lambda^2$ scaling to estimate intrinsic properties. However, physically motivated source and scattering models break all three of these assumptions. They also predict that refractive scattering effects will be significant, which have been ignored in standard model fitting procedures. We analyze radio observations of \sgra\ using a physically motivated scattering model, and we develop a prescription to incorporate refractive scattering uncertainties when model fitting. We show that an anisotropic Gaussian scattering kernel is an excellent approximation for \sgra\ at wavelengths longer than 1\,cm, with an angular size of $(1.380 \pm 0.013) \lambda_{\rm cm}^2\,{\rm mas}$ along the major axis, $(0.703 \pm 0.013) \lambda_{\rm cm}^2\,{\rm mas}$ along the minor axis, and a position angle of $81.9^\circ \pm 0.2^\circ$. We estimate that the turbulent dissipation scale is at least $600\,{\rm km}$, with tentative support for $r_{\rm in} = 800 \pm 200\,{\rm km}$, suggesting that the ion Larmor radius defines the dissipation scale. We find that the power-law index for density fluctuations in the scattering material is $\beta < 3.47$, shallower than expected for a Kolmogorov spectrum ($\beta=11/3$), and we estimate $\beta = 3.38^{+0.08}_{-0.04}$ in the case of $r_{\rm in} = 800\,{\rm km}$. We find that the intrinsic structure of \sgra\ is nearly isotropic over wavelengths from 1.3\,mm to 1.3\,cm, with a size that is roughly proportional to wavelength: $\theta_{\rm src} \sim (0.4\,{\rm mas}) \times \lambda_{\rm cm}$. We discuss implications for models of \sgra, for theories of interstellar turbulence, and for imaging \sgra\ with the Event Horizon Telescope. | The compact radio source at the Galactic Center, Sagittarius A$^\ast$ (\sgra), was discovered in 1974 \citep{Balick_Brown_1974}. Within two years, observers had deduced that the radio image was dominated by scatter broadening caused by the ionized interstellar medium (ISM) based on an observed scaling of image size with the squared observing wavelength, $\theta \propto \lambda^2$ \citep{Davies_1976}. In the decades since the initial discovery of \sgra, knowledge of its scattering properties has continually improved, but scattering uncertainties remain the primary limitation in determining the intrinsic structure of \sgra\ at wavelengths longer than a few millimeters. Motivated by the $\theta \propto \lambda^2$ scaling and approximately Gaussian image, many observers have sought to accurately measure the image of \sgra\ at a wide range of radio wavelengths, seeking to constrain the scattering law at long wavelengths (where the scattering dominates) and then to deconvolve its effects at shorter wavelengths to estimate the intrinsic source parameters. An advantage of treating both the source and the scattering as Gaussian is that the scattered image is then also a Gaussian because the time-averaged scattering acts as a convolution \citep[see, e.g.,][]{Coles_1987,GoodmanNarayan89,Johnson_Gwinn_2015}. Consequently, many techniques have been developed to accurately estimate Gaussian image parameters for \sgra\ from interferometric data, including image-domain parameter estimation \citep{Bower_2006}, model fitting using only closure quantities \citep{Bower_2004,Shen_2005,Bower_2014b,Ortiz_2016,Brinkerink_2016}, and self-calibration \citep{Doeleman_2001,Lu_2011,Ortiz_2016}. In addition, many techniques have been applied to ensure conservative estimates of parameter uncertainty, including standard exploration of the chi-squared hypersurface \citep[e.g.,][]{Bower_2014b}, Monte Carlo approaches \citep{Ortiz_2016}, and bootstrap approaches using multi-epoch data \citep[e.g.,][]{Lu_2011}. Nevertheless, the reported sizes and position angles still have significant unresolved discrepancies \citep[see][]{Psaltis_2015}. In addition to the simplified scattering model, a major missing component from all these previous studies has been refractive scattering effects. Refractive scattering will distort the instantaneous image, giving systematic departures from the ensemble-average image that are independent of observing quality \citep{Blandford_Narayan_1985}. Refractive scattering also introduces substructure in the image, which contributes additional ``refractive noise'' to interferometric measurements on baselines that resolve the image \citep{NarayanGoodman89,GoodmanNarayan89,Johnson_Gwinn_2015}. Recently, refractive noise was discovered in 1.3\,cm observations of \sgra\ \citep{Gwinn_2014}, suggesting that it may contribute significantly to the error budget when fitting Gaussian models. Refractive noise is especially problematic because the longer baselines, which are most affected, are also the most sensitive to compact structure; their measurements are what dominate Gaussian model fits. Because refractive noise tends to bias long-baseline visibility amplitudes upward, detections interpreted without a noise budget for refractive substructure will tend to imply artificially compact structure \citep[see, e.g.,][]{Johnson_2016RA,Pilipenko_2018}. Thus, refractive scattering effects are essential to include when fitting models to interferometric data, and they contribute in multiple ways, both by modulating the ``true'' instantaneous image size and orientation and by adding a new type of ``noise'' to interferometric measurements. Here, we analyze archival observations of \sgra\ at wavelengths from 1.3\,mm (EHT) to 30\,cm (VLA). We develop a framework to efficiently incorporate refractive noise into parametric model fitting, and we show how to isolate components of the refractive noise that may be absorbed into fitted model parameters (e.g., refractive flux modulation and image wander). We constrain a physically motivated scattering model \citep{Psaltis_2018}, which generically produces Gaussian scatter-broadening that scales as $\lambda^2$ in the limit $\lambda \rightarrow \infty$, but which differs at short wavelengths because of a finite inner scale $r_{\rm in}$ of the interstellar turbulence with an associated power-law index $\alpha$. In addition to these two parameters, the scattering model depends on the Gaussian scatter broadening in the long-wavelength limit, which we parameterize via the major axis full width at half maximum (FWHM) $\theta_{\rm maj,0}$, minor axis FWHM $\theta_{\rm min,0}$, and major axis position angle $\phi_{\rm PA}$, all specified at a reference wavelength $\lambda_0$ (we use $\lambda_0 \equiv 1\,{\rm cm}$). We estimate uncertainties in our parameter estimates by fitting representative ensembles of synthetic datasets that match the baseline coverage and sensitivity of the observations. These synthetic datasets are created using numerical simulations of the scattering and also include wavelength-dependent systematic gain calibration uncertainties to simulate imperfect amplitude and phase calibration. This approach allows us to incorporate thermal noise, refractive uncertainties, and systematic calibration errors in the overall error budget, and to verify that our model fitting is not biased by any of these effects or by the anisotropic baseline coverage. Using our estimated scattering model, we compute the wavelength-dependent intrinsic size of \sgra. We begin, in \S\ref{sec::Scattering_Model}, with a brief review of scattering theory. Next, in \S\ref{sec::Model_Fitting_Procedure}, we describe our procedure to fit individual observations and motivate how we can use the full set of observations to constrain the scattering model. In \S\ref{sec::Observations}, we provide details about the observations used to constrain the scattering model and give the results of Gaussian fits to each. In \S\ref{sec::Composite_Constraints}, we derive our parameter estimates and uncertainties for the scattering model, describe the expected scattering properties, and estimate the intrinsic source size of \sgra. In \S\ref{sec::Discussion}, we discuss implications for models of \sgra, implications for theories of interstellar turbulence, consequence of unmet assumptions in our approach, and prospects for continued study of \sgra. We summarize our findings in \S\ref{sec::Summary}. | \end{figure*} \begin{figure*}[t] \centering \includegraphics[width=\textwidth]{Kband_MultiIF_fits-crop.pdf}\\ \caption { Gaussian parameters for the four sub-bands of the $\lambda=1.26\,{\rm cm}$ observations fitted independently. The scatter among sub-bands when fitting the major axis scattering law is approximately $2\,\mu{\rm as}$, or roughly $0.08\%$ of the image size. However, the uncertainty from refractive fluctuations of the image size (which will give nearly identical bias for each sub-band) is approximately ${\approx}\, 20\,\mu{\rm as}$, or $1\%$ of the image size. Thus, the estimated major axis uncertainty relative to the ensemble average size is dominated by refractive image distortion. The close agreement with a $\lambda^2$ scaling law (shown in red) strongly suggests that intrinsic structure is heavily subdominant to scatter broadening at this wavelength and also that the inner scale must be larger than the diffractive scale, $ r_{\rm in} \gsim 300\,{\rm km}$, because otherwise the wavelength dependence of the scattering kernel steepens, $\theta \propto \lambda^{1+\frac{2}{\alpha}}$. } \label{fig::Kband_MultiIF} \end{figure*} \begin{figure}[t] \centering \includegraphics[width=\columnwidth]{KBand_Images_wResidual.pdf}\\ \caption { (left) Reconstructed image at $\lambda=1.3\,{\rm cm}$. The color scale is linear and ranges from $0 - 0.33\,{\rm Jy}/{\rm mas}^2$. The dashed blue ellipse shows the Gaussian half-maximum contour from model fitting; the solid black line shows the half-maximum contour of the reconstructed image. Substructure is apparent through the subtle distortions from a smooth Gaussian image. (right) Residual image after subtracting the best-fit Gaussian image. The color scale is linear, and the range extends over $\pm 0.033\,{\rm Jy}/{\rm mas}^2$. } \label{fig::KBand_Image} \end{figure} \begin{figure*}[t] \centering \includegraphics[width=0.98\textwidth]{Kband_Data_Summary-crop.pdf} \caption { VLBA+GBT observations at $\lambda=1.3\,{\rm cm}$ (see \S\ref{sec::Obs_Kband}). Panels are as described for Figure~\ref{fig::Xband_Data_Summary}. Larger points in the left panel denote baselines to the GBT. For clarity, we only show long baselines to GBT on this plot (omitting long baselines to NL and HN, which sample similar $(u,v)$ coordinates but with less sensitivity). } \label{fig::Self-cal_wRefractive} \end{figure*} \subsection{VLBA Observations at 1.3cm} \label{sec::Obs_Kband} We analyzed observations at $\lambda=1.3\,{\rm cm}$ taken with the VLBA+GBT in 2014. These observations were analyzed in \citet{Gwinn_2014}, which reported the initial discovery of refractive substructure in \sgra. As with the 3.6\,cm data, these observations recorded four contiguous $128\,{\rm MHz}$ channels, spanned approximately 3.5\,hours, and used NRAO~530 as a calibration source. They include strong detections to the VLBA antennas at North Liberty (NL) and Hancock (HN) in addition to the sites noted in \S\ref{sec::Obs_Xband}. After a global fringe search in AIPS \citep{Greisen_2003}, we averaged the data in frequency and in 30-second intervals before Gaussian fitting. However, we averaged the data to four 128\,MHz sub-bands and separately analyzed each. For these data, the overall baseline coverage is well matched to the scattered image, and we can tightly constrain the Gaussian parameters separately within each sub-band. For each, the thermal uncertainty on the major axis size is less than $0.1\%$, and we clearly identify the $\lambda^2$ scaling of image size across the four sub-bands (see Figure~\ref{fig::Kband_MultiIF}). However, the uncertainty from refractive distortion is an order of magnitude larger, so we can only constrain the ensemble-average FWHM to within approximately 1\%. For these data, there is sufficient baseline coverage to reliably synthesize an image. Figure~\ref{fig::KBand_Image} shows a maximum entropy image reconstruction using the \texttt{eht-imaging} library \citep{Chael_2016}. The effects of refractive substructure are evident in substructure of the image. However, the effects of substructure are most striking in the visibility domain, where long baselines from GBT to the inner VLBA give strong detections on baselines for which the Gaussian visibility contribution is negligible. Figure~\ref{fig::Self-cal_wRefractive} shows the final self-calibrated visibilities, following the procedure described in \S\ref{sec::Obs_Xband}. \subsection{KaVA Observations at 7mm} \label{sec::KaVA_Observations} The KaVA array has been conducting regular monthly monitoring of \sgra\ at 7\,mm since September 2014 as part of the KaVA AGN large program \citep{Kino_2015,Zhao_2017}. The KaVA baselines range from 300 to 2300\,km and provide excellent $(u,v)$ coverage for \sgra\ observations (Figure~\ref{fig::KaVA_Data}; see also~\citet{Akiyama_2014}). In particular, the KaVA coverage along the North-South direction is significantly better than VLBA coverage at this frequency, so the KaVA data are better suited to estimate the minor axis size. We analyzed data from the experiment r14308a, which were obtained in November 2014. The data were recorded with 256\,MHz total bandwidth, spanned 5.5~hours, and had an on-source time for \sgra\ of 220~minutes. NRAO~530 and two nearby SiO masers (OH 0.55-0.06, VX~Sgr) were observed as calibrators \citep{Cho_2017}. The correlated data were analyzed with AIPS in a standard pipeline. Most stations had good fringe detections in this experiment. After a global fringe search, the data were averaged in 30-second intervals and across the entire bandwidth for Gaussian model-fitting. See \citet{Zhao_2018} for more details of the monitoring and data analysis. Figure~\ref{fig::KaVA_Data} shows the data from this observation, after our Gaussian model fitting and self-calibration. For these observations, the contribution of renormalized refractive noise is insignificant and is only comparable to the thermal noise on the longest baselines. \begin{figure*}[t] \centering \includegraphics[width=0.98\textwidth]{KaVA_Data_Summary-crop.pdf} \caption { KaVA observations at $\lambda=0.7\,{\rm cm}$ (see \S\ref{sec::KaVA_Observations}). Panels are as described for Figure~\ref{fig::Xband_Data_Summary}. } \label{fig::KaVA_Data} \end{figure*} \subsection{VLBA+LMT Observations at 3.5mm} For $\lambda=3.5\,{\rm mm}$, we analyzed data from the first VLBI observations using the LMT in concert with the VLBA. These observations recorded 480\,MHz of bandwidth and spanned 7.5~hours (with approximately 3.4~hours on \sgra). We averaged the data in 10-second intervals and across the full bandwidth before Gaussian model fitting. For additional details about these observations, see \citet{Ortiz_2016}, who originally reported and analyzed them. Using our Gaussian fitting procedure, we found values and uncertainties very close to those reported in \citet{Ortiz_2016} using self-calibration. This agreement is expected because the only significant adaptation in our current approach is to include refractive noise, and the renormalized refractive noise is less than the thermal noise for all points. More recent data, reported by \citet{Brinkerink_2016}, also includes the GBT and shows marked non-Gaussianity in the closure phases. For these data, which achieve significantly better sensitivity, including refractive noise in the error budget for model fitting may be significant. \subsection{EHT Observations at 1.3mm} \label{sec::EHT} Since 2007, the Event Horizon Telescope (EHT) has observed \sgra\ using a 1.3\,mm VLBI array with stations in California, Arizona, and Hawaii. Recently, \citet{Lu_2018} reported observations that included a fourth station (in Chile). With only 3-4 stations, there is insufficient baseline coverage to create an image. In addition, the measured visibilities are markedly non-Gaussian \citep{Johnson_2015,Fish_2016,Lu_2018}, as expected because of complex intrinsic structure from optically thin emission near the black hole. Nevertheless, even under these circumstances, the image FWHM is still a meaningful quantity that is reliably constrained by sparse coverage because it represents universal behavior of the interferometric visibility function on short baselines (see \S\ref{sec::Gaussian_Model_Assumption}). Thus, we will now estimate this characteristic FWHM and its uncertainty at 1.3\,mm using previously published EHT data. Because the EHT baseline joining CARMA-SMT (California-Arizona) does not significantly resolve \sgra\ at 1.3\,mm, current EHT measurements primarily constrain the source size in the direction of the California-Hawaii and Arizona-Hawaii baselines, close to East-West (i.e., roughly along the major axis of the scattering kernel). Early detections were consistent with a Gaussian image having a FWHM of approximately $40\,\mu{\rm as}$ \citep{Doeleman_2008,Fish_2011}. However, more recent measurements with improved sensitivity and calibration find visibility amplitudes on the shortest Hawaii baselines (California-Hawaii) that are strongly inconsistent with the Gaussian model \citep{Johnson_2015,Lu_2018}. Thus, the appropriate FWHM is not that of the Gaussian fits, which are incompatible with the data, but can instead be estimated by computing the characteristic FWHM of models that do fit the short- and intermediate baseline visibility amplitudes. One such model is an annulus. The fitted annulus in \citet{Johnson_2015} gives a characteristic FWHM of $58.5\,\mu{\rm as}$ for the intrinsic source (as defined in Eq.~\ref{eq::fwhm_maj}). For comparison, the annulus model from \citet{Doeleman_2008} gave $51.5\,\mu{\rm as}$. Two-Gaussian model fits that also include closure phase measurements and baselines to APEX give FWHMs of $55.2\,\mu{\rm as}$ and $60.4\,\mu{\rm as}$ along the East-West direction or $62.5\,\mu{\rm as}$ and $60.5\,\mu{\rm as}$ along the major axis of the scattering \citep[for Models A and B of][]{Lu_2018}. Because the East-West scatter-broadening is ${\lsim}\,20\,\mu{\rm as}$ at this frequency, our revisions to the scattering model and remaining uncertainties have little effect on the estimated intrinsic FWHM. The uncertainties are instead dominated by the sparse baseline coverage, and we estimate a plausible range of $51-63\,\mu{\rm as}$ for the FWHM of the intrinsic source along the major axis of the scattering based on the span of these fitted models. Note that this range extends beyond the expected diameter of the black hole shadow ($51 \pm 3\,\mu{\rm as}$), so it does not necessitate that the accretion flow is viewed at large inclination. Accounting for both source and scattering uncertainties, we adopt a plausible range of $53-66\,\mu{\rm as}$ for the FWHM of the scattered image of \sgra\ at 1.3\,mm along the major axis of the scattering kernel. The North-South FWHM at 1.3\,mm is comparatively poorly constrained. \citet{Krichbaum_1998} reported detections of \sgra\ at $\lambda=1.4\,{\rm mm}$ on the baseline joining Pico Veleta and an antenna of the IRAM interferometer at Plateau de Bure. These observations had a baseline length $|\mathbf{u}| \approx 0.7 \times 10^9$, but the baseline was aligned close to East-West (position angle approximately $70^\circ$ East of North). \citet{Lu_2018} have recently reported detections at $\lambda=1.3\,{\rm mm}$ on baselines from APEX to California and Arizona, which are oriented close to North-South, but these heavily resolve the source. Thus, they are unreliable for estimating a FWHM using Eq.~\ref{eq::fwhm_maj} or for computing the second moment of a fitted model. Instead, we estimate a maximum size of the source along the scattering minor axis by requiring that the SMT-CARMA baseline amplitude be at least 80\% of the zero-baseline flux density over the GST range from $-0.5 - 4.0$~hours, as is supported by both a priori calibration \citep{Lu_2018} and polarization arguments \citep{Johnson_2015}. For a major axis FWHM of ${\sim}60\,\mu{\rm as}$, this requirement gives an upper limit to the minor axis FWHM of approximately $90\,\mu{\rm as}$. To obtain a corresponding lower limit, we require that the correlated flux density on the SMT-APEX baselines for the scattered image never exceeds 10\% of the zero-baseline flux density over the GST range from $0.0-2.5$~hours \citep[otherwise it would exceed measurements on this baseline;][]{Lu_2018}. This constraint only requires that the scattered source have a minor axis FWHM that exceeds $25\,\mu{\rm as}$. Combining these limits, we obtain a plausible range for the FWHM along the scattering minor axis direction of $25-90\,\mu{\rm as}$ (of course, the scattering position angle need not correspond to that of the scattered or unscattered image at 1.3\,mm). Finally, we note that the EHT has detected persistent non-zero closure phases of \sgra\ on the California-Arizona-Hawaii triangle, demonstrating that the scattered image structure is not point symmetric \citep{Fish_2016}. However, these results do not imply that the intrinsic or scattered FWHM is asymmetric because the non-zero closure phases may be produced by image substructure. For instance, Model~B in \citet{Lu_2018} fits both the visibility amplitudes and closure phases but has little asymmetry in the FWHM, with major and minor axes FWHMs of $60.5\,\mu{\rm as}$ and $60.3\,\mu{\rm as}$, respectively. \label{sec::Discussion} \subsection{The Intrinsic Structure of Sgr A*} Our estimates for the approximately linear wavelength dependence of intrinsic size are typical for stratified emission in synchrotron self-absorbed systems \citep[e.g.,][]{Blandford_Konigl_1979,Falcke_Markoff_2000,Davelaar_2018}, and they are plausible for both disk- and jet-dominated models for the radio emission of \sgra. However, the lack of asymmetry in the inferred intrinsic size and the stable position angle of the scattered image both argue against intrinsic structure that is highly asymmetric for $3.5\,{\rm mm} \lsim \lambda \lsim 1.3\,{\rm cm}$. For instance, the model of \citet{Falcke_Markoff_2000} predicts an image asymmetry of roughly $4{:}1$, and recent GRMHD simulations of jets show asymmetry of ${\sim}2{:}1$ at 7\,mm \citep{Davelaar_2018}. The intrinsic size of \sgra\ we find is qualitatively consistent with RIAF models \citep[e.g.,][]{Ozel_2000,Yuan_2003,Yuan_Narayan_2014}, which are more plausible for producing a nearly isotropic image at wavelengths as long as 1.3\,cm. Recent GRRMHD simulations show good agreement with our estimated size at 1.3\,mm and also show a similar size trend, but they generally underpredict the size at 7\,mm and 1.3\,cm \citep{Chael_2018}, perhaps highlighting the contribution from a non-thermal population of electrons. Note that there has not been consistency in how image FWHM from simulations is defined. For comparison with Gaussian image sizes reported here and elsewhere, simulations should compute the image FWHM from the second moment along principal axes of the image brightness distribution (see \S\ref{sec::Gaussian_Model_Assumption}). While some simulation papers have adopted this convention for comparisons \citep[e.g.,][]{Moscibrodzka_2009,Moscibrodzka_2012,Davelaar_2018,Chael_2018}, others have developed ad hoc definitions for the reported image size \citep[e.g.,][]{Ozel_2000,Falcke_Markoff_2000,Psaltis_2015b,Chan_2015} or do not state their procedure for estimating the size. An alternative is to fit or compare simulations directly to measured interferometric visibilities \citep[e.g.,][]{Broderick_2009,Dexter_2010,Pu_2016,Kim_2016,Broderick_2016,Gold_2016}. \subsection{Implications for Interstellar Scattering} We now reevaluate our assumptions for the scattering of \sgra, and we discuss implications of our findings. \subsubsection{The Outer Scale of Turbulence} \label{sec::Outer_Scale} All of our calculations and model fits have assumed that the outer scale of turbulence is effectively infinite. We now evaluate this assumption a posteriori. In particular, \citet{Goldreich_2006} estimated that $r_{\rm out} \lsim 10^{11}\,{\rm cm} \times \left( \frac{R}{130\,{\rm pc}} \right)^{5/2} \times \left(\frac{T}{10^4\,{\rm K}} \right)^{3/4}$. For the previously assumed value of $R = 130\,{\rm pc}$ \citep{Lazio_Cordes_1998}, they noted that this scale is unacceptably small, as it produces too much heating and it does not correspond to a reasonable astronomical scale for nonlinear density fluctuations. In addition to these objections, our measurements of refractive noise give a lower bound for the outer scale because refractive noise will be suppressed on angular scales larger than ${\sim}\,r_{\rm out}/D$. Thus, our measurements of refractive noise on baselines with $|\mathbf{u}| \sim 10^7$ at 3.6\,cm show that $r_{\rm out} \gsim (2\pi)^{-1} D/10^7 \sim 10^{14}\,{\rm cm}$. With the modified distance to the scattering (see \S\ref{sec::Scattering_Geometry} and \citet{Bower_Magnetar_2014}), the problems identified by \citet{Goldreich_2006} are mitigated, as we now discuss in detail. Specifically, an upper limit on the outer scale can be estimated as the scale on which the scattering power spectrum requires density fluctuations of order unity. Suppose that the scattering material is statistically homogeneous over a region of length $z$ along the line of sight. Electron density fluctuations $\delta n_{\rm e}(\ell)$ on a scale $\ell$ then introduce corresponding screen phase fluctuations of $\delta \phi(\ell) \sim r_{\rm e} \lambda \sqrt{\ell z} \delta n_{\rm e}(\ell)$ (because of the random walk through $z/\ell$ regions; see \S\ref{sec::Background}). Taking $\delta n_{\rm e}(\ell)/n_{\rm e} \sim (\ell/r_{\rm out})^{(\alpha-1)/2}$, we obtain, \begin{align} \delta \phi(\ell) \sim n_{\rm e} \lambda r_{\rm e} \sqrt{z \ell} \left(\frac{\ell}{r_{\rm out}} \right)^{(\alpha-1)/2},\\ \nonumber \Rightarrow r_{\rm out} \lsim \left(n_{\rm e} \lambda r_{\rm e} \sqrt{z r_{\rm diff}} \right)^{2/(\alpha-1)} r_{\rm diff}, \end{align} where $r_{\rm diff} \approx \lambda/((1+M) \theta_{\rm scatt})$ is the diffractive scale (i.e., $\delta \phi(r_{\rm diff}) \sim 1$). For \sgra, $r_{\rm diff} \approx 10^8\,{\rm cm}$ at $\lambda = 1\,{\rm cm}$. With our characteristic value $\alpha = 1.38$, we then find \begin{align} \label{eq::rout_max} r_{\rm out} \lsim (10\,{\rm pc}) \times \left( \frac{n_{\rm e}}{10\,{\rm cm}^{-3}} \right)^{5.26} \left( \frac{z}{10\,{\rm pc}} \right)^{2.63}\!. \end{align} For comparison, \citet{Armstrong_1995} estimate $r_{\rm out} \gsim 30\,{\rm pc}$ for the scattering material within 1\,kpc. For Eq.~\ref{eq::rout_max} to violate our measured lower limit for the 3.6\,cm refractive noise would require much lower electron densities $n_{\rm e} \lsim 0.1$ or larger values of $\alpha$ (approximately $\alpha \gsim 5/3$, for the characteristic values of $z$ and $n_{\rm e}$ given in Eq.~\ref{eq::rout_max}), although these results are highly sensitive to the screen thickness, $z$. From the dispersion measure of the Galactic Center magnetar \citep[e.g.,][]{Eatough_2013,Kravchenko_2016}, we can only estimate an upper bound on the plasma density, $n_{\rm e} < (180\,{\rm cm}^{-3}) / (z / 10\,{\rm pc})$. Regardless, the outer scale required by our scattering model is not implausible. \citet{Goldreich_2006} have also provided an alternative model for interstellar scattering from folded magnetic field structures. Their proposed model reproduces the $\lambda^2$ scaling and Gaussian scatter-broadening for \sgra, but it also predicts significantly suppressed refractive scintillation. Furthermore, intrinsic structure would be blurred out on small angular scales from the scattering in this model. Thus, our measurements of image substructure at 1.3\,cm conclusively reject the folded field model for the scattering of \sgra\ in its simplest form. However, we will demonstrate later that this model is compatible with our measurements if the inner scale (corresponding to the thickness of current sheets in this model) is significantly larger than expected: $r_{\rm in} \sim 2\times 10^6\,{\rm km}$. In this case, the \citet{Goldreich_2006} spectrum would instead produce significantly \emph{enhanced} refractive effects at millimeter wavelengths. \subsubsection{The Inner Scale of Turbulence} \label{sec::Inner_Scale} Our measurements constrain the inner scale of turbulence, both through plausibility arguments related to the $\lambda^2$ dependence of the angular broadening and image Gaussianity and by relating the scattering power on large scales (refractive noise) to that on small scales (the diffractive blurring). Ultimately, our most stringent lower limit on the inner scale comes from the image Gaussianity at 1.3\,cm, giving $r_{\rm in} \gsim 600\,{\rm km}$. Likewise, the 7\,mm data show a statistically significant departure from a Gaussian image, with a preference for $r_{\rm in} \lsim 1000\,{\rm km}$, although we regard this upper limit as tentative (see \S\ref{sec::gaussian_image_constraint}). Thus, we have adopted a recommended characteristic value of $r_{\rm in} = 800\,{\rm km}$. While the scattering of \sgra\ is anomalously strong, the dissipation mechanism for turbulence in the ISM may be universal. Thus, we now compare our estimate for $r_{\rm in}$ with previous theoretical and observational estimates. Using VLBI measurements of the angular broadening for several heavily scattered objects, \citet{Spangler_Gwinn_1990} estimated an inner scale of $50-200\,{\rm km}$. Based on weak scintillation measurements at centimeter wavelengths, \citet{Armstrong_1995} constrained the inner scale for the nearby ISM (within 1\,kpc) to be less than ${\sim}5\,{\times}\,10^4\,{\rm km}$. \citet{Rickett_2009} estimated $r_{\rm in} = 70\,{-}\,100\,{\rm km}$ from the pulse broadening of PSR~J1644-4559. \citet{Smirnova_2010} estimated $r_{\rm in} = 350\pm 150\, {\rm km}$ from the pulse broadening of PSR~B2111+46. Each of these studies has its own limitations. For instance, the pulsar analyses assumed isotropic scattering, and \citet{Rickett_2009} noted that a (finely-tuned) anisotropy would allow an arbitrarily large inner scale. Perhaps the most significant difficulty in our study of \sgra\ is that intrinsic source structure becomes significant for the baselines and wavelengths that are sensitive to a direct estimate of the inner scale for \sgra. Nevertheless, our lower limit on the inner scale is quite robust. \citet{Goldreich_Sridhar_1995} suggest that the inner scale in the ISM may approach the ion Larmor radius, and \citet{Spangler_Gwinn_1990} proposed that the inner scale corresponds to the larger of the ion inertial length and the ion Larmor radius in the scattering medium. The ion inertial length is $\ell_{\rm i} = V_{\rm A}/\Omega_{\rm i} \approx 230/\sqrt{n_{\rm e}/{\rm cm}^{-3}}\,{\rm km}$, where $V_{\rm A} = B/\sqrt{4\pi n_{\rm e} m_{\rm i}}$ is the Alfv\'en speed and $\Omega_{\rm i} = e B/(m_{\rm i} c)$ is the ion cyclotron frequency. The ion Larmor radius is $r_{\rm i} = v_{\rm th}/\Omega_{\rm i} \approx 930\,{\rm km} \times \left( \frac{B}{1\,\mu{\rm G}} \right)^{-1} \left( \frac{T}{10^4\,{\rm K}} \right)^{1/2}$, where $v_{\rm th} = \sqrt{k T/m_{\rm i}}$ is the ion thermal speed. Given the strong scattering of \sgra, it is likely that the ion Larmor radius will then determine the inner scale in this model. The required $B \sim 1\,\mu{\rm G}$ is somewhat lower than expected for magnetic fields in the ISM at the galactocentric distance $R \sim 5.5\,{\rm kpc}$ \citep[e.g.,][]{Han_2006}, and it may suggest that the inner scale is a few times larger than $r_{\rm i}$. In terms of a specific model for the scattering of \sgra, \citet{Sicheneder_2017} have proposed that the scattering may arise in a single \ion{H}{2} region along the line of sight, with density $n_{\rm e} \sim 200\,{\rm cm}^{-3}$ and radius ${\sim} 3\,{\rm pc}$. They note that this region can also produce the observed rotation measure of the Galactic center magnetar SGR~J1745-2900, if the field strengths in the scattering material are $15-70\,\mu{\rm G}$. In this model, the ion inertial length is only $\ell_{\rm i} \sim 10-20\,{\rm km}$ and the ion Larmor radius is $r_{\rm i} \lsim 60\,{\rm km}$. Thus, our estimates of an inner scale that is significantly higher than either of these values support the scenario in which the large RM of the magnetar arises from a local contribution near the Galactic Center \citep{Eatough_2013,Desvignes_2018}. For smaller magnetic fields, $B \sim 1\,\mu{\rm G}$, the parameters identified by \citet{Sicheneder_2017} remain plausible for the scattering. \subsubsection{The Power-Law Index of Turbulence} \label{sec::Power-Law} \begin{figure}[t] \centering \includegraphics[width=\columnwidth]{Q_Constraints-crop.pdf}\\ \caption { Summary of constraints on the power spectrum of phase fluctuations, $Q(\mathbf{q})$. Refractive noise on long baselines at 3.6 and 1.3\,cm constrains the power in wavenumbers $q^{-1}\,{\sim}\,10^{13} - 10^{14}\,{\rm cm}$ (red diamonds). Asymptotic Gaussian angular broadening constrains the power in wavenumbers $q\,{\sim}\,r_{\rm in}^{-1}$; the corresponding constraint on $Q$ depends strongly on $r_{\rm in}$ and weakly on $\alpha$. The orange band shows the angular broadening constraint as a function of $q=r_{\rm in}^{-1}$ over the range $1 < \alpha < 5/3$. Three models are plotted: a Kolmogorov spectrum with our minimum allowed inner scale (purple; $\alpha=5/3$, $r_{\rm in}=600\,{\rm km}$), our recommended characteristic model (blue; $\alpha=1.38$, $r_{\rm in}=800\,{\rm km}$), and a \citet{Goldreich_2006} spectrum (dashed gray; $\alpha=0$, $r_{\rm in}=2 \times 10^6\,{\rm km}$). Corresponding colored circles show the constraint on $Q$ from angular broadening for each model. While a Kolmogorov spectrum is compatible with the refractive noise measurements taken alone, it would require spectral flattening or additional power near the inner scale to be compatible with the measured angular broadening, as shown by the horizontal purple dashed line. The green shaded region shows the range of modes that contribute refractive noise to EHT images of \sgra\ ($\sigma_{\rm ref} \propto \sqrt{Q}$). Note that refractive noise predictions from our model for the EHT are rather insensitive to possible generalizations that would allow $\alpha=5/3$. However, the \citet{Goldreich_2006} spectrum would increase refractive noise by a factor of ${\approx}10$ relative to our characteristic model. } \label{fig::Q_Constraints} \end{figure} Figure~\ref{fig::Q_Constraints} shows our constraints on the power spectrum of phase fluctuations, $Q(\mathbf{q})$, along the direction of the scattering major axis. Refractive noise on a long baseline $\mathbf{u}$ is dominated by refractive modes with $\mathbf{q} \sim 2\pi \mathbf{u}/D$. Thus, our measurements of refractive noise at 3.6 and 1.3\,cm constrain the power in wavenumbers $q^{-1}\,{\sim}\,10^{13} - 10^{14}\,{\rm cm}$. In addition, our measurements of the asymptotic Gaussian angular broadening constrain the power in wavenumbers $q^{-1}\,{\sim}\,r_{\rm in}$, with the exact constraint also weakly dependent on $\alpha$: $Q(r_{\rm in}^{-1}) \propto r_{\rm in}^4/\Gamma(1-\alpha/2)$. As is evident from Figure~\ref{fig::Q_Constraints}, larger values of the inner scale require a flatter power spectrum for the measured refractive noise to be compatible with the measured angular broadening (see also Figure~\ref{fig::alpha_rin_limits}). Allowing arbitrarily small inner scales, we find $\alpha \lsim 1.6$, while including our derived constraints on the inner scale, we obtain $\alpha < 1.47$. Thus, a Kolmogorov spectrum ($\alpha=5/3$) is incompatible with our measurements, as is an $\alpha=3/2$ spectrum \citep[see, e.g.,][]{Iroshnikov_1964,Kraichnan_1965,Sridhar_1994,Goldreich_Sridhar_1995}. Our results are at tension with measurements for the local ISM \citep[e.g.,][]{Armstrong_1995}, the wavelength dependence of pulsar temporal broadening \citep[e.g.,][]{Lohmer_2001,Bhat_2004,Lewandowski_2013}, and VLBI of heavily scattered sources \citep{Spangler_Gwinn_1990}, all of which tend to infer somewhat larger values of $\alpha$. We now outline possible generalizations to our scattering model that might render higher values of $\alpha$, including a Kolmogorov spectrum, feasible. The first possibility is an outer scale of turbulence that is similar to the scales probed at 3.6\,cm, thereby reducing the 3.6\,cm refractive noise but perhaps not the 1.3\,cm refractive noise (which probes smaller scales). The 3.6\,cm refractive noise corresponds to scattering modes with a transverse scale of ${\sim}4\,{\rm AU}$ on the scattering screen, so the required outer scale is $r_{\rm out} \sim 1\,{\rm AU}$. This value is somewhat smaller than the lower limit estimated by \citet{Armstrong_1995}. Moreover, this is a finely-tuned constraint, requiring the outer scale to be precisely matched to our observing parameters --- a smaller outer scale would be inconsistent with the observed 1.3\,cm refractive noise, and a larger outer scale would not affect the 3.6\,cm noise. A second possibility is that there is extra power located near the dissipation scale. Spectral flattening near the inner scale has been seen in the solar wind \citep[e.g.,][]{Neugebauer_1975,Celnikier_1983,Coles_1991} and possibly also in the ISM \citep{Smirnova_2010}. A pile-up in power by a factor of ${\sim}15$ would reconcile our measurements with a Kolmogorov spectrum (see Figure~\ref{fig::Q_Constraints}); a factor of ${\approx}\,2$ would be needed for an $\alpha = 3/2$ spectrum. A more radical possibility, which is not excluded by our data, is that the spectrum is extremely shallow and the inner scale is correspondingly large. In particular, the model of \citet{Goldreich_2006} produces a power spectrum with $\alpha=0$, which would be consistent with all our measurements if $r_{\rm in} \sim 2\times 10^6\,{\rm km}$. This inner scale is a factor of 20 larger than the characteristic value used by \citet{Goldreich_2006} and requires an outer scale that is 400 times larger than expected, or a few kpc. Nevertheless, our measurements are insufficient to rule out this type of power spectrum, and it would produce refractive signatures for the EHT that are approximately 10 times stronger than those predicted by our recommended characteristic model. Thus, we expect continued studies at 1.3\,mm (and possibly 3.5\,mm) will be able to conclusively confirm or reject this scattering model for \sgra. \subsection{Sensitivity to the Assumed Scattering Model} \label{sec::Sensitivity_to_Model_Assumptions} We have analyzed all our data in the context of a single scattering model. The anisotropy in this model is determined by the magnetic field wander along the line of sight relative to its preferred orientation (which determines the minor axis of the scattering ellipse). \citet{Psaltis_2018} give three representative models for the field wander: ``von Mises,'' ``Dipole,'' and ``Boxcar.'' The von~Mises model represents the angular field wander using a generalized Gaussian distribution for circular quantities; the Dipole model uses a change of variables to rescale the principal axes of the power spectrum; and the Boxcar model has a power spectrum that is isotropic across a restricted range of angles and is zero elsewhere. Because of the efficient computational tools developed in Appendix~\ref{sec::Efficient_Computation}, all of our results have used the Dipole model. We now evaluate how sensitive our conclusions are to this choice. \citet{Psaltis_2018} show that the shape of the scattering kernel is almost independent of the choice of scattering model. However, the refractive noise along the minor axis is sensitive to the scattering model. Because our measurements of refractive noise are predominantly along the major axis, our results are not strongly affected by the choice of scattering model. For example, the mean refractive noise on our long baselines at 3.6\,cm changes by less than $\pm 2\%$ among the three scattering models (well within our uncertainty from sampling only a few elements of the refractive noise). Likewise, the 95\% confidence intervals are almost identical for the three models. For the average refractive noise on long baselines at 1.3\,cm, the von Mises and Dipole models agree to within 1\%, but the Boxcar model differs by 5\%. Again, these differences are negligible within our error budget. Thus, we conclude that our specific choice of scattering model is irrelevant for our results. Equivalently, our current measurements provide no firm guidance for discriminating among these models for the magnetic field wander. Future measurements of refractive noise on long baselines along the minor axis could immediately rule out the Boxcar model and would be sensitive to differences between the von Mises and Dipole models \citep[see, e.g., Figures~9 and 13 in][]{Psaltis_2018}. \subsection{Implications for Continued Studies of Sgr A*} For the scattering parameters that we have identified, a Gaussian scattering kernel is likely a good approximation for \sgra\ for centimeter wavelengths, although the full non-Gaussian kernel shape should be used for continued studies at millimeter wavelengths. The full kernel shape is especially important for scattering mitigation in imaging with EHT data \citep{Fish_2014,Stochastic_Optics}. We have shown that refractive noise is a critical component of the error budget when model fitting to observations of \sgra. In addition, we caution that inferences of intrinsic size should account for refractive image distortion. At long wavelengths, stochastic changes from refractive distortion can exceed the contribution of intrinsic structure (see Figure~\ref{fig::Fractional_Size_Effects}). Our results suggest that intrinsic structure can only be securely decoupled from refractive distortion for $\lambda \lsim 3.6\,{\rm cm}$ (minor axis) or $\lambda \lsim 1.3\,{\rm cm}$ (major axis). Note that these are fundamental limitations; they would apply even if the sensitivity and baseline coverage of the observations were perfect. \begin{figure}[t] \centering \includegraphics[width=\columnwidth]{Fractional_Size_Effects-crop.pdf}\\ \caption { Comparison of the wavelength-dependent fractional effects from intrinsic image structure and from refractive image distortion. The intrinsic curves show the fractional increase in the scattered image size because of intrinsic structure: $1 - \theta_{\rm scatt}/\theta_{\rm ea}$, where $\theta_{\rm scatt}$ is the size of a scattered point source and $\theta_{\rm ea}$ is the angular size of an extended source using our estimated wavelength-dependent size of \sgra. The refractive distortion curves show the expected fractional fluctuations of image size among different observing epochs because of refractive scattering (see \S\ref{sec::Image_Size_Fluctuations}). Intrinsic structure can only be reliably estimated when refractive jitter is significantly smaller than the intrinsic contribution, irrespective of the observing sensitivity or baseline coverage. Requiring that the fractional increase from intrinsic size must be at least three times the rms distortion, we estimate that intrinsic properties for the major axis can only be reliably constrained for observations with $\lambda \lsim 1.3\,{\rm cm}$, while intrinsic properties for the minor axis can only be constrained for observations with $\lambda \lsim 3.6\,{\rm cm}$. } \label{fig::Fractional_Size_Effects} \end{figure} Our work has two significant implications for imaging \sgra\ with the EHT. First, we have shown that the scattering kernel may be much smaller than has been estimated, so the blurring effects of scattering may be less severe than have been assumed (see Figure~\ref{fig::Scattering_Kernel}). Second, we find $\alpha \lsim 1.47$, which produces significantly less refractive noise than the standard Kolmogorov picture; it predicts that the renormalized refractive noise is at most ${\sim}1\%$ of the zero-baseline flux density (see Figure~\ref{fig::EHT_Refractive_Noise}). This estimate reinforces the conclusions of \citet{Fish_2016} and \citet{Lu_2018} that refractive noise is unlikely to be a significant component of the error budget for past EHT observations. Both these implications improve the prospects for horizon-scale imaging at 1.3\,mm. However, continued observations must also account for the possibility of strong refractive effects from shallow spectra, such as from the model with $\alpha=0$ and $r_{\rm in} = 2\times 10^6\,{\rm km}$. While this model remains speculative and lacks support from other lines of sight, it would produce striking differences from our characteristic model for EHT observations (an increase in refractive noise by a factor of 10) and should be tested further with long baseline measurements at 1.3\,mm and 3.5\,mm. \begin{figure}[t] \centering \includegraphics[width=\columnwidth]{EHT_Refractive_Noise-crop.pdf}\\ \caption { Expected rms of refractive noise (dashed) and renormalized refractive noise (solid) for \sgra\ at 1.3\,mm (i.e., for EHT observations). The curves correspond to our recommended characteristic scattering model and an isotropic intrinsic Gaussian source with FWHM $\theta_{\rm src} = 52\,\mu{\rm as}$. To express the refractive noise in units of flux density, we assume a total flux density of 3.5\,Jy for \sgra. } \label{fig::EHT_Refractive_Noise} \end{figure} \ \\ We have analyzed observations of \sgra\ at wavelengths from 1.3\,mm to 30\,cm using a physically motivated model for its scattering \citep[developed in][]{Psaltis_2018}. At long wavelengths, the angular broadening from scattering is an anisotropic Gaussian, and its size scales as $\theta_{\rm scatt} \propto \lambda^2$. At shorter wavelengths, the shape and wavelength dependence of the scattering depend on the inner scale of turbulence, $r_{\rm in}$, and on the power-law index of the scattering, $\alpha$. Using a new prescription to perform model fitting with refractive noise in the error budget, we are able to estimate the asymptotic Gaussian scattering parameters to excellent accuracy (see Table~\ref{tab::GaussianParameters}). In addition, we show that $\alpha \lsim 1.47$ and $r_{\rm in} \gsim 600\,{\rm km}$ (our recommended characteristic values for these parameters are $\alpha = 1.38$ and $r_{\rm in} = 800\,{\rm km}$). Our recommended scattering parameters are summarized in Table~\ref{tab::ScatteringParameters}. After deconvolving the effects of scattering from our estimated sizes of the scatter broadened images, we find that the intrinsic image of \sgra\ is nearly isotropic, with FWHM $\theta_{\rm src} \sim (40\,\mu{\rm as}) \times \lambda_{\rm mm}$ from 1.3\,mm to 1.3\,cm. While this linear wavelength dependence for the emission size is natural for both disk- and jet-dominated models, the nearly isotropic image shape strongly favors disk models. At 1.3\,mm, where the emission region is expected to be largely optically thin, our estimated image size is consistent with predictions from recent GRRMHD simulations \citep[e.g.,][]{Chael_2018}. For ISM scattering, our most surprising conclusion is that a Kolmogorov spectrum ($\alpha=5/3$) in the inertial range is incompatible with our observations. However, our constraints on $\alpha$ are somewhat indirect, as they relate the refractive power in large scattering modes to the diffractive power from small scattering modes (see Figure~\ref{fig::Q_Constraints}). For a generalized scattering model, larger values of $\alpha$ are possible but would require continuous injection of energy on ${\lsim}{\rm AU}$-scales, a pile-up of energy near the inner scale (${\sim}10^3\,{\rm km}$), or a small outer scale for the turbulence (${\sim}1\,{\rm AU}$). We have also shown that the inner scale cannot be smaller than $600\,{\rm km}$ and is likely $r_{\rm in} \approx 800\,{\rm km}$. This scale is comparable to the ion Larmor radius for regions of the ISM with weak magnetic fields $B \sim 1\,\mu{\rm G}$, thereby supporting identification of the ion Larmor radius (or a few times this radius) with the dissipation scale of ISM turbulence. This estimate also suggests that the rotation measure associated with the scattering material is modest, and hence that the rotation measure of the Galactic Center magnetar is dominated by local contributions \citep{Eatough_2013} rather than from the scattering material \citep{Sicheneder_2017}. Our estimated $r_{\rm in}$ is also comparable to the ion inertial length for $n_{\rm e} \sim 0.1\,{\rm cm}^{-3}$ and requires $n_{\rm e} \gsim 0.1\,{\rm cm}^{-3}$ if the inner scale is determined by the larger of these characteristic two plasma length scales \citep{Spangler_Gwinn_1990}. However, we cannot conclusively rule out much shallower spectra with correspondingly larger inner scales. While our primary objective has been to constrain the parameters of our specific scattering model, our observations also constrain alternative theories for the scattering of \sgra. For example, \citet{Goldreich_2006} have proposed a model in which the scattering is caused by an ensemble of folded current sheets in the ISM. While this model naturally reproduces the $\lambda^2$ scaling of angular broadening and the Gaussian image at long radio wavelengths, it predicts an absence of refractive scattering effects. This model would not produce scattering substructure in images, and any intrinsic substructure would be blurred out by small-scale scattering modes. Hence, the pronounced long-baseline refractive noise at 1.3\,cm enables us to firmly reject this alternative scattering model for \sgra\ in its simplest form \citep[see also][]{Gwinn_2014}. However, the model is compatible with our measurements if the thickness of the current sheets is significantly larger than expected, corresponding to $r_{\rm in} \sim 2\times 10^6\,{\rm km}$, in which case it would instead produce strongly \emph{enhanced} refractive effects at millimeter wavelengths. Thus, we expect that continued observations with the GMVA and EHT will be sufficient to firmly support or reject this model. Our results highlight the importance of including refractive noise when fitting models to radio observations of \sgra. Refractive uncertainties can plausibly explain many of the discrepancies in past measurements of the size of \sgra, such as those identified by \citet{Psaltis_2015}. In addition, we have shown that refractive effects likely prohibit a meaningful study of intrinsic structure at wavelengths longer than $1.3\,{\rm cm}$ (or $3.6\,{\rm cm}$ for the minor axis; see Figure~\ref{fig::Fractional_Size_Effects}). Nevertheless, our results also show that both the blurring and substructure from scattering may be significantly smaller at 1.3\,mm than expected. Thus, the prospects for deeper study of \sgra\ at millimeter wavelengths, including imaging with the EHT, are excellent. | 18 | 8 | 1808.08966 |
1808 | 1808.04379_arXiv.txt | We present COSMOS-Drift And SHift (DASH), a Hubble Space Telescope WFC3 imaging survey of the COSMOS field in the $H_{160}$ filter. The survey comprises 456 individual WFC3 pointings corresponding to an area of 0.49 deg$^2$ (0.66 deg$^2$ when including archival data) and reaches a $5\sigma$ point-source limit of $H_{160}=25.1$ (0\farcs3 aperture). COSMOS-DASH is the widest HST/WFC3 imaging survey in $H_{160}$ filter, tripling the extragalactic survey area in the near-infrared at HST resolution. We make the reduced $H_{160}$ mosaic available to the community. We use this dataset to measure the sizes of \totdash\ galaxies with $\log(M_{\star}/M_{\odot})>11.3$ at $1.5<z<3.0$, and augment this sample with \totacs\ galaxies at $0.1<z<1.5$ using archival ACS imaging. We find that the median size of galaxies in this mass range changes with redshift as $\langle{}r_{\rm eff}\rangle=(10.4\pm 0.4)\times(1+z)^{(-0.65\pm 0.05)}$\,kpc. Separating the galaxies into star forming and quiescent galaxies using their restframe $U-V$ and $V-J$ colors, we find no statistical difference between the median sizes of the most massive star-forming and quiescent galaxies at $\langle z\rangle =2.5$: they are $4.9 \pm 0.9$ kpc and $4.3 \pm 0.3$ kpc respectively. However, we do find a significant difference in the S\`ersic index between the two samples, such that massive quiescent galaxies have higher central densities than star forming galaxies. We extend the size-mass analysis to lower masses by combining it with the 3D-HST/CANDELS sample of \citet{VanderWel2014a}, and derive empirical relations between size, mass, and redshift. Fitting a relation of the form $r_{\rm eff}=A \times m_{\star}^{\alpha}$, with $m_{\star}=M_{\star}/5\times 10^{10}\,M_{\odot}$ and $r_{\rm eff}$ in kpc, we find $\log A=-0.25 \log$ $(1+z)+0.79$ and $\alpha=-0.13\log (1+z)+0.27$. We also provide relations for the subsamples of star forming and quiescent galaxies. Our results confirm previous studies that were based on smaller samples or ground-based imaging. | The sizes of galaxies reflect their assembly histories and their connection to their dark matter halos \citep{Mo1997TheDisks,Kravtsov2012TheGalaxies,Jiang2018IsSize}. Different modes of assembly of stars in galaxies lead to a different growth of their radii: passive evolution will cause no significant growth in size or mass, but only a maturation of the existing stellar population; dry major mergers lead to a proportional growth in size and mass as the two bodies come to dynamic equilibrium; and dry minor mergers increase the size of galaxies more rapidly by building an outer envelope \citep{Bezanson2009,Naab2009MinorGalaxies}. When gas physics are considered the evolution can be more complex; e.g., ``wet" gas-rich mergers may trigger compact starbursts leading to larger post-merger disks \citep{Hernquist1989TidalGalaxies,Robertson2005AFormation}, while gas flows to the central regions may both form compact bulges and feed a central black hole \citep{Efstathiou1982TheGalaxies,Dekel2013WetNuggets,Barro2017SpatiallyObservations}. The size of a galaxy may also hold information on the properties of the dark matter halo; galaxy size may be proportional to the halo virial radius as a result of conservation of angular momentum during the collapse and cooling of a galaxy \citep{Mo1997TheDisks,Dutton2006AExpansion,Shankar2011SizeUniverse,Kravtsov2012TheGalaxies,Porter2014UnderstandingGalaxies,Somerville2017a}, although it is unclear whether this expected correlation is actually preserved in the galaxy formation process \citep{DeFelippis2017TheSimulation, Jiang2018IsSize}. Observationally, the sizes of galaxies have been found to vary significantly with galaxy mass, color (or star formation activity) and redshift. Generally, the sizes are larger for galaxies that are more massive, galaxies that are forming stars, and galaxies at lower redshift \citep{Kormendy1996,Shen2003TheSurvey,Ferguson2003TheGalaxies,Trujillo2005TheFIRES,Elmegreen2007ResolvedRedshift,Williams2009TheZ=2,Mosleh2017Connection=2,Ono2012Evolution7,VanderWel2014a,Bernardi2012SystematicMorphology,Carollo2013,Lange2014GalaxyZ}. At intermediate masses, the slope of the size-mass relation is found to be shallow for star forming galaxies ($r_{\rm eff} \propto M^{0.2}_{\star}$) and steeper for quenched galaxies ($r_{\rm eff} \propto M^{0.8}_{\star}$), where $r_{\rm eff}$ is the half-light radius. Both galaxy types exhibit a large intrinsic scatter in the size-mass relation at all redshifts \citep{VanderWel2014a}. The slope of the size-mass relation of star forming galaxies is similar to the growth track of individual galaxies, both in observations and simulations \citep{Lilly1997HubbleZ=1,Ravindranath2004EvolutionDistribution,Trujillo2005TheFIRES,VanDokkum2015_compactmassive}. Following quenching galaxies follow a steeper growth track in the size-mass plane, probably because dry minor mergers rapidly increase the size \citep[see][]{Hilz2012HowFraction,Carollo2013,VanDokkum2015_compactmassive}. Physically, the central stellar density has been proposed as a key parameter connecting galaxy morphology and star formation histories \citep{Bezanson2009,Carollo2013,Fang2013AGalaxies,VanDokkum2014_densecore,Whitaker2016}. Galaxies with high central densities are found to be redder with lower specific star formation rate than bluer galaxies at a given redshift. A possible explanation is that feedback mechanisms that shut off star formation are more effective when the central density becomes high \citep[e.g.,][]{Croton2005TheGalaxies,Conroy2014PreventingHeating}. Whether the same processes operate in the most massive galaxies, here defined as galaxies with $M_{\rm\star}>2\times 10^{11}$\,M$_{\odot}$, is still an outstanding question. \citet{Carollo2013} present a comprehensive analysis of the sizes of galaxies at 0.2$<$z$<$1 in the COSMOS field, measured from the ACS F814W imaging. Out to $z\sim 1$ nearly all such galaxies are found to be quiescent \citep{Hahn2014PRIMUS:0.8}. At higher redshift this mass range is not commonly studied, as their number is low in the fields that have been observed so far with \textit{Hubble Space Telescope} (HST) in the near-IR. \citet{VanderWel2014a} studied the mass-size relation in the 3D-HST/CANDELS fields, and finds that the most massive star forming and quiescent galaxies at $z\sim 2.5$ have similar sizes. That is, the ``rule'' that star forming galaxies are larger than quiescent galaxies appears to not apply at the highest masses and redshifts. However, given the small number of galaxies with masses $>2\times 10^{11}$\,M$_{\odot}$ found within extragalctic pencil-beam studies, this is largely driven by the extrapolation of trends seen at lower masses; essentially, the fitted relations to lower mass galaxies intersect at $M_* \sim 5\times 10^{11}$\,M$_{\odot}$. Recently \citet{Faisst2017} studied the sizes of galaxies in this mass range out to $z\sim 2$ using ground-based imaging, calibrated with HST data in smaller fields. They find similar results as \citet{VanderWel2014a}. Similarly, \citet{Hill2017} study the size evolution of galaxies since $z\sim 5$ using number density-matched samples, again consistent with previous size measurements in smaller fields. Here we build on these previous studies by studying the most massive galaxies out to $z\sim 3$ with a new wide-field HST survey, COSMOS-Drift And SHift (DASH). COSMOS-DASH provides the large area and high resolution needed for structural study of massive galaxies at $1.5 \leq z \leq 3.0$. It is a wide and medium depth survey using the near-infrared channel of Wide Field Camera 3 (WFC3) on HST, utilizing a novel drift-and-shift technique. COSMOS-DASH covers 0.49 deg$^2$ of the UltraVISTA \citep{McCracken2012} deep stripes in the COSMOS field down to $H_{160}=25.1$, or 0.66 deg$^2$ when archival data are included \ref{tab:archival}, tripling the extragalactic survey area observed by HST in the near-IR \citep{Momcheva2016a}. The paper is structured as follows. In Section \ref{sec:dash} we give a brief description of the COSMOS-DASH survey. In Section \ref{sec:sample}, the selection of the massive galaxy sample and the separation of quenched galaxies from the star forming galaxies are described. Section \ref{sec:size} goes into the details of the size measurement of galaxies using COSMOS-DASH images. Analysis of the evolution of size-mass relation is described in Section \ref{sec:evolution}, while in Section \ref{sec:discussion} we interpret the results in the context of the termination of star formation in the most massive galaxies. In this paper, we assume a $\Lambda$CDM cosmology with $\Omega_{\rm m}=$ 0.3, $\Omega_{\Lambda}=$ 0.7, and $H_0=$ 70 km s$^{-1}$ Mpc$^{-1}$. | \label{sec:discussion} \subsection{Comparison to Previous Studies} We present the first comprehensive measurements of the sizes of the most massive galaxies with $M_{\star}>2\times 10^{11}$ M$_{\odot}$ within $0.1<z<3.0$ measured at HST resolution. As shown in Figure \ref{fig:size_med}, we confirm that the galaxies in this study are larger in size than the less massive galaxies \citep{Shen2003TheSurvey, Carollo2013,VanderWel2014a} at the same epoch. The bottom-left panel of Figure \ref{fig:size_med} shows the ratio of median sizes of star-forming galaxies to that of quiescent galaxies. For intermediate to massive galaxies, with stellar masses between $ 10^{9}$ and $10^{11} M_{\odot}$, quiescent galaxies are, on average, smaller than star-forming galaxies \citep{Corollo2013,VanderWel2014a}. At z$\sim$2.25, for stellar mass of $M_{\star}\sim 5 \times 10^{10}$ M$_{\odot}$ the median size of star-forming galaxy is 3.39$\pm$0.08 kpc and the median size of quiescent galaxy is 1.20$\pm$0.03 kpc, which is almost a third of the size of the star-forming galaxies. However, for galaxies with stellar masses $\geq 2 \times 10^{11} M_{\odot}$ the two classes of galaxies are found to have similar sizes of 5.1$^{+0.6}_{-0.1}$ kpc and 3.6$^{+1.9}_{-0.4}$ kpc. Our results are consistent with the measurements of \citealt{VanderWel2014a} in the same mass range, although their sample is smaller by a factor of 3--4. More precisely, they confirm the extrapolated size-mass relations of \citet{VanderWel2014a}, which were dominated by less massive galaxies. Our results are also in agreement with the ground-based measurements of \citet{Faisst2017} and \citet{Hill2017}. \citet{Faisst2017} studied galaxies with $\log(M_{\star}/M_{\odot})>11.4$, slightly more massive than galaxies in our sample. Within the uncertainties, the sizes of these galaxies are fully consistent with our median sizes. \subsection{Central Density of the Most Massive Galaxies} \label{sec:transition} \begin{figure}[ht] \centering \includegraphics[width=0.46\textwidth]{Sersic_evolution.png} \caption{Evolution of median S\`ersic indices with redshift of all, star-forming and quiescent galaxies.} \label{fig:Sersic} \end{figure} \begin{figure}[ht] \centering \includegraphics[width=0.5\textwidth]{vdisp_frac.png} \caption{Top panel: Evolution of the fraction of the most massive star forming and quiescent galaxies. The fraction of star-forming galaxies increases with redshift while that of quiescent galaxies decreases with the cross-over happening at z$\sim$2. Bottom panel: Evolution of the central mass density of star-forming and quiescent galaxies. The central mass density is the density of the central 1 kpc of the galaxy, calculated using Eq.\ \ref{eq:vsig}. The blue circles and red dots show the median central densities of the star-forming and quiescent populations, and the hatched areas show 1-$\sigma$ spread. The black line is the median velocity dispersion of all quenched galaxies with $9.0<\log(M_{\star}/M_{\odot})<12.0$ measured by \citet{Whitaker2015}. The dotted black line is the assumed constant threshold in velocity dispersion above which galaxies at $1.5<z<3.0$ quench from \citet{VanDokkum2015_compactmassive}. The dashed black line is the predicted quenching threshold from \citet{Voit2015Precipitation-RegulatedGalaxies} normalized to 300 km/s at z$=$2. } \label{fig:core} \end{figure} Studies of the central galaxy density out to z$=$3 \citep{Cheung2012The0.5,Saracco2012OnGalaxies,Barro2012CANDELS:Z2,Barro2015} find that the innermost structure of galaxies correlates with the star formation rate. Quenching of galaxies (whether by AGN feedback or other mechanisms) is thought to become very efficient when the central density reaches a certain threshold \citep{VanDokkum2015_compactmassive,Whitaker2016}. As stellar density and velocity dispersion are closely related, observations therefore indicate that galaxies are statistically more likely to be quiescent once they have surpassed a threshold in either density or velocity dispersion. This has been studied in detail by \citet{Whitaker2016} who found an abrupt cessation of star formation when galaxies reach a threshold central stellar density. We can use the information of the two dimensional light profiles of the galaxies and their total stellar mass to infer the central stellar density of the galaxy, assuming spherical symmetry. We follow the prescription from \citet{Bezanson2009} and \citet{Whitaker2016} to calculate the central stellar density and central velocity dispersions of the galaxies. We perform an Abel transform to deproject the circularized, three-dimensional light profile using: \begin{equation} \rho(x) = \frac{b_n}{\pi} \frac{I_0}{r_{\rm eff}} x^{1/n-1} \int^{\infty}_1 \frac{{\rm exp}(-b_nx^{1/n}t)}{\sqrt{t^{2n}-1}} dt, \end{equation} with $\rho$ the three-dimensional (3D) luminosity density in a particular filter, $x\equiv r/r_{\rm eff}$, $r_{\rm eff}$ the circularized effective radius, the $n$ the S\`ersic index, and $b_n$ the $n$-dependent normalization parameter of the S\`ersic profile. We note that this methodology may lead to errors for galaxies that are far from spherical symmetry, in particular for flat disks. The mass within $r=1$\,kpc is calculated by integrating the 3D luminosity profiles. Assuming mass follows light and neglecting color gradients, we convert the central luminosity to central stellar mass using the corrected total stellar masses from the UVISTA catalog. The central density is calculated by numerically integrating the following equation: \begin{equation} \rho_{\rm 1kpc} = \frac{\int^{\rm 1kpc}_0 \rho(r)r^2 dr }{\int^{\infty}_0 \rho(r)r^2 dr}\frac{M_{\rm tot}}{\frac{4}{3}\pi \rm(1 kpc)^3}. \label{eq:vsig} \end{equation} Figure \ref{fig:core} shows the central stellar densities of star-forming and quiescent galaxies in our sample. The central velocity dispersion is calculated assuming $v_{\rm disp, 1kpc}=v_{\rm circ, 1kpc}/\sqrt{2}$, where $v_{\rm circ, 1kpc}=\sqrt{4/3 G \rho_{\rm 1kpc}}$ and $G$ is the gravitational constant. As can be seen, the central stellar density of both populations of galaxies are lower at $z=0$ than at $z\sim3$, but the central density of quiescent galaxies is always higher than that of star-forming galaxies: the equivalent velocity dispersion of star-forming galaxies decreased from 181 km/s to 67 km/s whereas that of quiescent galaxies decreased from 285 km/s to 180 km/s between z$\sim$2.63 and z$\sim$0.13. This difference in central density between star forming and quiescent galaxies may seem surprising, given the fact that they have very similar sizes at fixed mass. However, the Sersic indices also enter the calculation of the central densities, and they are significantly different between the two samples: we find $\langle n\rangle=4.0 \pm 0.$ for quiescent galaxies at $z\sim 2.75$ and $\langle n\rangle=2.2 \pm 0.2$ for star forming galaxies, as shown in Fig. \ref{fig:Sersic}. The difference in central density between the two samples is driven by this difference in the profile shape. We show quenching threshold velocity dispersions from previous studies in Fig.\ \ref{fig:core}. As can be seen the median central density of quenched galaxies are always above the threshold density while that of star-forming galaxies is above the threshold density only at the highest redshifts. This analysis thus provides evidence that it is indeed the central density, rather than the average surface density within the effective radius \citep[e.g.,][]{Franx2008,Maier2009The0.5}, that determines whether a galaxy is star forming or quenched. We note here that quenching may not be a single event in the life of a galaxy; specifically, many of the massive star forming galaxies at low redshift may be rejuvenated by fresh gas infall. These transitions change the fraction of quiescent and star forming galaxies, with young quiescent galaxies being added to the sample (Fig.\ \ref{fig:core} upper panel, also see, e.g., \citet{VanDokkum2001_morphology, Carollo2013}). | 18 | 8 | 1808.04379 |
1808 | 1808.03286_arXiv.txt | HST/ACS observations along the major axis of M33 show that the mean age of its stars decreases with increasing distance from the galaxy center. Such a behavior is consistent with an inside-out growth of the disc. However, in the outermost observed field, at $r\simeq$11.6 kpc, a reversal of this gradient is detected, with old stars found in high percentages beyond this radius. In this work we investigate the origin of such a reversal in stellar age gradient, by using a simulated M33 analogue from the Constrained Local UniversE Simulations (CLUES). The simulated M33 is similar to the observed one in terms of mass, rotation velocity, surface brightness and, similar to what has been reported in observations, shows a stellar age turnaround at large radii. We demonstrate that this reversal is mostly a result of stellar accretion from old satellite galaxies and, to a lesser extent, of stellar migration of in-situ stars. The old accreted stars, with formation times t$_{f}<4$ Gyrs, are kinematically hot and can be differentiated from the in-situ stars by their high velocity dispersion and the fact that they do not have rotationally-supported orbits. In the future, obtaining kinematic information of the stars in the outskirt of M33 will help to verify this scenario. | \label{sec:intro} In a $\Lambda$CDM universe, spiral galaxies consist of a disc component made of stars, cold gas and dust, a central bulge and a stellar halo, all embedded in a dark matter halo \citep{white78}. The disc component can be separated into two different parts: the thin disc, and the thick disc \citep{Burstein79, Gilmore83}. These two components are defined by examining the vertical scale height of stars when separated by age (e.g. \citealt{Haywood13,Bensby14}) or metallicity (e.g. \citealt{Fuhrmann08,Bensby14}). The stars in the thin disc component are formed by gas accretion at the later stages of galaxy formation and they have a wide range of ages \citep{Yoachim06}. The stars in the thick disc, however, are older and their origin is still debated (e.g. \citealt{Brook04, Villalobos08, Minchev15}). The distribution of stars in galactic discs is also an ongoing research area. One of the favorite modes for the mass assembly of a galaxy is the ``inside-out'' scenario \citep{Chiappini97, Mo98, Brook12, Pilkington12, Bird13}. In the inside-out growth proposal, the inner disc is thought to assemble first as a consequence of the high density of accreted gas residing in the center of the galaxy's potential well. Thus, the fraction of young stars is expected to increase with galactocentric radius. Several galaxies have been found to be compatible with such a growth model \citep{Perez13, SanchezBlazquez14, Tacchella15}. Recent observations regarding the ages of stars in the neighboring galaxy M33 indicate that this galaxy is compatible with an inside-out disc growth scenario, in which old stars are detected in the inner region of the galaxy, while young, disc stars tend to naturally be found in the outskirts of the disc \citep{Williams09, Barker11}. Specifically, these observations made use of the \textit{Hubble Space Telescope Advance Camera for Surveys} (HST/ACS), to derive the cumulative star formation history (SFH) along M33's major axis and for different radii. The SFH was derived using the synthetic color-magnitude diagram (CMD) fitting method. CMDs were obtained by measuring resolved stellar photometry using the ACS module of the \code{DOLPHOT} software package \citep{Dolphin00}. Assuming an initial mass function and stellar evolution isochrones, a fitting is performed on the CMD to obtain the star formation rate at their respective ages and metallicities. \citet{Williams09} and \citet{Barker11} showed that within $\approx 9$ kpc from M33's center, the mean age of stars decreases as one moves further out from the galactic center. They also showed that at radii greater than $\approx 9$ kpc, however, this \textit{age gradient} reverses, such that the mean age of stars increases as one approaches the outer region of M33. The age gradient thus reverses from decreasing mean stellar age with radius (within 9kpc) to increasing mean stellar age with radius (beyond 9kpc). Note that the age gradient reversal is accompanied by a surface brightness and stellar mass surface density break beyond 8 kpc \citep{Ferguson07,Barker11}, whose physics remains contentious (see \citealt{trl17} for a recent review of the subject using simulations). Similar age profiles have been seen in both simulations (e.g. \citealt{Roskar08a,Roskar08b, SanchezBlazquez09,MartinezSerrano09,RuizLara16a}) and observations of disc galaxies \citep[e.g.,][]{Bakos08, Yoachim12, zheng15,RuizLara16b}, yet the origin of the reversal is not clear. Several explanations for the reversal have been proposed: stellar migration, in which the inner disc forms inside-out and the region beyond the upturn radius is populated with stars that migrated from the inner disc \citep{Roskar08a,Roskar08b,RuizLara16a}; projection effects, that cause a contamination and overlap of stars from different galactic components \citep{Barker11}; a decrease in the gas volume density in the disc, which induces a break in the star formation density which itself coincides with the radius where the gas disc begins to warp \citep{SanchezBlazquez09}; or old stars coming from mergers that, due to their significant energy, remain in orbits at large, outer radii \citep{Gill05,Sales2007,Brook12,RuizLara16a}. In this paper we explore the age gradient of a simulated M33 analogue galaxy, formed in a constrained Local Group environment as part of the CLUES\footnote{\url{www.clues-project.org}} project \citep{Gottloeber10,Carlesi16}. The initial conditions have been constrained by observational data such that the z=0 cosmography is forced to reproduce the real local environment \citep{Libeskind15,Sorce16}. The simulated M33 analogue shares many properties with the observed M33 and was formed in a similar environment. This means that our analysis of the origin of the M33 analogue may provide insights into the mechanisms driving the age gradient, in the real M33, in particular the reversal of the age gradient that is observed. The paper is organized as follows. In Section \ref{sec:m33_counterpart} we present the simulated M33. In Section \ref{sec:m33_counterpart.subsec:construct_numerical_M33} we present the simulation's properties. In Section \ref{sec:m33_counterpart.subsec:validate_numerical_M33} we focus on the features of our candidate galaxy. The reversal of the age gradient in the SFH of M33 is presented in Section \ref{sec:m33_age_grad}. In Section \ref{sec:m33_age_grad.subsec:presentation} we present the adopted methods to analyze the age reversal, in Section \ref{sec:m33_age_grad.subsec:explanation} we discuss the implications of our study, and in Section \ref{sec:m33_age_grad.subsec:observation_predict} we give observational predictions. Finally, in Section \ref{sec:conclusions} we summarize our results. \begin{figure*} \hspace{-0.5cm} \includegraphics[width=1.\textwidth]{NEWAHF_M33_starsANDgas_gasoline_faceANDside_init80_0Zrot.pdf} \caption{Stellar and gas mass densities of the simulated M33 galaxy at z=0. \textit{Left}: stellar mass density, face-on and edge-on. \textit{Right}: gas density, face-on and edge-on. Spiral features and a thin warped disc of gas are visible, in agreement with observational data (e.g. \citealt{Corbelli1997,Kam17}).} \label{fig:density_gasANDstars} \end{figure*} | \label{sec:conclusions} We presented properties of a M33-analogue galaxy, simulated within the framework of Constrained Local UniversE Simulations \citep{Gottloeber10,DiCintio12,Carlesi16}, run with the code \code{GASOLINE} \citep{gasoline} and including supernova feedback \`{a} la \cite{Stinson06}, which allows for an efficient regulation of star formation within galaxies. The properties of the simulated M33 are in fair agreement with observational data from \citet{Ferguson07,Corbelli14} and \citet{Kam17}, in terms of mass, rotation velocity, and surface brightness. Our simulated M33 has a total virial mass of $2.7\cdot10^{11}$M$_{\odot}$ with a stellar mass of $5.1\cdot10^{9}$M$_{\odot}$, placing the galaxy on the correct expectations from abundance matching prediction, a thin extended disc with scale length value of $\sim 3.2$ kpc, and a rotation curve whose maximum value of $V=127.6$ km/s is reached at a radius of $r\sim17$ kpc, similar to what is reported in observations. The M33 simulated candidate does, however, have a small bulge in its inner region, not observed in the Triangulum galaxy: consequently we have avoided the region of the galaxy $r<3$ kpc in our analysis. In the CLUES-M33 analogue, we observe a trend of decreasing stellar age as we move towards outer radii, a sign of an inside-out formation of the disc (e.g. \citealt{Pilkington12}), in which old stars are found within the inner most regions of a galaxy and young stars in the outskirt of the disc, as a result of newly accreting gas. In order to compare the observational results with our simulations we selected stellar particles within concentric annuli regions of similar thickness as the field of view of the HST/ACS camera, that has been used to derive the cumulative star formation history (SHF) of M33 along its major axis \citep{Williams09,Barker11}. Interestingly, similarly to what was found in observations, the age gradient of stars in the simulated M33 shows a turnaround at large radii, r$\geq$ 17-20 kpc, with the percentage of old stars ($t_{f}<4$ Gys) increasing from $\sim20$ to $\sim80$ per cent, moving from radii 14 to 30 kpc. Several proposals appear in the literature for explaining such age profile turn-around: stellar migration \citep{Roskar08a,Roskar08b,RuizLara16a}, projection effects \citep{Barker11}, a decrease in gas volume density possibly related to the warps in the disc \citep{SanchezBlazquez09}, and accretion events from mergers \citep{Gill05,Sales2007,Brook12,RuizLara16a}. In this work, we demonstrate that this age reversal is mostly a result of accretion of old ($t_{f}\leq4$ Gyrs) stars from merging satellite galaxies into the main host galaxy and, to a much lesser extent, of stellar migration of old in-situ stars from the central regions towards the outskirts of M33. This result is in agreement with previous work by \cite{RuizLara16a}, using more massive galaxies. Indeed, at large radii where the age turnaround is found, about $93$ per cent of the old stars come from accretion events, while only a mere $7$ per cent were formed within the M33 galaxy disc (i.e. are in-situ): the reversal in the stellar age gradient disappears when considering only in-situ star particles. This suggests that accretion from mergers are the origin of the turnaround in our simulated M33. This scenario could be verified observationally, studying the kinematic of stars in the outer fields of M33: in-situ stars should be co-rotating with the galactic disc, and should have a relatively small velocity dispersion $\sigma$, while accreted stars, which are kinematically hot, are expected to have a random distribution in their line-of-sight velocity, and to show a large velocity dispersions (in our model, more than two times higher than the $\sigma$ of in-situ stars at the same radius). Moreover, the median age of the rotationally supported, in-situ stars, should indicate that this stellar population is young (median age of 4-5 Gyrs), unlike the pressure supported, accreted stars, causing the age turnaround, which should be all old (median age $\sim$ 11 Gyrs). We highlight that the method used is sensitive to projection effects. While changing the inclination of the galaxy did not induce an apparent change in the turnaround radius region observed in the cumulative star formation history of the galaxy (i.e. $17\leq r\leq20$ kpc), since it only provides a rough estimation of the turnaround radius; considering the correlation between the age reversal radius in the median formation time profile and the break radius of the surface brightness and stellar mass density profiles we are able to trace the projection effects through the break radius. Therefore, projection effects must be thought of carefully since they might play an important role in the determination of the true age turnaround radius. Similar to what was found in observations \citep{Ferguson07}, a break in the surface brightness profile of our M33 candidate in its inclined configuration appears at $r=17.6$ kpc ($5.3$ times its disc scale length), coinciding with the radius at which the age turnaround is found. Moreover, following the \citet{Martin-Navarro12,Martin-Navarro14} classification, we detect a truncation coexisting with an up-bending of the surface brightness profile associated with the stellar halo component of the simulated galaxy at $r=25$ kpc. Similar results are obtained from the stellar surface mass density profile of the M33 candidate, i.e. a comparable disc scale length, and a break and a truncation at the same radii. Thus, both the radial mass distribution of the star particles and their age/metallicity contributes to the reversal of the age gradient at the outskirts of the galaxy. Recently, \citet{trl17} showed, using simulations, that breaks are a consequence of the combined effects of outward-moving and accreted stars, in good agreement with our results. Finally, we note that \citet{RuizLara16a} found similar results when studying Milky Way-mass galaxies in the RADES \citep[\code{RAMSES} \textit{Disc Environment Study} simulations,][]{Few12}. In those simulations, the age reversal appears due to a combination of an inside-out growth of the disc, stellar migration (both inwards and outwards) of disc stars and accretion from old satellites: interestingly, as in our model, the age reversal was still recovered after suppressing stellar radial motion, indicating the minor relevance of stellar migration in generating the age upturn observed at large radii in massive galaxies. In the future we intend to verify if the accretion phenomenon causing the age turnaround is dependent on the specific mass accretion history of each galaxy: in order to shed light on this we would need a large statistical sample of hydrodynamically simulated halos of M33's mass. The recently developed Local Group Factory \citep{Carlesi16} could be used to this aim. | 18 | 8 | 1808.03286 |
1808 | 1808.04373_arXiv.txt | {Debris disks or belts are important signposts for the presence of colliding planetesimals and, therefore, for ongoing planet formation and evolution processes in young planetary systems. Imaging of debris material at small separations from the star is very challenging but provides valuable insights into the spatial distribution of so-called hot dust produced by solid bodies located in or near the habitable zone. We report the first detection of scattered light from the hot dust around the nearby ($d=28.33$~pc) A star HD 172555.} {We want to constrain the geometric structure of the detected debris disk using polarimetric differential imaging (PDI) with a spatial resolution of 25~mas and an inner working angle of about 0.1$''$.} {We measured the polarized light of HD 172555, with SPHERE-ZIMPOL, in the very broad band (VBB) or RI filter ($\lambda_c=735$~nm, $\Delta\lambda=290$~nm) for the projected separations between $0.08''$ (2.3 au) and $0.77''$ (22 au). We constrained the disk parameters by fitting models for scattering of an optically thin dust disk taking the limited spatial resolution and coronagraphic attenuation of our data into account.} {The geometric structure of the disk in polarized light shows roughly the same orientation and outer extent as obtained from thermal emission at 18 $\mu$m. Our image indicates the presence of a strongly inclined ($i\approx 103.5^\circ$), roughly axisymmetric dust belt with an outer radius in the range between 0.3$''$ (8.5 au) and 0.4$''$ (11.3 au). An inner disk edge is not detected in the data. We derive a lower limit for the polarized flux contrast ratio for the disk of $(F_{\rm pol})_{\rm disk}/F_{\rm \ast}> (6.2 \pm 0.6)\cdot 10^{-5}$ in the VBB filter. This ratio is small, only $\sim$9~\%, when compared to the fractional infrared flux excess ($\approx 7.2\cdot 10^{-4}$). The model simulations show that more polarized light could be produced by the dust located inside $\approx 2$ au, which cannot be detected with the instrument configuration used.} {Our data confirm previous infrared imaging and provide a higher resolution map of the system, which could be further improved with future observations.} | Young stars are often surrounded by circumstellar dust debris disks or rings. These consist of solid bodies, such as planetesimals and comets, as well as large amounts of dust and small amounts of gas. Small dust grains with a broad size distribution are generated in steady collisions of solid bodies and perhaps by the evaporation of comets. Debris disks are usually recognized by infrared (IR) excess on top of the spectral energy distribution (SED) of the stellar photosphere because of the thermal emission of the heated dust grains \citep[see, e.g.,][for a review]{Wyatt2008, Matthews2014}. In scattered light, debris disks are usually very faint, i.e., $> 10^3$ times fainter than the host star, and therefore they are difficult to image. To date, several dozen debris disks have been spatially resolved in various wavelength bands, from visible to millimeter, using various space and ground-based telescopes. These data provide important constraints on the debris disk morphologies \citep[e.g.,][]{Moerchen2010, Schneider2014, Choquet2016, Olofsson2016, Bonnefoy2017}. In the absence of imaging data, the location of the bulk of dust in the system can be inferred from the grain temperature derived by SED modeling. Based on modeling results, most debris disks have been found to harbor warm dust, meaning that the temperature of small grains lies between 100 and approximately 300 K \citep[e.g.,][]{Moor2006, Trilling2008, Chen2011, Morales2011}. Cold debris reside far away from the host star where their temperature does not exceed 100 K. The prominent analogs to warm and cold dust belts are the main asteroid belt at 2-3.5 au and the Kuiper belt between 30 and 48 au in the solar system \citep{Wyatt2008}. Aside from this, several stars with hot disks have been found \citep{Wyatt2007, Fujiwara2009} based on the IR excess with an estimate temperature above 300 K. The stars that possess hot dust are very interesting targets because the submicron-sized grains at a small small distance from a star could be a signpost for transient collision events. HD 172555 (HIP 92024, HR 7012) is one of the most studied objects among these special stellar systems \citep{Cote1987, Schuetz2005, Chen2006, Wyatt2007a, Lisse2009, Smith2012, Riviere-Marichalar2012, Johnson2012, Kiefer2014, Wilson2016, Grady2018}. \begin{table*} \caption[]{Log of observations with the atmospheric conditions for each run.} \centering \label{Settings} \renewcommand{\arraystretch}{1.3} \begin{tabular}{cccccccccc} \hline \hline Date&Observation & Field & & \multicolumn{4}{c}{Observing conditions (on average)} \\\cline{5-8} &identification$^1$ & offset [$^\circ$] & & Airmass & Seeing ["] & Coherence time [ms] & Wind speed [ms$^{-1}$] \\ \hline \hline \noalign{\smallskip} 2015-06-21&OBS172\_0004-0036 & 0 & &1.32& 1.39 &1.1 & 7.9\\ 2015-06-26&OBS177\_0010-0042 & 120 & &1.31 & 0.78 &3.5 & 6.7\\ 2015-07-12&OBS193\_0049-0081 & 60 & &1.33 & 1.41 &0.9 & 1.8\\ 2015-09-06&OBS249\_0028-0060 & 0 & &1.35 & 0.93 &3.8 & 9.0\\ 2015-09-06&OBS249\_0061-0094 & 60 & &1.47 & 1.11 &3.4 & 8.7\\ 2015-09-07&OBS250\_0002-0034 & 120 & &1.55 & 2.01 &1.0 & 16.1\\ \hline \hline \noalign{\smallskip} \end{tabular} \begin{flushleft} {\bf Notes.} $^{(1)}$ The observation identification corresponds to the fits-file header keyword ``origname'' without prefix ``SPHERE\_ZIMPOL\_''. The first three digits give the day of the year followed by the four-digit observation number. \\ Instrument setup for the deep imaging mode: Fast polarimetry, VBB filter, derotator mode P2, coronagraph V\_CLC\_MT\_WF, and total exposure time for each run 42.7 min. \\ Instrument setup for the flux calibration mode: Fast polarimetry, VBB filter, density filter ND2, derotator mode P2, and total exposure time for each run 14.4 s.\end{flushleft} \end{table*} HD 172555 is a V = 4.8$^{m}$, A7V star \citep{Hog2000, Gray2006} at a distance of $28.33\pm 0.19$ pc \citep{GaiaCollaboration2016} and a member of the $\beta$ Pictoris moving group. The deduced age for the group is $23\pm3$ Myr \citep{Mamajek2014}. HD 172555 has a K5Ve low mass companion CD-64 1208, separated by more than 2000~au or 71.4$''$ \citep{Torres2006}. The Infrared Astronomical Satellite (IRAS) identified HD 172555 as a Vega-like star that has a large IR excess with an effective temperature of 290~K \citep{Cote1987}. Mid-infrared (mid-IR) spectroscopy from the ground \citep{Schuetz2005} and 5.5-35 $\mu$m spectroscopy obtained with the \textit{Spitzer} Infrared Spectrograph (IRS) show very pronounced SiO features \citep[][their Fig. 2(d)]{Chen2006} indicating large amounts of submicron sized crystalline silicate grains with high temperatures (> 300 K). \cite{Lisse2009} proposed, based on the spectral analysis of mineral features, that the IR-excess emission originates from a giant collision of large rocky planetesimals similar to the event that could have created the Earth-Moon system. Until now, only one resolved image of the HD 172555 debris disk existed that was obtained in the Qa band ($\lambda_c=18.30\, \mu$m, $\Delta\lambda=1.52\, \mu$m; hereafter Q band) with TReCS (Thermal-Region Camera Spectrograph) at Gemini South telescope and described by \citet{Smith2012}. They detected extended emission after subtraction of a standard star point spread function (PSF), which is consistent with an inclined disk with an orientation of the major axis of $\theta=120^{\circ}$, an inclination of $\approx75^{\circ}$, and a disk outer radius of 8 au. In Si-5 filter ($\lambda_c=11.66\, \mu$m, $\Delta\lambda=1.13\, \mu$m; hereafter N band) data, the disk was not detected. Combining these results with the visibility functions measured with the MID-infrared Interferometric instrument (MIDI), \cite{Smith2012} concluded that the warm dust radiating at $\sim$10 $\mu$m is located inside 8 au from the star or inside the detected 18 $\mu$m emission. The proximity of the host star and strong excess in the mid-IR \citep[$L_{\rm {IR}}/L_{\ast}=7.2\cdot 10^{-4}$;][]{Mittal2015} from hot dust, makes HD 172555 an excellent, yet challenging, object for the search of scattered light. In this paper, we report the first detection of scattered light from the hot debris disk around HD 172555 with polarimetric differential imaging (PDI). We present the PDI data from 2015 taken with ZIMPOL (Zurich IMaging POLarimeter; Schmid et al. 2018, submitted) at the Very Large Telescope (VLT) in Chile. The ZIMPOL instrument provides high contrast imaging polarimetry, coronagraphy with a small inner working angle of $\sim$0.1$''$, and high spatial resolution of $\sim$25 mas. The next section presents the available data together with a brief description of the instrument. Section 3 describes the data reduction and Section 4 the obtained results. Section 5 includes analysis of the observed disk morphology and photometry. For the interpretation of the data we calculated three-dimensional (3D) disk models for spatial distribution of dust to put constraints on the geometric parameters of the disk and to estimate the flux cancellation effects in the Stokes $Q$ and $U$ parameters. In Sect.~\ref{Discussion}, we compare our results with previous TReCS and MIDI observations \citep{Smith2012}. We also investigate the possible presence of a very compact dust scattering component that would not be detectable in our data. We conclude and summarize our findings in Sect. \ref{s_Summary}. | \label{Discussion} \subsection{Comparison with thermal light detection and interferometry} \label{s_Smith} The disk geometry is consistent with the $18~\mu$m image presented in \cite{Smith2012}. They found two lobes of extended emission along PA $= 110^\circ$ with a separation of about $0.4''$ from the star and a flux of $105$~mJy. On top of this, they measured a roughly seven times stronger (732~mJy) unresolved dust component and a stellar flux of 202~mJy. In the N-band image the dust was not resolved by \citet{Smith2012}, but their N-band MIDI interferometry resolved the dust fully, putting it at separations between $>0.035''$ and $<0.27''$ ($1-8$~au). Therefore, we also expect that a lot of polarized light from the dust scattering contributes to the signal inside the inner working angle of the SPHERE/ZIMPOL observations at $0.12''$, which cannot be measured. The disk modeling of the infrared flux by \citet{Smith2012} yields a disk with radius $r=0.27''$ and width $dr= 1.2r,$ which is equivalent to the ring with constant surface brightness between $0.1''$ and $0.4''$, inclined at 75$^\circ$ to the line of sight and at PA$= 120^\circ$ for the disk major axis. We confirm these parameters with our observation and can provide more stringent parameters for the inclination and the disk orientation because of the higher spatial resolution of our data. Also the disk extension agrees, but it is not clear whether the scattered light emission and thermal emission should show the same radial flux distributions. In Fig.~\ref{f_SED_LF}, we compare the flux distribution $\lambda F_{\lambda}$ from the stellar photosphere with the thermal emission of the disk and the polarized flux distribution measured in this work. Photometric data points of HD 172555 are listed in Table~\ref{t_SED} and plotted as green diamonds in Fig.~\ref{f_SED_LF}, which also includes a $Spitzer$/IRS spectrum\footnote{downloaded from http://www.stsci.edu/~cchen/irsdebris.html}. The stellar spectral flux density is approximated by a Planck function for the star temperature of 7800 K \citep{Riviere-Marichalar2012}. A black asterisk (labeled VBB) denotes the stellar magnitude measured in the VBB (this work). The near-IR aperture photometry of HD 172555 in the N and Q bands (green diamonds at 11.7 and 18.3 $\mu$m), performed by \citet{Smith2012} using TReCS imaging data, is indicated by capital letters `$N$' and `$Q$'. The thermal flux from the disk is shown as a blackbody emission with the temperature of 329 K found by \cite{Riviere-Marichalar2012} (orange solid line). The net polarized flux measured in the $Q_\varphi$ image (Sect.~\ref{Contrast}) is indicated by a light blue asterisk (labeled with '$I_{pol}=Q_\varphi$'), and the polarized flux corrected for the polarimetric cancellation effects and aperture size is indicated by a blue asterisk in this figure. The blue dotted line shows an approximate curve for the SED of the polarized light obtained by scaling down the SED of the star. Using our best-fitting grid model we roughly estimated the scattered flux from the disk (averaged over full solid angle) in the VBB filter. This conversion neglects the contribution from the diffracted light and assumes that the maximum polarization fraction of the phase function for the polarized flux is $p_{\rm max}= 0.3$. This value is denoted with a yellow asterisk and labeled '<$I_{sca}$>'. We used this magnitude as a proxy for the scattered flux of the disk to compare it with the maximum thermal flux of the disk at $\lambda = 10$ $\mu$m and thus to estimate a kind of the disk albedo. We obtain <$I_{sca}$>$\,(p_{\rm max}= 0.3)$/$\lambda F_{\lambda}\,$($\lambda = 10$ $\mu$m) = 0.54, which should be considered as a very rough estimate for this ratio. The lower and upper limits on the scattered flux $I_{sca}$, shown in Fig.~\ref{f_SED_LF}, are obtained assuming maximum polarization fraction $p_{\rm max}= 0.1$ (for the upper limit) and $p_{\rm max}= 0.7$ (for the lower limit). \begin{table} \caption[]{Photometry of HD 172555. \label{t_SED}} \centering \begin{tabular}{lcccc} \hline \hline \noalign{\smallskip} Filter & Wavelength & Flux density & Error & Ref.\\ & ($\mu$m) & (Jy) & (Jy) & \\ \hline \noalign{\smallskip} Johnson $B$ & 0.44 & 53.87 & 0.94 & 1 \\[5pt] Johnson $V$ & 0.55 & 38.45 & 0.53 & 1 \\[5pt] Johnson $R$ & 0.7 & 33.40 & 0.71 & 1\\[5pt] Johnson $I$ & 0.9 & 35.74 & 0.82 & 1 \\[5pt] Si-5 ($N$) & 11.7 & 1.120 & 0.067 & 2 \\[5pt] Qa ($Q$) & 18.3 & 1.039 & 0.085 & 2 \\[5pt] PACS70 & 70 & 0.191 & 0.005 & 3 \\[5pt] PACS100 & 100 & 0.089 & 0.003 & 3 \\[5pt] PACS160 & 160 & 0.036 & 0.02 & 3 \\[5pt] \noalign{\smallskip} \hline \hline \noalign{\smallskip} \end{tabular} \begin{flushleft} {\bf References.} (1) \citet{Johnson1966}; (2) \cite{Smith2012}; (3) \cite{Riviere-Marichalar2012}. \end{flushleft} \end{table} \begin{figure*} \centering \includegraphics[width=16cm]{model_hidden_v3_minmax_corr.pdf} \caption{Model of the small debris disk that could be hidden behind the coronagraphic mask. The model parameters correspond to those of the mean model (Table \ref{t_results}) except the radius of the belt and scale height of the dust vertical distribution, which are reduced by a factor of 5. ({\bf a}) $Q_\varphi$ image of the debris disk model showing the polarized intensity before being convolved with the instrumental PSF. ({\bf b}) $Q_\varphi$ image of the debris disk model showing the polarized intensity after convolution with the instrumental PSF. The black circle denotes the edge of the coronagraphic mask. ({\bf c}) $U_\varphi$ image of the model demonstrating nonzero $U_\varphi$ signal appearing after convolution of the Stokes $Q$ and $U$ parameters with the instrumental PSF. The color bars show counts per pixel. \label{f_hidden}} \end{figure*} \subsection{Hidden source of thermal emission } According to the mid-IR observations of \cite{Smith2012}, most of the dust is located inside 8~au, down to 1~au. Therefore, we investigate with model calculations whether a strong but compact source of scattered light could be present, which is not detectable in our data because of the limited resolution and inner working angle limit of the used coronagraph. The occulting spot in our coronagraphic data has a radius $r\sim0.08''$ (2.27 au). We model a small disk with a radius of 2 au (0.07$''$) with exactly the same geometric morphology as the best-fit grid model derived in Sect~\ref{Modelling}, but just smaller by a factor of 5. This small disk is shown als $Q_\varphi$ image in Fig.~\ref{f_hidden}(a), together with the convolved $Q_\varphi$ image with the size of the coronagraphic mask indicated in Fig.~\ref{f_hidden}(b). First, we notice a large difference between $Q_\varphi$ intrinsic polarized flux and the convolved polarized flux of a factor of 2.5 because of the very strong polarimetric cancellation for such a compact disk. Then, there is another factor 2.67 between the total convolved polarized flux from the disk and the convolved polarized flux that falls into the rectangular measuring areas of our observations shown in Fig.~\ref{Qphi}(a) because a major region of the polarization signal is hidden by the coronagraphic mask. Thus, only about 15~\% of the polarized flux produced by an inner compact disk would contribute to our measurement. We measure a net flux of 1540 ct/s in the measuring area of our observation and estimate that it is not possible to recognize an inner disk in our data, which contributes less than 150 ct/s to our measurement. Thus, an unseen compact disk with the geometry as described above and an intrinsic polarized flux of 6.675 x 150 ct/s could be present without being in conflict with our observations and even more scattered light could be hidden for a more compact inner disk. \subsection{Comparison between polarized flux and thermal emission} \label{s_Lambda parameter} To characterize the scattering albedo of the dust, we compute the $\Lambda$ parameter describing the ratio of the fractional polarized light flux to fractional infrared luminosity excess of the disk \citep{Engler2017}, i.e., \begin{displaymath} \Lambda = {(F_{\rm pol})_{\rm disk}/F_{\rm \ast} \over L_{\mathrm{IR}}/L_{\ast}}, \end{displaymath} where the ratio of total polarized flux of the disk to the stellar flux for HD 172555 is $(F_{\rm pol})_{\rm disk}/F_{\rm \ast} \geqslant (6.2 \pm 0.6)\cdot 10^{-5}$ (Sect. \ref{Contrast}). Adopting the ratio of the disk infrared luminosity to stellar luminosity $L_{\rm {IR}}/L_{\ast}$ equal to $7.2\cdot 10^{-4}$ \citep{Mittal2015}, we obtain a lower limit for $\Lambda \geqslant 0.086$. This value is slightly smaller than the $\Lambda$ parameters we have estimated for the F star HIP 79977 and the M star AU Mic ($\Lambda_{\rm HIP\, 79977}=0.11$ and $\Lambda_{\rm AU\, Mic}=0.55$) in \citet{Engler2017}. In this paper, we presented images of polarized scattered light from the debris disk around HD 172555 obtained in the VBB filter using differential polarimetry. We found that the observed polarized intensity is consistent with an axisymmetric dust distribution, which can be interpreted as a parent belt of planetesimals or disk with a radius in the range between 0.3$''$ (8.5 au) and 0.4$''$ (11.3 au). We analysed the disk structure and obtained the following results: \begin{itemize} \item The PA for the disk major axis is $\theta_{\rm disk} = 112.3^{\circ} \pm 1.5^{\circ}$ on the sky. \item The disk emission is slightly shifted in NNE direction from the star indicating that the front side or forward-scattering part of the disk is on the NNE side. The observed dust distribution can be described with the HG asymmetry parameter $g\approx0.7$ and disk inclination $\approx103.5^{\circ}$. \item Data analysis and modeling results do not suggest the clearing of the inner disk regions (inside $r=0.3''$). \item The total disk magnitude in polarized flux in the VBB filter is $mp_{\mathrm{disk}}(\mathrm{VBB})$ = 15.20$^m \pm$ 0.37$^m$ and the stellar flux is $m(\mathrm{VBB})$ = 4.68$^m \pm 0.03^m$. This gives a disk to star contrast $(F_{\rm pol})_{\rm disk}/F_{\rm \ast}$ of $(6.2 \pm 0.6)\cdot 10^{-5}$. The measured peak surface brightness of the polarized light is ${\rm SB} \mathrm{_{peak}(VBB)} = 13.3^m \pm 0.3^m$ arcsec$^{-2}$. This corresponds to a surface brightness contrast of ${\rm SB} \mathrm{_{peak}(VBB)} - m \mathrm{_{star}(VBB)} = 8.62$ mag arcsec$^{-2}$. \item When compared with the fractional infrared luminosity of the disk using the $\Lambda$ parameter, the fractional polarized light flux in the VBB filter makes up $\sim 9\%$. \end{itemize} Our data demonstrate high sensitivity and ability of ZIMPOL to resolve the polarized light from the hot debris around a nearby A star as close as $0.1''$ or $\sim$3 au. It would be worth reobserving HD172555 in the I band without the coronagraph. Additional observations of this target under excellent observing conditions, i.e., seeing $\leqslant 0.7$, airmass $\leqslant 1.4$, coherence time $ \geqslant 3.5$ ms, and wind speed $>3$ m/s, to avoid the low wind effect (see SPHERE Manual) and with different offsets of the sky field on the detector would allow us to better distinguish between the instrumental features in the PSF and the real astrophysical signal. In the VBB filter the speckle ring extends from $0.30''$ to $0.45''$ and overlaps the region of interest. In the I band, the speckle ring is further outside of $\sim0.40''$ and the achievable contrast is higher under optimal observing conditions even if the amount of photons received in the I band is a factor of 0.6 lower than in the VBB. Moreover, noncoronagraphic data taken with I band are more suitable for the measurement of the beam shift between the two orthogonal polarization states and, hence, the correction of the beamshift effect (Schmid et al. 2018, submitted) is much easier. Such observation may also allow us to probe the innermost dust in the HD 172555 system. | 18 | 8 | 1808.04373 |
1808 | 1808.06140_arXiv.txt | Scaling relations among structural and kinematical features of 79 late-type spiral and dwarf irregular galaxies of the SPARC sample are revisited or newly established. The mean central surface brightness $\mu_0$ $<$$\mu_{0,[3.6]}$$>$$= 19.63\pm0.11\, \mathrm{mag\,arcsec^{-2}}$ allows for a clear-cut distinction between low and high surface brightness galaxies. At a given luminosity, LSB galaxies are more extended than HSB galaxies and the rotation curves have smaller inner circular velocity gradients $dv(R_d)/dr$ at one disk scale length $R_d$. Irrespective of luminosity, the geometry of rotation curves is characterized by the relation $dv(R_d)/dr \approx v_\mathrm{max}/R_\mathrm{max}$, with $v_\mathrm{max}$ being the maximum circular velocity reached at $R_\mathrm{max}$. For the rotation curve decompositions disk mass-to-light ratios are restricted to have constant, but semi-free best-fit values 0.2, 0.5, or 0.8 $M_\odot/L_\odot$ at [3.6]; they exhibit an asymmetric bimodal distribution with the dominant peak located at the median value of 0.2 (minimum disks) and with the subdominant peak at 0.8 (maximum disks). Assuming dark matter halos of Burkert and of pseudo-isothermal (PITS) type, the former provide better fits for about two thirds of all galaxies. While the halo core densities $\rho_0$ are about equal, the core radii $r_0$ of PITS halos are systematically lower by a factor of about 0.6 as compared with those of the Burkert type. Focussing on the Burkert halo, the baryonic mass fraction at intermediate radii is included to address both an adjusted baryonic Tully-Fisher relation and the significance of deviations from the mean radial acceleration relation. The average radial decrease of the baryonic mass fraction within galaxies is quantified. The Burkert halo parameters obey $\rho_0$$\,\propto\,$$r_0^{-1.5\pm 0.1}$ with considerable scatter, but allowing $v_\mathrm{max}$ as a third variable we find $\rho_0$$\,\propto\,$$r_0^{-1.84\pm0.07} v_\mathrm{max}^{2.00\pm 0.11}$ with small scatter. The halo central surface density $\rho_0 r_0$, with a sample median $<$$\rho_0 r_0$$>$$\,\approx 121\, \mathrm{M_\odot pc^{-2}}$ ($\sigma = 112$), weakly correlates with $\mu_0$ and with compactness C and strongly correlates with the observed radial acceleration $g_{\mathrm{obs}}=v^2_{\mathrm{obs}}(r)/r$ at different galactocentric radii. Consequently, because $R_\mathrm{max} \propto r_0$, we have a tight central halo column density versus maximum circular velocity relation $v_\mathrm{max}^2 \propto \rho_0 r_0^2$. Halo cores barely extend over the luminous disk, but their sizes do not correlate with the optical radii. We introduce an alternative to a prominent conventional universal rotation curve; it is based on the non-singular total matter density profile $\rho_\mathrm{total}(r)= (v_\mathrm{max}^2/4\pi G r^2)\left( 1 - (1-r/r_c) \exp(-r/r_c) \right)^2$, with the scaling parameter $r_c$ correlating with the halo core size $r_0$. Fitting the synthetic URC to a selection of galaxies, the co-added doubly-normalized rotation curves exhibit a high degree of similarity. A couple of analytic URC decompositions into a baryonic disk and a dark matter component is accomplished. | Late-type spiral and dwarf irregular galaxies are composed of, firstly, bulgeless and non-barred gaseous and luminous stellar disks, where accretion, star formation, and stellar evolution takes place. Within classical Newtonian physics, disks are recognized to be prominently surrounded by, secondly, extended nonluminous accumulations of invisible matter, called dark matter (DM) halos. The presence of DM is inferred only indirectly due to a luminous mass (LM)-deficiency or mass-degeneracy. Historically, the first dynamical evidence for Galactic non-luminous matter goes back to Lord Kelvin who applied the kinetic theory of gases to stellar systems and in particular to a handfull of galactic-plane stars with known velocity dispersion \citep{Thomson04}. In extragalactic astronomy, the mass-degeneracy was first put on firm quantitative grounds by \citet{Zwicky33} on applying the virial theorem to the Coma cluster of galaxies, realizing an indisputable lack of luminous matter and coining the term "dark matter". For disk galaxies of all corresponding Hubble types, DM is typically inferred since the late 1950ies from the amplitude and shape of their rotation curves (RC) that disagree with Keplerian motion, in particular, the flatness observed way beyond the optical radius of the luminous disk indicates a hidden mystery. The first accurate RC was determined for M31 in the radio and goes back to \citet{vandeHulst57} and \citet{Schmidt57}. For a historical review see \citet{Bertone16} and references therein. One path to an understanding of the nature of DM leads over finding model relations between observed or infered features of luminous and dark matter. Luminous matter particularly in late-type spiral and dwarf irregular disk galaxies typically goes with exponential surface brightness profiles at intermediate radii (with central surface brightness and disc scale length as structural parameters) and for the dynamics rotating disks close to centrifugal equilibrium are adopted \citep{VanderKruit11}. The corresponding disk mass models predict at most radii circular velocities way below to what is observed, even if maximum disks with high mass-to-light ratios are assumed. The excess of the observed total circular velocity may be attributed to non-detected or dark matter \citep[e.g.,][]{Freeman70, Rubin70,Rubin83, vanAlbada85}. Parts of these halos are composed not of some enigmatic type of matter, but simply are baryonic due to the speedy outflow of stellar debris due to supernovae, an ongoing process within and around each disk galaxy since the early times of its formation. There is a vast literature on suggested parametrizations for total DM halo density profiles \citep[e.g.,][for double power laws and the Einasto profile]{An13}. For late-type spiral and low-luminous disk galaxies the rotation curves are particulary well described by corresponding halo density profiles with, on one hand, non-singular and non-steep ("constant") density cores, as with the \citet{Burkert95} profile, the PSS profile \citep{Persic96, Borriello01}, or pseudo-isothermal spherical (PITS) profiles \citep{Kormendy04, Kormendy16}. On the other hand, simulations of galaxies within a lambda cold dark matter ($\Lambda$CDM) universe with hierarchical clustering predict cuspier, steep core density profiles that even become singular at the galaxy center, as with the widely used three-parameter double-power law by \citet{Zhao96}. This model includes the isothermal profile (ISO), the Navarro-Frenk-White \citep[NFW,][]{Navarro96} profile and its generalization \citep[gNFW, e.g.,][]{Moore99}, and the DC14 profile by \citet{DiCintio14b}. The term "constant core" is owing to $\log \rho(r)-\log r$-diagrams wherein singular profiles show everywhere a falling straight line while cored non-singular profiles exhibit a nearly flat part in the central region. Overall, there is a prevailing degeneracy with respect to halo density profiles. A review on the cusp-core-controversy is given in \citet{deBlok10}. Comparing cored halo models, the results do not allow to conclusively select a single superior one, \citet{Breddels13} even suspect that more than one type of parametric halo profile is necessary for realistic assessments. To give an idea, while \citet{Adams14} compares fits to the RCs of dwarf galaxies using Burkert and gNFW profiles and finds the latter to slightly perform better, \citet{Pace16} finds the PITS halo and the DC14 profiles to perform equally well and to outperform the Burkert and the NFW profiles. The result of \citet{DiCintio14b} is motivated by hydrodynamical N-body simulations and includes as special cases both the isothermal and the gNFW halos, favored within the $\Lambda$CDM scenario. These hydrodynamical numerical approaches as well as semi-analytic models manage to transform cusped DM halos into cored halos by means of galaxy evolution with supernovae feedback or due to interaction of baryon clumps with dark matter as secular dynamical processes \citep[e.g.,][]{Governato10, Oh11b, DiCintio14a, Katz17, DelPopolo18}. For the sake of analytical simplicity, halo models used for RC decompositions usually assume spherical symmetry, but most N-body simulations agree in their prediction that evolving halos in a CDM universe with hierarchical clustering are more accurately described by triaxial systems \citep[e.g.,][]{Bryan13, Despali14}. For dwarf galaxies, however, the asphericity seems to be very moderate \citep{Trachternach09}. According to their star-formation history, late-type spirals and dwarf irregulars are capable to reproduce their stellar mass over the cosmic time, with bursty episodes \citep[e.g.,][]{Karachentsev18}. Star-forming regions within rotationally supported low-mass disk galaxies are systematically found out to the outer disk until more than two and a half optical radii (corresponding to about eight exponential disk scale lengths and about four times the size of typical halo cores) \citep{Parodi03, Hunter16}. This is more than twice as far as the onset of the flat part of typical RCs. Stellar truncation is due to reaching a critical star formation threshold density \citep{vandenBosch01b}. Disk supported galaxies of all luminosities are subject to the baryon-halo conspiracy: they obey both the \citet{Tully77} relation (TFR), linking the maximum circular velocity with disk luminosity, and the baryonic Tully-Fisher relation (BTFR), linking the total baryonic mass with the observed rotational dynamics. These structure-dynamics relations are indirectly verified by the recently established mass-deficiency versus radial acceleration relation (MDAR) or equivalenty by the two-radial accelerations relation (RAR) of rotationally supported galaxies including dwarf disc galaxies \citep{McGaugh14a,McGaugh16}. Typically the observed rotation curves (RC) exhibit a linear or nearly linear increase at small galactocentric radii before they curve at intermediate radii to become flat or, in some cases, start to moderately decrease after reaching a maximum circular velocity. Further out, any RC will eventually drop. Irrespective of the maximum circular velocity the central velocity gradients may vary from galaxy to galaxy, from steep to rather flat. In general, there's a variety of possible amplitudes, from slow to fast, and of possible shapes of RCs, from narrow to extended curvatures. Despite of the diversity of observed RCs \citep{Oman15}, based on cored halo models the increasingly successful composition of universal rotation curves (URC), in particular for dwarf disc galaxies \citep{Karukes17, DiPaolo18, Lapi18}, continuously adds to the confidence concerning a deeper understanding of DM halos. A panopticum of other scaling relations among various variables describing structural and kinematical properties of disk galaxies are known and kept being refined, such as LM parameter correlations \citep{Giovanelli99,Courteau07}, DM correlations among halo parameters (e.g., the supposed uniformity of halo central surface density), or DM-LM relations \citep[e.g.,][and references therein]{Kormendy04, Kormendy16, Lapi18}. In this paper we address a few known and investigate some new structural and kinematical scaling relations for a sample of 79 late-type spiral and dwarf irreguar galaxies, based on a homogenous data set. The photometry and preprocessed kinematic data are taken from the SPARC database and consistently postprocessed by us with respect to halo mass model decompositions. We proceed as follows: in Sect. 2 we present the data (Sect. 2.1), include an absolute magnitude independent distinction of high-surface and low-surface brightness at 3.6$\mu$m (Sect. 2.2), and describe the decomposition procedure applied by us to the galaxy RCs in order to obtain Burkert and PITS halo parameters (Sect. 2.3). We thereby present a peculiarity concerning the distribution of best-fit mass-to-light ratios. Section 3 presents and discusses the investigated scaling relations, which are the RAR (Sect. 3.1), the halo core density versus size relation subject to a third variable (thereby providing an original new result, Sect. 3.2), some kinematic dependencies of the central halo surface density (and refuting its supposed uniformity, Sect. 3.3), an inner circular velocity gradient (based a new simple derivation, Sect. 3.4), an adjusted BTFR for varying radii (Sect. 3.5), the central halo column density versus maximum velocity relation (Sect. 3.6), and finally the URC in a conventional as well as in an alternative new form, the latter being based on a genuinely proposed total matter density profile (Sect. 3.7). Section 4 summarizes the results and takes a glance at further topics. Finally, the appendix contains a couple of preliminary formal decompositions of the alternative URC (App. A) as well as three tables with selected photometric, kinematic, and halo structural data (App. B). | In this study, we confirmed, refined, or controversely discussed some details concerning well-known scaling relations or laws for late-type spiral and dwarf irregualar galaxies, linking the structure and the dynamics of these rotationally supported systems. We additionally contributed some new scaling relations that were not yet (to the best of our knowledge) discussed in the literature. The following results were obtained. \begin{itemize} \item Rotation curve (RC) mass model decompositions for 79 late-type spiral and dwarf irregular galaxies of the SPARC sample are provided, assuming spherically-symmetric dark matter halos of Burkert and of pseudo-isothermal (PITS) type. \item Disk mass-to-light ratios are restricted to have constant, but semi-free best-fit values 0.2, 0.5, or 0.8 $M_{\sun}/L_{\sun}$ at [3.6]. They exhibit an asymmetric bimodal distribution with the dominant peak located at the median value of 0.2 (minimum disks) and the subdominant peak at 0.8 (maximum disks). It remains to be checked, whether this is artefactual or really related to some physical parameters like baryonic mass fraction or galaxy color and metallicity. \item As compared to decompositions with PITS halos, those with Burkert halos provide better fits for about two thirds of all galaxies. Comparing the best-fit halo parameters for either approach, the central densities $\rho_0$ have similar values, but the scale lengths $r_0$ of the Burkert model are systematically higher by a factor of 1.80$\pm$0.05. This is of relevance for the calculation of, e.g., mean central surface densities for a given sample of galaxies. Within galactocentric radii smaller than a few halo scale lengths, $r_0$-scaled versions of the Burkert and the PITS model represent the generalized Burkert-PITS halo density profile $\rho (r)=\rho_0 (1+r/r_0)^{-\eta}$ $(1+(r/r_0)^2)^{-1}$ with 0$\le$$\eta$$\le 1$. \item Focussing on the Burkert halo, we use the two-radial accelerations relation (RAR) diagram for three purposes: based on the mass model relation $g_{\mathrm{obs}}=g_{\mathrm{bar}} (1+(v_{\mathrm{halo}}/ v_{\mathrm{bar}})^2)$, with the factor within the outer brackets being the inverse baryonic mass fraction that quantifies the mass deficiency, we (i) illustrate the numerical accuracy of our RC decompositions; (ii) we emphasize the information content in the spreaded distribution of the data points; and (iii) we confirm some decrease of the baryonic mass fraction when going from HSB to LSB galaxies. \item The latter distinction is anchored on the observation that for any given absolute magnitude (i) the central surface brightnesses have a uniform mean value $<$$\mu_{0,[3.6]}$$>$$= 19.63\pm0.11\, \mathrm{mag\,arcsec^{{-2}}}$, and (ii) larger-than-mean disk scale lengths (lower compactness) correspond to LSB galaxies, while smaller-than-average scale lengths (higher compactness) correspond to HSB galaxies. Consistently, given the baryonic Tully-Fisher relation (BTFR), some expected increase of the inner velocity-gradient (inferred here at $R_d$) with brighter central surface brightness is indeed observed on average. At a given luminosity, more compact (i.e., HSB) galaxies do have steeper-than-mean circular velocity gradients than less compact (i.e., LSB) galaxies. \item The average inner circular velocity gradient at one disk scale length roughly follows $\nabla v(R_d) \approx v_\mathrm{max}/R_\mathrm{max}$ (equation \ref{nablav2}). This relation may be considered as a constraint concerning the universal shape of the RC of rotationally supported low-mass galaxies. It links the kinematics within the luminous inner part of a galaxy with variables taken from the outer part, where DM dictates the motion. \item We formulate some minimum disk hypothesis as follows: DM rich disk galaxies with low mass-to-light ratios fullfill the velocity-ratio criterium $v_\mathrm{bary}(2.2\,R_d)/v_\mathrm{obs} (2.2\,R_d)\,\vline\,_{R=2.2R_d} =0.5\pm0.1$ (equation \ref{vrc}). \item The luminous and the dark matter content at inner radii are both closely linked with the observed kinematic behaviour. We report on an adjusted baryonic Tully-Fisher relation (aBTFR), that however incorporates the baryonic mass fraction as an essentially non-baryonic ingredient. Taking the DM content explicitely into account, the spread accompaning the usual BTFR relation is highly reduced. \begin{table}\centering \small \begin{tabular}{lll} \hline scaling rel.: & & \\ aBTFR & adjusted baryonic TF relation & 3.5 \\ CDC & core density vs. compactness & 3.3 \\ CDMV & column density vs. maximum velocity & 3.2 \\ CRC & core radius vs. compactness & 3.6 \\ CRRC & core radius vs. rotation curve shape & 3.6\\ MDAR & mass discrepancy - radial acceleration rel.& 3.1\\ RAR & two-radial accelerations relation & 3.1 \\ SDR & two-surface densities relation & 3.3.2 \\ VGMV & velocity gradient vs. maximum velocity & 3.4.3 \\ VRC & velocity-ratio criterium & 2.3.2\\ \hline others: & & \\ DM / LM & dark / luminous matter & \\ LSB / HSB & low / high surface brightness & 2.2 \\ OLSB & ordinary-least squares bisector & \\ RC & individual rotation curve & \\ URC & universal rotation curve & 3.7\\ \hline \hline \end{tabular} \caption{\small Overview on acronyms for some scaling relations and other terms. The columns give the acronym, its meaning, and the number of the section in the text where the corresponding scaling relation or phenomenon is introduced or discussed in more depth. \normalsize} \label{Table2} \end{table} \item The baryonic mass fraction within individual galaxies decreases with growing galactocentric distance according to the statistical relation $f_\mathrm{bary}\propto v_\mathrm{obs}^{1.18}(R)/R$ (equation \ref{fbary2}). \item The Burkert halo parameters roughly follow $\rho_0$$\,\propto\,$$r_0^{-1.5}$. More subtle, allowing for the maximum circular velocity as a third variable, the tight relation $\rho_0$$\,\propto\,$$r_0^{-1.84\pm0.07}v_\mathrm{max}^{2.00\pm 0.11}$ with very small scatter emerges. As with the velocity gradient, this links the inner region a galaxy (actually, the halo structure) with a kinematic variable taken from the outer part, where DM dictates the motion. \item The halo central surface density $\rho_0 r_0$, with a sample median $<$$\rho_0 r_0$$>$$\,\approx 121\, \mathrm{M_\odot pc^{-2}}$, weakly correlates with the disk central surface brightness $\mu_0$ and strongly correlates with the observed radial acceleration $g_{\mathrm{obs}}=v^2_{\mathrm{obs}}(r)/r$ at different galactocentric radii. Consistently, and even more pronounced, the maximum velocity $v_\mathrm{max}$, that typically represents the flat part of the RC, is tightly proportional to the halo component of circular velocity at $r_0$, i.e. $v_\mathrm{max} \propto v_\mathrm{halo}(r_0) \propto \rho_0^{1/2} r_0$. This is equivalent to a maximum circular velocity versus central halo column density (CDMV) relation $v_\mathrm{max}^2 \propto \rho_0 r_0^2$ (equation \ref{CDMV}). Given some $v_\mathrm{max}$, larger halo cores not only correspond to smaller central halo densites but systematically go with fainter central surface brightness or less compact disks, too. \item Overall, halo cores are small (we have a median $r_0 \approx 2 R_d \approx 0.64 R_\mathrm{opt}$, hence close to a Freeman disk peak radius or good half an optical radius, to be compared with the average $R_\mathrm{max}\approx 3.9 R_d \approx 1.22 R_\mathrm{opt}$), but their size does actually not correlate with the optical radius. We did not yet look at the transition of low mass galaxies to very luminous galaxies that usually exhibit high values of $v_\mathrm{max}$ and sometimes a decreasing outer RC. Whether or not the scaling relations discussed in this paper will change due to Hubble type was not investigated here. \item Certainly, some of the scatter in our figures and the spreads in our relations could be reduced if we would restrain to a more restrictive selection of galaxies. This potentially may alter some of our fitted scaling relations, too. Whether or not the frequent occurence of a difference between observed central surface brightness and extrapolated central surface brightness plays a role in the interplay between LM and DM remains an open question. The simple ordinary least-squares bisector (OLSB) fitting procedure adopted throughout the paper allows for rudimentary inclusion of errors in both variables and seems to be a working alternative to more sophisticated routines that consider individual errors for all data points. However, the question on how to deal best with errors has no simple answer, despite its enormous influence on the produced relations. The overall consistency of the results presented above makes us nevertheless feel confident with the findings of our study. In order to attempt to further reduce the amount of uncertainty in the context of halo related relations, it seems worthwhile to adopt the hybrid Burkert-PITS (BP) model given by equation (\ref{halodens}) for the decompositions, with the fitting parameter $\eta$ possibly being somehow related to observable quantities. \item Some of our best-fit model RCs only accidentally match with the conventional universal rotation curve (URC). A better conversion is provided by the cuspy (but non-singular) total matter density profile $\rho_\mathrm{total}(r)=$ $(v_\mathrm{max}^2/4\pi G r^2)$ $( 1 - (1-r/r_c)\exp(-r/r_c))^2$, with the scaling parameter $r_c$ linearly scaling with $r_0$ of the Burkert model or being roughly estimable via a quadratic regression of $r_\mathrm{max}$, and with $v_\mathrm{max}$ being replaceable by means of the CDMV relation. Consistently, for galaxies with comparable luminosity this RC-shape parameter $r_c$ anticorrelates with the inner circular velocity gradient. We provide the corresponding formulae for the total radial mass distribution and the total velocity profile. The co-added doubly-normalized velocity profiles of a selection of nine galaxies exhibit a high degree of similarity. This is strongly encouraging us to continue the quest for a synthetic URC along this path. It is strengthened by the successful application of analytic RC decompositions of this synthetic URC into a baryonic disk and a halo component to a couple of example galaxies (Appendix A). An inquiry of this novel density profile by testing it against many more galaxies than done here is a postboned endeavour. \item A plethora of acronyms is used for brief reference to the various scaling relations and phenomena (Table \ref{Table2}). \end{itemize} If the usefullness of the alternative URC model density will be confirmed, a new playground would be provided. For example, if $r_c$ and eventually $v_\mathrm{max}$ or the parameters related to the decompositions are \emph{time-dependent}, evolutionary aspects could be addressed within our simple parameterization approach. Figs. \ref{Fig13} (right-hand panel) and \ref{UGCA442_figs} (lower right-hand panel) invite to be interpreted in terms of such a toy scenario. Of course, the astrophysical reason for the proposed total matter density profile would have to be addressed in any case, as well as its deeper interrelationship with the scaling relations or laws picked up and investigated in this study. Irrespective of any URC model these genuine scaling relations among structural and dynamical parameters stand alone and help to shape our understanding of late-type spiral and dwarf irregular galaxies. Explaining their origin will lead to refined and maybe surprising insights concerning theories of galaxy formation and on the nature of DM. | 18 | 8 | 1808.06140 |
1808 | 1808.03415_arXiv.txt | We present a systematic study of ionized gas outflows based on the velocity shift and dispersion of the [O III] $\lambda 5007$\AA \, emission line, using a sample of $\sim 5000$ Type 1 AGNs at $z<0.3$ selected from Sloan Digital Sky Survey. This analysis is supplemented by the gas kinematics of Type 2 AGNs from \citet{2016ApJ...817..108W}. For the majority of Type 1 AGNs (i.e., $\sim 89$\%), the [O III] line profile is best represented by a double Gaussian model, presenting the kinematic signature of the non-virial motion. Blueshifted [O III] is more frequently detected than redshifted [O III] by a factor of 3.6 in Type 1 AGNs, while the ratio between blueshifted to redshifted [O III] is only 1.08 in Type 2 AGNs due to the projection and orientation effect. The fraction of AGNs with outflow signatures is found to increase steeply with [O III] luminosity and Eddington ratio, while Type 1 AGNs have larger velocity dispersion and more negative velocity shift than Type 2 AGNs. The [O III] velocity $-$ velocity dispersion (VVD) diagram of Type 1 AGNs expands towards higher values with increasing luminosity and Eddington ratio, suggesting that the radiation pressure or wind is the main driver of gas outflows, as similarly found in Type 2 AGNs. In contrast, the kinematics of gas outflows is not directly linked to the radio activity of AGN. | Mass-accreting supermassive black holes are manifested as active galactic nuclei (AGNs), which are classified into two categories based on the unification model: Type 1 AGNs, of which the central engine is directly viewed, and Type 2 AGNs, of which the central engine is obscured by the torus \citep{1993ARA&A..31..473A,1995PASP..107..803U}. The observed correlation between black hole mass and host galaxy properties is typically interpreted as that feeding and feedback work together in self-regulation of black hole growth and galaxy evolution \citep{2013ARA&A..51..511K}. During the process of galaxy-galaxy interaction, for example, strong inflows supply a vast amount of gas to the supermassive black hole helping it to evolve and power AGN. The radiation emitted from AGN drives out a large amount of gas, quenching star formation and also the growth of the black hole \citep[see review by][]{2015ARA&A..53..115K}. The feedback of AGNs is considered to be traced by frequently seen gas outflows in various scales \citep[see review by][]{2012ARA&A..50..455F}. Though various mechanisms, i.e., the disc wind \citep[e.g.,][]{2003ARA&A..41..117C,2009ApJ...702L.187R,2015Natur.519..436T}, radiation pressure on dust \citep{1998MNRAS.294L..47B,2002ApJ...572..753D,2010MNRAS.402.2211A}, interaction of radio jet with clouds \citep{2005MNRAS.359..781S,2008A&A...491..407N,2012ApJ...747...95G} etc. have been considered, the main driver of AGN feedback remains largely debatable \citep[see][]{2015ARA&A..53..115K}. AGN-driven gas outflows have been detected in various energy bands, e.g., optical, UV and X-ray \citep[e.g.,][]{2003ApJ...593L..65R,2013MNRAS.436.3286A} allowing to probe AGN feedback in different scales, while the kpc-scale outflows observed in the narrow line region (NLR) are particularly important in understanding AGN feedback since they are extended to galactic scales, where the outflows may interact with interstellar medium and suppress star formation. The [O III] $\lambda 5007$\AA \, emission line being strong in AGN spectra is a good tracer of outflows and consequently subjected to various studies. On one hand, spatially resolved spectroscopy based on [O III] kinematics mapped the velocity structure in the NLR \citep[e.g.,][]{2000ApJ...532L.101C,2013ApJS..209....1F}. More detailed studies of gas kinematics and star formation became possible thanks to the integral field spectroscopy \citep[e.g.,][]{2006ApJ...650..693N,2014MNRAS.441.3306H,2016ApJ...819..148K,2016ApJ...833..171K,2017ApJ...837...91B,2017MNRAS.467.2612W,2018ApJ...858...48M,2018MNRAS.476.2760F}. On the other hand, more systematic studies of [O III] kinematics were based on the spatially integrated spectra obtained from large surveys since a large sample of AGNs can be utilized, revealing that outflows are prevalent, particularly among luminous AGNs \citep[e.g.,][]{ 2013MNRAS.433..622M,2014ApJ...795...30B,2016ApJ...817..108W,2016ApJ...831....7S,2017ApJ...839..120W,2017MNRAS.468..620Z,2018ApJ...852...26W,2018ApJ...856...76D}. The kinematics of [O III] are mainly caused by AGN outflows, while the virial motion due to the gravitational potential of the host galaxy is partly responsible for the broadening of the [O III] line \citep{1984ApJ...281..525H,1996ApJ...465...96N,2008ApJ...680..926K,2014JKAS...47..167W,2017ApJ...839..120W,2017ApJ...842....5E}. A positive correlation between mid-infrared luminosity and velocity width of [O III] suggests that the gas outflows are mainly radiation driven \citep{2014MNRAS.442..784Z}. The studies of gas outflows based on a large sample of Type 1 and Type 2 AGNs support this idea \citep[e.g.,][]{2016ApJ...817..108W, 2018ApJ...852...26W}. \citet{2016ApJ...817..108W} studied outflow kinematics of $\sim 39,000$ Type 2 AGNs at $z<0.3$ by carefully estimating velocity dispersion and shift of [O III] with respect to the stellar velocity dispersion and systemic velocity of the host galaxies. Their combined analysis of velocity dispersion and velocity shift exhibits the presence of outflow signatures in the majority of the high luminous AGNs and the strong dependence of the outflow properties on the radiation emitted by AGNs. In a series of papers on ionized gas outflows in AGNs, the demography of ionized gas outflows in type 2 AGNs was reported, respectively, based on [O III] \citep{2014ApJ...795...30B,2016ApJ...817..108W} and H$\alpha$ \citep{2017ApJ...845..131K}, while \citet{2016ApJ...828...97B} constrained the physical properties of the outflows, e.g., launching velocity, dust extinction, and the opening angle of the cone, based on the kinematical modeling of the outflows and Monte Carlo simulations. Based on these studies \citet{2017ApJ...839..120W} showed that AGNs with strong outflows tend to have a regular star formation rate similar to the star-forming galaxies in the main sequence, while AGNs with weak or no outflows have on average much lower specific star formation rate, suggesting that the effect of AGN-driven outflows is delayed. In this paper, we focus on the gas outflows of Type 1 AGNs based on the kinematics of [O III]. Compared to Type 2 AGNs, Type 1 AGNs have a number of merits. First, since the direction of the outflows is closer to the line-of-sight than that of Type 2 AGNs, the projection effect is less problematic and the measured velocity is expected to be higher. Second, the main physical parameters of AGNs, i.e., black hole mass ($M_{\mathrm{BH}}$) and Eddington ratio ($\lambda_{\mathrm{Edd}}$) can be properly estimated in Type 1 AGNs, while for Type 2 AGNs, mass and bolometric luminosity is much more difficult to estimate due to the lack of broad emission lines and AGN continuum. On the other hand, the downside of type 1 AGNs includes the difficulty of measuring the systemic velocity (e.g., based on stellar absorption lines), which is required to measure the velocity shift of outflows based on gas emission lines, and host galaxy mass or stellar velocity dispersion, which are needed to calculate the host galaxy gravitational potential for removing the effect of the virial motion in the width of gas emission lines. Therefore, it is important to combine Type 1 and Type 2 AGNs for better understanding gas outflows and their connection to AGN energetics. We investigated outflow kinematics based on [O III] for a large sample of Type 1 AGNs at $z<0.3$. By combining these Type 1 AGNs with the Type 2 AGNs from \citet{2016ApJ...817..108W}, we perform a demography of ionized gas outflows over a luminosity range of $\sim 5$ orders of magnitude. The data and spectral analysis method are presented in Section \ref{sec:data}. The main results are given in Section \ref{sec:results} and discussion in Section \ref{sec:discussion}. The summary and conclusions are presented in Section \ref{sec:conclusion}. A cosmology with $H_0=70 \, \mathrm{km \, s^{-1} Mpc^{-1}}$, $\mathrm{\Omega_m}=0.3$ and $\Omega_{\Lambda}=0.7$ is used throughout the paper. | \label{sec:conclusion} We have investigated ionized gas outflows based on the [O III] kinematics using a large sample of $\sim 5000$ Type 1 AGNs at $z<0.3$. For comparison, we combined Type 1 AGNs with the sample of $\sim 39,000$ Type 2 AGNs from \citet{2016ApJ...817..108W}. Our main findings are summarized as follows. \begin{itemize} \item For the majority of Type 1 AGNs ($\sim$89\%), the [O III] line profile presents a broad wing component, representing a non-virial motion, i.e., outflows. Compared to Type 2 AGNs, of which $\sim$43\% shows broad [O III] fitted with a double Gaussian model, outflow signature is more easily detected in Type 1 AGNs. This is partially due to the luminosity effect since the mean luminosity of Type 1 AGNs is much higher than that of Type 2 AGNs. The fraction of AGNs with double Gaussian [O III] steeply increases with AGN luminosity and Eddington ratio in Type 1 AGNs as similarly found in Type 2 AGNs. \item The velocity dispersion of [O III] strongly correlates with [O III] luminosity while Type 1 AGNs have on average higher velocity dispersion than Type 2 AGNs. \item Although many AGNs show $\sim$zero velocity shift, a significant fraction of AGNs presents strong velocity shift, suggesting various effects on the observed kinematic signatures, i.e., the inclination and opening angle of the cone and dust obscuration. The average velocity shift increases with the [O III] luminosity, while the velocity shift is larger in Type 1 AGNs than in Type 2 AGNs, reflecting the orientation and projection effect. \item The VVD diagram expands toward higher values with increasing AGN luminosity and Eddington ratio in both Type 1 and Type 2 AGNs, suggesting that outflows are radiation-driven. \item Blueshifted [O III] is more frequently detected than redshifted [O III] in Type 1 AGNs as expected from the biconical outflow model combined with dust obscuration. The ratio between AGNs with blueshifted [O III] and AGNs with redshifted [O III] is higher in Type 1 AGNs than in Type 2 AGNs. These results are consistent with the unification model of AGNs and well reproduced by the 3D bicone models. \item The apparent trend of increasing $\sigma_{\mathrm{[O\,III]}}$ with radio luminosity or radio-loudness vanishes when normalized by the velocity dispersion of [O III] core component, which is a proxy for the galaxy gravitational potential. This is in agreement with the previous findings for Type 2 AGNs where no trend of radio luminosity and dispersion have been found once [O III] velocity dispersion was normalized with stellar velocity dispersion, suggesting that outflows are not directly linked to the radio activity of AGNs. \item Mass outflow, energy injection and momentum flux rates increase with AGN luminosity. A majority of the Type 1 AGNs in our sample have mass outflow rates between $1-10 \,L_{\mathrm{bol}}/\eta c^2$ indicating powerful mass loading by AGN outflows to the interstellar medium. Energy conversion efficiency is estimated to be smaller than 0.01\% of bolometric luminosity and momentum flux is ranging between $0.01-0.1 \,L_{\mathrm{bol}}/c$. \end{itemize} | 18 | 8 | 1808.03415 |
1808 | 1808.00620_arXiv.txt | We present observations of disc-bearing stars in Upper Scorpius (US) and Upper Centaurus-Lupus (UCL) with moderate resolution spectroscopy in order to determine the influence of multiplicity on disc persistence after $\sim5-20\,\mathrm{Myr}$. Discs were identified using infra-red (IR) excess from the Wide-field Infra-red Survey Explorer (WISE) survey. Our survey consists of 55 US members and 28 UCL members, using spatial and kinematic information to assign a probability of membership. Spectra are gathered from the ANU 2.3m telescope using the Wide Field Spectrograph (WiFeS) to detect radial velocity variations that indicate the presence of a companion. We identify 2 double-lined spectroscopic binaries, both of which have strong IR excess. We find the binary fraction of disc-bearing stars in US and UCL for periods up to 20 years to be $0.06^{0.07}_{0.02}$ and $0.13^{0.06}_{0.03}$ respectively. Based on the multiplicity of field stars, we obtain an expected binary fraction of $\sim0.12^{0.02}_{0.01}$. The determined binary fractions for disc-bearing stars does not vary significantly from the field, suggesting that the overall lifetime of discs may not differ between single and binary star systems. % | \label{sec:introduction} In the age of Kepler and other planet finding surveys a number of planets have been discovered around binary star systems. Some of these planets are in S-Type orbits where the planet orbits one star in the binary system (e.g. $\gamma$ Cephei Ab \citep{neuhauser_direct_2007} and HD 196885 Ab \citep{chauvin_characterization_2007}) and others are in P-Type, or circumbinary orbits, where the planet orbits both stars (e.g. Kepler-47b and c \citep{orosz_kepler-47:_2012}, PH-1 \citep{schwamb_planet_2013} and ROXs 42Bb \citep{kraus_three_2014}). When considering planet formation, current models mostly focus on single star systems. This picture is, however, insufficient, given that a large fraction of stars are in binary star systems \citep{raghavan_survey_2010} and therefore a complete understanding of planet formation must take these environments into account. Current planet occurrence rates around binary stars show that the frequency drops for systems with separations $a\sim40\,\mathrm{AU}$ \citep{kraus_impact_2016}, which might indicate that it is difficult for planets to form around binaries of this separation. Could this be because the disc material from which the planets would form is destroyed faster in binary star systems than around single stars? There are various mechanisms that contribute to the dispersal of protoplanetary discs such as accretion, photo-evaporation and outflows. How these mechanisms change and affect the discs in binary star systems has been investigated using simulations (e.g. \citealt{artymowicz_dynamics_1994, kuruwita_binary_2017}) which find that circumstellar discs are truncated to $r_{disc} \sim a/3$ and outflows from binary star systems are less efficient at carrying away mass and momentum compared to single star counterparts. \cite{harris_resolved_2012} and \cite{cox_protoplanetary_2017} found that circumstellar discs in binaries were smaller and fainter, suggesting that they may be dispersed faster. On the surface it seems that discs have a shorter lifetime around stars in binary systems, and as a result shorten the time in which planets may form. However, they also found that circumbinary discs have at least an order of magnitude higher millimetre flux densities compared to circumstellar discs around binaries of the same separation, suggesting they are larger and have more material in the circumbinary disc to form planets. There is also a handful of other considerably old circumbinary discs (e.g. AK Sco \citep[$18\pm1\Myr$,][]{czekala_disk-based_2015}, HD 98800 B \citep[$10\pm5\,\mathrm{Myr}$,][]{furlan_hd_2007}, V4046 Sgr \citep[$12$--$23\,\mathrm{Myr}$,][]{rapson_combined_2015} and St 34 \citep[also known as HBC 425, $\sim$25$\,\mathrm{Myr}$,][]{hartmann_accretion_2005} compared to the typical lifetime of protoplanetary discs of $3\,\mathrm{Myr}$ (\citealt{haisch_disk_2001,mamajek_initial_2009}). If circumbinary discs always have a significantly longer lifetime than discs around single stars, it would provide a greater opportunity for planets to form in these systems. In this work we aim to determine what fraction of disc-bearing stars are in binary star systems in the OB associations Upper Scorpius (US) and Upper Centaurus-Lupus (UCL). These OB associations are part of the larger Scorpius-Centaurus-Lupus-Crux (Sco-Cen) association which is the nearest region of recent massive star formation at a distance of $\sim140$~pc \citep{zeeuw_hipparcos_1999}. The subgroups of the Sco-Cen association show a well known age gradient from $\sim5$~Myr at high galactic latitudes to $\sim26$~Myr in the galactic plane \citep{pecaut_star_2016}, making this region an ideal place to study evolution of pre-main sequence stars and protoplanetary discs. The ages of US and UCL sub-groups in the Sco-Cen association are $\sim$11$\,\mathrm{Myr}$ \citep{pecaut_revised_2012} and $\sim$17$\,\mathrm{Myr}$ \citep{mamajek_post-t_2002} respectively. These regions are expected to have $\sim10^4$ G/K/M pre-main sequence stars based on any IMF. Many studies have looked to discover and characterise this low mass population (e.g. \citealt{zeeuw_hipparcos_1999, preibisch_history_1999, mamajek_post-t_2002, rizzuto_multidimensional_2011, rizzuto_new_2015}). Previous work on the presence of discs around these members (\citealt{luhman_disk_2012, rizzuto_wise_2012, pecaut_star_2016}) show a general increase in disc fraction with later spectral types. This is consistent with previous studies showing that protoplanetary disc lifetime is short around higher mass stars \citep{ribas_protoplanetary_2015}. However, none of these works investigated the binary fraction of these disc-bearing members. Work on binary fractions in the Sco-Cen region have primarily focused on the the higher mass B/A/F stars (\citealt{kouwenhoven_primordial_2005,kouwenhoven_primordial_2007}) finding binary fractions of $>70\%$. In order to investigate the influence of binarity on disc lifetime, we look for radial velocity variation in selected disc-bearing stars in US and UCL to determine the binary fraction. The Upper Scorpius and Upper Centaurus-Lupus regions are relatively old considering the typical protoplanetary disc lifetime, but it is at these older ages that we believe any variation in disc fractions between single and binary stars would be amplified. \Cref{sec:method} describes identification of US and UCL members, the target selection criteria, observations and how radial velocities are determined. In \Cref{sec:resultsanddiscussion} we discuss our estimated binary fractions and Bayesian analysis. In \Cref{sec:discussion} we discuss our results and caveats of our work. | \label{sec:conclusion} We determine the multiplicity of disc-bearing stars in the $\sim$11$\,\mathrm{Myr}$ Upper Scorpius (US) and $\sim$17$\,\mathrm{Myr}$ Upper Centaurus-Lupus (UCL) for periods up to 20 years. Our sample consists of 55 US members and 38 UCL members that are shown to have an IR excess, indicating the presence of a disc. Targets are observed with the Wide Field Spectrograph (WiFeS) on the Australian National University 2.3m telescope to search for radial velocity variation. Using Bayesian statistics we determine the binary fraction of our disc-bearing stars in Upper Scorpius and Upper Centaurus-Lupus to be $0.06^{0.07}_{0.02}$ and $0.13^{0.06}_{0.03}$ respectively. When compared to the expected binary fraction of $\sim0.12^{0.02}_{0.01}$ based on our observational limits and the multiplicity of field stars, these results are consistent with disc lifetime around binary stars being similar to single stars. This implies that planet formation is equally likely around binary stars as around single stars. | 18 | 8 | 1808.00620 |
1808 | 1808.09471_arXiv.txt | Birefringence in ionized, magnetized media is usually measured as Faraday rotation of linearly polarized radiation. However, pulses propagating through regions with very large Faraday rotation measures (RMs) can split into circularly polarized components with measurable differences in arrival times $\propto \nu^{-3}\,\RM$, where $\nu$ is the radio frequency. Differential refraction from gradients in DM (dispersion measure) and RM can contribute a splitting time $\propto \vert \grad_\perp \DM \vert \vert \grad_\perp\RM \vert\,\nu^{-5}$. Regardless of whether the emitted pulse is unpolarized or linearly polarized, net circular polarization will be measured when splitting is a significant fraction of the pulse width. However, the initial polarization may be inferable from the noise statistics of the bursts. Extreme multipath scattering that broadens pulses can mask splitting effects. We discuss particular cases such as the Galactic center magnetar, J1745$-$2900, and the repeating fast radio burst source, FRB 121102. Both lines of sight have $\vert \RM \vert \sim 10^5~\RMunits$ that yields millisecond splittings at frequencies well below $\sim 1 \text{ GHz}$. We also consider the splitting of nanosecond shot pulses in giant pulses from the Crab pulsar and the minimal effects of birefringence on precision pulsar timing. Finally, we explore the utility of two-dimensional coherent dedispersion with DM and RM as parameters. | Radio waves propagating through the interstellar medium (ISM) display dispersive arrival times following the cold plasma dispersion relation. Traditional pulse dedispersion techniques used in pulsar data analysis typically correct for this inverse-frequency squared delay that scales linearly with the dispersion measure (DM). As most pulsars are characterized by low Faraday rotation measures (RMs), time-of-arrival (TOA) corrections introduced by birefringence effects are usually negligible. However, recent discoveries of objects with large RMs, such as the repeating FRB 121102 \citep{Spitler+14,Michilli2018} and the Galactic center magnetar J1745$-$2900 \citep{Eatough2013}, have revealed the presence of extreme magnetoionic environments in galaxies. Magnetized plasma offers different refractive indices for the two opposite modes of circular polarization. These modes propagate with different speeds and hence, develop a propagation-induced relative time delay in such media. If this delay is long enough, a linearly polarized or unpolarized wave emitted at a transient astrophysical source may manifest to an observer as a sequence of two pulses with opposite circular polarizations. Here, we study birefringence effects arising via different mechanisms, and evaluate their significance at different observing frequencies. Sections~\ref{sec:arrival_time} and \ref{sec:obs_effects} discuss TOA corrections due to dispersive group velocities, differential refraction and plasma lensing. In \S~\ref{sec:FRB_split}, we investigate the frequency dependence of burst shapes and polarizations for FRB 121102 by taking into account pulse broadening and scattering. Observable birefringence effects for the Galactic center magnetar are examined in \S~\ref{sec:Galcenter_magnetar}. We find that arrival time perturbations due to differential dispersion and refraction of nanoshot pulses through filaments in the Crab nebula could possibly generate the nanosecond-duration pulses \citep{Hankins2003,Hankins2007} seen from the Crab pulsar. This is discussed in \S~\ref{sec:Crab}. Accurate dedispersion of pulses received from sources embedded in extreme magneto-active environments requires knowledge of the source RM in addition to its DM. In \S~\ref{sec:deDMRM}, we consider the plausibility of a two-dimensional coherent dedispersion technique that operates to simultaneously optimize two parameters, i.e., the DM and RM. In \S~\ref{sec:noise}, we highlight the importance of exploiting burst noise statistics for inferring initial polarization states of observed pulses. Finally, we summarize our findings and present the conclusions from our study in \S~\ref{sec:summary}. | \label{sec:summary} In this paper, we have considered observable manifestations of radio wave propagation through magnetized plasmas besides the standard Faraday rotation of the polarization ellipse. Arrival time variations arising from birefringence are typically too small to measure directly for most pulsars. But, in extreme conditions where RMs are very large or pulses are extremely narrow, TOA effects can become measurable. We have identified two distinct contributions to pulse TOAs, one due to the difference in group velocity of RHCP and LHCP pulses and the other due to differential refraction. For FRB 121102, the Galactic center magnetar, and the Crab pulsar, we have quantified the spectral regimes where refraction dominates group-velocity effects and vice versa. Birefringent TOAs may be relevant to fast radio bursts observed at low frequencies. The repeating FRB121102 has a high enough RM that polarization splitting at $100-300$ MHz can be larger than (intrinsic) burst widths, though multi-path scattering likely will broaden bursts significantly at those frequencies. Splitting is in principle relevant to the high-RM line of sight to the Galactic-center magnetar J1745$-$2900, though intense scattering will prevent measurement of splitting times at sub-GHz frequencies. The Crab pulsar, however, shows nanoshot pulses that occasionally display both RHCP and LHCP that may arise from splitting due to propagation through dense filaments in the Crab Nebula. To aid further study of this possibility, we have outlined a two-parameter, coherent dedispersion method that we will explore further. | 18 | 8 | 1808.09471 |
1808 | 1808.09192_arXiv.txt | We present experimental data on H$_2$ formation processes on gas-phase polycyclic aromatic hydrocarbon (PAH) cations. This process was studied by exposing coronene radical cations, confined in a radio-frequency ion trap, to gas phase H atoms. Sequential attachment of up to 23 hydrogen atoms has been observed. Exposure to atomic D instead of H allows one to distinguish attachment from competing abstraction reactions, as the latter now leave a unique fingerprint in the measured mass spectra. Modeling of the experimental results using realistic cross sections and barriers for attachment and abstraction yield a 1:2 ratio of abstraction to attachment cross sections. The strong contribution of abstraction indicates that H$_2$ formation on interstellar PAH cations is an order of magnitude more relevant than previously thought. | The formation of molecular hydrogen in the interstellar medium (ISM) is a topic of ongoing debate (see recent review \cite{wakelam2017}). As molecular hydrogen is the most abundant molecule in the universe and plays a key role in many astrophysical processes, a proper understanding of the processes that lead to its formation is of great interest. A number of potential formation processes have already been explored. While gas phase routes to H$_2$ formation have been found to be inefficient, formation on interstellar dust grains has been identified as one possible mechanism \citep{oort1946,gould1963}. Another possible route to molecular hydrogen formation is on interstellar polycyclic aromatic hydrocarbons (PAHs) \citep{bauschlicher1998,hirama2004,lepage2009,mennella2012,boschman2012,thrower2012}. A variety of objects inside and outside our galaxy exhibit spectra crowded with unidentified lines. These lines can be seen in absorption in the optical (diffuse interstellar bands: DIBS) and in emission in the infrared (aromatic infrared bands: AIB). Ubiquitous interstellar AIB are now commonly considered to be carried by PAHs. These PAHs are typically at least partially hydrogenated \citep{schutte1993,bernstein1996}. Assuming DIBS to be due to PAH-based species as well, it is most likely that part are present in protonated form or as syperhydrogenated cations, which feature the observed transitions in the visible \citep{lepage1997,snow1998,hammonds2009}. Atomic hydrogen impacting on PAH may undergo an attachment reaction and become attached to the molecule. Sequences of these addition reactions, with s additional H atoms, are of the type \begin{equation} [\mbox{PAH}+s\mbox{H}]+\mbox{H}\rightarrow [\mbox{PAH}+(s+1)\mbox{H}], \;\; \Delta M=+1 \label{eq:H-attachment} \end{equation} \noindent are referred to as hydrogenation sequences, and leave the PAH in a state of superhydrogenation and increase its mass $M$ by 1. Recent experimental and theoretical research on trapped coronene radical cations C$_{24}$H$_{12}^+$ revealed that in the gas-phase and at $T=300$ K, hydrogenation proceeds through a specific sequence of well-defined atomic sites. Reaction barriers and binding energies lead to odd - even oscillations in the observed superhydrogenation states, and magic numbers of particularly high intensity for the attachment of $n=$5, 11, and 17 extra H atoms \citep{boschman2012,cazaux2016}. Superhydrogenation dramatically alters the response of neutral and ionic gas-phase PAHs in various astrochemically relevant interaction processes. Attachment of small numbers of H atoms to coronene cations can for instance quench photoionization-induced H loss from a C$_{24}$H$_{12}^+$ precursor cation \citep{reitsma2014,reitsma2015}. For superhydrogenation of neutral pyrene molecules, an opposite effect was observed.The C-backbone is weakened and fragmentation upon ion collisions or photoionization is increased \citep{wolf2016,gatchell2015}. Attachment of a single H atom to a PAH radical cations has a dramatic influence on its IR spectrum \citep{knorke2009} and substantially decreases the HOMO-LUMO gap \citep{pathak2008}. \begin{figure*} \includegraphics[width=170mm]{sketch.jpg} \caption{Schematic representation of the reaction sequences for subsequent H (top row) or D (bottom row network) interactions with a coronene cation. H atoms are marked in blue and D atoms in red. Only a selection of possible attachment (+H, +D) and abstraction (+H-H$_2$, +D-D$_2$, +D-HD) processes are indicated. The third atom is attached to an inner edge site, as indicated by recent IR spectroscopy data \citep{schlatholter2018} and not an outer edge site as previously predicted \citep{cazaux2016}. The masses of systems that possess an odd total number of H+D atoms are given in bold letters.} \label{fig:sketch} \end{figure*} In 1998 a second type of H reaction with PAH molecules was proposed \citep{bauschlicher1998}, which was experimentally confirmed in 2012 through hydrogenation experiments on supported PAH thin films by Mennella et al. \citep{mennella2012}. In these so-called direct abstraction reactions of the Eley-Rideal type \begin{equation} [\mbox{PAH}+s\mbox{H}]+\mbox{H}\rightarrow [\mbox{PAH}+(s-1)\mbox{H}]+\mbox{H}_2, \;\; \Delta M=-1 \label{eq:H2-abstraction} \end{equation} \noindent the incoming H atom does not get bound to the PAH molecule, but rather reacts with an H atom already present in the hydrogenated PAH, to directly desorb as an H$_2$ molecule. In these experiments on solid PAH films the abstraction channel is determined to be more than an order of magnitude weaker than H attachment, with the reaction cross sections for abstraction and attachment being respectively 0.06 \AA$^2$ and 1.1 \AA$^2$ corresponding to a ratio between abstraction and attachment of $\approx 1:20$ \citep{mennella2012}. In the following, we study H abstraction reactions on gas phase coronene cations, C$_{24}$H$_{12}^+$. It should be noted that coronene is not a major species in the interstellar medium \citep{hirama2004}, but it is used as one of the prototypical PAHs in related astrolaboratory research \citep{boschman2012,cazaux2016,rauls2008,boschman2015,jochims1994,ling1998} because it is fairly large and has a compact shape, making it relatively easy to work with, and it is commercially available in large quantities. By comparing the mass spectra obtained from C$_{24}$H$_{12}^+$ exposure to $T=300$\;K H and D beams, respectively, direct evidence for the occurrence of abstraction reactions is observed. The modeling of the measured mass spectra with a time-dependent rate equation model indicates a relative cross section for abstraction that is about an order of magnitude larger than previously thought. | The main conclusion from our study is therefore the 2:1 ratio between attachment and abstraction. This implies that H abstraction from gas phase coronene cations is at least more than 7 times more efficient than H abstraction from neutral coronene thin films. This could have implications for H$_2$ production in the ISM. For instance, Andrews {\it et al.} \citep{andrews2016} have recently shown that in photodissociation regions, PAHs only contribute to H$_2$ formation via photodissociation channels and not via abstraction mechanisms. However, their calculations were based on the low abstraction cross sections from \citep{mennella2012}. An increase of the abstraction cross section by one order of magnitude could make PAHs a very important route for the formation of H$_2$ in space. | 18 | 8 | 1808.09192 |
1808 | 1808.00599_arXiv.txt | In this paper, we compile the very-high-energy and high-energy spectral indices of 43 BL Lac objects from the literature. Based on a simple math model, $\Delta\Gamma_{obs}=\alpha {\rm{z}}+\beta $, we present evidence for the origin of an observed spectral break that is denoted by the difference between the observed very-high-energy and high-energy spectral indices, $\Delta\Gamma_{obs}$. We find by linear regression analysis that $\alpha\ne 0$ and $\beta\ne 0$. These results suggest that the extragalactic background light attenuation and the intrinsic curvature dominate on the GeV-TeV $\gamma$-ray energy spectral break of BL Lac objects. We argue that the extragalactic background light attenuation is an exclusive explanation for the redshift evolution of the observed spectral break. | BL Lac objects show that continuum emission, arising from the jet emission taking place in an AGN whose jet axis is closely aligned with the observer's line of the sight, is dominated by non-thermal emission, as well as rapid and large amplitude variability. Their spectral energy distributions (SEDs) are characterized by two distinct broad bumps, implying two emission components (e.g., \citealt{phdthesis2009}). It is widely believed that the first bump is produced by electron synchrotron radiation. The second bump is probably produced by inverse Compton scattering off the relativistic electrons, either on the synchrotron photons \citep{Maraschi1992} or on some other photon populations(\citealt{Dermer1993Model}; \citealt{Sikora1994Comptonization}). These processes should result in a smooth spectrum in the $\gamma$-ray band. However, the SED shows a tendency with steepening toward higher energies (\citealt{costamante2013gamma-rays}; \citealt{dwek2013the}). This tendency indicates a spectral break in different bands. Extragalactic background light (EBL) is the diffuse and isotropic radiation fields from ultraviolet to far-infrared wavelengths \citep{Hauser2001THE}. It plays an important role in the formation and evolution of stellar objects and galaxies. Since the very-high-energy (VHE, 100 GeV$\le$ E $<$ 10TeV) $\gamma$-ray photons propagate through cosmic space, they can interact with EBL photons producing the electron-positron pairs (i.e., $\gamma_{VHE} + \gamma_{EBL}\to e^{+} + e^{-}$, e.g., \citealt{Stecker1992TeV}; \citealt{Stecker1998Absorption}; \citealt{Stecker2006Intergalactic}; \citealt{A2008Extragalactic}), which will steepen the observed spectrum (\citealt{Ackermann2012the}; \\\citealt{Sanchez2013Evidence}). If the $\gamma$-ray radiation from BL Lacs cannot be attenuated by either the secondary cascade (\citealt{Essey2012On}; \citealt{Zheng2013Evidence}; \citealt{Zheng2016Bethe}) or the axion-like particle conversion(\citealt{Horns2012Indications}), we should expect a distinctive EBL attenuation characteristic (e.g., \citealt{Kneiske2004Implications}; \citealt{Stecker2006Intergalactic}; \\ \citealt{Imran2008Detecting}; \citealt{A2008Extragalactic}; \citealt{Tavecchio2009Intrinsic}) at high energies in gamma-ray sources. Taking the EBL attenuation into account during the $\gamma$-ray propagation, we expect a difference in the observed high energy (HE, 100 MeV $\le$ E $<$ 300GeV) and VHE spectrum of BL Lacs. In order to determine the spectral break, \\ \cite{Dom2015Spectral} selected 128 extragalactic sources from the second catalog of Fermi-LAT sources (2FHL), which is observed in HE bands extending in energy from 50 GeV to 2 TeV. Nevertheless, the study of the spectral break mechanism with the 2FHL catalog cannot explain the spectrum detected with energies higher than 2 TeV. Owing to the MAGIC, HESS, and VERITAS broadening of the 2FHL spectra (\citealt{Funk2012The}), we can study the relativistic particle acceleration (\citealt{Holder2012TeV}) and the spectral energy break. Additionally, the sources of the 2FHL catalog used by \cite{Dom2015Spectral} included the flat spectrum radio quasars (FSRQs), whose GeV-TeV spectrum would be contributed by the photon populations of the broad line region (BLR) or of the accretion disk (\citealt{Poutanen2010GeV}; \citealt{Ackermann2012the}). This results from without a significant EBL attenuation in the SED of FSRQs. Therefore, we focus on the spectra of BL Lacs. In order to clarify the spectral break mechanism, in this paper we focus on the analysis of the GeV-TeV energy spectral index of the BL Lac objects. Our goal is to determine whether the EBL attenuation is the exclusive origin of the energy spectral break of the BL Lac objects. Since the observed spectrum is attenuated by EBL, we should obtain the absorption-corrected spectrum. In Section 2, we describe the sample; in Section 3, we compare the observed spectrum and intrinsic spectrum; in Section 4, we analyze the origin of spectral break; and in Section 5, we provide our conclusions and a discussion. | In this paper, we compile the GeV-TeV energy spectral indices of 43 BL Lac objects to analyze the scatter of the $\gamma$-ray observed spectra and intrinsic spectra. We found that the mean observed index is significantly higher than the mean intrinsic index (i.e., $\Gamma _{0, obs}= 3.10$ for observed spectrum and $\Gamma _{0, int}= 2.05 $ for the intrinsic spectrum), implying that the intrinsic spectrum is a hard spectrum. It was also proved that the observed spectrum is attenuated by EBL, which would reduce the mean spectral index. We focus on a relationship, $\Delta\Gamma_{obs}=\alpha\rm{z} + \beta $, between the observed spectral break and redshift. Three cases exist for this mathematical model, as follows: \begin{enumerate}% \item When $\alpha \ne 0$ and $\beta=0$, the equation can be replaced by the relationship between EBL attenuation and redshift, $\Delta {\Gamma _{EBL}(E,z)} = \Gamma _{VHE,obs}- \Gamma _{VHE, int}$, which is not in accord with the observed spectral break. This shows that the EBL attenuation will change with redshift, especially for the high-redshift spectrum. \item When $\alpha \ne 0$ and $\beta\ne 0$, the origin of the observed spectral break can be affirmed by linear regression and the simulation of theoretical models (i.e., SSC and EBL models). The observed spectral break is determined by the EBL attenuation and the intrinsic curvature (some blazar physics processes) due to the facts that $\alpha \ne 0$ and $\beta=0$. \item When $\alpha=0 $ and $\beta\ne 0$, the equation relates to the difference between the the intrinsic spectral index at VHEs and spectral index at HEs (intrinsic spectral break), ${\Delta\Gamma _{int}} = {\Gamma _{VHE,int }} - {\Gamma _{HE,obs}}$. In this case, we only study the evolution of the intrinsic spectral break with redshift. The EBL attenuation is removed due to the property of the intrinsic spectral break. From the distribution of intrinsic break and the simulation, the intrinsic curvatures play a crucial role in the intrinsic spectral break, and they cannot evolve with redshift. \end{enumerate} Owing to the statistical results, $\alpha = 3.60 \pm 0.72 \ne 0$ and $\beta = 0.83\pm 0.13 \ne 0$, it is suggested that the observed spectral break is dominated by EBL attenuation and the intrinsic curvature (some blazar physics processes). This confirms that EBL attenuation is an origin for the observed spectral break that has evolved with redshift. Although we have verified that the EBL attenuation is an origin for the observed spectral break of the BL Lac objects, according to the statistical results $\beta \ne 0$ the study of $\beta$ will become indispensable. Some different methods can be used to simulate the spectrum, \cite{Tramacere2011Stochastic} employed the log-parabolic-law (log-parabolic-law electron-energy distribution) SSC model to obtain the intrinsic curvature. However, our research employed the broken-power law SSC model to obtain it. Different methods will lead to different electron spectra, which could affect the intrinsic curvature. Additionally, the intrinsic curvature also can be interpreted by the Klein-Nishina suppression (emission effects) or a turnover in the distribution of the underlying emitting particles (acceleration effects, e.g., \citealt{Sanchez2013Evidence}.) The redshift evolution of the observed spectral break can be explained solely by EBL attenuation (without the secondary cascades or axion-like particle conversion), and there is also no evidence of evolution with redshift of the physics that drives the photon emission (\citealt{Dom2015Spectral}). \acknowledgment We thank the anonymous referee for valuable comments and suggestions. This work is partially supported by the National Natural Science Foundation of China under grants 11463007, 11673060 and the Natural Science Foundation of Yunnan Province grants 2017FD072. | 18 | 8 | 1808.00599 |
1808 | 1808.10495_arXiv.txt | {This is a proceeding of a rapporteur talk given on ground-based gamma-ray astronomy at the $35^\mathrm{th}$ International Cosmic-Ray Conference (ICRC) held in 2017 in Busan, Republic of Korea. A total of $\sim$300 contributions were presented during the ICRC over 17 gamma-ray sessions. Here, I summarize the contributions mainly focusing on the source observations performed by ground-based gamma-ray instruments and the connection between gamma rays and cosmic rays. Any such summary must necessarily be incomplete. However, I have attempted to provide a glance into recent progress that has been made in using ground-based gamma-ray observations to understand the nature of high energy particles in our Universe.} \FullConference{35th International Cosmic Ray Conference --- ICRC2017\\ 10--20 July, 2017\\ Bexco, Busan, Korea} \begin{document} | Very high energy (VHE, E$>$100 GeV) gamma-ray observations reveal the extreme environments in our Universe where particles can be accelerated above TeV energies and allow us to study the dynamics of these high energy particles. We can learn how particles are accelerated in different classes of sources in various evolutionary stages and environments. This also opens a window to understand the origin and acceleration of high energy particles observed locally on Earth as cosmic rays. Since the first detection of VHE gamma-ray emission from the Crab Nebula in 1989 by the Whipple Observatory~\cite{1989Weekes}, VHE gamma-ray astronomy has been rapidly progressing. By the summer of 2017, the number of VHE gamma-ray sources has reached 198 with at least 8 broad source classes. The science topics studied by VHE gamma-ray observatories cover a wide range from gamma-ray astronomy to indirect dark matter detection, measurements of the local cosmic-ray fluxes and anisotropy, and particle physics. In this $35^{\textnormal{th}}$ International Cosmic-Ray Conference (ICRC), there were in total 17 gamma-ray sessions, which includes 289 contributions. Among these, about 225 contributions are closely related specifically to VHE gamma-ray observations. The considerable overlap with contributions from space-based gamma-ray observations highlights the scientific value of the multiwavelength approach that has become standard for many studies. Also, there are about 20 contributions from VHE gamma-ray observations to the other sessions, including the indirect cosmic-ray session and the dark matter session, and these contributions emphasize the multidisciplinary nature of the field. For this review, I will focus on the updates from the observations of gamma-ray sources in particular. Since this is a review of the VHE gamma-ray studies presented during the cosmic-ray conference, it is worthwhile to point out the connection between cosmic rays and VHE gamma rays. VHE gamma rays are produced by the interaction of high-energy particles that will eventually become cosmic rays. Thus, understanding VHE gamma-ray sources is connected to understanding the sources of cosmic rays. Gamma-ray observations can provide what cosmic-ray observations cannot. Namely, because gamma rays are neutral particles their directional information is preserved and they can be used to find the source. This allows us to study sources that accelerate particles to high energies. We can study different source classes and the evolution of sources as accelerators. However, the gamma-ray studies have their own difficulties because the observed emission may come from different species of particles. Usually understanding the nature of the accelerator requires detailed models of the source region and evolution to disentangle leptonic and hadronic contributions. Also, unlike direct cosmic-ray measurements, hadronic gamma-ray emission is mostly dominated by protons and is not sensitive to different compositions. Compared to this, cosmic-ray observations generally cannot be used to detect the individual sources, and what we measure at Earth is likely contributed by multiple sources over millions of years. However, cosmic-ray measurements can provide detailed composition information, which allows us to study the source site composition, potential acceleration differences between elements, and propagation of the particles in our Galaxy by looking at the ratio between primary and secondary particles. We know that there should be high energy accelerators in our Universe to explain the remarkably smooth distribution of cosmic rays from several GeV up to $10^{21}$ eV. Answering what is the origin and acceleration mechanism for these high-energy particles is one of the most important tasks of VHE gamma-ray astronomy. Thus, the best way to understand the dynamics of high energy particles in our Universe is to combine the knowledge gained from cosmic-ray and gamma-ray observations. | After more than ten years of observations with the current generation's IACTs, we are still finding large numbers of new VHE gamma-ray sources in the sky. Since the last ICRC, the number of VHE gamma-ray sources has increased by 36, and the total number of VHE gamma-ray sources counts 198. About half of the new sources were detected by HAWC as they accumulated more than two years of data with a full detector configuration. HAWC's new sources present different characteristics than sources that were detected by IACTs. Because HAWC has a large FoV and covers 2/3 of the sky, HAWC can provide a population of very extended sources including local sources that are closer than 1 kpc to us. These are particularly interesting in connection to the leptonic cosmic-ray fluxes we measure at Earth. For IACTs, we achieve one more milestone for the technique by a firm detection of the extension of the Crab Nebula, which was treated as a point-like source up to now. The first binary from outside of our Galaxy has been detected, and new flaring episodes observed in wide electromagnetic bands allow us detect new VHE gamma-ray sources and study the extreme environments around AGNs. Detailed studies of known sources reveal the particle acceleration and interaction in different environments. Currently we know more about the VHE accelerators than ever before with the best sky coverage in history. There will be another leap in our understanding of the VHE gamma-ray universe once CTA comes online. We will learn about the in-depth details of sources we are studying now. More importantly, CTA will reveal more new sources, which likely will include at least one new class of accelerators we have not yet detected. The construction of CTA will start in 2018, so the future is not very far away. | 18 | 8 | 1808.10495 |
1808 | 1808.03659_arXiv.txt | {}{ We explore the wealth of high quality photometric data provided by data release 2 of the \Gaia mission for long period variables (LPVs) in the Large Magellanic Cloud. Our goal is to identify stars of various types and masses along the Asymptotic Giant Branch. }{ For this endeavour, we developed a new multi-band approach combining Wesenheit functions $W_{RP,BP-RP}$ and $W_{K_s,J-K_s}$ in the \Gaia BP, RP and 2MASS J, K$_\mathrm{s}$ spectral ranges, respectively, and use a new diagram ($W_{RP,BP-RP}-W_{K_s,J-K_s}$) versus $K_s$ to distinguish between different kinds of stars in our sample of LPVs. We used stellar population synthesis models to validate our approach. }{ We demonstrate the ability of the new diagram to discriminate between O-rich and C-rich objects, and to identify low-mass, intermediate-mass and massive O-rich red giants, as well as extreme C-rich stars. Stellar evolution and population synthesis models guide the interpretation of the results, highlighting the diagnostic power of the new tool to discriminate between stellar initial masses, chemical properties and evolutionary stages. } {} | Asymptotic Giant Branch (AGB) stars represent a decisive part within the evolution of low and intermediate mass stars. These high luminosity, cool objects contribute significantly to the metal content of their host galaxy by various nucleosynthesis processes in their interior, deep mixing events, and strong mass loss. The first two points modify the surface composition -- even changing the chemistry from O-rich to C-rich -- which affects the mass loss. All these contributions depend on the stellar parameters, in particular the stars' masses. To study the role of AGB stars for the evolution of galaxies, it is important to reliably identify and characterize (mass, chemistry) them within a stellar population. AGB stars show a typical variability pattern with long periods of several ten to several hundred days and moderate to large light amplitudes in particular in the visual range. These variables form the class of the long period variables (LPVs) which also includes other cool, evolved stars like red supergiants or stars at the tip of the Red Giant Branch (RGB). This kind of variability is easily detectable even in other galaxies. Thus, LPVs define an excellent sample of AGB stars. Major advances in the understanding of AGB stars were often related to large surveys in various photometric bands. Monitoring surveys like MACHO \citep{1999IAUS..191..151W} and OGLE \citep[e.g.][]{2009AcA....59..239S} and others delivered excellent samples of AGB stars for various galaxies and the Galactic Bulge. With the advent of the \Gaia mission \citep{2016A&A...595A...1G}, a new era is starting. The all-sky nature of the \Gaia survey, the high spatial resolution of $\sim$0.4~arcsec, the provision of parallaxes, the availability of three photometric bands providing wide \gmag, blue \gbp and red \grp magnitudes for more than 1.5~billion stars, from very bright (few magnitudes in \gmag) to as faint as $\gmag \simeq 20.7$~mag, and the provision of time series in all three bands for all the stars, not mentioning the expected provision of radial velocity and astrophysical parameters for stars brighter than $\gmag \simeq 16$~mag, are some of the unique advantages of the survey. In \Gaia data release 2 \citep[DR2,][]{Brown_etal18}, the first \Gaia catalog of LPVs with \gmag-band variability amplitudes larger than 0.2~mag (measured by the 5--95\% quantile range of the \gmag time series) has been published \citep{2018LPV_DR2}. It contains 151\,761 candidates. As a first step in the analysis of this unique dataset, we study in this Letter the \Gaia DR2 LPVs in the Large Magellanic Cloud (\object{LMC}). The above mentioned task of identification and distinction of various groups of stars among the LPVs is challenging because the upper giant branch represents a mixture of objects of various masses and evolutionary stages. In this letter we present a new method for an efficient identification of stars of different mass and chemistry based on \Gaia and 2MASS photometry. | \label{Sect:conclusions} We demonstrated that the new tool for the investigation of stellar populations of AGB-stars presented in this paper offers an excellent possibility to study the large \Gaia dataset of evolved stars. It allows to distinguish stars of different mass, chemistry and evolutionary stage very efficiently. The high diagnostic capability of the new diagram is corroborated with the help of state-of-the-art stellar evolution models. We plan a more extensive exploration of this tool in the future. As a next step, we are going to use this tool for a similar analysis of other stellar populations and for an investigation on the pulsational behaviour of LPVs. | 18 | 8 | 1808.03659 |
1808 | 1808.03145_arXiv.txt | We study the narrow emission line properties and stellar populations of a sample of 1385 AGN selected to have strong excess emission at mid-infrared wavelengths based on comparing Wide-field Infrared Survey Explorer W1-W2 band colours with optical stellar absorption line indicators. Our goal is to understand whether the physical conditions in the interstellar medium of these objects differ from those of local AGN selected by their optical emission line ratios. To enable this comparison, we construct a {\em control sample} of 50,000 optically-selected AGN with the same redshifts that do not have strong mid-IR excess emission. The mid-IR excess and control samples differ strongly in [OIII] line luminosity, ionized gas excitation mechanism, ionization state and electron density. We show that the radio-detected, mid-IR excess AGN constitute the most luminous and highly ionized AGN in the local Universe and they contribute primarily to the growth of black holes in the most massive galaxies. At least half of this black hole growth is occurring in galaxies with recent starbursts. The morphologies of these systems indicate that the starbursts have probably been triggered by galaxy-galaxy mergers and interactions. The most luminous AGN in our mid-IR excess sample have properties that are similar to the Type II quasars studied at higher redshifts. In contrast, the control sample constitute a class of lower ionization, less luminous AGN in more quiescent galaxies that contribute primarily to the growth of low mass black holes. | 18 | 8 | 1808.03145 |
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1808 | 1808.05625_arXiv.txt | Blazars are known for their energetic multiwavelength flares from radio wavelengths to high-energy $\gamma$-rays. In this work, we study radio, optical, and $\gamma$-ray light curves of 145 bright blazars spanning up to 8~yr, to probe the flaring activity and interband correlations. Of these, 105 show $>1\sigma$ correlations between one or more wavebands, 26 of which have a $>3\sigma$ correlation in at least one wavelength pair, as measured by the discrete correlation function. The most common and strongest correlations are found between the optical and $\gamma$-ray bands, with fluctuations simultaneous within our $\sim 30$~d resolution. The radio response is usually substantially delayed with respect to the other wavelengths with median time lags of $\sim 100$--160~d. A systematic flare identification via Bayesian block analysis provides us with a first uniform sample of flares in the three bands, allowing us to characterise the relative rates of multiband and ``orphan'' flares. Multiband flares tend to have higher amplitudes than ``orphan'' flares. | \label{sec:intro} Blazars are luminous and highly variable across the entire electromagnetic spectrum. Their emission, dominated by relativistically boosted jets, shows a spectral energy distribution (SED) with two distinct humps. The low-energy peak (radio to ultraviolet, and in some cases X-rays) is thought to be produced by synchrotron radiation, and the high-energy hump (X-rays to high-energy $\gamma$-rays) is likely Compton emission. However, the source of the seed photons that produce the observed high-energy emission is not fully understood. Some models also attribute $\gamma$-ray flux to proton-mediated emission rather than inverse-Compton scattering. Traditionally, blazars are classified by their optical spectral line properties as BL Lacertae objects (BL Lacs; broad lines typically lost against a bright continuum) and flat-spectrum radio quasars (FSRQs; strong, broad lines). Blazars exhibit violent flaring, in which the Earth-directed flux can increase by orders of magnitude over short timescales. Flares appear in multiple wavebands either simultaneously or with a time delay ranging from days to months. These delays can provide clues to the emission process and the relative location of the emitting region at different wavelengths. Flaring events are often accompanied by other phenomena such as the ejection of jet components mapped at very long baseline interferometry (VLBI) scales (e.g., \citealp{Marscher2008}) and rotations of the optical polarization plane (e.g., \citealp{Blinov2017}). The origin of the $\gamma$-rays in particular is poorly understood. The most likely mechanism appears to be inverse-Compton (IC) scattering of lower-energy photons by the relativistic jet electrons. However, the nature of the incident photon field and the location of the upscattering electrons are poorly constrained. If the incident photon field is external to the jet (accretion disk, broad-line region, etc.), the IC scattering is referred to as external Compton (EC, e.g., \citealp{Dermer1992}), while if the incident photons are from the jet itself it is called synchrotron self-Compton (SSC, e.g., \citealp{Abdo2010-II}). One can probe these uncertainties by studying the correlated variability between different frequency bands. Several monitoring programs have followed $\gamma$-ray-loud blazars at various wavelengths since the launch of the {\it Fermi} gamma-ray space telescope, with the primary goal of constraining the mechanism and location of the $\gamma$-ray emission through time-series analysis. This is usually accomplished through the cross-correlation of radio and $\gamma$-ray light curves (e.g., \citealp{Fuhrmann2014,Max-Moerbeck2014,Ramakrishnan2015}), or optical and $\gamma$-ray light curves (e.g., \citealp{Pati2013,Cohen2014,Hovatta2014-II,Ramakrishnan2016}). An alternative approach is to organise intensive multiwavelength campaigns of individual sources (e.g., \citealp{Rani2013,Karamanavis2016}). In this work, we revisit the optical--$\gamma$-ray \citep{Cohen2014} and $\gamma$-ray--radio correlations \citep{Max-Moerbeck2014} as well as the underexplored optical--radio correlation using long-term monitoring of a large sample of $\gamma$-ray-bright blazars. Our goal is to apply statistical analysis to reveal correlation trends and compare them across blazar classes. We present the sample and the results of basic cross-correlation analysis in Section \ref{sec:samp-cros}. In section \ref{cor_flares} we use the correlations to associate individual flares seen in different wavelengths and examine the statistical properties of the correlated and non-correlated ``orphan'' events. Section \ref{sec:disc-conc} discusses our findings and summarises our conclusions. | \label{sec:disc-conc} In this work we have investigated the temporal correlations between the optical, radio, and $\gamma$-ray emission of a large sample of blazars. Out of the 145 sources, 105 revealed at least one $>1\sigma$ significant correlation coefficient between wavebands. Out of these, 38 showed a correlation in only one pair of wavelengths, 35 in two pairs, and 32 in all three pairs. Based on the significance of DCF peaks, we estimated that only 9.8\% of the optical--radio, 6.4\% of the optical--$\gamma$-ray, and 11.2\% of the $\gamma$-ray--radio cross correlations will be false positives. For the sources with a $>2\sigma$ significant correlation the false-positive rate is roughly 2\% in any wavelength pair. It should be noted that blazars exhibit a wide variety of flaring patterns and behaviours which could be the result of different mechanisms operating even in the same source (e.g., \citealp{Chen2012,Liodakis2017}). It is possible for a source to show both correlated and uncorrelated events with other wavelengths, which can impact the significance of DCF. This could explain why many of our sources show only a 1$\sigma$ or 2$\sigma$ significant DCF peaks. Our results are qualitatively consistent with those of previous studies, but our longer time base has allowed us to identify a larger number of, and more significant, correlations. Continued monitoring would doubtless improve the situation further. We find that the optical--$\gamma$-ray time lags are generally small, while both lead the radio by $\sim30$--150~d. However, there are three sources (J0433+2905, J1849+6705, and J2236+2828) in which all three wavelengths are correlated with lags $<30$~d. In fact, unusually, J0433+2905 shows the optical emission coming 20--30~d earlier than both radio and $\gamma$-rays (the other two are more conventional). This suggests that the emission regions are located in close proximity. The fact that there is a strong connection between optical and $\gamma$-ray variations favours leptonic (i.e., inverse-Compton scattering of the optical photons as the production mechanism for the $\gamma$-ray emission) over hadronic processes. On the other hand, the fact that the radio usually lags all other wavelengths suggests that it is typically downstream from both the optical and $\gamma$-ray emission regions. Since the radio variations are connected to the ejection of new components from the radio core seen in VLBI maps \citep{Savolainen2002}, this would place the $\gamma$-ray emission regions between the supermassive black hole and the radio core. Combined with the generally longer timescale variations seen at radio wavelengths (e.g., \citealp{Hovatta2007}), our results favour emission scenarios of an expanding disturbance propagating in the jet and becoming optically thin at higher frequencies before becoming transparent at radio wavelengths (e.g., \citealp{Marscher1985,Max-Moerbeck2014}). Using Bayesian block decomposition \citep{Scargle2013}, we have studied a large number of individual flares and their multiwavelength properties. We have relied on the DCF analysis to align these light curves, allowing cross-band identification of individual flaring events. This has also for the first time allowed a robust determination of orphan (to our sensitivity) flares in all three wavebands. Overall, sources showing a lower flaring rate tend to have less significant interband correlations, but this may be a simple selection effect. Comparing BL Lacs and FSRQs, the former show a higher flaring rate in the optical (3.6--3.7 vs. 2.3--2.5 flares yr$^{-1}$ source$^{-1}$ on average), while showing lower radio (0.9--1 vs. 1.1--1.3 flares yr$^{-1}$ source$^{-1}$) and $\gamma$-ray (0.3--0.5 vs. 0.6--0.8 flares yr$^{-1}$ source$^{-1}$) activity. Two main effects limit the present analysis. First, seasonal gaps in the optical light curves limit the number of flare identifications and make long-term trends difficult to follow. Although observational gaps in the optical are unavoidable because of the Sun, in many cases these are enlarged (or even induced) by weather and/or technical related constraints. Multisite monitoring could help minimise the extent of said gaps. Second, we have used only publicly accessible $\gamma$-ray light curves, with coarse 30~d sampling. As shown in Figures \ref{plt:30-7_comp} and \ref{plt:30-7_comp2} (see also the discussion in Section \ref{cor_flares}), our overall statistical results and conclusions should not be strongly affected by the choice of bin size; however, finer sampling such as the adaptive binning used by \citet{Cohen2014} can find more flares and probe shorter timescales . For this reason, our results on the flare rates and flare associations with respect to the $\gamma$-ray light curves should be treated as limits. We are currently pursuing a more detailed analysis of the $\gamma$-ray light curves that will allow us to probe time delays and flare associations on shorter timescales between optical and $\gamma$-rays. In summary, our results show the following. \begin{itemize} \item {The radio emission generally lags the optical/$\gamma$-rays, suggesting that the higher energy radiation arises inward of the radio cores of the jets. } \item {The optical emission is closely connected to the $\gamma$-ray emission, with roughly half the sources having time lags consistent with zero.} \item {A few sources seem to have all three bands colocated (e.g., J0433+2905, J2236+2828).} \item {Low radio and $\gamma$-ray activity likely explains the lack of significant correlation for many of our sources.} \item We found no significant difference between associated and orphan $\gamma$-ray flares in sources with a significant $\gamma$-ray--radio correlation. In all other cases (and wavelengths), flares have higher flux when associated with the other band than when they are orphans. \end{itemize} | 18 | 8 | 1808.05625 |
1808 | 1808.02143_arXiv.txt | We have used the Australia Telescope Compact Array (ATCA) to undertake the first high angular resolution observations of 37.7-GHz ($7_{-2} \rightarrow 8_{-1}E$) methanol masers towards a sample of eleven high-mass star formation regions which host strong 6.7-GHz methanol masers. The 37.7-GHz methanol sites are coincident to within the astrometric uncertainty (0.4 arcseconds) with the 6.7-GHz methanol masers associated with the same star formation region. However, spatial and spectral comparison of the 6.7- and 37.7-GHz maser emission within individual sources shows that the 37.7-GHz masers are less often, or to a lesser degree co-spatial than are the 12.2-GHz and 6.7-GHz masers. We also made sensitive, high angular resolution observations of the 38.3- and 38.5-GHz class II methanol transitions ($6_{2} \rightarrow 5_{3}\mbox{A}^-$ and $6_{2} \rightarrow 5_{3}\mbox{A}^+$, respectively) and the 36.2-GHz ($4_{-1} \rightarrow 3_{0}E$) class I methanol transition towards the same sample of eleven sources. The 37.7-, 38.3- and 38.5-GHz methanol masers are unresolved in the current observations, which implies a lower limit on the brightness temperature of the strongest masers of more than $10^6$~K. We detected the 38.3-GHz methanol transition towards 7 sources, 5 of which are new detections and detected the 38.5-GHz transition towards 6 sources, 4 of which are new detections. We detected 36.2-GHz class~I methanol masers towards all eleven sources, 6 of these are new detections for this transition, of which 4 sources do not have previously reported class~I methanol masers from any transition. | The details of the physical processes that govern the earliest stages of high-mass star formation remain an area of intense research and much debate \citep[e.g.][]{Tan+14}. Class~II methanol masers are widely recognised as one of the most reliable tracers of the earliest stages of high-mass star formation \citep{Ellingsen06} and have the advantage that they are exclusively associated with this phenomenon \citep{Breen+13b}. The rich microwave spectrum of the methanol molecule contributes to the large number of different transitions which show maser emission. Methanol masers are empirically divided into two classes, known as class~I and class~II \citep{Menten91b}. The population inversion for class~I methanol masers is due to collisions\citep{Cragg+92,Sobolev+07,Leurini+16}, with the 44- and 36-GHz transitions the most commonly observed and generally strongest transitions \citep[e.g.][]{Voronkov+14}. The class~I methanol masers are associated with molecular outflows \citep[e.g.][]{Plambeck+90,Voronkov+06} and other phenomenon which produce low-velocity shocks in molecular gas, such as expansion of \ionhy regions \citep[e.g.][]{Voronkov+10a} and cloud-cloud collisions \citep[e.g.][]{Sobolev92,Sjouwerman+10b}. Searches for class~I methanol masers were initially limited by insufficient target sites known to host outflows. However, searches towards class~II methanol maser sites \citep[e.g.][]{Slysh+94,Valtts+99,Ellingsen05} and more recently sources showing excess 4.5-\micron\/ emission in {\em Spitzer} images \citep[known as extended green objects - EGOs][]{Cyganowski+08}, have been very successful \citep[e.g.][]{Chen+11,Chen+13a} and there are now more than 500 class~I methanol maser sites known throughout the Milky Way \citep{Yang+17}. The limited sensitive and unbiased searches for class~I methanol masers undertaken to date \citep[e.g.][]{Jordan+15,Jordan+17,Yusef-Zadeh+13}, suggest that the ATCA Legacy survey of the Galactic Plane at 7mm which commenced in 2017 will more than double the number of known sources over the next few years. Infrared radiation creates the population inversion in class~II methanol masers \cite[e.g.][]{Sobolev+97a,Cragg+05}, with the 6.7- and 12.2-GHz transitions the most commonly observed and generally strongest. The majority of studies to date have focused on the class~II methanol masers, as the strong transitions are at lower frequencies, (hence more readily observed) and early targeted searches towards other maser transitions and infrared sources were very successful \citep[e.g.][]{Menten91a,MacLeod+92,Caswell+95a,Walsh+98}. There are now more than 1000 class~II methanol maser sites known in the Milky Way, with a number of large-scale, complete searches having been undertaken \citep[e.g.][]{Ellingsen+96b,Pandian+07}, of which the Methanol Multibeam (MMB) is the most comprehensive \citep{Green+10,Caswell+10,Caswell+11,Green+12a,Breen+15}. A total of 18 different class~II methanol maser transitions or transition series have been observed toward high-mass star formation regions \citep[see][and references therein]{Ellingsen+12b}. The majority of these transitions have only been detected in a small number of sources and have significantly weaker emission than is observed from the 6.7-GHz transition. Targeted searches have been made at 12.2-GHz towards all 6.7-GHz methanol masers detected in the MMB \citep{Breen+12a,Breen+12b,Breen+14a,Breen+16}, with detections towards 45.3 per cent of the sources. However, the only other class~II methanol maser transitions for which there have been sensitive searches towards moderately large samples of sources are the 19.9-GHz transition \citep[][22 sources]{Ellingsen+04}, the 23.1-GHz transition \citep[][50 sources]{Cragg+04}, the 37.7-, 38.3- and 38.5-GHz transitions \citep[][106 sources]{Haschick+89,Ellingsen+11a,Ellingsen+13a}, the 85.5-, 86.6- and 86.9-GHz transitions \citep[][22 sources]{Ellingsen+03} and the 107- and 156.6-GHz transitions \citep[][80 sources]{Caswell+00}. Where multiple class~II methanol maser transitions are observed to be co-spatial there is the potential to use them as probes of the physical conditions within the star formation regions \citep[e.g.][]{Cragg+01,Sutton+01}. However, for many of the maser transitions observations have only been made with single-dish telescopes with angular resolutions of order arcminutes. Interferometric observations are required to determine whether the different transitions are truly co-spatial, and hence whether the underlying assumption of the multi-transition maser modelling is valid \citep[e.g.][]{Krishnan+13}. Here we present the first interferometric observations of the 37.7-, 38.3- and 38.5-GHz class~II methanol maser transitions. All of the sources in our sample have interferometric observations at 6.7-GHz available in the published literature, and many have also been observed at arcsecond angular resolution in the 12.2- and 19.9-GHz transitions, enabling us to test the degree to which the emission is co-spatial over these transitions. | We have made the first high resolution observations of the 37.7-, 38.3- and 38.5-GHz class~II methanol maser transitions towards a sample of eleven sources. We find that to within the absolute astrometric accuracy (around 0.4 arcseconds), these class~II transitions arise from the same location as the strong 6.7-GHz class~II methanol masers. The spectral and imaging data both suggest that the 37.7-GHz masers are less co-spatial with the 6.7-GHz maser spots than has previously been observed for the 12.2-GHz class~II transition. The peak velocity of the 37.7-GHz masers is found to be preferentially blueshifted with respect to the systemic velocity, suggesting that these maser arise in molecular gas which is outflowing from the region in the general direction of the observer. The detection of class~I methanol masers towards what are hypothesised to be star formation regions towards the end of the class~II methanol maser phase, combined with the results of recent untargetted searches for class~I methanol masers suggest that the duration of the class~I methanol maser phase exceeds that of class~II methanol masers. | 18 | 8 | 1808.02143 |
1808 | 1808.00977_arXiv.txt | {% } {The potential similarity of the powering mechanisms of relativistic SNe and GRBs allowed us to make a prediction that relativistic SNe are born in environments similar to those of GRBs, that is, ones which are rich in atomic gas. Here we embark on testing this hypothesis by analysing the properties of the host galaxy NGC 3278 of the relativistic {\sn}. This is the first time the atomic gas properties of a relativistic SN host are provided and the first time resolved 21\,cm-hydrogen-line ({\hi}) information is provided for a host of an SN of any type in the context of the SN position. } {We obtained radio observations with ATCA covering the {\hi} line, and optical integral field unit spectroscopy observations with MUSE. Moreover, we analysed archival carbon monoxide (CO) and multi-wavelength data for this galaxy. } {The atomic gas distribution of {\ngc} is not centred on the optical galaxy centre, but instead around a third of atomic gas resides in the region close to the SN position. This galaxy has a few times lower atomic and molecular gas masses than predicted from its star formation rate (SFR). Its specific star formation rate ($\mbox{sSFR}\equiv\mbox{SFR}/\mstar$) is approximately two to three times higher than the main-sequence value. {\sn} exploded close to the region with the highest SFR density and the lowest age ($\sim5.5$\,Myr). Assuming this timescale was the lifetime of the progenitor star, its initial mass would have been close to $\sim36\,\msun$. } {As for GRB hosts, the gas properties of {\ngc} are consistent with a recent inflow of gas from the intergalactic medium, which explains the concentration of atomic gas close to the SN position and the enhanced SFR. Super-solar metallicity at the position of the SN (unlike for most GRBs) may mean that relativistic explosions signal a recent inflow of gas (and subsequent star formation), and their type (GRBs or SNe) is determined either {\it i)} by the metallicity of the inflowing gas, so that metal-poor gas results in a GRB explosion and metal-rich gas results in a relativistic SN explosion without an accompanying GRB, or {\it ii)} by the efficiency of gas mixing, or {\it iii)} by the type of the galaxy. } % | \label{sec:intro} \begin{figure*} \begin{center} \includegraphics[width=\textwidth,clip]{sed.ps} \end{center} \caption{Spectral energy distribution of the {\ngc} (red points). The {\sc Grasil} and {\sc Magphys} models are shown as black solid and blue dotted lines, respectively. For comparison we show the models for the hosts of GRB 980425 \citep[][]{michalowski14} and 111005A \citep{michalowski18grb}, scaled to approximately match near-IR fluxes. } \label{fig:sed} \end{figure*} The gas inflow from the intergalactic medium is predicted to be an important process providing the fuel for star formation (see e.g.~\citealt{sancisi08}, \citealt{spring17}). It has been studied mostly from indirect diagnostics because compiling a sample of galaxies for which this process can be observed directly is difficult. Based on the analysis of gas properties in long gamma-ray burst (GRB) host galaxies, we have recently proposed that the progenitors of GRBs are preferentially born when a galaxy accretes fresh gas from the intergalactic medium \citep{michalowski15hi,michalowski16}. This is based on a high abundance of atomic gas in GRB hosts and its concentration close to the GRB position \citep{arabsalmani15b,michalowski15hi}. This may also imply that a fraction of star formation is fuelled directly by atomic, not molecular, gas. The majority of star formation in the Universe is fuelled by molecular gas, as shown by many observations \citep[e.g.][]{carilli13,rafelski16}. However, {\hi}-fuelled star formation has been shown to be theoretically possible \citep{glover12,krumholz12,hu16} and it was supported by the existence of {\hi}-dominated, star-forming regions in other galaxies \citep{bigiel08,bigiel10,fumagalli08,elmegreen16}. If the connection between GRBs and recent inflow is confirmed, this will allow the use of GRB hosts to study gas accretion and/or {\hi}-fuelled star formation. On the other hand, relativistic supernovae (SNe) without detected $\gamma$-rays are thought to be powered by similar engines to those of GRBs, but with the jet failing to break out from the exploding star \citep{paragi10,lazzati12,margutti14,chakraborti15,milisavljevic15}. The potential similarity of this powering mechanism to that of GRBs allowed us to make a prediction that relativistic SNe are born in environments similar to those of GRBs, that is, those rich in atomic gas. Here we embark on testing this hypothesis by analysing the properties of the host of the relativistic {\sn}. \object{SN 2009bb} was discovered by the galaxy-targeted survey, the CHilean Automatic Supernova sEarch (CHASE; \citealt{chase}) on 21 March 2009 \citep{pignata09cbet} at the position of 10:31:33.8762, $-$39:57:30.022 \citep{bietenholz10} and was a broad-line type-Ic supernova \citep{stritzinger09cbet}. Radio and optical behaviour, and the relativistic ejecta velocity of {\sn} were very similar to those of low-$z$ GRBs, especially GRB\,980425 \citep[SN\,1998bw;][]{soderberg10,bietenholz10,pignata11}. {\sn} exploded within a spiral galaxy type Sa \citep{devaucouleurs91} \object{NGC 3278} (\object{ESO 317-G 043}, \object{PGC 031068}) at a redshift of $0.009877\pm0.000123$ \citep{strauss92}. It has an inclination to the line of sight of 41\,deg \citep{hyperleda}.\footnote{\hyperleda.} The {\sn} explosion site was reported to have super-solar metallicity \citep{levesque10d}. The objectives of this paper are: {\it i)} to provide the first resolved measurement of the atomic gas properties of a relativistic SN host, {\it ii)} to test whether these properties are consistent with a recent inflow of atomic gas from the intergalactic medium, and {\it iii)} to derive the properties of {\ngc} to assess the possible implications regarding the nature of the progenitor of {\sn}. We use a cosmological model with $H_0=70$ km s$^{-1}$ Mpc$^{-1}$, $\Omega_\Lambda=0.7$, and $\Omega_m=0.3$, so {\sn} at $z= 0.009877$ is at a luminosity distance of 42.6 Mpc and $1\arcsec$ corresponds to 203 pc at its redshift. We also assume the \citet{chabrier03} initial mass function (IMF), to which all star formation rate (SFR) and stellar masses were converted (by dividing by 1.8) if given originally assuming the \citet{salpeter} IMF. | \label{sec:conclusion} We obtained 21\,cm hydrogen line ({\hi}) and optical integral field unit spectroscopy observations of {\ngc}, the host galaxy of the relativistic {\sn}. This is the first time the atomic gas properties of a relativistic SN host have been provided and the first time resolved 21\,cm-hydrogen-line ({\hi}) information is analysed for the host of an SN of any type in the context of the SN position. The atomic gas distribution of {\ngc} is not centred on the optical galaxy centre, but instead around a third of the atomic gas resides in the region close to the SN position. This galaxy has a few times lower atomic and molecular gas masses than predicted from its star formation rate (SFR). Its specific star formation rate ($\mbox{sSFR}\equiv\mbox{SFR}/\mstar$) is approximately two to three times higher than the main-sequence value, placing it at the higher end of the main sequence towards starburst galaxies. {\sn} exploded close to the region with the highest SFR density and the lowest age, as evident from a high H$\alpha$ EW, corresponding to the age of the stellar population of $\sim5.5$\,Myr. Assuming this timescale was the lifetime of the progenitor star, its initial mass would have been close to $\sim36\,\msun$. The gas properties of {\ngc} are consistent with a recent inflow of gas from the intergalactic medium, which explains the concentration of atomic gas close to the SN position and the enhanced SFR. Super-solar metallicity at the position of the SN (unlike for most GRBs) may mean that relativistic explosions signal a recent inflow of gas (and subsequent star formation), and their type (GRBs or SNe) is determined by either {\it i)} {\it } the metallicity of the inflowing gas, so that metal-poor gas results in a GRB explosion and metal-rich gas (for example a minor merger with an evolved galaxy or re-accretion of expelled gas) in a relativistic SN explosion without an accompanying GRB, {\it ii)} the efficiency of gas mixing (efficient mixing for SN hosts leading to quick disappearance of metal-poor regions), or {\it iii)} the type of the galaxy (more metal-rich galaxies would result in only a small fraction of star formation being fuelled by metal-poor gas). % | 18 | 8 | 1808.00977 |
1808 | 1808.05887_arXiv.txt | We present high resolution transmission spectra, calculated directly from a 3D radiative-hydrodynamics simulation that includes kinetic cloud formation, for HD~209458b. We find that the high opacity of our vertically extensive cloud deck, composed of a large number density of sub-$\mu$m particles, flattens the transmission spectrum and obscures spectral features identified in observed data. We use the PandExo simulator to explore features of our HD~209458b spectrum which may be detectable with the James Webb Space Telescope (JWST). We determine that an 8 -- 12\,$\mu$m absorption feature attributed to the mixed-composition, predominantly silicate cloud particles is a viable marker for the presence of cloud. Further calculations explore, and trends are identified with, variations in cloud opacity, composition heterogeneity and artificially scaled gravitational settling on the transmission spectrum. Principally, by varying the upper extent of our cloud decks, rainout is identified to be a key process for the dynamical atmospheres of hot-Jupiters and shown to dramatically alter the resulting spectrum. Our synthetic transmission spectra, obtained from the most complete, forward atmosphere simulations to--date, allow us to explore the model's ability to conform with observations. Such comparisons can provide insight into the physical processes either missing, or requiring improvement. | \label{sec:intro} Clouds are expected to be ubiquitous in exoplanetary atmospheres \citep{marley13}. Transmission spectra obtained observationally from hot--Jupiters, highly--irradiated Jovian--like giant exoplanets, often contain a number of gas--phase atomic and molecular absorption features \citep[e.g.][]{charbonneau02,snellen10,sing11,birkby17}, possible small--particle condensates or photochemical `haze' that appears in spectra as non--H$_2$/He Rayleigh scattering \cite[e.g.][]{etangs08,pont08,nikolov15,kirk17} and evidence of clouds in the form of a large multi--wavelength opacity that can weaken water and other gaseous signatures \citep{deming13,sing16,iyer16}. The ability for clouds and haze to mute or mask entirely the underlying chemical composition and thermal structure of their host atmospheres, means an improved understanding of their atmospheric feedback may be critical in correctly interpreting observations. \begin{figure*} \centering \includegraphics[scale=0.4]{trans_spec2.pdf} \caption{Spherical shell geometry used for the calculation of transmission spectra. Left-hand plot shows the view perpendicular to the transit. Right-hand plot shows the view from the night side in line with the transit. Parameters are shown for a model column located in the position of the dotted line in each plot, giving a transmission spectrum at the point where the arrow leaves the top of the atmosphere (indicated by the dot in the right-hand plot). $\zeta$ denotes the stellar zenith angle, $b$ the impact parameter, and $ds$ the path length element for the layer bounded by radii $r_1$ and $r_2$. Note the path of the beam will pass through each layer twice, except for the layer in which the impact parameter is found.} \label{fig:trans_calc} \end{figure*} Since cloud formation is dependent upon the local thermochemical conditions, mapping clouds is important to infer the underlying atmospheric properties. The distribution of clouds across a diverse range of planetary types is made more complex by the flow or advection, particularly from super--rotating equatorial or general zonal jets \citep[see e.g.][]{showman02,cho06,menou09,showman11,heng11,dobbs13,rauscher14,mayne14,carone15,heng15,carone16,kataria16,mayne17}, meridional advection \citep[from jet--momentum coupling, see][]{showman11,mayne13,mayne17,lines18a} and vertical mixing from a combination of mean flow (circulation) and atmospheric turbulence \citep{parmentier13}. The swift increase in our knowledge of the dynamics and structure of hot-Jupiter atmospheres has been due to, in part, the development and adaptation of 3D atmosphere and Global Circulation Models (GCMs) which can capture both the full vertical and horizontal dynamics \citep[e.g.][]{showman02,menou09,rauscher13,mayne14}. Since the atmospheres of hot-Jupiters are heated by both a convective flux at the base of the radiative zone, and intense stellar irradiation, one important model consideration is the treatment of radiative transfer. Cloud--free simulations have produced results that closely match observations, e.g. the prediction of kilometre--per--second wind velocities from the super--rotating jets \citep{snellen10,louden15,brogi16}, and the agreement with the observed day--side emission \citep{showman09,amundsen16}. Aerosols are prevalent in the atmospheres of planets within our solar system, and the influence on their host atmospheres \citep[e.g.][]{zhang17} demonstrates clearly that neglecting the radiative feedback from clouds on their host atmospheres is not always a suitable approximation. In the last few years, a number of cloudy-GCMs, of varying complexity, have been developed \citep{parmentier13,oreshenko16,lee16,parmentier16,roman18} and advanced our understanding of cloud dynamics, radiative feedback and their effect on observables. Recently, in \citet{lines18a}, we continued this effort by coupling the Met Office GCM, the {\emph{Unified Model}} (UM) to a sophisticated kinetic, non-equilibrium cloud formation model \citep{woitke03,woitke04,helling06,helling08,lee16}. The coupled cloud-GCM model considers the homogeneous nucleation of seed particles and subsequent heterogeneous surface growth (condensation) and evaporation. The model also allows for the advection of cloud and depleted/enriched gas with the bulk atmospheric flow, gravitational settling (precipitation) and both gas and solid phase interaction with planetary and stellar radiation via absorption and scattering. \begin{figure*} \begin{subfigure}{0.48\textwidth} \includegraphics[scale = 0.56, angle = 0]{nd_1mbar.pdf} \caption{} \end{subfigure}\hspace*{\fill} \begin{subfigure}{0.48\textwidth} \includegraphics[scale = 0.56, angle = 0]{nd_01mbar.pdf} \caption{} \end{subfigure} \medskip \begin{subfigure}{0.48\textwidth} \includegraphics[scale = 0.56, angle = 0]{aod_1mbar_10u.pdf} \caption{} \end{subfigure}\hspace*{\fill} \begin{subfigure}{0.48\textwidth} \includegraphics[scale = 0.56, angle = 0]{aod_01mbar_10u.pdf} \caption{} \end{subfigure} \medskip \begin{subfigure}{0.48\textwidth} \includegraphics[scale = 0.56, angle = 0]{aod_1mbar_1u.pdf} \caption{} \end{subfigure}\hspace*{\fill} \begin{subfigure}{0.48\textwidth} \includegraphics[scale = 0.56, angle = 0]{aod_01mbar_1u.pdf} \caption{} \end{subfigure} \caption{Cloud particle number density (upper), aerosol total optical depth for $\lambda$ = 6.45 - 10.40\,$\mu$m (middle) and aerosol total optical depth for $\lambda$ = 0.96 - 1.01\,$\mu$m (lower) in HD~209458b. Data obtained from the `hot' simulation at t$_{\textrm{cloud}}$ = 100 days in \citet{lines18a}. The substellar point is at $\lambda$ = 180$^o$.} \label{fig:h209_cloud} \end{figure*} While theoretical modelling provides a valuable opportunity to understand atmospheric chemical processes which are not directly observable in the atmospheres of extrasolar planets, in order to verify their accuracy, comparisons to observed datasets are necessary. With respect to observations of cloud in substellar atmospheres, a combination of powerful methods have been used, including analysis of emission spectra \citep{knutson08,line16,evans17} and phase curves \citep{armstrong16,demory13,stevenson14,stevenson17}. Transmission spectra, however, remain one of the most valuable sources of atmospheric data, allowing for the direct identification of gas-phase species via their interaction with the stellar photons. This claim is supported by the wealth of transmission spectra obtained from forward models of hot-Jupiter atmospheres \citep[e.g.][]{seager00, brown01, dobbs12, showman13, wakeford15, goyal18}. With hot-Jupiters expected to contain similar chemical species, albeit in variable quantities due to their vast thermal range, the detection or absence of signatures of key species, such as the alkali metals and water vapour, can help us to infer the presence of haze and/or condensate cloud. One of the foremost issues with atmosphere characterisation however, is the continued absence of directly detected clouds; haze and clouds are currently indirectly inferred from non-H$_2$/He Rayleigh scattering and weakened spectral signatures from their expected broad opacity. Many studies have discussed potential condensate species operating across the JWST wavelength window \citep[e.g.][]{helling06b,helling08,min08,zeidler13,lee14,marley15,wakeford15,zeidler15,helling16,kitzmann18}. Of particular interest for mineral dust clouds forming in hot-Jupiters is the Si-O bond which has active vibrational modes operating between $\lambda$ $\sim$ 8 -- 12\,$\mu$m \citep{lee14,wakeford15}. Plotting the complex refractive index (see Figure \ref{fig:nk}) for our included dust species displays, via the extinction coefficient, the strong attenuating properties of the silicate species in the 10\,$\mu$m region can be seen. Cloud-coupled GCMs can indicate the expected 3D distribution of cloud, precipitation efficiency and help constrain which condensate species are important for a given planet. Comparisons with observational data can only be made by producing synthetic model observables. \citet{hubbard01}, \citet{seager00} and \citet{brown01} presented some of the first studies into how model transmission spectra can be used as diagnostics to characterise giant exoplanet atmospheres. Atmospheric properties, such as the temperature and cloud cover, were investigated to reveal their impact on the resulting spectrum. While most transmission models have been analytically prescribed for 1D Pressure-Temperature (PT) profiles, \citet{fortney03} presented a 2D study and furthered this in \citet{fortney10} for application to a 3D atmosphere. Since then, various models have considered the effect of a 3D atmosphere on the transmission spectrum \citep{burrows10, dobbs12, kempton12, showman13, drummond18b, powell18}. In this letter, we present for the first time, synthetic transmission spectra of a cloudy hot-Jupiter atmosphere which are computed directly from a state-of-the-art prognostic (predictive) and radiatively-active cloudy 3D simulation. This methodology is applied to the existing radiatively-active HD~209458b cloudy atmosphere simulation performed in \citet{lines18a} and we explore within this atmosphere, how cloud particle opacity, composition and gravitational settling affect the transmitted flux. In Section \ref{sec:theory} we introduce the details of our transmission spectrum calculation, initial conditions and methodology, in Section \ref{sec:results} we present the transmission spectra and potential for cloud detection with the James Webb Space Telescope (JWST), and in Section \ref{sec:discussion} we discuss and summarise the implications from our findings. | \label{sec:discussion} Since our simulations produce vertically extensive cloud condensate decks with large cloud particle number densities ($>$ 3 x 10$^4$ cm$^{-3}$) even at the lowest pressures, it is not surprising that the blanket opacity across the visual and infra--red mask, almost completely, the atomic and molecular signatures in the gas--phase. Despite this, the obtained synthetic transmission spectrum for the full cloud opacity model still contains potentially detectable features. Firstly there is the flat nature, compared to a clear-skies atmosphere, of the spectrum. This indicates the presence of high opacity aerosols that can efficiently obscure the spectral signatures of the majority of gas phase absorbers, unless the cloud opacity is scaled down or the cloud deck upper boundary or `cloud top' is forced to higher (deeper) pressures. Secondly, there remains a tiny signature from CH$_4$ absorption which is amplified from the cooler atmospheric conditions produced by highly scattering cloud in our \citet{lines18a} simulations. Finally, the 9\,$\mu$m cloud absorption feature, shown to be broadened by mixed-composition cloud particles, is potentially detectable by JWST, providing a reduced cloud opacity. The previously predicted, Si-O vibrational absorption from cloud particles \citep{lee14,wakeford15}, leads to strong absorption (shown via the imaginary component of the material refractive indices) at around 10\,$\mu$m for the SiO$_2$, TiO$_2$, MgSiO$_3$ and Mg$_2$SiO$_4$ condensates considered in our work. Further understanding of how this feature changes in response to the cloud composition will likely prove critical in determining cloud properties (e.g. composition, particle/droplet size) of observed atmospheres. We address the bias introduced by an approximation in our transmission model, and find that the offset in the transit radius ratio when considering atmospheric properties about both the night and daysides of the terminator, is minimal. The offset, which for our case lies below the JWST instrument noise, is a function of the cloud longitudinal asymmetry level. To address planetary atmospheres which could feature stronger asymmetries in the cloud distribution, future work may need to improve upon our approximation. One aspect of cloud formation and evolution that could exert a strong influence on the cloud vertical structure and observable properties, is the gravitational settling of cloud and/or haze particles. While cloud particles can obtain large precipitation velocities \citep[see][]{lines18a}, this fall-speed can be offset by a combination of vertical mean winds that result in a slow net downwards advection of cloud. Cloud particles may also be transported by turbulence, but this process is not included in our simulations since the model does not include a sub-grid turbulence parametrisation nor is at a suitable resolution to capture these processes. Future studies will need to consider the implications of turbulence, since this may have a significant effect on the cloud vertical structure. As cloud particles settle to deeper pressures, the total cloud cross-section reduces in the upper atmosphere, leading to a reduced opacity in the transmission region. While the timescales involved are typically too long to capture with current 3D simulations, it will be necessary in future studies to address this mechanism. The results from our enhanced precipitation tests, which simulate the effects of a more vertically evolved cloud deck, indicate that the ability to detect gaseous absorption signatures through the grey opacity, is strongly dependent on the cloud top pressure. The retrieval analysis of \citet{barstow17} on the \citet{sing16} dataset have previously indicated the importance of this connection, stressing the potential wide range of cloud top pressures across their small sample of hot-Jupiters. While we can, for the cloud structural and compositional conditions determined in \citet{lines18a}, constrain the cloud top pressure to be P $>$ 15 mbar for HD~209458b, the motivation of this result is clearly to more accurately pin-point the vertical equilibrium of the cloud (and hence upper boundary of the cloud opacity). It is also worth considering and cautioning that by parametrising the cloud top, the complex feedback between the cloud radiative transfer and the atmosphere's thermal, chemical and dynamical properties is circumvented. Thus, the resulting synthetic observations do not necessarily represent an atmosphere with a PT profile that has converged to the cloud vertical extent. \begin{figure*} \includegraphics[scale=0.4]{h209_pandexo.pdf} \caption{JWST detectability with PandExo simulator showing each simulation from Figure \ref{fig:h209_tran} with the same format and colour scheme.} \label{fig:h209_jwst} \end{figure*} Observations of HD~209458b have revealed, repeatedly, the presence of the sodium \citep{charbonneau02,sing08,snellen08,jensen11} as well as water vapour and carbon monoxide \citep{snellen10,deming13}. It is interesting therefore to consider that while the literature typically considers HD~209458b to have a cloudless atmosphere due to the prominence of the aforementioned gaseous chemical signatures, our \citet{lines18a} simulations are remarkably cloudy in terms of the opacity and both horizontal and vertical distribution of cloud particles. As a result, our transmission spectrum probes that of an optically thick atmosphere, with many of the radiative interactions occurring within the first few vertical layers of dense silicate particles; cloud is radiatively dominant in the most upper layer. The result is an inevitably featureless profile, since there is no layer of gas absorbers above the cloud that can imprint on our spectra. Our opacity scaling tests indicate that a reduction in both the absorption and scattering coefficients of at least three orders of magnitude is required to see CO and water bands in the model spectrum, although CH$_4$ is visible in the full cloud opacity spectrum. We note that even with a strong reduction in the opacity, we are still unable to unmask the alkali metal spectral signatures. Our inability to emulate the observed spectrum of HD~209458b is an indication that we are missing physical processes or chemical constituents in our model\footnote{Caution must always be taken when making comparisons between transmission spectra from theoretical models and observations due to the presence of degeneracies, for example, those existing between baseline pressure, planetary radius and absorber abundances \citep[see][for more information]{benneke12,heng17b}.}. Our current model setup considers only five condensate materials and therefore may underestimate the atmospheric `cloudiness' by mass. The inclusion of a wider variety of important condensing species may lead to increased cloud particle sizes which will modify the cloud particle optical properties; the significance of this consideration is posited by the 1D microphysical study of \citet{powell18} who use a mass-binning technique for their cloud particles to reveal an irregular distribution of silicate particles and explore its importance. The current omission of iron itself, as a strong absorber, means we are likely underestimating cloud warming and this could lead to a situation whereby the atmospheric temperature increases, instead of the net cooling mechanism from a silicate dominant one. Aside from the aforementioned increase in condensate species, particle growth via coagulation could also play a role in producing heavier particles. Coagulation has been shown to be a requirement to drive precipitation in Earth's atmosphere \citep{pruppacher78}, with condensational growth alone not able to provide particle sizes large enough to initiate rainout. However, \citet{helling08-b1} shows that for sub-micron silicate grains in substellar atmospheres, coagulation can operate on a timescale orders of magnitude longer than chemical growth and therefore may not be an essential model component. Our chosen initial metallicity may also play a significant role in cloud abundance, with non--solar values either reducing or increasing the total condensate mass and giving rise to a weakened or increased cloud opacity. This has been demonstrated in \citet{helling17b} for varying C/O ratios, and in \citet{mahapatra17} for rocky versus solar element abundances. Whatever the precise mechanism behind condensate growth, larger particles will effectively sediment out of the atmosphere and support the importance of investigating the role of precipitation. We also acknowledge the omission of hydrocarbons formed photochemically in the gas--phase (e.g., Polycyclic Aromatic Hydrocarbons, or PAHs), which may play a large role in the atmospheric opacity and hence spectrum. Even if PAHs cannot themselves survive the large UV flux on the day side, their precursors radicals have been found to be quite abundant in the upper atmospheric layers \citep{venot15}. That situation may change as the radical--rich gas is advected by zonal winds from the hot/irradiated day side to the cold/dark night side, leading to the efficient formation of PAHs and, subsequently, photochemical hazes. It has been demonstrated however, that the abundances of PAHs can be very low, despite their abundances increasing in cloud forming regions due to the reduction of oxygen \citep{bilger13}. These few physically motivated model adjustments alone may be enough to better reproduce the observed data, by way of increasing the detectability of gas--phase atomic and molecular absorption features. The benefit and power of using this physically motivated model is the ability to isolate and identify such specific physical processes and details which can be difficult for parametrised models which may obscure the underlying physical mechanisms at play.Therefore, although unlikely to represent the current conditions on HD~209458b, our simulations do provide an insight into those atmospheres which contain optically thick cloud, potentially from suspended silicate condensates. For example, our results are consistent with those of WASP-101b, which has a flat spectrum (there is no WFC3 H$_2$O feature) despite the planet possessing similar atmospheric properties to HD~209458b \citep{wakeford17b}. This is a possible indication that subtle differences in the atmospheric circulation, metallicity, C/O ratio etc. and interplay with the cloud radiative feedback could lead to large changes in the cloud's ability to impact on the observations. Additionally, despite the grey cloud opacity for our simulations, we are still able to identify transmission features which indicate cloud coverage and which have the potential for detection with JWST. Finally, we acknowledge the convenient ability to obtain synthetic transmission spectra directly from our 3D simulations which will enhance our ability to connect with the latest observations. | 18 | 8 | 1808.05887 |
1808 | 1808.07772_arXiv.txt | We present an improved Minimal Variance (MV) method for using a radial peculiar velocity sample to estimate the average of the three-dimensional velocity field over a spherical volume, which leads to an easily interpretable bulk flow measurement. The only assumption required for this interpretation is that the velocity field is irrotational. The resulting bulk flow estimate is particularly insensitive to smaller scale flows. We also introduce a new constraint into the MV method that ensures that bulk flow estimates are independent of the value of the Hubble constant $H_o$; this is important given the tension between the locally measured $H_o$ and that obtained from the cosmic background radiation observations. We apply our method to the \textit{CosmicFlows-3} catalogue and find that, while the bulk flows for shallower spheres are consistent with the standard cosmological model, there is some tension between the bulk flow in a spherical volume with radius $150$\hmpc\ and its expectations; we find only a $\sim 2\%$ chance of obtaining a bulk flow as large or larger in the standard cosmological model with \textit{Planck} parameters. | \label{sec:intro} One of the central ideas of the standard model of cosmology is the gravitational instability paradigm, in which small amplitude density fluctuations in the early Universe were amplified by gravity into the large-scale structure that we see today \citep{MukChi81}. In this picture, growth of fluctuations continues today, taking the form of flows of galaxies on scales of 1 Mpc and larger. On scales of order 100 Mpc, these flows are still in the linear regime, and thus can be directly related to density fluctuations on these scales. Large-scale flows are thus an important cosmological probe that can potentially provide confirmation of many features of the standard model. Given that measurements of peculiar velocities of individual galaxies have large uncertainties, researchers have traditionally focused on velocity field statistics that average over many galaxies; by far the most common statistic is the dipole moment of the velocity field, also called the bulk flow. Early attempts to measure the bulk flow \citep[\eg][]{RubThoForRob76,DreFabBurDav87,LP94,RiePreKir95} were plagued by small sample sizes and difficult to control biases. Recently, however, the compilation of large catalogues of distance measurements with $\sim 10,000$ objects \citep[\eg][]{CF3} has made it possible to accurately estimate the bulk flow on scales $\sim 100$\hmpc\ \citep{ WatFelHud09,FelWatHud10,DavNusMas11,NusDav11,Nusser14,WatFel15,ScrDavBlaSta15,FeiBraNus17,HelNusFeiBil17,HelBilLib18}. For recent discussions of bulk flow determination methods, see \eg, \citet{DavScr14,WatFel15a,Nusser16}. Bulk flows have additional significance as a large-scale flow statistic in being, at least in principle, independent of the value of the Hubble constant, $H_o$. While calculated peculiar velocities depend strongly on the Hubble constant, since bulk flows measure the dipole of the velocity field, they are theoretically insensitive to phantom monopole flows introduced by using an incorrect value of $H_o$. This characteristic of bulk flows is particularly appealing in an era where there is tension between cosmic background radiation measurements of $H_o$ and more local measurements using the Hubble diagram (see, \eg, \citet{RieMarHofSco16,BerVerRie16,PlanckXLVIII16,Freedman17,RieCasYuaMar18}). This paper is organized as follows: in Section~\ref{stats} we discuss some of the different ways of defining the bulk flow and the difficulty of comparing bulk flow estimates reported in different studies. In Section \ref{sec:theory} we follow \citet{Nusser14} and show how the assumption of an irrotational flow allows the full, three-dimensional bulk flow to be written in terms of a weighted average of the radial component of the peculiar velocity, the only component of the velocity that is actually observable. In Section~\ref{sec:MV}, we show that the resulting weighted average is well suited for estimation via the Minimal Variance (MV) method. We also show how a new constraint can be added to the MV method to ensure that the bulk flow estimate is independent of the value of the Hubble constant. Section \ref{sec:data} discusses the data on which we apply our method. In Section \ref{sec:results} we present the results of our analysis. Finally, in section \ref{sec:discussion}, we discuss our results and put them in the context of other estimates of the bulk flow. | \label{sec:discussion} We have introduced a new method for calculating the bulk flow from a catalogue of peculiar velocities that has several advantages over other methods. First and foremost, it is much easier to interpret than estimates obtained using other methods. In particular, with the minimal assumption that the velocity field is irrotational, the bulk flow estimated with the MV method using a $r^{-2}$ weighted ideal survey corresponds to the integral of the full three dimensional velocity field over a well defined spherical volume. In cases where the data cannot provide an accurate estimate of the bulk flow on a particular scale, it is clearly indicated by the window function of the bulk flow estimate not matching the ideal survey window function. In contrast, bulk flows calculated by methods such as maximum likelihood are very difficult to interpret, as they are particularly sensitive to a given survey geometry, and sample and error distributions and thus do not correspond to a well-defined volume. Methods such as maximum likelihood give the most weight to nearby galaxies that have smaller uncertainties in their velocities, which can result in bulk flow estimates that probe smaller scales than those of the survey probes. A second advantage of the MV method is that it is very effective at averaging out flows on scales smaller than the volume of interest, as can be seen from the window functions in Fig.~\ref{fig:wincomp}. Other methods of estimating the bulk flow, including other weighting schemes, can give window functions with wider central peaks and with non-negligible side lobes, resulting in bulk flow estimates that contain significant contributions from smaller scale power. This is a very important consideration. Bulk flow components are different from some other cosmological probes in that models do not predict their mean, which is zero, but rather their variance. The strongest possible constraint from bulk flow measurements comes from minimizing the predicted variance; a smaller predicted variance shrinks the acceptable range of bulk flow component values for a given model. The primary way of reducing this variance is by increasing the depth of peculiar velocity surveys. This has the effect of reducing the width of the central peak of the window function so that a larger fraction of bulk flow is coming from large scales, where the power spectrum vanishes. However, a deep peculiar velocity survey does not necessarily guarantee a small predicted variance. If a bulk flow analysis gives more weight to the inner part of the survey, where there is more information, then the central peak of the window function can be wider than expected. Furthermore, additional variance can come from the incomplete cancellation of smaller scale flows. The MV method with $r^{-2}$ weighting can provide bulk flow estimates with minimum predicted variance both by having very small side lobes and thus allowing one to simply determine the maximum radius for which a given survey can accurately determine the bulk flow. An additional advantage of the MV method is that it allows for constraints to be easily placed on the bulk flow moments using Lagrange multipliers. Here we have imposed a constraint that ensures that the bulk flow moments are independent of the value of the Hubble constant $H_o$. For the case where CF3 catalog velocities are calculated using $H_o=75$ km/s/Mpc, a value for which radial flows are roughly minimized, the constraint changes the bulk flow components by a few \kms and reduces the $\chi^2$ value from 10.49 to 9.59. The probability of finding as large or larger a bulk flow is increased to 2.2\% from the 1.5\% obtained without the constraint. We note that the current tension in the value of the Hubble constant makes it difficult to assess the magnitude of the radial flows in our local volume. The effect of the constraint could be somewhat larger in a peculiar velocity catalog that has more significant radial flows or a more anisotropic distribution. While the result of our analysis, that there is only a $\sim 2\%$ chance of obtaining the observed bulk given the parameters of the cosmological standard model, is similar to that of WFH, it is important to note the differences in both the data and the method used. First, the quantity of peculiar velocity measurements has increased dramatically; whereas WFH had $\sim$4,500 peculiar velocities of groups and individual galaxies in the COMPOSITE catalogue, the current analysis uses nearly 12,000, an increase of more than a factor of two. The addition of new data hasn't significantly changed the direction of the bulk flow, but it has reduced it's magnitude; the gaussian weighted bulk flow with $R=150$\hmpc\ (corresponding to $R=50$\hmpc\ in WFH) went from about $400$\kms\ in WFH to less than $300$\kms\ in the current analysis. The addition of new data has also resolved an unexpected result from WFH; they saw the bulk flow sharply \textit{increase} as the scale $R$ became large, a result that is difficult to provide a physical explanation for. In Fig.~\ref{fig:bf} we see that in the current work the bulk flow decreases and then remains roughly constant with increasing $R$. Second, we have introduced a $r^{-2}$ weighting scheme which results in a window function with smaller side lobes than in the Gaussian weighting, indicating that bulk flows calculated with this weighting are less susceptible to velocity modes on scales smaller than the survey. Given that they are only sensitive to scales as large or larger than $R$, they necessarily have smaller \textit{expectations} for the bulk flow, as seen in Fig.~\ref{fig:bf}. However, we also see in Fig.~\ref{fig:bf} that the bulk flow for the $r^{-2}$ weighting is actually larger than that found using Gaussian weighting, so that the $\chi^2$ for the $r^{-2}$ weighting is significantly larger than for the Gaussian case. In fact, the $\chi^2$ we calculate for the $r^{-2}$ weights are similar to those obtained by WFH for Gaussian weighting. It is important to note that disagreement with the standard cosmological model only occurs at the largest scales probed in this study, $\sim 150$\hmpc. At smaller scales, the bulk flow magnitude is consistent with expectations. Thus the question of whether the bulk flow is inconsistent with the standard cosmology depends strongly on precisely how the bulk flow is calculated. It is not surprising, then, that there is disagreement on this question in the community, where different analyses can weigh information in different ways, even when using the same peculiar velocity catalogue. Our result by itself is only suggestive of continuing tension with the standard model, but is not conclusive. Ultimately, resolving the question of whether large scale flows are consistent with expectations will require additional measurements of peculiar velocities, particularly of objects with distances $\sim 150$\hmpc. \noindent{\bf Acknowledgements:} SP and RW have been supported in part by a grant from the Murdock Charitable Trust. We would like to thank Andrew Jaffe for his helpful comments. | 18 | 8 | 1808.07772 |
1808 | 1808.00478_arXiv.txt | Cosmic sources of gamma-ray radiation in the GeV range are often characterized by violent variability, in particular this concerns blazars, gamma-ray bursts, and the pulsar wind nebula Crab. Such gamma-ray emission requires a very efficient particle acceleration mechanism. If the environment, in which such emission is produced, is relativistically magnetized (i.e., that magnetic energy density dominates even the rest-mass energy density of matter), then the most natural mechanism of energy dissipation and particle acceleration is relativistic magnetic reconnection. Basic research into this mechanism is performed by means of kinetic numerical simulations of various configurations of collisionless relativistic plasma with the use of the particle-in-cell algorithm. Such technique allows to investigate the details of particle acceleration mechanism, including radiative energy losses, and to calculate the temporal, spatial, spectral and angular distributions of synchrotron and inverse Compton radiation. The results of these simulations indicate that the effective variability time scale of the observed radiation can be much shorter than the light-crossing time scale of the simulated domain. | The high-energy gamma-ray sky (0.1-10 GeV), as observed by the Fermi Large Area Telescope \citep{2009ApJ...697.1071A}, is dominated by two classes of sources: blazars and pulsars. Gamma-ray pulsars are concentrated along the Galactic plane, while extragalactic blazars are distributed uniformly. Both of these sources are strongly variable, and their gamma-ray emission is non-thermal, indicating efficient mechanisms of particle acceleration. Gamma-ray variability has been observed in certain blazars and radio galaxies on time scales of several minutes, much shorter than the light crossing time of their supermassive black holes (typically a few hours), e.g.: PKS~2155-304 \citep{2007ApJ...664L..71A}, Mrk~501 \citep{2007ApJ...669..862A}, PKS~1222+216 \citep{2011ApJ...730L...8A}, IC~310 \citep{2014Sci...346.1080A}, 3C~279 \citep{2016ApJ...824L..20A}. It has been argued that such short variability time scales require a highly localized dissipation mechanism, not directly related to variations in the jet structure induced at the central black hole. In addition, such rapid variations observed at very high gamma-ray luminosity impose a potential problem of intrinsic absorption of the gamma-ray radiation in a photon-photon pair creation process. Such absorption can be avoided by postulating very high Doppler or Lorentz factor $\mathcal{D} \sim \Gamma \sim 100$ \citep{2008MNRAS.384L..19B}. In the case of luminous quasars, like PKS~1222+216 and 3C~279, additional absorption of gamma rays can be expected at subparsec scales due to the external radiation field that includes broad emission lines and direct accretion disk radiation. Relativistic magnetic reconnection has been proposed as a solution to these challenges in the form of the minijets model, in which reconnection produces additional relativistic bulk outflows in the jet co-moving frame, increasing the effective Doppler factor \citep{2009MNRAS.395L..29G}. A semi-analytical model of minijets has been applied directly to the case of PKS~2155-304 \citep{2011MNRAS.413..333N}. However, that model was highly simplified, and over the last several years numerical simulations showed that relativistic magnetic reconnection is a much more complex phenomenon. Understanding of magnetic reconnection has been developing slowly since the first ideas were formulated in the 1950s in the context of solar physics. Analytical models have difficulty in describing the reconnection process in detail, as it necessarily involves plasma physics beyond the standard MHD or force-free regimes. Kinetic numerical simulations, in particular the particle-in-cell (PIC) algorithm, are our best tools for studying magnetic reconnection in both non-relativistic and relativistic regimes. Particle acceleration in relativistic reconnection has been studied with PIC simulations since the work of \cite{2001ApJ...562L..63Z}. At that time, it was not even clear whether relativistic reconnection is an efficient dissipation mechanism, as solutions based on smooth Sweet-Parker current layers predicted very low reconnection rates (slow inflow velocities, and hence weak electric fields). It has been known that long current layers are prone to tearing-mode instability, which produces chains of plasmoids, but PIC simulations were necessary to demonstate that formation of plasmoids accelerates the reconnection rate \citep{2004PhPl...11.1151J}. With increasing computational power, PIC simulations showed that relativistic reconnection is very efficient in accelerating particles, producing power-law particle energy distributions with indices $N(\gamma) \propto \gamma^{-p}$ approaching $p \to 1$ in the limit of relativistic background magnetization $\sigma_0 = B_0^2/(4\pi w_0) \gg 1$, where $w_0$ is relativistic enthalpy including the rest-mass energy \citep{2014ApJ...783L..21S,2014PhRvL.113o5005G,2016ApJ...816L...8W}. Recent simulations of Harris layers with open boundaries showed how reconnection can operate as a steady-state mechanism producing plasmoids of specific size distribution accelerated to relativistic bulk motions \citep{2016MNRAS.462...48S}. Introduction of synchrotron radiative losses to the PIC algorithm \citep{2013ApJ...770..147C} allowed to study particle acceleration under severe radiative losses, with direct application to the gamma-ray flares from the Crab Nebula \citep{2011Sci...331..736T}, which are interpreted in terms of synchrotron emission exceeding the radiation reaction photon energy limit of $\sim 100\;{\rm MeV}$. Finally, alternative initial conditions are being explored, i.e., ``ABC fields'' that allow in addition to study the formation and dynamics of current layers \citep{2016ApJ...826..115N,2016ApJ...828...92Y,2017JPlPh..83f6302L}. One of the most interesting properties of radiation produced by relativistic magnetic reconnection is its rapid variability. This is illustrated with the example of 2-dimensional simulation of a Harris layer described in detail in \cite{2015ApJ...815..101N}. These results have been obtained by us from numerical simulations performed with the PIC code Zeltron \citep{2013ApJ...770..147C} created by Beno{\^i}t Cerutti\footnote{http://benoit.cerutti.free.fr/Zeltron/}. The space-time diagrams of synchrotron emissivity and the corresponding lightcurves are presented here for the first time. \begin{figure}[ht] \centering \includegraphics[width=0.7\textwidth]{xymap_ne.png} \caption{Snapshots (x,y-maps) from the evolution of Harris layer undergoing tearing instability and hierarchical mergers of plasmoids. Color scale indicates the particle (electron-positron) density $n_{\rm e}$, and cyan lines indicate the magnetic field lines. Simulation time for each panel is given on the left side. The dashed white lines mark the regions, from which the x-profiles used to build the (x,t) space-time diagrams were extracted.} \label{fig_xy} \end{figure} | 18 | 8 | 1808.00478 |
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1808 | 1808.07896_arXiv.txt | The Simons Observatory (SO) will provide precision polarimetry of the cosmic microwave background (CMB) using a series of telescopes which will cover angular scales from arc-minutes to tens of degrees, contain over 60,000 detectors, and observe in frequency bands between 27 GHz and 270 GHz. SO will consist of a six-meter-aperture telescope initially coupled to roughly 35,000 detectors along with an array of half-meter aperture refractive cameras, coupled to an additional 30,000+ detectors. The large aperture telescope receiver (LATR) is coupled to the SO six-meter crossed Dragone telescope and will be 2.4 m in diameter, weigh over 3 metric tons, and have five cryogenic stages (80 K, 40 K, 4 K, 1 K and 100 mK). The LATR is coupled to the telescope via 13 independent optics tubes containing cryogenic optical elements and detectors. The cryostat will be cooled by by two Cryomech PT90 (80 K) and three Cryomech PT420 (40 K and 4 K) pulse tube cryocoolers, with cooling of the 1 K and 100 mK stages by a commercial dilution refrigerator system. The secondo component, the small aperture telescope (SAT), is a single optics tube refractive cameras of 42 cm diameter. Cooling of the SAT stages will be provided by two Cryomech PT420, one of which is dedicated to the dilution refrigeration system which will cool the focal plane to 100 mK. SO will deploy a total of three SATs. In order to estimate the cool down time of the camera systems given their size and complexity, a finite difference code based on an implicit solver has been written to simulate the transient thermal behavior of both cryostats. The result from the simulations presented here predict a 35 day cool down for the LATR. The simulations suggest additional heat switches between stages would be effective in distribution cool down power and reducing the time it takes for the LATR to reach its base temperatures. The SAT is predicted to cool down in one week, which meets the SO design goals. | \label{sec:intro} The cosmic microwave background (CMB) has become one of the most powerful probes of the early universe. Measurements of its temperature anisotropy on the level of $\sim$ ten parts per million, which have brought cosmology into a precision era, have placed tight constraints on the fundamental properties of the universe. Beyond the temperature anisotropy, CMB polarization anisotropy not only enriches our understanding of our cosmological model, but could potentially provide clues to the very beginning of the universe via the detection (or non-detection) of primordial gravitational waves. A number of experiments have made and are continuing to refine measurements of the polarization anisotropy. However, these experiments are typically dedicated to a relatively restricted range of angular scales, e.g., large angular scales (tens of degrees) or high resolution/small angular scales ($\sim$~1 arcminute). To provide a complete picture of cosmology, both large and small angular scales are important. Ideally these measurements would be made from the same observing site so that the widest range of angular scales can be probed, at multiple frequencies, on the same regions of the sky. This is the goal of the Simons Observatory (SO). The Simons Observatory will comprise a combination a single large aperture telescope receiver (LATR) and and array of small aperture telescopes (SAT) for observing both large and small angular scales. The observatory will be located in Chile's Atacama Desert at an altitude of 5190 m. The LATR is designed with a large FOV capable of supporting a cryostat with up to 19 optics tubes. To limit the development risk, the LATR is designed to accommodate up to 13 optics tubes. We plan to deploy 7 optics tubes with 3 detector wafers in each for a total of roughly 35,000 detectors, primarily at 90/150 GHz in the initial SO deployment. We note that each optics tube could be upgraded to support 4.5 wafers for a~50\% increase in the number of detectors per optics tube. With this upgrade and the deployment of 19 optics tubes, the LAT could support roughly 145,000 detectors at 90/150 GHz. In addition, the SAT cameras will be deployed with more than 30,000 detectors. With the current scale and potential to upgrade, SO will serve as a valuable step to advance to the next-generation CMB experiment, CMB-S4 \citenum{Technology, Science}. To achieve the scientific goal of Simons Observatory, it is important to reduce the non observing time and in particular the time required for cooling down the cryostat as part of the preparation. In this paper, we discuss a model developed to estimate the cooldown of both the LATR and the SAT. In sections~\ref{sec:theory}\textendash\ref{sec:contacts}, the theory for the model is introduced and its mathematical implementation is explained. Sections \ref{sec:cooling} and \ref{sec:setup} are dedicated to introduce the general setup of the system that we simulated including the cooling capacity available and the external load. Finally, in section \ref{sec:results} we discuss the results from the simulations. The model developed in this paper can be also useful in developing the future large CMB experiments, like CMB-S4. | \label{sec:conclusion} The code developed has allowed us to estimate the cooldown of both the LATR and SAT of the Simons Observatory. The cool down time is particularly important for the LATR. Indeed, given the dimensions of the LATR, a proper study of the cool down of the cryostat is necessary. The results from this simulation show that the required time to cool down to 4K is approximately 35 days. This time is dominated by cooling the $1$ K and $100$ mK stages. A possible solution to reduce this time is the use of switches connecting different stages to redistribute the cooling power. For example, nitrogen heat pipes for thermally connecting the $40$ K stage with the $4$ K stage are in the testing phase. Whereas, for connecting colder stages, such as the $4$ K stage with the $1$ K stage more classical heat switches, like gas-gap, are considered. Future tests will give a complete characterization of both nitrogen heat pipes and gas-gap switches so that it is possible to include in the cool down simulation. The other stages take less time to thermalize, therefore connecting these colder stages will provide additional cooling power for the $1$ K and $100$ mK stages. Instead, for the SAT, the cooling time is in the order of $1$ week, which is within SO defined technical requirements. This study could serve as a useful reference for the design of the next generation CMB experiments, like CMB-S4. | 18 | 8 | 1808.07896 |
1808 | 1808.02011_arXiv.txt | On hot Jupiter exoplanets, strong horizontal and vertical winds should homogenize the abundances of the important absorbers CH$_4$ and CO much faster than chemical reactions restore chemical equilibrium. This effect, typically neglected in general circulation models (GCMs), has been suggested as explanation for discrepancies between observed infrared lightcurves and those predicted by GCMs: On the nightsides of several hot Jupiters, GCMs predict outgoing fluxes that are too large, especially in the Spitzer 4.5 $\mu$m band. We modified the SPARC/MITgcm to include disequilibrium abundances of CH$_4$, CO and H$_2$O by assuming that the CH$_4$/CO ratio is constant throughout the simulation domain. We ran simulations of hot Jupiter HD 189733b with 8 CH$_4$/CO ratios. In the more likely CO-dominated regime, we find temperature changes $\geq$50-100 K compared to the equilibrium chemistry case across large regions. This effect is large enough to affect predicted emission spectra and should thus be included in GCMs of hot Jupiters with equilibrium temperatures between 600K and 1300K. We find that spectra in regions with strong methane absorption, including the Spitzer 3.6 and 8 $\mu$m bands, are strongly impacted by disequilibrium abundances. We expect chemical quenching to result in much larger nightside fluxes in the 3.6 $\mu$m band, in stark contrast to observations. Meanwhile, we find almost no effect on predicted observations in the 4.5 $\mu$m band, as the opacity changes due to CO and H$_2$O offset each other. We thus conclude that disequilibrium carbon chemistry cannot explain the observed low nightside fluxes in the 4.5 $\mu$m band. | Close-in extrasolar giant planets, known as hot Jupiters, are the best characterized exoplanets to date. Due to their proximity to their host stars, they are expected to be tidally locked. This creates strong temperature contrasts between the permanent day side and the night side. For the temperature ranges of typical hot Jupiters, assuming chemical equilibrium, these temperature contrasts translate to large horizontal gradients in the abundance of methane (CH$_4$) and carbon monoxide (CO), two important infrared absorbers in the atmospheres of hot Jupiters. Carbon is preferentially found in CH$_4$ at high pressures and low temperatures, while at low pressures and high temperatures, CO is the dominant carbon-bearing species. Most models of hot Jupiters, especially three-dimensional general circulation models (GCMs) allowing for realistic representation of opacities and the effect that chemical composition exert on them, assume equilibrium chemistry as a basis for calculating opacities. However, the stark day-night temperature contrast also drives a vigorous atmospheric circulation, with a strong eastward equatorial jet advecting thermal energy from the day side to the night side and strong vertical mixing \citep[e.g.][]{ShowmanGuillot2002,ShowmanPolvani2011,Dobbs-DixonLin2008,HengEtAl2011b,ThrastarsonCho2011,RauscherMenou2012,PernaEtAl2012,Dobbs-DixonAgol2013,ParmentierEtAl2013,MayneEtAl2014,KatariaEtAl2016,MendoncaEtAl2016}. In addition, at low pressures ($\lesssim 1$ bar), the chemical time scale at which the interconversion between CH$_4$ and CO acts, becomes very long. Therefore, an air parcel is advected much faster than its CH$_4$ and CO abundances can adapt to the local equilibrium values. This process, called quenching, is expected to vertically and horizontally homogenize the abundances of CH$_4$ and CO (as well as of many other species, including N$_2$ and NH$_3$) in the near-infrared photospheres (located roughly between 1 bar and 1 mbar) of hot Jupiters \citep{CooperShowman2006,AgundezEtAl2012,AgundezEtAl2014,DrummondEtAl2018,DrummondEtAl2018HD189733b}. These disequilibrium abundances can have a significant effect on the opacities and thus the radiative transfer. Including this effect in general circulation models (GCMs) could potentially impact the thermal structure and atmospheric circulation, as well as the predicted spectra and phase curves. In fact, this has been proposed as a solution for the observed discrepancy between phase curves predicted by state-of-the-art GCMs assuming equilibrium chemistry and observations with the Spitzer Space Telescope \citep{KnutsonEtAl2012}. GCMs over-predict the night side fluxes in the Spitzer 4.5 $\mu$m band for the hot Jupiters HD 189733b \citep{KnutsonEtAl2012}, HD 209458b \citep{ZellemEtAl2014} and WASP-43b \citep{StevensonEtAl2017}. Within this wavelength bandpass, CO has a strong absorption band. For relatively cool hot Jupiters, like HD 189733b, transport-induced disequilibrium chemistry is expected to enhance the CO abundance on the night side compared to the equilibrium chemistry value \citep{CooperShowman2006,AgundezEtAl2014}. \citet{KnutsonEtAl2012} argued that this would lead to increased opacity in the 4.5 $\mu$m band, such that the outgoing radiation in that band would probe higher, cooler regions of the atmosphere, decreasing the flux emitted in this band. In addition to changing the pressure level from which the outgoing radiation is emitted, the altered opacities can also affect the thermal structure and the atmospheric circulation in a GCM, an effect not considered in the argumentation of \citet{KnutsonEtAl2012}. So far, this effect has only been taken into account by \citet{DrummondEtAl2018,DrummondEtAl2018HD189733b}. The goal of this study is to better quantify how the combination of these two effects of disequilibrium carbon chemistry (the change in the level from which radiation escapes to space and the change in thermal structure) affects the thermal emission spectra predicted from GCMs. We assume that the abundance ratio of CH$_4$ to CO is constant throughout the entire simulation domain and treat the CH$_4$/CO abundance ratio as free parameter. This simple approach allows us to explore a broader parameter range than \citet{DrummondEtAl2018,DrummondEtAl2018HD189733b} and to focus on the radiative effects. Our approach is justified by the findings of previous studies: Coupling a simple chemical relaxation scheme to a GCM of HD 209458b, both \citet{CooperShowman2006} and \citet{DrummondEtAl2018} found that the CO and CH$_4$ abundances are homogenized everywhere above the $\sim 3$ bar level. In a later study (published while this paper was in the peer-review process), \citet{DrummondEtAl2018HD189733b} find that the same is true for HD 189733b. \citet{AgundezEtAl2014} instead used a full kinetical network in a pseudo-2D framework that was able to capture vertical and horizontal transport, and included photochemistry on the day side. They found that quenching homogenized the CO and CH$_4$ abundances at pressures between $\sim1$ and $\sim10^{-4}$ bars on HD 189733b. (At lower pressures, photochemical processes destroy CH$_4$ on the day side.) While these studies disagree on the relative importance of vertical versus horizontal quenching, all of them conclude that the abundances of CO and CH$_4$ should be homogeneous in the near-infrared photosphere, justifying our assumption. However, coupling a chemical relaxation scheme to their GCM, \citet{MendoncaEtAl2018_chemistry} found that on WASP-43b, the CH$_4$ abundances were only homogenized horizontally but not vertically. A similar behavior is seen in the HD 209458b case of \citet{AgundezEtAl2014} (though only for their nominal eddy diffusion profile). While there are several differences in the models and planet parameters used that could contribute to this different outcome, the perhaps most important factors are the hotter day side and the weak vertical mixing in both of these models. If horizontal transport dominates over vertical transport in setting the disequilibrium abundances, as these two papers find, whether abundances are homogenized only horizontally or vertically and horizontally depends on whether vertical quenching happens in the hottest region of the day side. Abundances thus are homogenized vertically and horizontally if in the hottest regions of the day side the vertical mixing time scale is shorter than the chemical time scale. If vertical mixing is very weak or the day side is too hot (leading to a shorter chemical time scale), this condition is not fulfilled and abundances are only homogenized horizontally but not vertically. In the case of \citet{AgundezEtAl2014}, the hotter day side compared to \citet{DrummondEtAl2018} at low pressures is largely due to their assumption of a thermal inversion on the day side of HD 209458b. \citet{MendoncaEtAl2018_chemistry} look at WASP-43b, for which \citet{KatariaEtAl2015} also found a larger day-night contrast and hotter day side compared to HD 209458b due to its shorter orbital period and larger gravity. Our study looks at HD 189733b, which has a significantly lower zero-albedo equilibrium temperature than HD 209458b and WASP-43b. Since the chemical time scale dramatically increases with decreasing temperature, it is likely that abundances are homogenized horizontally and vertically on this planet, and both published studies looking at HD 189733b confirm this \citep{AgundezEtAl2014,DrummondEtAl2018HD189733b}. With the focus of these previous studies being on chemistry, most of them did not self-consistently calculate the radiative transfer, and thus were not able to quantify the effect of the changed opacities on the temperature structure in their model. In their GCM, \citet{CooperShowman2006} used a Newtonian cooling scheme, in which the temperature relaxes towards a prescribed temperature profile at each point in the atmosphere. Like typical kinetical networks, \citet{AgundezEtAl2014} used a prescribed background pressure-temperature profile (in this case derived from a GCM assuming equlibrium chemistry). \citet{MendoncaEtAl2018_chemistry} use a double-grey radiative transfer scheme. Only \citet{DrummondEtAl2018} and \citet{DrummondEtAl2018HD189733b}, whose GCM uses state-of-the-art radiative transfer with correlated k-coefficients, included the effect of the changed opacities on the temperature structure in a GCM. Our approach complements previous studies by focusing on the radiative transfer rather than on chemistry. The photospheric disequilibrium CH$_4$ and CO abundances found in coupled chemistry-circulation models such as \citet{DrummondEtAl2018} strongly depend on the temperature profile in transition region between equilibrium chemistry and disequilibrium chemistry ($\sim 1$ to 10 bars)\citep{MosesEtAl2011,VenotEtAl2014}. However, the temperature profile in this region depends on the initial temperature profile used in the GCM \citep{AmundsenEtAl2016}, assumptions about the interior heat flux \citep[e.g.,][]{GuillotShowman2002,BurrowsEtAl2003,FortneyEtAl2008} % and dissipation in deep layers \citep[e.g.,][]{GuillotShowman2002,TremblinEtAl2017, KomacekYoudin2017} . Furthermore, typical GCM simulations only resolve dynamic mixing through the large-scale circulation. If unresolved sub-grid-scale turbulence is relevant near the quench level \citep[e.g., ][]{Menou2019}, it might further alter the quenched abundances. Uncertainties in reaction rates can also affect the resulting abundances \citep[e.g., ][]{VisscherMoses2011}. Overall, although it is well established that the CH4/CO ratio should be reasonably homogenized throughout the photosphere, many unconstrained factors will determine the actual value of this ratio. Studying how the thermal response and the resulting emission spectra and phase curves depends on the CH$_4$/CO ratio, as is our goal, thus adds to the overall understanding of the effects of disequilibrium chemistry. \begin{figure*} \plotone{temp_composition.pdf} \caption{Left: horizontal temperature maps of the reference simulation (equilibrium chemistry) at pressures of 1 mbar, 30 mbars and 1 bar, respectively. The substellar point (0 $^\circ$ longitude, 0 $^\circ$ latitude) is at the center of each panel. The arrows represent the velocities of the horizontal component of the wind. Right: horizontal maps of the abundance ratio of methane to carbon monoxide assuming chemical equilibrium. Methane dominated regions are magenta while carbon-monoxide dominated regions are green. The contours are evenly spaced in log space and mark log$_{10}$(CH$_4$/CO)=(-4,-3,-2,-1,0,1, 2), respectively. The zero contour (corresponding to CH$_4$/CO=1) is thickened. \label{fig:temp_composition}} \end{figure*} We choose to focus on HD 189733b. Due to its cooler temperature compared to HD 209458b and WASP-43b, disequilibrium effects are expected to be stronger on this planet. The latter two planets are hot enough that in simulations assuming equilibrium chemistry, CO is the dominant carbon species at all longitudes and latitudes \citep{ShowmanEtAl2009,DrummondEtAl2018}. In contrast, on the cooler HD 189733b, in chemical equilibrium one expects that the day side is dominated by CO, while the night side is dominated by CH$_4$ (see Figure \ref{fig:temp_composition}). HD 189733b thus occupies an interesting point in the parameter space: Regardless of the quenched CH$_4$ and CO abundances, including disequilibrium chemistry will change which of the two species is dominant on about half of the planet. The radiative effects of disequilibrium chemistry are thus expected to play a larger role than on hotter planets. In addition, unlike for HD 209458b, for HD 189733b there exist phase curve observations in multiple Spitzer bandpasses \citep{KnutsonEtAl2007,KnutsonEtAl2009_HD189,AgolEtAl2010,KnutsonEtAl2012}, providing stronger observational constraints. | We have included the radiative effect of transport-induced disequilibrium CH$_4$, CO and H$_2$O abundances in a GCM to study the effect on the atmospheric structure, phase curves and emission spectra. We have assumed that the ratio of CH$_4$ to CO is constant throughout the entire simulation (an assumption that is expected to be well fulfilled at pressures between $\sim10^{-4}$ bars and 1 bar) and treat the CH$_4$/CO ratio as free parameter. The water abundance is updated accordingly, such that the total number of oxygen atoms is preserved. It is important to include this change in the water abundance, as in equilibrium chemistry the water abundance varies by a factor of $\sim2$ between the CO-dominated day side and the CH$_4$-dominated night side. Assuming vertical and horizontal quenching, however, the water abundance is expected to be homogenized between day and night side. We ran simulations of hot Jupiter HD 189733b with eight different quenched CH$_4$/CO ratios. We find that in the CO dominated case, which is the case favored by chemical kinetics models, the temperature changes locally by up to 150 K, with cooler temperatures compared to equilibrium chemistry on the day side and warmer temperatures on part of the night side. In the less plausible CH$_4$ dominated case, the addition of greenhouse gases leads to hotter temperatures everywhere at pressures higher than a few tens of mbars. When comparing the predicted phase curves from GCM simulations including disequilibrium CH$_4$ and CO abundances to phase curves obtained from an equilibrium chemistry GCM simulation that has been post-processed assuming quenched CH$_4$/CO abundances, we find that the eastward offset of the phase curve maximum can differ by up to $10^\circ$. We thus conclude that it is important to self-consistently include the effect of disequilibrium abundances of CH$_4$ and CO on the opacities in GCMs rather than including disequilibrium abundances only in the post-processing while continuing to use opacities based on equilibrium chemistry abundances in the GCM. This is in contrast to \citet{DrummondEtAl2018} who find in their study of HD 209458b that the effect of radiative feedback of disequilibrium abundances on the temperature and wind fields is only $\sim 1 \%$, but agrees with their more recent findings for HD 189733b \citep{DrummondEtAl2018HD189733b}. These seemingly conflicting results can be understood when considering the difference in the equilibrium temperatures of the planets: On the hotter HD 209458b, the CH$_4$ abundance remains low compared to the CO and H$_2$O abundances in both equilibrium and disequilibrium chemistry even on the night side. In contrast, on the cooler HD 189733b, the night side is cool enough to be dominated by CH$_4$ in equilibrium chemistry. Including disequilibrium chemistry thus changes the dominant carbon species on half of the planet, resulting in much larger changes. In addition, in the regions where disequilibrium chemistry changes the dominant carbon species, the water abundance is also altered by a factor of $\sim2$, further contributing to the effect on temperatures. Furthermore, we show that disequilibrium CH$_4$ and CO abundances have only a small effect on the Spitzer 4.5 $\mu$m phase curve despite CO having a prominent absorption band within this wavelength band. This is because the change in opacity due to CO is offset by a change in water opacity in the opposite direction. In wavelength regions dominated by CH$_4$ opacity, including the Spitzer 3.6 $\mu$m and 8 $\mu$m bands, the phase curve amplitude decreases significantly, resulting in a much worse fit to the observed Spitzer 3.6 $\mu$m and 8 $\mu$m phase curves. We thus conclude that disequilibrium carbon chemistry cannot explain the observed low night side fluxes in the 4.5 $\mu$m band, in contrast to the interpretation of \citet{KnutsonEtAl2012}. Other effects, for example night side clouds, must be responsible for the observed shape of the phase curve. While disequilibrium chemistry does not explain existing observations of HD 189733b, it may be detectable on other hot Jupiters with a similar equilibrium temperature. Therefore, we examine the effect of disequilibrium carbon chemistry on emission spectra and simulated JWST observations. We find that in the expected CO dominated regime, spectral regions dominated by methane absorption bands are most suitable to observe disequilibrium abundances. Assuming that the spectral signatures of disequilibrium carbon chemistry are not obscured by clouds or other effects not considered in our model, it will be possible to distinguish between different quenched ratios with JWST in both secondary eclipse and phase curve observations. | 18 | 8 | 1808.02011 |
1808 | 1808.00014_arXiv.txt | We present two new \nustar\ observations of the narrow line Seyfert 1 (NLS1) galaxy Mrk~766 and give constraints on the two scenarios previously proposed to explain its spectrum and that of other NLS1s: relativistic reflection and partial covering. The \nustar\ spectra show a strong hard ($>15$\,keV) X-ray excess, while simultaneous soft X-ray coverage of one of the observations provided by \xmm\ constrains the ionised absorption in the source. The pure reflection model requires a black hole of high spin ($a>0.92$) viewed at a moderate inclination ($i=46_{-4}^{+1}\,^\circ$). The pure partial covering model requires extreme parameters: the cut-off of the primary continuum is very low ($22_{-5}^{+7}$\,keV) in one observation and the intrinsic X-ray emission must provide a large fraction (75\%) of the bolometric luminosity. Allowing a hybrid model with both partial covering and reflection provides more reasonable absorption parameters and relaxes the constraints on reflection parameters. The fractional variability reduces around the iron~K band and at high energies including the Compton hump, suggesting that the reflected emission is less variable than the continuum. | A common feature of the X-ray spectra of many non-jetted active galactic nuclei (AGN) is the hard excess, a strong increase in flux above $\sim15$~keV. This was first detected in stacked spectra from {\it Ginga} \citep{Pounds90} and measured for individual sources with {\it BeppoSAX} \citep{perola02}. Prior to the launch of \nustar , detailed measurements had only been made in a handful of AGN, using \swift-BAT or the \suzaku\ PIN detector, which were interpreted either as evidence of Compton thick absorption \citep{Risaliti09,Turner09} or the Compton hump of reflected emission \citep{Walton10}. X-ray reflection \citep{lightman88,George91} occurs when the primary X-ray source, known as the corona, illuminates the accretion disc or other relatively cold material such as the torus. This illumination triggers the emission of fluorescent lines at low energies and is scattered into a `Compton hump' at high energies. When the reflection spectrum originates from the parts of the accretion disc close to the innermost stable circular orbit (ISCO) of the black hole, the narrow features are blurred out by relativistic effects \citep{Fabian89, Laor91}. Since the launch of \nustar\ \citep{Harrison13}, the Compton hump has been more definitively detected in many AGN \citep{Risaliti13,Walton14,Marinucci14,Brenneman14,Parker14c,Balokovic15,Kara15}. The sensitivity of \nustar\ at high energies means that it can be used to differentiate between reflection and absorption models for the hard excess \citep[e.g.][]{vasudevan14}. One object which has had both these processes proposed to explain its spectrum is Mrk~766. Mrk~766 is a nearby ($z=0.013$) narrow line Seyfert 1 (NLS1) galaxy. NLS1s are thought to be rapidly accreting ($\dot{m}\sim0.01-1$), relatively low mass (typically $M_{\rm BH}\sim10^{6}-10^{7}M_{\odot}$) AGN, and are distinguished by narrow optical Balmer lines, weak [O \textsc{iii}] and strong Fe \textsc{ii} emission \citep[see review by][]{Komossa08}. In the X-ray band, NLS1s are spectrally soft, and are thus easily detected by low energy instruments. They frequently show complex, rapid variability and non-trivial spectral shapes, so are of great interest for study. The supermassive black hole in the nucleus of Mrk~766 has a mass of 1--6$\times10^6 M_\odot$ \citep{Bentz09,Bentz10} and the host is a barred spiral galaxy. Spectrally, the evidence for a relativistically-broadened iron line in Mrk~766 is tentative. A broad line was claimed with {\it ASCA} by \citet{Nandra97asca}. However, later analysis of a more sensitive \xmm\ spectrum by \citet{Pounds03} showed that the line profile could instead be described by ionized reflection alone, with no need for relativistic blurring. Based on \xmm\ and \suzaku\ observations of Mrk~766, \citet{Miller07} and \citet{Turner07} proposed a model where the bulk of the spectral variability is due to variations in multiple complex (partially-covering, ionized) absorbing zones. A recent re-analysis of the archival \xmm\ data by \citet{Liebmann14} showed that the spectra and variability could be well described by a composite model, containing both partial-covering absorption and relativistic reflection. More robust evidence for the presence of relativistic reflection in Mrk~766 comes from the detection of a reverberation lag \citep{Emman11,DeMarco13}, thought to be caused by the time delay induced in the reflected signal due to the light travel time from the corona to the disc. \citet{Emman11} found almost identical reverberation lags in Mrk~766 and MCG--6-30-15, the first source in which a broad iron line was discovered \citep{Tanaka95}. Mrk~766 is included in the sample of objects studied by \citet{Emman14}, who found that by modelling the time lag spectra they could precisely determine the mass ($M_{\rm BH}=1.6_{-1.2}^{+1.4}\times 10^{6}\,M_{\sun}$) and constrain other physical parameters (e.g. the dimensionless spin, $a>0.56$). The discovery of iron K lags in some sources \citep[e.g.][]{zoghbi12,Kara13,kara16}, which have so far only been explained by invoking relativistic reflection, have reinforced the interpretation of these high-frequency time lags as originating from reverberation close to the black hole. However, iron K reverberation lags have not yet been detected in Mrk~766 \citep{kara16}. In this paper, we present the results of recent \nustar\ observations of Mrk~766, where we examine the hard X-ray spectrum using the sensitivity and high-energy spectral resolution of \nustar\ to enable us to constrain the different physical models for the hard excess. The observations and methods of data reduction are presented in Section~\ref{section_datareduction}; results of the analysis are given in Section~\ref{sec:results}; these results are discussed in Section~\ref{section_discussion}; and conclusions are made in Section~\ref{section_conclusions}. | \label{section_conclusions} We have presented two new observations of Mrk~766 taken by \nustar, providing a detailed view of its hard X-ray spectrum. With simultaneous coverage in soft X-rays by \xmm\ or \swift, we are able to exploit the high spectral resolution of \xmm-RGS to take account of warm absorption and so produce better constraints on the broadband spectrum. We can model the spectrum with reflection or partial covering to generate the iron~K feature and Compton hump. In the reflection model, the system has a high spin black hole ($a>0.92$) viewed at intermediate inclination ($i=46_{-2}^{+1}\,^\circ$). The best-fitting partial covering model is questionable as it requires a very low cut-off energy and the intrinsic X-ray luminosity is high compared to the bolometric luminosity. A hybrid model including reflection and partial covering allows less extreme conditions for each component of the model. | 18 | 8 | 1808.00014 |
1808 | 1808.02820_arXiv.txt | {The second data release (DR2) from the European Space Agency mission \mbox{{\it Gaia}} took place on April 2018. DR2 included photometry for more than $1.3 \cdot 10^9$~sources in the three bands \GG, \GBP, and \GRP. Even though the \mbox{{\it Gaia}} DR2 photometry is very precise, there are currently three alternative definitions of the sensitivity curves that show significative differences.} {The aim of this paper is to improve the quality of the input calibration data to produce new compatible definitions of the \GG, \GBP, and \GRP\ bands and to identify the reasons for the discrepancies between previous definitions.} {We have searched the HST archive for STIS spectra with G430L+G750L data obtained with wide apertures and combined them with the CALSPEC library to produce a high quality SED library of 122 stars with a broad range of colors, including three very red stars. This library defines new sensitivity curves for \GG, \GBP, and \GRP\ using a functional analytical formalism.} {The new sensitivity curves are significantly better than the two previous attempts we use as a reference, REV and WEI. For \GG\ we confirm the existence of a systematic bias in magnitude and correct a color term present in REV. For \GBP\ we confirm the need to define two magnitude ranges with different sensitivity curves and measure the cut between them at $\GGp = 10.87$~mag with a significant increase in precision. The new curves also fit the data better than either REV or WEI. For \GRP, our new sensitivity curve fits the STIS spectra better and the differences with previous attempts reside in a systematic effect between ground-based and HST spectral libraries. Additional evidence from color-color diagrams indicate that the new sensitivity curve is more accurate. Nevertheless, there is still room for improvement in the accuracy of the sensitivity curves because of the current dearth of good-quality red calibrators: adding more to the sample should be a priority before \mbox{{\it Gaia}} data release 3 takes place.} {} | $\,\!$\indent The second data release (DR2) of the {\it Gaia} mission \citep{Prusetal16} took place in April 2018 \citep{Browetal18}. {\it Gaia} DR2 includes photometry for over $1.3 \cdot 10^9$~sources in the three bands \GG, \GBP, and \GRP. The \GG\ photometry was extracted using PSF fitting and has formal uncertainties under 1 mmag for most stars brighter than $\GG = 16$. The \GBP\ and \GRP\ magnitudes were obtained through aperture photometry and have larger formal uncertainties, of the order of a few mmag for stars brighter than $\GG = 15$, larger than those for \GG\ because they are measured just once per transit as opposed to the nine measurements per transit for \GG. {\it Gaia} DR2 constitutes the first all-sky multiband high-precision deep optical photometric survey and as such is likely to be considered an astronomical milestone that will be used as a reference and a calibration source for many studies. However, a high formal precision does not necessarily imply a high accuracy, as one needs to read the ``fine print'' of how the photometry was obtained to determine the applicability of the published magnitudes and uncertainties. For example, the different nature of the photometry (PSF vs. aperture) leads to \GG\ being more accurate than \GBP\ and \GRP\ in crowded (where multiple sources can be included more easily) or nebular (where the background model can be biased) regions \citep{Evanetal18}. Another accuracy issue, which is the main subject of this paper, is the comparison between the observed magnitudes ($m_{\rm phot}$) and the synthetic ones ($m_{\rm synth}$) derived from the spectral energy distributions (SEDs) of the sources. In this paper we use Vega magnitudes, as customary for Gaia photometry, and the reader is referred to the Appendix to see how we define the relevant quantities, including the zero points (ZPs) that are one of our results. An accurate definition of the sensitivity curves is especially important for the {\it Gaia} photometric system because the three passbands are very broad: \GG\ has an effective width\footnote{There are different ways to measure the center and width of a passband (see e.g. section 5.1 in \citealt{synphot}) but that does not affect the argument here.} around 2900~\AA\ (centered around 6400~\AA) while those of \GBP\ and \GRP\ are close to 1900~\AA\ (with that of \GRP\ slightly larger) and centered around 5100~\AA\ and 7800~\AA, respectively. For comparison, the widths of the Johnson $UBV$ system are 500-700~\AA. When doing broad-band photometry of sources with very different intrinsic SEDs and degrees of extinction one needs to integrate each SED to calculate the magnitudes, as a simple evaluation of the flux at a central wavelength does not work. Already for the Johnson $UBV$ system the classical $Q$ approximation to calculate extinction \citep{JohnMorg53} breaks down in many practical situations (see Appendix~B in \citealt{MaizBarb18}) due to the non-linearity of the extinction trajectories in the $U-B$ + $B-V$ plane induced by this effect. For {\it Gaia} photometry such extinction non-linearities in a color-color plane are even larger and more dependent on the precise definition of the passbands, as we will show later on in this paper. The first sensitivity curves for the three {\it Gaia} passbands were published by \citet{Jordetal10} but those were based on pre-launch data that were later modified. In one of the {\it Gaia} DR1 calibration papers, \citet{Carretal16} noted that if one used those curves a color term was present in the \GG\ photometry and \citet{Maiz17} published a modified sensitivity curve that was able to correct for it. An independent analysis by \citet{Weiletal18} found a very similar sensitivity curve. The {\it Gaia} DR1 photometry was affected by a contamination effect caused by water freezing in some optical elements \citep{Prusetal16} so the {\it Gaia} DR2 \GG\ data were expected to be characterized by a different sensitivity curve. Indeed, \citet{Evanetal18} published not a set but two sets of sensitivity curves for the \GG, \GBP, and \GRP\ photometry in the second data release: one they called DR2 and another one they called REV (for revised, that set was considered the preferred one by the authors). Later on, \citet{Weil18} provided a third set that differed from the other two, and that we refer to as WEI. All three by now published sets of {\it Gaia} DR2 passbands are based on the same set of calibration sources, the ''Spectrophotometric Standard Stars'' (SPSS, \citealt{Pancetal12, Altaetal15}) with the only exception of the WEI \GBP\ passbands, which were derived using CALSPEC \citep{Bohletal14,Bohletal17}. The SPSS set of calibration spectra is being constructed for the calibration of {\it Gaia}, and a first subset of 92 stars was made available for deriving {\it Gaia} DR2 passbands. Other spectral libraries, namely CALSPEC, the Next Generation Spectral Library (NGSL, \citealt{HeapLind07}), and the library by \citet{Strietal05} have been used for validation purposes by \citet{Evanetal18} and \citet{Weil18}. The DR2, REV, and WEI results are similar (but not identical) for $G$: the three sensitivity curves show few differences and agree in their overall shape. They all require a correction for a drift in the zero point of the observed \GG\ photometry, as discussed further in section~3 (the corrected \GG\ magnitude is denoted \GGc\ here). The results are more different for \GBP, as \citet{Weil18} found that bright and faint stars follow different sensitivity curves and that there is a jump of 20~mmag in zero point between the two. The WEI sensitivity curves for \GBP\ for the bright and the faint stars both differ strongly in their overall shape from the DR2 and REV passbands. For \GRP\ the differences in shape between the DR2 and REV sensitivity curves and the WEI curve are large, too, although resulting in a small improvement for the SPSS calibration spectra only. Furthermore, \citet{Weil18} noted that, while the WEI sensitivity curve for \GRP\ improves the results for the SPSS, Stritzinger, and NGSL libraries, it yields a worse result than the REV passband for the CALSPEC spectra. \citet{Weil18} also compared synthetic color-color relationships with observed relationships to test the consistency of a set of sensitivity curves for the three different {\it Gaia} passbands. This consistency test showed that the REV set of sensitivity curves fails to reproduce the observed color-color relationships. On the other hand, the WEI sensitivity curves have been designed not only to result in a good reproduction of the observed photometry for each passband individually, but also to reproduce the observed color-color relationships even outside the range in colors covered by the calibration spectra. In this work, we first compile a new set of calibration spectra based on high-quality HST/STIS optical observations. This set of calibration spectra extends the CALSPEC set significantly, both in number and coverage of different spectral types. In section~2 we describe this set of calibration spectra in detail. We then use the new set of calibration spectra to derive refined sensitivity curves (that will be referred to as MAW from our last names) for all {\it Gaia} passbands in section~3. Finally, in section~4 we demonstrate that the calibration data of this work is superior in quality to existing sets of calibration spectra. We also compare synthetic color-color relationships with observed ones, both for main sequence stars and for highly reddened stars, demonstrating that the sensitivity curves derived in this work with a new set of calibration spectra are the most accurate available to date. | 18 | 8 | 1808.02820 |
|
1808 | 1808.04771_arXiv.txt | \iftoggle{endauthorlist}{ }{ \newpage } We present the first Advanced LIGO and Advanced Virgo search for \compact{} binary systems with component masses between \MinSearchMass{} -- \MaxSearchMass{} using data taken between \ObsOneStart{} and \ObsOneEnd{}. We find no viable gravitational wave candidates. Our null result constrains the coalescence rate of monochromatic (delta function) distributions of non-spinning (\MinSearchMass{}, \MinSearchMass{}) \compact{} binaries to be less than \ZeroPointTwoRate{} and the coalescence rate of a similar distribution of (\MaxSearchMass{}, \MaxSearchMass{}) \compact{} binaries to be less than \OneRate{} (at $90\%$ confidence). Neither black holes nor neutron stars are expected to form below $\sim 1 M_\odot$ through conventional stellar evolution, though it has been proposed that similarly low mass black holes could be formed primordially through density fluctuations in the early universe. Under a particular primordial black hole binary formation scenario, we constrain monochromatic primordial black hole populations of 0.2 $M_\odot$ to be less than \DMConstraintPointTwo{} of the total dark matter density and monochromatic populations of 1.0 $M_\odot$ to be less than \DMConstraintOne{} of the dark matter density. The latter strengthens the presently placed bounds from micro-lensing surveys of MAssive Compact Halo Objects (MACHOs) provided by the MACHO and EROS collaborations. | Introduction} The era of gravitational wave astronomy began with the observation of the binary black hole merger GW150914~\cite{Abbott:2016blz}. Since then, four additional binary black hole mergers~\cite{Abbott:2016nmj, Abbott:2017vtc, Abbott:2017gyy, Abbott:2017oio} and one binary neutron star merger~\cite{TheLIGOScientific:2017qsa} have been announced as of November 2017. Thus far, Advanced LIGO and Advanced Virgo searches have targeted binary systems with total masses from \SearchedMassRange{}~\cite{TheLIGOScientific:2016pea, Abbott:2017iws}, but the LIGO and Virgo detectors are also sensitive to \compact{} binaries with components below 1 $M_\odot$ if the compactness (mass to radius ratio) is close to that of a black hole. White dwarf binaries, while often formed with components below one solar mass, are not sufficiently compact to be a LIGO/Virgo gravitational wave source. Neutron stars or black holes are sufficiently compact as would be other exotic compact objects. Previous gravitational wave searches for sub-solar mass \compact{} binaries used data from initial LIGO observations from \STwoFourDates{}~\cite{Abbott:2005pf, Abbott:2007xi}. Advanced LIGO~\cite{TheLIGOScientific:2014jea} presently surveys a volume of space approximately \VolumeIncrease{} times larger than the previous search for sub-solar mass \compact{} objects therefore improving the chances of detecting such a binary \VolumeIncrease{}-fold. In conventional stellar evolution models, the lightest \compact{} objects are formed when stellar remnants exceed $\sim 1.4 \Msun{}$, the Chandrasekhar mass limit~\cite{Chandrasekhar:1935zz, Chandrasekhar:1931ih}. Beyond the Chandrasekhar mass limit, electron degeneracy pressure can no longer prevent the gravitational collapse of a white dwarf. The lightest remnants that exceed the Chandrasekhar mass limit form neutron stars ~\cite{glendenning2012compact}. When even the neutron degeneracy pressure cannot prevent collapse, heavier stellar remnants will collapse to black holes. Some equations of state predict that neutron stars remain stable down to $\sim 0.1 \Msun$~\cite{Potekhin:2013qqa}; there is no widely accepted model for forming neutron stars below $\sim 1 M_\odot$, though a recent measurement does not exclude the possibility of $0.92 M_\odot$ neutron star~\cite{Martinez:2017jbp}. Observationally, black holes appear to have a minimum mass of $\sim 5\Msun$ with a gap between the observed neutron star masses and black hole masses~\cite{2010ApJ...725.1918O,2011ApJ...741..103F,2012ApJ...757...36K}. Detecting \compact{} objects below one solar mass could challenge our ideas about stellar evolution or possibly hint at new, unconventional formation scenarios. Beyond conventional stellar evolution, one of the most prolific black hole formation models posits that primordial black holes (PBHs) could have formed in the early universe through the collapse of highly over-dense regions~\cite{Zeldovich1967,Hawking1971,Carr1974,Meszaros1974,Chapline1975}. It has been suggested that PBHs could constitute a fraction of the missing dark matter~\cite{Chapline1975}, though this scenario has been constrained~\cite{Carr2016}. LIGO's detections have revived interest in black hole formation mechanisms and, in particular, the formation of primordial black holes (PBHs)~\cite{Bird:2016dcv, Sasaki:2016jop}. Though there are proposals on how to distinguish a primordial black hole distribution from an astrophysical one~\citep{Kovetz:2016kpi}, disentangling them is challenging when the populations overlap in mass. Hence, detection of sub-solar mass \compact{} objects would provide the cleanest signature for determining primordial formation. Still, recent proposals for non-baryonic dark matter models can produce sub-solar mass black holes either by allowing a lower Chandrasekhar mass in the dark sector~\cite{Shandera:2018xkn}, or by triggering neutron stars to collapse into $\sim 1 M_\odot$ black holes ~\cite{Kouvaris:2018wnh}. This letter describes a gravitational wave search for \compact{} binary systems with component masses between \MinSearchMass{} and \MaxSearchMass{} using data from Advanced LIGO's first observing run . No viable gravitational wave candidates were identified. We briefly describe the data analyzed and the anticipated sensitivity to sub-solar mass \compact{} objects, as well as the search that was conducted, which led to the null result. We then describe how the null result constrains the merger rate of sub-solar mass binaries in the nearby universe. We consider the merger rate constraints in the context of binary merger rate estimates most recently given by Sasaki et al~\cite{Sasaki:2016jop} thereby constraining the fraction of dark matter density made up of PBHs between \MinSearchMass{} and \MaxSearchMass{}. Finally, we conclude with a discussion of future work. | We presented the first Advanced LIGO and Advanced Virgo search for \compact{} binary mergers with components below 1 $M_\odot$. No viable gravitational wave candidates were found. Therefore, we were able to constrain the binary merger rate for monochromatic mass functions spanning from 0.2 -- 1.0 $M_\odot$. Using a well-studied model from the literature~\cite{Nakamura:1997sm,Ioka1998,Sasaki:2016jop}, we constrained the abundance of primordial black holes as a fraction of the total dark matter for each of our nine monochromatic mass functions considered. This work was only the first step in constraints by LIGO on new physics involving sub-solar mass \compact{} objects. The constraints presented in Fig. \ref{fig:rates} (and consequently those that arise from the model of binary formation we consider shown in Fig. \ref{fig:fraction}) may not apply if the \compact{} binary components have non-negligible spin since the waveforms used for signal recovery were generated only for non-spinning binaries. Future work may either quantify the extent to which the present search could detect spinning components, or expand the template bank to include systems with spin. Third, we should consider more general distributions of primordial black hole masses; extended mass functions allow for the possibility of unequal mass binaries, and the effect of this imbalance on the predicted merger rate has not been quantified. We also stress that our present results do not rule out an extended mass function that peaks below 0.2 $M_\odot$ and extends all the way to LIGO's currently detected systems at or above 30 $M_\odot$. Each model would have to be explicitly checked by producing an expected binary merger rate density that could be integrated against Advanced LIGO and Advanced Virgo search results. The first two areas of future work are computational challenges. Lowering the minimum mass and including spin effects in the waveform models could easily increase the computational cost of searching for sub-solar mass \compact{} objects by an order of magnitude each, which would be beyond the capabilities of present LIGO data grid resources. Advanced LIGO and Advanced Virgo have not reached their final design sensitivities. The distance to which Advanced LIGO will be sensitive to the mergers of \compact{} binaries in this mass range should increase by a factor of three over the next several years~\cite{Aasi:2013wya}. Furthermore, at least a factor of ten more data will be available than what was analyzed in this work. These two facts combined imply that the merger rate constraint should improve by $\gtrapprox 2$ orders of magnitude in the coming years. | 18 | 8 | 1808.04771 |
1808 | 1808.09009_arXiv.txt | Starspots, plages, and activity cycles cause radial velocity variations that can either mimic planets or hide their existence. To verify the authenticity of newly discovered planets, observers may search for periodicity in spectroscopic activity indices such as Ca H \& K and H$\alpha$, then mask out any Doppler signals that match the activity period or its harmonics. However, not every spectrograph includes Ca H \& K, and redder activity indicators are needed for planet searches around low-mass stars. Here we show how new activity indicators can be identified by correlating spectral line depths with a well-known activity index. We apply our correlation methods to archival HARPS spectra of $\epsilon$~Eri and $\alpha$~Cen~B and use the results from both stars to generate a master list of activity-sensitive lines whose core fluxes are periodic at the star's rotation period. Our newly discovered activity indicators can in turn be used as benchmarks to extend the list of known activity-sensitive lines toward the infrared or UV. With recent improvements in spectrograph illumination stabilization, wavelength calibration, and telluric correction, stellar activity is now the biggest noise source in planet searches. Our suite of $> 40$ activity-sensitive lines is a first step toward allowing planet hunters to access all the information about spots, plages, and activity cycles contained in each spectrum. | Although NASA is investing tremendous resources in space missions designed to find habitable planets and detect biomarkers \citep{gardner04, ricker14}, the search for Earthlike planets orbiting Sunlike stars still requires ground-based radial velocity measurements to identify or confirm targets for space-based observations \citep[e.g.][]{mayor14}. Of the main radial velocity (RV) noise sources which obscure the signals of low-mass planets, stellar noise is one of the most intractable. Indeed, there is concern that stellar jitter may create an intrinsic RV precision floor near the 0.5-1~m~s$^{-1}$ precision of current-generation RV instruments, which would limit the minimum mass of a detectable planet in the habitable zone of $\alpha$~Cen~B, a K1 dwarf, to $2.5 M_{\oplus}$ \citep[e.g.][]{dumusque11}. For next-generation RV instruments expected to reach instrumental precisions of 10-20 cm/s, stellar jitter is widely anticipated to be the dominant source of RV noise \citep{pepe13, jurgenson16}. Yet stellar jitter is not white noise, but is instead structured in time: p-modes have timescales of minutes \citep{christensendaalsgard04,kjeldsen05}, granulation and supergranulation have timescales of hours to days \citep{derosa04,brandt08}, and magnetic activity shows timescales ranging from the stellar rotation period \citep{saar97} to multi-year cycles \citep{baliunas95}. Methods that either suppress stellar jitter \citep{pepe02, angladaescude12} or flag it using stellar activity signals and model it out \citep{aigrain12, dumusque12, dumusque14, haywood14, rajpaul15, giguere16, meunier17} allow us to push below the stellar noise floor toward true earth analogs. Although RV corrections for stellar activity are often based on only one or a few lines \citep[such as Ca$_{\rm II}$ H \& K, Na D, or H$\alpha$;][]{lovis11, barnes14, robertson13, robertson15}, hundreds or even thousands of spectral lines are present in most planet-search spectra. Including information from as many lines as possible is an important step toward improving activity diagnostics \citep{giguere16}. Here we present a new method of identifying all activity-sensitive lines in the HARPS (High Accuracy Radial velocity Planet Searcher) spectrograph's wavelength range of $3800 \AA - 6900 \AA$. The methods presented here are not particular to any wavelength range, so they could be replicated to find activity-sensitive lines in any wavelength range that contains a previously known stellar activity indicator. Our final product is a new list of $> 40$ activity-sensitive spectral lines whose depths vary with the same periodicity as $\log(R'_{HK})$, an index measuring activity-induced emission in the cores of the Ca$_{\rm II}$ H \& K absorption lines. Our new activity indicators can be measured using the same high-resolution spectra that are used to make precise RV measurements. The paper is organized as follows. In Section~\ref{finding}, we give a step-by-step explanation of how we derived our list of activity-sensitive lines. In Section~\ref{additional}, we describe the tests we performed to validate our criteria for labeling these lines as activity sensitive. Finally, in Section~\ref{discussion}, we discuss how our results might be used to improve RV techniques for finding exoplanets. | \label{discussion} We searched archival HARPS spectra of $\epsilon$~Eri and $\alpha$~Cen~B for optical spectral lines whose variation correlates with the Mt.\ Wilson S-index, a measure of chromospheric activity. After we visually identified lines of interest based on $|({\rm A-Q})|$ correlating with S-index, we measured three properties for each line of interest: core flux, half-depth range, and center of mass (see \S~\ref{AlphaCenB}). We found that core flux was the most robust quantitative indicator of stellar activity, as measured by the Kendall's $\tau$ correlation coefficient with S-index, for our visually identified lines of interest. In addition to our methods for identifying activity-sensitive lines, we also present a list of 39 lines whose core fluxes strongly correlate with S-index (see Table~\ref{lines}). We also note that Fe~$5250 \AA$, which Zeeman splits in sunspots, shows a strong correlation between half-depth range and S-index. Our list of activity-sensitive lines may be used to improve methods of disentangling the RV signals of exoplanets from RV signals of stellar activity. \cite{giguere16} note that their HH' method, an H$\alpha$-based variation on the FF' method \citep{aigrain12} for removing RV variations due to stellar activity from planet-search data, could be improved by incorporating more activity-sensitive lines. Similarly, the multiple activity indicators used to gauge stellar activity in \cite{rajpaul15} could be expanded to include activity-sensitive line depths. However, these methods of trying to separate components of RV time series will continue to perform poorly when an exoplanet period is similar to the stellar activity signal period. One potential solution is directly calculating convective blueshift using stellar effective temperatures and activity levels \citep{meunier17}. Another possibility is using photospheric activity indicators, not just chromospheric activity indicators, to improve the RV measurement process \citep{davis17}. Half-depth range, which we believe probes Zeeman splitting in some lines, may be a promising photospheric activity metric. Although the activity-sensitive lines we have identified are at optical wavelengths ($4376 \AA$ to H$\alpha$), the methods we used to find them are not particular to only one wavelength range. A similar procedure that unites visual inspection with automated line property measurements could be used to find stellar activity-sensitive lines in any wavelength range that contains a previously known stellar activity indicator. | 18 | 8 | 1808.09009 |
1808 | 1808.08525_arXiv.txt | { In order to study the state of gas in galaxies, diagrams of the relation of optical emission line fluxes are used allowing one to separate main ionization sources: young stars in the H\,II regions, active galactic nuclei, and shock waves. In the intermediate cases, when the contributions of radiation from OB stars and from shock waves mix, identification becomes uncertain, and the issue remains unresolved on what determines the observed state of the diffuse ionized gas (DIG) including the one on large distances from the galactic plane. Adding of an extra parameter --- the gas line-of-sight velocity dispersion --- to classical diagnostic diagrams helps to find a solution. In the present paper, we analyze the observed data for several nearby galaxies: for UGC\,10043 with the galactic wind, for the star forming dwarf galaxies VII\,Zw\,403 and Mrk\,35, for the galaxy Arp\,212 with a polar ring. The data on the velocity dispersion are obtained at the 6-m SAO RAS telescope with the Fabry-Perot scanning interferometer, the information on the relation of main emission-line fluxes --- from the published results of the integral-field spectroscopy (the CALIFA survey and the MPFS spectrograph). A positive correlation between the radial velocity dispersion and the contribution of shock excitation to gas ionization are observed. In particular, in studying Arp\,212, ``BPT--$\sigma$ relation'' allowed us to confirm the assumption on a direct collision of gaseous clouds on the inclined orbits with the main disk of the galaxy. | Diagrams of ratios of optical emission-line fluxes are widely used for diagnostics of gas-ionization sources in galaxies. In the classic work by \citet*{Bald1981} the two-dimension diagram of line fluxes of \OIII$\lambda5007$/\Hb\, and \linebreak\NII$\lambda6583$/\Ha\ was suggested for separation of objects with different ionization sources. The method became popular due to the use of measurements of lines bright in the visible range that are close in wavelengths and, consequently, with a weak dependence of their intensity ratio on the interstellar extinction. Later, this method was extended by adding the relations \SII/\Ha\footnote{Hereinafter, for short we will designate \OI$\lambda6300$\, as \OI, \OIII$\lambda5007$\, as \OIII, \NII$\lambda6583$\, as \NII, and \SII$\lambda6717$+\SII$\lambda6731$\, as \SII.} and \OI/\Ha~\citep{Veilleux1987,Kewl2001} as the second parameter. All the mentioned diagrams are frequently called in the literature the ``BPT diagrams'' after the authors of the method. Using them, it is possible to distinguish the regions, where the largest contribution to gas ionization is made by young massive stars (hereinafter, the H II type) and the regions of dominating hard ionizing radiation of the active galactic nucleus (AGN). At the same time, the regions ionized by shock waves, the asymptotic giant branch (AGB) stars or nuclei of galaxies of the LINER type mix in the diagrams \footnote{Low-Ionization Narrow Emission-line Region in which the shock ionization of gas can be associated both with a burst of star formation and with a weak nuclear activity.}. Various variants of the demarcation lines were suggested~\citep{Mon2006,Ho2014}, but it is often problematic to separate the contribution of ionizing sources with the soft spectrum. Addition of one more parameter -- the velocity dispersion of the ionized gas along the line of sight ($\sigma$) -- to the classic diagnostic diagrams allows one to escape uncertainty in the cases, when the increase of $\sigma$ indicates the increase of turbulent velocities of gas beyond the front of a shock wave. However, to estimate $\sigma$ reliably, the spectral resolution is necessary that is noticeably better than that usually required for measuring the fluxes of radial velocities of separate spectral lines. Thus, until recently, the dependence of the relation of line fluxes characterizing the shock ionization from $\sigma$ was rarely considered and mainly for the objects with $\sigma>100$--$200\km$ such as galaxies with intense star formation~\citep{Mon2006,Ho2014}. Such an approach has not previously been used to study the ionization of diffuse gas in dwarf galaxies, around separate star-forming regions, or at some distance from the plane of the galactic disk. There is a discussion about the sources of ionization of this diffuse ionized gas (DIG) in galaxies whose role is assigned to an old stellar population, leakage of Lyman quanta from H\,II regions, and also possibly to shock fronts caused by star formation processes \citep[see references and discussion in][]{Jone2017,Egorov2017}. The most effective methods for studying the extended low-brightness structures in galaxies are panoramic spectroscopy also called integral-field, or 3D. In a recent paper based on the results of the SDSS MaNGA survey by \citep{Zhang2017}, it was concluded that DIG is associated mainly with the evolved stellar population (AGB stars, etc.). At the same time, in Section 6.2 of the paper cited, it was noted that shock waves can be the cause of the observed increase in the flux ratio of the forbidden and Balmer lines. It is difficult to verify, since the spectral resolution of the MaNGA survey is about two times poorer than that required to be able see the effects of moderate shock waves (with a velocity of less than 500~\km) in the observed kinematics of the ionized gas. Unfortunately, most of the available observed data on spectrophotometry and kinematics of the gas of nearby galaxies are obtained with the spectral resolution $FWHM>5$~\AA~which corresponds to values greater than 100~\km\ in terms of radial velocity dispersion or greater than 230 \km\ in terms of the $FWHM$ in the \Ha line. Observations with such resolution are a compulsory compromise in the study of low-surface-brightness objects. In the SAMI survey of galaxies \citep{Ho2014} with the 3D spectroscopy at the 3.9-m Anglo-Australian Telescope (AAT), the ``line ratio--velocity dispersion'' diagrams were built for the galaxies with active star formation. A positive correlation of the ionized gas $\sigma$ with a characteristic emission lines ratios was noticed, which was interpreted as an increase of shock waves contribution with velocities of about 200-300 \km\ accompanying a burst of star formation. The spectral resolution of the SAMI survey is greater than that of MaNGA and equals $R\approx4500$. A significant limitation of these two most massive today 3D spectroscopy surveys of galaxies is rather low spatial resolution (more than 1 kpc). It is related to the fact that the field of view of integral field unit (IFU) is small and is about $15\arcsec$ in SAMI~\citep{SAMI2012} and $12\arcsec$--$32\arcsec$ in SDSS MaNGA~\citep{MANGA2015}. In these surveys, relatively distant ($z>0.01$) galaxies are studied. At the same time, the largest contribution to the kinematics of interstellar medium from motion due to supernovae and winds of young stars in star-forming regions is made on considerably smaller spatial scales (from tens to hundreds of parsec). Consequently, any observed manifestations of shock fronts in star-forming regions become unevident, when averaging over a scale of one kpc or more. The examples of decreasing of peak velocity dispersion of the ionized gas in dwarf galaxies with degradation of a spatial resolution are presented in~\cite{Mois2012}. \cite{Vasiliev2015} considered the same effect in simulations of multiple supernova explosions. Therefore, for observational studies of the relation between an ionization state of gas and dispersion of its radial velocities in galaxies without an active nucleus and with a moderate star formation rate, 3D spectroscopic data are required simultaneously with a considerably high spectral and spatial resolution. In this paper, we consider this relation for several nearby galaxies using a combination of two spectroscopic methods with similar spatial resolution and quite a large field of view. Velocity dispersion map are derived from the observations with a scanning Fabry-Perot interferometer (FPI) at the 6-m SAO RAS telescope. Information on the main emission lines ratios are taken from open data on the integral-field spectroscopy with low spectral resolution. In order to show the relation between the velocity dispersion and the lines ratios characterizing the ionization state, we use various methods through our paper: coloring in BPT diagrams, the ``$\sigma$--line-flux relation'' diagrams, and ``$\sigma$--distance from the H\,II/AGN demarcation line.'' \textit{As a general name for the dependencies under study, we use the term ``BPT relation--$\sigma$''}. Classical BPT--diagrams are two-dimensional plots, where the axes represent relations of line fluxes. The inclusion of the velocity dispersion in the analysis is equivalent to the transition to three-dimensional plots, where a coordinate axis $\sigma$ is added to each diagram. The more familiar two-dimensional plots given in our paper and in the above papers are a projection of the BPT--$\sigma$ common relation to the selected plane. \begin{table*}[] \caption{Characteristics of the galaxies under study and parameters of their observations with different methods} \label{tab_1} \begin{tabular}{cccccccccc} \hline Galaxy & $D$, & $M_B$ &\multicolumn{4}{c}{Integral-field spectroscopy} &\multicolumn{3}{c}{Scanning FPI} \\ & Mpc & &Instrument &$\Delta\lambda$, \AA &$\delta\lambda$, \AA &$\theta\, ''$ &$\Delta\lambda$ &$\delta\lambda$, \AA &$\theta\,''$ \\ \hline UGC\,10043 &34.9 &$-17.6$ &PPAK &3750--7500 &5--9 &2.7 &\NII &1.7 &1.5 \\ VII\,Zw\,403~(UGC\,6456) &4.34 &$-13.87$ &MPFS &4250--7200 &8 &2.0 &\Ha &0.8 &2.2 \\ Mrk\,35~(NCG\,3353) &15.6 &$-17.75$ &PPAK &3750--7500 &5--9 &2.7 &\Ha &0.8 &2.1 \\ Arp\,212~(NGC\,7625) &23.5 &$-18.9$ &PPAK &3750--7500 &5--9 &2.7 &\Ha &0.8 &2.7 \\ \hline \end{tabular} \end{table*} | For an observational study of the relation between turbulent motions of the ionized gas in nearby galaxies and the state of its ionization, it is required to have panoramic spectroscopy data together with a large field of view and quite high spectral resolution. Since it is necessary to observe the low-surface-brightness region with an angular resolution of about $1''$, then an optical telescope of a large (\mbox {$D>3$--$5$}~m) diameter is needed. All these requirements are implemented together probably only in the unique MUSE instrument at the 8-m VLT telescope~\citep{MUSE2010}. Our idea is to combine the ionization ionized gas velocity dispersion maps obtained in the observations with the scanning FPI and the panoramic spectrophotometry data for low-spectral-resolution galaxies. The observed line-of-sight velocity dispersion characterizing the turbulent motions of the ionized gas can be due to various causes such as virial motions in the galaxy's gravitational potential, the effect of expanding shells on the gas, or, more generally, energy injected into the interstellar medium by star-forming processes~\citep[see discussion and references in][]{Mois2015,Krum2016}. Various factors influence the value of line flux relations with different excitation mechanisms. Observational information fusion makes it possible, in certain cases, to draw unambiguous conclusions about contribution of shock waves to the gas ionization in low-surface-brightness regions. From the lines flux ratio in the conical nebula in UGC 10043 only, it cannot be definitely concluded what leads the growth of the relative intensity of the forbidden lines: ionization by shock waves from the central burst of star formation or the old stellar population of the thick disk, in which it is located. Additional information on the gas kinematics allows to say that there is a galactic wind. For Arp 212, our approach allowed to confirm the previously assumption in \cite{Mois2008} on the direct collision of gaseous clouds on inclined orbits with the main disk of the galaxy generating shock fronts. Thus, the use of the BPT--$\sigma$ diagram together with the classical diagnostic methods based on lines ratios helps us better understand of ionization of the galactic interstellar medium in each specific case. The only galaxy in which we did not find a correlation between $\sigma$ and characteristic line flux relations (or $\rho$ parameter)~is VII\,Zw\,403. The ongoing star-formation rate here is the lowest in our sample \citep[$\sim0.015\,M_{\odot}$\,yr$^{-1}$,][]{Loz2006}. Apparently, for this reason, the contribution of shock waves to gas ionization is practically invisible. We plan to conduct further expansion of the sample of the objects under study in two ways. The first is new observations with a high spectral resolution of galaxies, for which there are the CALIFA survey data already, with the scanning FPI. The second is the creation of images in the emission lines of galaxies, for which we already have maps of the velocity dispersion of the ionized gas. Here it is proposed to use a tunable--filter photometer, the first observations with which are already being conducted by our team\footnote{\url{https://www.sao.ru/Doc-en/Events/2017/Moiseev/moiseev_eng.html }}. | 18 | 8 | 1808.08525 |
1808 | 1808.03637_arXiv.txt | 1I/'Oumuamua is the first interstellar object observed passing through the Solar System. Understanding the nature of these objects will provide crucial information about the formation and evolution of planetary systems, and the chemodynamical evolution of the Galaxy as a whole. We obtained the galactic orbital parameters of this object, considering 8 different models for the Galaxy, and compared it to those of stars of different ages from the Geneva-Copenhagen Survey (GCS). Assuming that the galactic orbital evolution of this object is similar to that of stars, we applied a Bayesian analyses and used the distribution of stellar velocities, as a function of age, to obtain a probability density function for the age of 'Oumuamua. We considered two models for the age-velocity dispersion relation (AVR): the traditional power law, fitted using data from the GCS; and a model that implements a second power law for younger ages, which we fitted using a sample of 153 Open Clusters (OCs). We find that the slope of the AVR is smaller for OCs than it is for field stars. Using these AVRs, we constrained an age range of 0.01--1.87 Gyr for 'Oumuamua and characterized a most likely age ranging between 0.20--0.45 Gyr, depending on the model used for the AVR. We also estimated the intrinsic uncertainties of the method due to not knowing the exact value of the Solar motion and the particularities of 1I/'Oumuamua's ejection. | The discovery of the interstellar object 1I/2017 U1 ('Oumuamua) by the Pan-STARRS survey \citep{Chambers+2016} in October, 2017 \citep{MPC2017a, MPC2017b}, and confirmation of its interstellar origin \citep{delaFuenteMarcos+delaFuenteMarcos2017}, sets a new field in astrophysics: the study of materials coming directly from other stellar systems. Crucial information about the Sun comes from the analysis of asteroids, like its chemical abundances \citep{Palme1988} and age with precision unmatched by any methods in the literature \citep{Bouvier+Wadhwa2010}. By analysing unbound asteroidal objects (\designation), like 'Oumuamua, we may be able to access this kind of precise information for other stellar systems as well. \begin{figure*} \centering \includegraphics[scale=0.57]{Figures/orbits_final_edited.pdf} \caption{'Oumuamua's Galactic orbit integrated in the past for several orbits ($\approx 3.5$ Gyr) for the eight models described in Table \ref{tab::orbital_parameters}. Although the shape of the orbit varies from one potential to the other, the orbital parameters ($e$, $z_\mathrm{max}$, $R_\mathrm{min}$ and $R_\mathrm{max}$) do not significantly change between them (in comparison to the range of these parameters observed for stars in the Solar Neighbourhood). Changing $R_0$ and $v_0$ does not affect the shape of the orbit, but causes a displacement. The model parameter that has the most significant effect in the derived orbital properties is the velocity of the Sun.} \label{fig::orbits} \end{figure*} Current efforts are concentrated into determining how common these objects are, and if we can identify its original stellar system. \citet{PortegiesZwart+2017} take into account the volume sampled by the Pan-STARRS survey to estimate the number density of \designations \, and find the value of $7.0 \times 10^{14}$ pc$^{-3}$. They estimate a number of 2-12 encounters a year. Regarding the identification of its original system, \citet{Mamajek2017} finds that its spatial velocity is not compatible with any nearby star, while \citet{Gaidos+2017} suggest its origin to be the Carina and Columba associations (since its velocity differs by $\lesssim 2 \, \mathrm{km\,s}^{-1}$). \citet{Feng+Jones2017} go beyond the comparison of spatial velocities and integrate the orbits of 0.23 million local stars as well as that of 'Oumuamua and provide a list of 109 stars which had encounters closer than 5 pc (17 of which had encounters distances of less than 2 pc). There is also interest in knowing the characteristics of the body, as it may provide further clues to its origin and provide constraints regarding the formation and evolution of planetary systems \citep[e.g.][]{Raymond+2017}. 'Oumuamua display an unusual elongated shape, with an axial ratio that may reach 10:1 and a rotational period of about 8 hr \citep{Knight+2017, Bolin+2017, Jewitt+2017}. Optical and spectral analysis show no evidence of cometary activity \citep{Ye+2017} but observed deviations in its predicted path indicates that some material is being expelled from the body \citep{Micheli+2018}. Its color is redder than asteroids but consistent with Kuiper Belt Objects \citep{Masiero2017}. \citet{Fitzsimmons+2017} suggested that its surface may be as organically rich as outer solar system bodies due to it displaying a similar red and featureless spectrum. However, one fundamental property of this body, that may provide clues about its nature and origin, remains unknown: the age. In this work, we aim to derive the age of this body from constraints provided by its galactic orbit. The formation age of \designations\, is important for many reasons: (i) it helps narrowing down the list of potential stellar systems from which the object came from; (ii) combined with the chemical abundances observed in this and future similar objects, it will provide another tool to understand the chemical evolution of our galaxy; (iii) provide information about the dynamics and evolution of proto-planetary disks; (iv) and, considering that we can measure its $UVW$ velocity with higher precision than any star, these objects shed light to the whole dynamics of the Galaxy and the orbital heating mechanisms. In this paper, we characterize the age of 'Oumuamua from a probability density function derived from the relation between stellar orbital parameters and age. Our methodology is explained in Section \ref{sec::methodology}. In Section \ref{sec::results} we show the results of the orbital integration and the obtained age. Finally, Section \ref{sec::conclusions} contains our conclusions. | \label{sec::conclusions} Based on dynamical arguments, we suggest that \designations\, will follow the same evolution of galactic orbits as that of stars. It is then expected that the same relations between kinematical parameters and age that we observe for star will also be observed for \designations. In this work, we have compared the orbital parameters derived for 'Oumuamua with the orbital parameters of the GCS as a function of age, to obtain kinematical constraints for the age of this \designation. We performed orbital integrations using the module \texttt{galpy} considering 8 different models for the Galaxy (two potentials, two sets of $R_0$ and $v_0$ and two sets of $UVW_\odot$). The estimated orbital parameters range are $e = 0.043$--$0.076$, $z_\mathrm{max} = 27.05$--$27.91$ pc, $R_\mathrm{min} = 6.89$--$7.53$ kpc, $R_\mathrm{max} = 8.00$--$8.23$ kpc. We find that the orbit is then very close to planar and circular, which as noted by \citet{Gaidos+2017} is an indication of a lower age. However, in this work a more robust statistical approach was used to characterize the age of 'Oumuamua. By performing a Bayesian analysis using the relation between velocity distribution and age, we estimated the age pdf for 'Oumuamua from its $UVW$ velocities. We have taken into account the bias in the point estimators, considered different values for the Solar motion, analysed the possible bias caused by the unknown ejection velocity from parental stellar system and investigated two different models of AVR. Finally, we estimated the age of $t_\mathrm{kin} = 0.50^{+1.37}_{-0.27}$ Gyr (with a most likely age of 0.45 Gyr and an expected age of 0.66 Gyr), when considering the extrapolated AVR for younger stars; and $t_\mathrm{kin} = 0.18^{+0.70}_{-0.17}$ (with a most likely age of 0.20 Gyr and an expected age of 0.11 Gyr) when considering the AVR of OCs for younger stars. From this estimations we constrained the possible age range for 'Oumuamua to be between 0.01--1.87 Gyr (0--2.14 Gyr, when considering the predicted additional uncertainty from 'Oumuamua's ejection process). The obtained age shows that, indeed, the kinematics of the object predict that it comes from a young stellar system, which may further constrain the list of parental system candidates. \begin{figure} \centering \includegraphics[scale=0.53]{Figures/age_combination.pdf} \caption{Characterized kinematical ages as well as their lower and upper limits for the scenarios using the extrapolated AVR of field stars (top, red), using the AVR of OCs for younger stars (middle, blue) and the combination of both results (bottom, black). The faded lines include an addition 20\% uncertainty estimated from not knowing the exact conditions of 'Oumuamua ejection.} \label{fig::age_comparison} \end{figure} During our analysis, we considered two models for the AVR, which differs for younger stars. In one of them, we extrapolate the AVR of field stars older than $\approx$ 1 Gyr, as done in A\&R18. For the other model, we considered that young stars are part of OCs and used the DAML sample to fit the AVR for these ages. We find that the slope of the AVR for OCs is smaller than for field stars and that the correlation between the $U$ and $V$ velocities can be neglected in this case. In both cases, the upper limit for the age of 'Oumuamua is well constrained, indicating that the object in fact had a space velocity of a younger body, before its encounter with the Sun. As more objects like 'Oumuamua are observed (\citealp{PortegiesZwart+2017}, predict 2 to 12 encounters each year), we believe this method to derive their ages will further help the understanding of the characteristics of these objects and what they imply for proto-planetary disk formation. | 18 | 8 | 1808.03637 |
1808 | 1808.04367_arXiv.txt | We determine the thermal evolution of the intergalactic medium (IGM) over \(\SI{3}{\giga\yr}\) of cosmic time \(1.8<z<5.4\) by comparing measurements of the Ly\(\alpha\) forest power spectrum to a suite of \(\sim70\) hydrodynamical simulations. We conduct Bayesian inference of IGM thermal parameters using an end-to-end forward modeling framework whereby mock spectra generated from our simulation grid are used to build a custom emulator which interpolates the power spectrum between thermal grid points. The temperature at mean density $T_0$ rises steadily from \(T_0\sim\SI{6000}{K}\) at \(z=5.4\), peaks at \(\SI{14000}{K}\) for \(z\sim 3.4\), and decreases at lower redshift reaching \(T_0\sim\SI{7000}{K}\) by \(z\sim1.8\). This evolution provides conclusive evidence for photoionization heating resulting from the reionization of \HeII{}, as well as the subsequent cooling of the IGM due to the expansion of the Universe after all reionization events are complete. Our results are broadly consistent with previous measurements of thermal evolution based on a variety of approaches, but the sensitivity of the power spectrum, the combination of high precision BOSS measurements of large-scale modes (\(k\lesssim\SI{0.02}{\skm}\)) with our recent determination of the small-scale power, our large grid of models, and our careful statistical analysis allow us to break the well known degeneracy between the temperature at mean density \(T_0\) and the slope of the temperature density relation \(\gamma\) that has plagued previous analyses. At the highest redshifts \(z\geq5\) we infer lower temperatures than expected from the standard picture of IGM thermal evolution leaving little room for additional smoothing of the Ly\(\alpha\) forest by free streaming of warm dark matter. | The \ac{Lya} forest \citep{gunn1965DensityNeutralHydrogen,Lynds1971Absorption-LineSpectrum} is the premier probe of diffuse baryons in the \ac{IGM} at high redshifts. Its fluctuations can be accurately described in the current \(\Lambda \mathrm{CDM}\) framework --- on large scales it is mostly sensitive to cosmological parameters such as the amplitude of fluctuations \(\sigma_8\), primordial power spectrum slope \(n_s\), baryon density \(\Omega_b\), number of neutrino species \(N_\mathrm{eff}\), and the sum of neutrino masses \(\sum m_\nu\) \citep{Palanque-Delabrouille2015Neutrinomasses,Rossi2017ImpactMassive}. On small scales, however, it is sensitive to the thermal state of the \ac{IGM}\footnote{Note that the small scale \ac{Lya} forest is also sensitive to the nature of dark matter \citep[like \ac{WDM}, see e.g.][]{Seljak2006CanSterile,viel2013Warmdarkmatter} which will not be the focus of this work, but leads to important applications of our results.}. This alters the observed spectra via the Doppler broadening of absorption features due to thermal motions, as well as pressure smoothing of the gas (sometimes called "Jeans" broadening), which affects the underlying baryon distribution and depends on the integrated thermal history of the \ac{IGM} \citep{gnedin1998ProbingUniverseLyalpha, kulkarni2015CharacterizingPressureSmoothing,Onorbe2017Self-consistentModeling}. The thermal evolution is largely driven by impulsive heating from cosmic reionization events and the cooling process due to adiabatic expansion and Compton cooling \citep{McQuinn2016intergalactictemperature-density}. Current constraints imply that hydrogen and \HeI{} were reionized at \(z_\mathrm{reion, 50}= 6.4\textrm{--}9.0 (95\%)\)\footnote{\(z_\mathrm{reion, 50}\) is the redshift at which \(x_\mathHI= 0.50\).} (see \citealt{PlanckCollaboration2018Planck2018}). Additionally, measurements of the \ac{Lya} forest optical depth show a strong increase close to \(z=6\) leading to complete Gunn-Peterson absorption \citep{fan2006Constrainingevolutionionizing, becker2015Evidencepatchyhydrogen, Bosman2018Newconstraints, Eilers2018OpacityIntergalactic} which reveals that reionization ends at \(z\sim 6\). As for \HeII{} reionization, which is driven by the hard \(>4\) Ryd photons emitted by luminous quasars, observations of the \HeII{} Ly\(\alpha\) forest indicate \HeII{} had to be reionized by \(z=2.7\) \citep{worseck2011EndHeliumReionization} and possibly as early as \(z=3.4\) \citep{Worseck2016EarlyExtended}, but the limited number of observational constraints imply that the exact timing remains largely uncertain. While it is observationally tricky to obtain direct higher redshift constraints on \HeII{} reionization through \HeII{} \ac{Lya} absorption measurements because the \HeII{} forest becomes more and more opaque, we can indirectly constrain it via its imprint on the thermal state of the IGM. In the standard picture of thermal evolution cold \ac{IGM} gas (few \(\si{\kelvin}\)) is strongly heated during \HI{} and \HeI{} reionization (by few times \(\SI{1e4}{\kelvin}\)), subsequently cools and then experiences additional heating during \HeII{} reionization \citep{mcquinn2009HeIIReionization,Compostella2013imprintinhomogeneous,puchwein2015photoheatingintergalacticmedium,Greig2015impacttemperature,uptonsanderbeck2016Modelsthermalevolution,McQuinn2016intergalactictemperature-density,Onorbe2017Self-consistentModeling,Puchwein2018Consistentmodelling}. The combined effects of photoionization heating, Compton cooling, and adiabatic cooling due to the expansion of the universe lead to a net cooling of intergalactic gas between and after the reionization phases which has so far not been conclusively observed. Another consequence of these effects is a tight power law \ac{TDR} for most of the \ac{IGM} gas \citep{hui1997Equationstatephotoionized, puchwein2015photoheatingintergalacticmedium, McQuinn2016intergalactictemperature-density} about \(\Delta z\approx 1\text{--}2\) after the impulsive heating from a reionization event: \begin{equation} \label{eq:t-rho-relation} T(\Delta) = T_0 \Delta^{\gamma-1}, \end{equation} where \(\Delta = \rho / \bar{\rho}\) is the overdensity, \(T_0\) is temperature at mean density \(T_0\), and the index \(\gamma\) is expected to approach \(\sim 1.6\) long after the completion of reionization. As we recently summarized in \citet{Walther2017NewPrecision} (hereafter \citetalias{Walther2017NewPrecision}) there have been many attempts to measure the IGM's thermal parameters \citep{haehnelt1998Probingthermalhistory,schaye2000thermalhistoryintergalactic,bryan2000DistributionLyForest,ricotti2000EvolutionEffectiveEquation,mcdonald2001MeasurementTemperatureDensity,theuns2002Temperaturefluctuationsintergalactic,bolton2008Possibleevidenceinverted,Viel2009Cosmologicalastrophysical,lidz2010MeasurementSmallscale,becker2011DetectionextendedHe,rudie2012TemperatureDensityRelation,garzilli2012intergalacticmediumthermal,rorai2013NewMethodDirectly,viel2013Warmdarkmatter, boera2014thermalhistoryintergalactic,bolton2014consistentdeterminationtemperature,Lee2015IGMConstraints,Rorai2017Exploringthermal,Rorai2017Measurementsmall-scale,Irsic2017Newconstraints,yeche2017Constraintsneutrinomasses,Garzilli2017CutoffLyman-a,Rorai2018newmeasurement,DAloisio2018Largefluctuations,Hiss2017NewMeasurement} based on different statistical techniques which typically constrain the smoothness of the \ac{Lya} forest as a whole via some summary statistics (e.g. wavelet amplitudes, spectral curvature or the power spectrum) or decompose the forest into individual absorption lines by Voigt profile fitting. While there were some notable discrepancies between some of the older measurements (e.g. low values of \(\gamma\) inferred from the \citealt{bolton2008Possibleevidenceinverted} flux PDF or the high \(T_0\) measurements from the \citealt{lidz2010MeasurementSmallscale} wavelet analysis), more recent measurements appear to be in better agreement. For example, temperature determinations from the curvature statistic \citep{becker2011DetectionextendedHe,boera2014thermalhistoryintergalactic} agree fairly well with those determined from Voigt profile fitting \citep{bolton2014consistentdeterminationtemperature,Rorai2018newmeasurement,Hiss2017NewMeasurement}. Note however, that different techniques have distinct systematics and parameter degeneracies, that often complicate detailed comparisons. In this work, we use the power spectrum of the \ac{Lya} forest to obtain an accurate self-consistent measurement of IGM thermal evolution over a large redshift range from \(z=5.4\) to \(z=1.8\). The power spectrum exhibits a cutoff at small scales (high \(k\)) beyond which there is no structure left in the \ac{Lya} forest. The reason for this is both the smoothness in the baryon density resulting from the finite gas pressure (often called Jeans pressure smoothing) as well as thermal Doppler broadening. The great advantage of the power spectrum compared to other methods, is its sensitivity to structure on a multitude of scales. Specifically, whereas other methods like the curvature \citep{becker2011DetectionextendedHe} and wavelets \citep{lidz2010MeasurementSmallscale} provide only a small-scale measurement of spectral smoothness, the overall shape of the power spectrum for scales between \(\sim \SI{500}{\kilo\parsec}\) and \(\sim \SI{10}{\mega\parsec}\) as well as small-scale (high-k) cutoff provides additional constraining power that breaks degeneracies between different thermal parameters\footnote{Note that this property can also be used to break degeneracies with cosmological parameters, e.g. the nature of dark matter \citep{viel2013Warmdarkmatter, Irsic2017Newconstraints, armengaud2017Constrainingmasslight} or the mass of neutrinos \citep{Palanque-Delabrouille2015Neutrinomasses, yeche2017Constraintsneutrinomasses, baur2017ConstraintsLybackslash}.}. For this work we consider \(T_0\), \(\gamma\) and the pressure smoothing scale \(\lambda_P\) as thermal parameters and the mean transmission \(\bar{F}\) as a further astrophysical parameter. We additionally marginalize over the strength of \SiIII{} correlations and the resolution of the X-SHOOTER spectrograph (see \autoref{sec:priors} for more detailed information about our prior assumptions). Our analysis is based upon our recent high-precision measurements of the the small-scale (high wavenumber \(k\)) the \ac{Lya} forest flux power spectrum in \citetalias{Walther2017NewPrecision} as well as other recent measurements from different instruments (\citealt{palanque-delabrouille2013onedimensionalLy}, hereafter \citetalias{palanque-delabrouille2013onedimensionalLy}; \citealt{viel2013Warmdarkmatter}; and \citealt{Irsic2017Newconstraints}) combined with the new \ac{THERMAL} grid\footnote{see \url{thermal.joseonorbe.com}} of hydrodynamical simulations. We then perform inference by employing fast interpolation of our model power spectra and performing an MCMC analysis with a Gaussian likelihood. This paper is organized as follows. The measurements we used in this work are summarized in \autoref{sec:measurements}. In \autoref{sec:sims} we present our grid of hydrodynamical simulations. We use modified versions of our forward modeling, interpolation and inference tools from \citetalias{Walther2017NewPrecision}, which we present in \autoref{sec:method}, to measure the thermal state of the IGM at each redshift. In \autoref{sec:results} we present these results and compare them to measurements from the literature as well as thermal evolution models. Finally, we discuss the results in \autoref{sec:discussion} and conclude with \autoref{sec:conclusions}. | 18 | 8 | 1808.04367 |
|
1808 | 1808.06648_arXiv.txt | The Simons Observatory (SO) will make precision temperature and polarization measurements of the cosmic microwave background (CMB) using a series of telescopes which will cover angular scales between one arcminute and tens of degrees, contain over 60,000 detectors, and sample frequencies between 27 and 270 GHz. SO will consist of a six-meter-aperture telescope coupled to over 30,000 detectors along with an array of half-meter aperture refractive cameras, which together couple to an additional 30,000+ detectors. SO will measure fundamental cosmological parameters of our universe, find high redshift clusters via the Sunyaev-Zeldovich effect, constrain properties of neutrinos, and seek signatures of dark matter through gravitational lensing. In this paper we will present results of the simulations of the SO large aperture telescope receiver (LATR). We will show details of simulations performed to ensure the structural integrity and thermal performance of our receiver, as well as will present the results of finite element analyses (FEA) of designs for the structural support system. Additionally, a full thermal model for the LATR will be described. The model will be used to ensure we meet our design requirements. Finally, we will present the results of FEA used to identify the primary vibrational modes, and planned methods for suppressing these modes. Design solutions to each of these problems that have been informed by simulation will be presented. | The cosmic microwave background (CMB) has become one of the most powerful probes of the early universe. Measurements of temperature anisotropies on the level of approximately ten parts per million have brought cosmology into a precision era, and have placed tight constraints on the fundamental properties of the universe. Beyond temperature anisotropies, CMB polarization anisotropies not only enrich our understanding of our cosmological model, but could potentially provide clues to the very beginning of the universe via the detection (or non-detection) of primordial gravitational waves. A number of experiments have made and are continuing to refine measurements of the polarization anisotropy. However, these experiments are typically dedicated to a relatively restricted range of angular scales, e.g., large angular scales (tens of degrees) or high resolution/small angular scales (on the order of 1\,arcminute). To provide a complete picture of cosmology, both large and small angular scales are important. Ideally these measurements would be made from the same observing site so that the widest range of angular scales can be probed, at multiple frequencies, on the same regions of the sky. This is the goal of the Simons Observatory (SO). SO will field a 6-meter large aperture telescope (LAT) coupled to the large aperture telescope receiver. During initial deployment, seven of the planned thirteen optics tubes will be installed in the LATR, containing over 30,000 detectors. The LAT is designed with a large FOV capable of supporting a cryostat with up to 19 LATR-like optics tubes. To limit the development risk of the large SO cryostat, the LATR is designed to accommodate up to 13 optics tubes. We plan to deploy 7 optics tubes with 3 detector wafers in each for a total of roughly 35,000 detectors, primarily at 90/150~GHz in the initial SO deployment. We note that each optics tube could be upgraded to support 4.5 wafers for a ~50\% increase in the number of detectors per optics tube. With this upgrade and the deployment of 19 optics tubes, the LAT could support roughly 145,000 detectors at 90/150~GHz. SO will also have an array of half-meter large angular scale cameras coupled to an additional 30,000 detectors. The unique combination of telescopes in a single CMB observatory, which will be located in Chile\textsc{\char13}s Atacama Desert at an altitude of 5190~m, will allow us to sample a wide range of angular scales over a common survey area. In this paper we will provide an overview of the simulations done in support of the LATR design process. Finite element analysis (FEA) simulations give us critical feedback about the performance of our components, which we use to refine their design. In Sec.~\ref{sec:mech} we cover the suite of mechanical simulations that we did in support of our design using the Solidworks Simulation module.\footnote{Dassault Syst\`emes, 10, Rue Marcel Dassault, 78140, V\'elizy-Villacoublay, FRANCE, https://www.solidworks.com/} Then, in Sec.~\ref{sec:therm} we described our combined thermal model and detail the thermal gradient simulations performed with the COMSOL software.\footnote{COMSOL, Inc., 100 District Avenue, Burlington, MA 01803, USA, https://www.comsol.com} The challenges we faced when designing the LATR are unique in that we are trying build the largest ground-based CMB receiver to date. However, the solutions to these challenges we developed with the assistance of FEA simulations will provide a critical stepping stone for the next generation CMB experiments, in particular, CMB-S4\cite{Abazajian2016,Abitbol2017}. | Computer assisted design and FEA simulations have been key for the design of the Simons Observatory LATR. As we iterated through designs of various components, FEA provided critical feedback on the performance of those components. This feedback allowed us to evaluate the performance of our design, determine what could be improved, and design the next iteration. Through this process, we were able to converge on our current and final design for the LATR. Lastly, FEA provided the final validation of each component, ensuring for us that they would meet the performance specifications that we set. The result of this process is a design which is validated and slated for manufacture in the near future, and which can inform the design of a future CMB-S4 instrument.\cite{Abitbol2017}\cite{Abazajian2016} \appendix % | 18 | 8 | 1808.06648 |
1808 | 1808.09578_arXiv.txt | We present an analysis of the nuclear variability of $\sim28,000$ nearby ($z<0.15$) galaxies with Sloan Digital Sky Survey (SDSS) spectroscopy in Stripe 82. We construct light curves using difference imaging of SDSS \textit{g-}band images, which allows us to detect subtle variations in the central light output. We select variable AGN by assessing whether detected variability is well-described by a damped random walk model. We find 135 galaxies with AGN-like nuclear variability. While most of the variability-selected AGNs have narrow emission lines consistent with the presence of an AGN, a small fraction have narrow emission lines dominated by star formation. The star-forming systems with nuclear AGN-like variability tend to be low-mass ($M_{\ast}<10^{10}~M_{\odot}$), and may be AGNs missed by other selection techniques due to star formation dilution or low-metallicities. We explore the AGN fraction as a function of stellar mass, and find that the fraction of variable AGN increases with stellar mass, even after taking into account the fact that lower mass systems are fainter. There are several possible explanations for an observed decline in the fraction of variable AGN with decreasing stellar mass, including a drop in the supermassive black hole occupation fraction, a decrease in the ratio of black hole mass to galaxy stellar mass, or a change in the variability properties of lower-mass AGNs. We demonstrate that optical photometric variability is a promising avenue for detecting AGNs in low-mass, star formation-dominated galaxies, which has implications for the upcoming Large Synoptic Survey Telescope. | Supermassive black holes (BHs; $M_{\rm BH}\gtrsim 10^{5}~M_{\odot}$) are ubiquitous in the centers of galaxies with stellar masses $\gtrsim10^{10}M_{\odot}$. Less is known about the population of BHs in the centers of low-mass galaxies (here defined as galaxies with $M_{\ast}\lesssim10^{10}M_{\odot}$). However, the population of BHs in low-mass galaxies has the potential to place constraints on the mechanisms by which the seeds of present day BHs formed. The occupation fraction (i.e., the fraction of galaxies containing BHs) is expected to differ depending on the seed formation mechanisms at play (see reviews by \citealt{2012NatCo...3E1304G, 2014GReGr..46.1702N}). In particular, the occupation fraction is sensitive to the seed formation mechanism for galaxies with stellar masses $M_{\ast}<10^{10}~M_{\odot}$. Detecting BHs in low-mass galaxies poses unique observational challenges. A BH with $M_{\rm BH}=10^{5}~M_{\odot}$ has a gravitational sphere of influence of just a few pc, i.e., largely unresolvable outside the Local Group even with the \textit{Hubble Space Telescope}. In recent years, an increasing number of actively accreting massive black holes have been discovered in low-mass galaxies, particularly using X-ray emission and optical spectroscopic signatures (\citealt{2004ApJ...610..722G, 2007ApJ...670...92G, 2008AJ....136.1179B,Reines:2011fr, Reines:2013fj, 2014ApJ...787L..30R, 2014AJ....148..136M, 2015ApJ...798...38S, 2015ApJ...805...12L, 2016ApJ...831..203P, 2016ApJ...817...20M, 2018MNRAS.478.2576M}). However, emission line ratio diagrams commonly used to identify AGN were developed using samples of massive galaxies and do not necessarily apply for lower-mass, lower-metallicity systems \citep{2006MNRAS.371.1559G, 2018arXiv180510874B}. Moreover, at low galaxy stellar masses, star formation can dilute the AGN emission-line signal, resulting in AGN potentially being missed \citep{2015ApJ...811...26T}. While a sufficiently bright, hard X-ray point source can be a relatively unambiguous signature of an AGN, X-ray imaging down to the relevant luminosities is observationally expensive for large samples. Motivated by the potential for identifying systems missed by other selection techniques, we take the approach of searching for AGN via low-level optical variability. AGN are known to vary at all wavelengths, and searching for optical variability has been a rather prolific tool for identifying quasars (e.g., \citealt{2003AJ....125....1G, 2007AJ....134.2236S, 2010ApJ...714.1194S, 2011ApJ...728...26M, 2011A&A...530A.122P, 2014ApJ...782...37C, 2014AJ....147...12B}). The origin of the variability remains uncertain, but is potentially related to thermal instabilities in the accretion disk (e.g., \citealt{1984ARA&A..22..471R, 2009ApJ...698..895K}). Since the advent of the Sloan Digital Sky Survey (SDSS), great advances have been made in understanding the characteristics of AGN variability, as well as in the identification of AGN through the presence of variability. Optical variations at the 0.03 magnitude level have been observed in at least $90\%$ of quasars in the SDSS Stripe 82 \citep{2007AJ....134.2236S}. Variable AGN can also be distinguished from other variable objects (such as variable stars) based on their variability properties \citep{2011AJ....141...93B}. In the last several years, an increasing number of time domain surveys have come online; an incomplete list includes the Palomar Transient Factory \citep{2009PASP..121.1334R}, Zwicky Transient Facility \citep{2014htu..conf...27B}, Pan-STARRS \citep{2016arXiv161205560C, 2016arXiv161205243F}, La Silla-QUEST \citep{2015ApJ...810..164C} and GAIA \citep{2016A&A...595A...1G}. Additionally, the Large Synoptic Survey Telescope (LSST), which will image the entire visible sky every three nights, is scheduled to begin operations in 2023. Given the wealth of available and upcoming time domain data, it is important to assess the utility of variability for identifying AGN in low-mass galaxies. Using \textit{g-}band imaging data from the SDSS Stripe 82, we construct light curves for $\sim28,000$ galaxies in the NASA-Sloan Atlas, with stellar masses spanning from $10^{7}-10^{12}~M_{\odot}$. Section 2 describes the sample and data. Section 3 describes the difference imaging analysis and selection criteria for AGN candidates. In Section 4, we present the full sample of variability-selected AGN. In Section 5, we discuss the low-mass systems with AGN variability, and present an analysis of the detection limits and expected number of detections for the low-mass end. | We analyze the light curves of $\sim28,000$ nearby (z<0.15) galaxies in the SDSS Stripe 82 field for AGN-like nuclear variability. We concentrate our analysis on galaxies in NASA-Sloan Atlas, which have spectroscopic redshifts and stellar mass estimates. We construct light curves using difference imaging, which allows us to detect small variations superposed on top of the stellar light from the galaxy. We then determine whether variability is AGN-like by assessing how well individual light curves are described by a damped random walk. We find 135 galaxies with AGN-like variability. The variability-selected sample spans roughly four orders of magnitude in stellar mass. Above $M_{\ast}=10^{10}~M_{\odot}$ (100 galaxies), almost all of the variability-selected AGNs also have narrow emission line ratios in the AGN or composite regions of the BPT diagram. Below $M_{\ast}=10^{10}~M_{\odot}$, half of the galaxies have narrow emission-line ratios dominated by star formation, indicating they may be AGNs missed by other selection techniques due to star formation dilution or metallicity effects, which are expected to be more important at lower stellar masses. We stress that the low-mass galaxies with AGN-like variability falling in the star forming region of the BPT diagram should be regarded as candidates; we are in the process of obtaining higher spatial resolution follow-up spectroscopy which better isolates emission from the nucleus. High spatial resolution X-ray and/or radio observations will also be valuable in the effort to confirm the presence of AGNs in these systems. Using this sample, we study the fraction of variable AGN as a function of stellar mass, and find that even when accounting for magnitude bias, there is a decline in the fraction of variable AGN with decreasing stellar mass. This could be attributed to a lower BH occupation fraction for galaxies with $M_{\ast}<10^{10}~M_{\odot}$, a change in the $M_{BH}-M_{\ast}$ relation for low mass galaxies, or a change in the variability properties of BHs in low-mass galaxies. Figure~\ref{agnvmstar} shows that the active fraction is lower for galaxies with $M_{\ast}<10^{10}~M_{\odot}$ than for galaxies with $M_{\ast}>10^{10}~M_{\odot}$, even after taking into account the bias introduced by different apparent magnitude distributions for the high and low-mass subsamples. There are several possible explanations for a decline in the active fraction with stellar mass, which we discuss below. A relatively straightforward interpretation of our results is that there is a drop in the \textit{occupation fraction} (i.e., the fraction of galaxies with BHs) which leads naturally to a drop in the fraction which are active. While it is well established that the occupation fraction for massive galaxies is 100\% (or extremely close; \citealt{1998AJ....115.2285M}), it is unclear whether that holds for low-mass galaxies. There are several nearby galaxies for which there are stringent limits on the mass of a central BH based on dynamical modeling; for example, \cite{2001AJ....122.2469G} find that the upper limit on a central BH in M33 is just $1500~M_{\odot}$. There are ongoing efforts to constrain the BH occupation fraction at low masses using X-ray observations; current estimates place a firm lower limit of 20\% on the low-mass occupation fraction \citep{2015ApJ...799...98M}. Another possibility is that $M_{\rm BH}/M_{\ast}$ is lower for low-mass galaxies than for more massive galaxies. Massive galaxies show a fairly constant (with some scatter) ratio between the mass of the central BH and the mass of the galaxy ($\sim1/1000$ or 0.001). Recent works have shown that BHs in low-mass galaxies may be under-massive with respect to scaling relations between BH mass and bulge/galaxy stellar mass defined for more massive galaxies \citep{Greene:2008qy, 2015ApJ...813...82R, 2017ApJ...850..196B}. As mentioned above, we estimate BH masses for the 16 low-mass galaxies with broad H$\alpha$ emission. The BH masses range from $\log(M_{\rm BH}/M_{\odot})$ = 6.0 to 7.9 (with uncertainties of $\sim0.3$ dex). The median BH mass is $\log(M_{\rm BH}/M_{\odot})$ = 6.8. These BH masses correspond to $M_{\rm BH}/M_{\ast}$ ratios of 0.0002 to 0.007 (median $M_{\rm BH}/M_{\ast}$=0.0013). Thus, we may only be able to detect variable AGN (with this data set) in the low-mass galaxies that have unusually massive BHs (as compared to other galaxies of similar stellar mass). Finally, it is possible that AGNs in low-mass galaxies are intrinsically less variable, or are less variable in optical wavelengths. Monitoring campaigns of known low-mass galaxies with AGNs will be important for studying whether low-mass galaxies vary on similar timescales and/or with similar amplitudes as more massive AGNs. The objects with AGN-like variability falling in the star forming region of the BPT diagram will require multi-wavelength follow-up to search for additional evidence for AGNs in these systems. Our analysis will also be extended to additional existing repeat imaging surveys to search for more low-mass systems with AGN-like variability. The Large Synoptic Survey Telescope will be ideal for continuing searches for low-level variability on days-to-months timescales and should be sensitive to less massive BHs and/or BHs accreting at lower Eddington fractions. | 18 | 8 | 1808.09578 |
1808 | 1808.02801_arXiv.txt | {The halos of disk galaxies form a crucial connection between the galaxy disk and the intergalactic medium. Massive stars, H{\sc ii} regions, or dwarf galaxies located in the halos of galaxies are potential tracers of recent accretion and/or outflows of gas, and are additional contributors to the photon field and the gas phase metallicity.} {We investigate the nature and origin of a star-forming dwarf galaxy candidate located in the halo of the edge-on Virgo galaxy NGC~4634 with a projected distance of 1.4 kpc and a H$\alpha$ star formation rate of $\sim 4.7 \times 10^{-3} \text{M}_\odot \text{yr}^{-1}$ in order to increase our understanding of these disk-halo processes.} {With optical long-slit spectra we measured fluxes of optical nebula emission lines to derive the oxygen abundance 12~+~log(O/H) of an H{\sc ii} region in the disk of NGC~4634 and in the star-forming dwarf galaxy candidate. Abundances derived from optical long-slit data and from Hubble Space Telescope (HST) r-band data, H$\alpha$ data, Giant Metrewave Radio Telescope (GMRT) H{\sc i} data, and photometry of SDSS and GALEX data were used for further analysis. With additional probes of the luminosity--metallicity relation in the $B$-band from the H$\alpha$-luminosity, the H{\sc i} map, and the relative velocities, we are able to constrain a possible origin of the dwarf galaxy candidate.} {The high oxygen abundance (12 + log(O/H) $\approx$ 8.72) of the dwarf galaxy candidate leads to the conclusion that it was formed from pre-enriched material. Analysis of auxiliary data shows that the dwarf galaxy candidate is composed of material originating from NGC~4634. We cannot determine whether this material has been ejected tidally or through other processes, which makes the system highly interesting for follow up observations.} {} | Star formation in spiral galaxies usually takes place close to the midplane of the disk. Nevertheless, star formation high above the midplane is detected in the Milky Way and in external galaxies. Extraplanar H{\sc ii} regions indicate an exchange of matter between the disk and the halo, for example NGC~3628 and the Virgo cluster galaxy NGC~4522 \citep{steinetal2017} as well as NGC~4402 \citep{corteseetal2004}. Some extraplanar H{\sc ii} regions are related to the intracluster or intergalactic medium of interacting galaxies. These H{\sc ii} regions probably formed in tidal debris \citep{mendesetal2004, ryanweberetal2004} or by ram pressure stripping, for example in the Virgo cluster \citep{gerhardetal2002,oosterloovangorkom2005}. Understanding the physical processes in interacting systems, groups, or clusters is essential to understanding galaxy evolution. Additionally, dwarf galaxies are apparent in external galaxy groups \citep[e.g.,][]{mulleretal2015} and in the Local Group \citep[e.g.,][]{mateo1998}. Based on their optical appearance, dwarf galaxies are classified into different categories. Mainly, they are divided into star-forming dwarf galaxies with gas, which are interesting to study as analogs for galaxies in the early Universe \citep[e.g.,][]{vaduvescumccallricher2007}; intermediate types with little star formation and less gas \cite[e.g.,][]{sandagehoffman1991}; and non-star-forming dwarf galaxies with just a little gas, which appear like small elliptical galaxies \cite[e.g.,][]{binggeli94}. Furthermore, stellar/gaseous tidal debris might recondense within the halo of a merger and form a tidal dwarf galaxy (TDG), e.g., due to local gravitational instabilities in the gaseous component \citep{ducetal2000}. As they are born in situ, their stellar populations are young in comparison to those in normal dwarf galaxies. TDGs can serve as a laboratory for star formation studies in low-density environments \citep{boquienetal2009}. The building material of TDGs used to belong to a larger parent galaxy and is chemically pre-enriched compared to normal dwarf galaxies \citep{Duc2012}; i.e., in comparison to dwarf galaxies, TDGs do not follow the luminosity--metallicity relation (L-Z relation). Here, we investigate the case of a dwarf galaxy candidate detected by \citet{rossaetal2008} close to the edge-on spiral galaxy NGC~4634. It is located at a projected distance of 1.4 kpc from the NGC~4634 midplane. NGC~4634 is a Virgo Cluster member with enhanced star-forming activity and a neighboring galaxy of NGC~4633. NGC~4634 has a distance of 19.1 Mpc \citep{teerikorpi1992} and it is 2.6\arcmin $\times$ 0.7\arcmin \ in size. The dynamical galaxy mass is 2.7 $\times$ 10$^{10}$~M$_\odot$ \citep{rossaetal2008}. Its H{\sc ii} regions in the disk are widely spread throughout the midplane. This could be due to the interaction with NGC~4633 as the two galaxies are a close binary pair \citep{rotv46341993, rossadettmar2000A&A}. The halo shows diffuse X-ray emission \citep{tuellmann2006} and has a bright diffuse ionized gas component \citep{rossadettmar2000A&A}.\\ \indent NGC~4634 was part of a large sample of H$\alpha$ observations performed by \citet{rossadettmar2000A&A}, who searched systematically for extraplanar diffuse ionized gas (DIG). Among several extraplanar H{\sc ii} regions an extraordinary emission patch was detected in NGC~4634, which showed, in contrast to the other regions, a distributed stellar population in addition to higher H$\alpha$ emission. The authors argued that it could be either part of the DIG or a dwarf galaxy in the halo of NGC~4634. Further study by \citet{rossaetal2008}, which is based on the $R$-band morphology, suggests that this emission patch is a gas-rich dwarf irregular galaxy that has been tidally disrupted by the gravitational forces of NGC~4634. However, this study only used optical imaging, and thus the kinematical state and metallicity of the dwarf galaxy candidate have not been taken into account.\\ \indent We further analyze and investigate the origin of this peculiar object, which we refer to as the dwarf galaxy candidate of NGC~4634, by using various data. We present optical long-slit spectroscopic observations with ESO Faint Object Spectrograph and Camera (v.2) (EFOSC2) attached to the ESO 3.6~m Telescope, Hubble Space Telescope (HST) r-band and H$\alpha$ data, and Giant Metrewave Radio Telescope (GMRT) H{\sc i} data. We also use photometry of SDSS and GALEX data.\\ \normalsize \indent The paper is structured as follows. The data reduction and observational strategy of the different observations are presented in Section~2. The analysis and the results of the optical spectroscopy and the other data are described in Section~3. In Section~4 the results are discussed, and in Section~5 the paper is briefly summarized. \begin{figure} \centering \includegraphics[width=0.45\textwidth]{N4634Halpha2_neu2.pdf} \caption{Slit position on the HST H$\alpha$ image of NGC~4634. Indicated in red (see inset) are the regions of interest.} \label{NGC 4634} \end{figure} | In this paper we present multifrequency data of a dwarf galaxy candidate and its edge-on host NGC~4634. With optical long-slit spectra we calculate the oxygen abundances of one disk H{\sc ii} region and of a dwarf galaxy candidate's H{\sc ii} region. The comparable oxygen abundances hint at the material of the dwarf galaxy candidate being pre-enriched and that it originates from the disk of NGC~4634. This is confirmed by its location in a L-Z diagram, which shows that its metallicity is much higher than expected for a dwarf galaxy that evolved in relative isolation. The images of r-band, $B$-band, and the H$\alpha$-luminosity show that the dwarf galaxy candidate is star forming and approximately 550\,pc in diameter. The heliocentric velocities of the dwarf galaxy and its host NGC~4634 are similar and indicate that the galaxies are connected. The H{\sc i} data give the hint that a spur of NGC~4634 is located at the same position as the dwarf galaxy candidate. We furthermore determined the total dynamical mass of NGC~4634 to be 1.3~$\times 10^{10}$~M$_{\odot}$ from the GMRT data analyzed here.\\ \indent With the above discussion one possible origin is the tidal dwarf scenario where the dwarf galaxy candidate formed in the spur of NGC~4634 from disk material. However, the data presented in this paper cannot rule out other forces that separated the gas in the star-forming region from the disk of NGC~4634. In any case, it is clear that the halo of NGC~4634 contains a star-forming object of significant size that is formed from gas stripped from NGC~4634, which is an excellent laboratory for further studies in terms of galaxy evolution. | 18 | 8 | 1808.02801 |
1808 | 1808.02030_arXiv.txt | In hierarchical triple systems, the inner binary is perturbed by a distant companion. For large mutual inclinations, the Lidov-Kozai mechanism secularly excites large eccentricity and inclination oscillations of the inner binary. The maximal eccentricity attained, $e_{\rm max}$ is simply derived and widely used. However, for mildly hierarchical systems (i.e. the companion is relatively close and massive), non-secular perturbations affect the evolution. Here we account for fast non-secular variations and find new analytic formula for $e_{{\rm max}}$, in terms of the system's hierarchy level, correcting previous work and reproducing the orbital flip criteria. We find that $e_{{\rm max}}$ is generally enhanced, allowing closer encounters between the inner binary components, thus significantly changing their interaction and its final outcome. We then extend our approach to include additional relativistic and tidal forces. Using our results, we show that the merger time of gravitational-wave (GW) sources orbiting massive black-holes in galactic nuclei is enhanced compared with previous analysis accounting only for the secular regime. Consequently, this affects the distribution and rates of such GW sources in the relevant mild-hierarchy regime. We test and confirm our predictions with direct N-body and 2.5-level Post-Newtonian codes. Finally, we calculate the formation and disruption rates of hot-Jupiters (HJ) in planetary systems using a statistical approach, which incorporates our novel results for $e_{{\rm max}}$. We find that more HJ migrate from further out, but they are also tidally disrupted more frequently. Remarkably, the overall formation rate of HJs remains similar to that found in previous studies. Nevertheless, the different rates could manifest in different underlying distribution of observed warm-Jupiters. | Three body systems are ubiquitous in astrophysics and appear in a plethora of configurations and scales, from moons and asteroids of planets, to multiple stars and binary compact object around supermassive black holes. The general three body problem is notoriously non-integrable \citep{1892mnmc.book.....P}, but some special cases allow useful analytic approximations that shed light on their features \citep{ValtonenBook2006}. Hierarchical triples are systems where an inner binary is perturbed by a third distant companion. Observations of exo-planets \citep{2011PASP..123..412W,Kn14,2015ARA&A..53..409W}, multiple stars \citep{Rag10,2014AJ....147...86T} and compact objects in extreme orbital inclinations and eccentricities call for better understanding of such hierarchical multiple systems. The key parameter in the study of evolution of hierarchical systems is the maximal eccentricity $e_{{\rm max}}$ of the (inner) binary. Under appropriate conditions large eccentricities can be induced in the inner binaries of hierarchical triples through secular processes. These, in turn, can result in close encounters of the inner binary components during their pericentre approach, giving rise to a plethora of astrophysical phenonema, depending of the astrophysical set-up. Such processes include tidal dissipation in triple stars \citep{1998MNRAS.300..292K,2001ApJ...562.1012E}, Hot-Jupiter (HJ) formation \citep{wu03,2007ApJ...669.1298F, Naoz11nat,Anderson_ps,PetHJ,munoz16}, secular evolution of planets and satellites \citep{Per+09,T9,Grishin17,Grishin18}, triple stellar evolution \citep{2009ApJ...697.1048P,perets12, hamers13, michaely14,frewen16, T16, stephan2018}, gravitational-wave (GW) emission and mergers \citep{wen03, 2012ApJ...757...27A,ant14,AMK14,sil17,LiuLai17, LL18,Randall_LK2,Randall_LK1,FK18,hvrr}, tidal disruption events \citep{fraglei18a,fraglei18b}, direct collisions and type Ia supernovae \citep{katz_sne} etc. The main approach in studying the long-term evolution of hierarchical triples is through a perturbative method. In hierarchical systems, the interaction potential is expanded in multipoles in the (small) ratio of the inner to outer separations, and than double-orbit-averaged (DA) on both orbits \citep[and references therein]{Kozai62,Lidov62,Naoz2016review} . The resulting leading DA quadrupole term is integrable and the system admits an exact analytic solution \citep{2007CeMDA..98...67K}. Lidov-Kozai (LK) oscillations occur if the mutual inclination is in the range of the well known critical values $i_{c}=\arccos\pm\sqrt{3/5}=39.2^{\circ}, 140.8^{\circ}$. During the LK cycle, the maximal eccentricity attained is \begin{equation} e_{{\rm max}}^{{\rm DA}}=\sqrt{1-\frac{5}{3}\cos^{2}i_{0}}, \label{eq:emax} \end{equation} where $i_{0}$ is the initial mutual inclination, if the initial eccentricity $e_{0}\ll1$ is low. Eq. (\ref{eq:emax}) can be derived using the conservation of the specific $\hat{z}$ component of the inner binary's angular momentum, $j_{z}=\sqrt{1-e^{2}}\cos i$ (in the limit where the outer angular momentum dominates, i.e. the test particle limit), where $e$ is the inner binary eccentricity. The typical (secular) timescale for change in the orbital elements is \citep{2007CeMDA..98...67K, antognini15} \begin{equation} \tau_{{\rm sec}}\approx\frac{1}{2\pi}\frac{m_{{\rm tot}}}{m_{{\rm out}}}\frac{P_{{\rm out}}^{2}}{P_{{\rm in}}}(1-e_{{\rm out}}^{2})^{3/2}, \label{eq:tsec} \end{equation} where $m_{{\rm out}}$ is the mass of the outer companion, $m_{{\rm tot}}=m_{{\rm out}}+m_{{\rm bin}}$ is the total mass in the system, $m_{{\rm bin}}$ is the mass of the inner binary, $e_{{\rm out}}$ is the outer eccentricity, $P_{{\rm in}}$ and $P_{{\rm out}}$ are the inner and outer orbital periods, respectively. The DA approximation neglects any osculating fluctuations of the orbital elements on timescales $t\ll \tau_{{\rm sec}}$. However, in mildly hierarchical systems, such shorter-term effects change the evolution of the triple, and can induce larger eccentricites than predicted by the DA approach, as first shown by \citet{2012ApJ...757...27A}, while keeping an overall ``quasi-secular'' evolution (Lidov-Kozai cycles) very similar to that expected in the DA regime. Accounting for the quasi-secular regime can be important for a wide variety of systems at all scales \citep{cuk04,2012ApJ...757...27A,katz_sne,ant14,AMK14,Grishin17}. The rapid oscillations identified near the maximal eccentricity have been considered in \citet{ant14,AMK14}, and recently, \citet{Katz16} have shown that the orbital elements can be decomposed into averaged and fluctuating parts, and computed the additional corrections due to single-averaged (SA) potential, providing consistent results with direct N-body integrations. When the eccentricity (pericenter) is large (small), additional short-range forces (e.g. tides, general-relativistic (GR) precession or tidal and rotational rotational bulges) could affect and constrain the maximal eccentricity attained. \citet{LML15} used conservation of the total potential energy and $j_{z}$ to find the maximal eccentricity. For large enough strengths of the extra forces, the eccentricity excitations can be suppressed \citep{LML15}. Here we calculate the maximal eccentricity in the quadrupole order level of approximation and test particle limit, taking into account the additional SA potential, the osculating oscillations of $j_{z}$ (and consequently in $e$) and the additional extra forces. Relaxing these limitations is discussion in sec. \ref{sec5}. We show that contrary to quenching due to short range forces, the maximal eccentricity is enhanced due to the dominating effect of fluctuations in $j_{z}$. The enhancement may be orders of magnitude larger than the widely used $e_{{\rm max}}^{{\rm DA}}$ (Eq. \ref{eq:emax}), and even unconstrained, depending the level of the hierarchy and the initial inclination. This, in turn have consequences for the mildly-hierarchical triples in all scales. Here we explore these effects and discuss their implications for two test cases - production of GW-sources near MBHs and the formation and evolution of HJs. Our paper is organized as follows: In sec. \ref{sec2} we review basic LK mechanism and its coupling to additional extra forces. In sec. \ref{sec3} we derive the new formula for the maximal eccentricity in the quasi-secular CDA regime, and compare and validate our results with N-body integrations. In sec. \ref{sec4} We extend our analysis to include extra forces. We apply our results to find the GW merger time for Black-Hole binaries in the Galactic Centre (sec. \ref{41}), and then compare the changes in the rate of HJ formation with the recent analytical models (sec. \ref{42}). Finally, in sec. \ref{sec5} we discuss the limitations of our model and summarize in sec. \ref{sec6}. | \label{sec6} In this paper we studied the effects of short-term perturbations on a mildly hierarchical triple system, together with already studied non-Keplerian perturbations (e.g. general relativity and tides). We focused on the maximum eccentricity, a key parameter that determines the result of many short-range interactions and subsequent evolution of the system. Our result can be summarized as follows: \begin{enumerate} \item The strength of the perturbations and typical corrections are encapsulated in the hierarchy strength (single-averaged, SA) parameter $\epsilon_{{\rm SA}}$ (Eq. \ref{eq:epssa}). \item The critical inclinations for the onset of the Lidov-Kozai mechanism change according to Eq. (\ref{eq:inc_crit}), and the overall maximal eccentricity is increased according to Eqns. (\ref{eq:emaxsa}), (\ref{eq:ecorr}) and (\ref{eq:de3}). The new formula is robust, reproduces the orbital flip criteria, is in good agreement with N-body integrations, and corrects previous work, which has overestimated the eccentricity fluctuations. The main advantage of our calculation is retaining the secular approach, allowing for an efficient computaional approach and better analytic understanding. The double-averaging approximation is not breaking down, but rather is corrected for, such that lesser hierarchies are correctly accounted for. \item When general relativistic effects are included, they tend to add extra precession and quench the secular Lidov-Kozai eccentricity excitations. We find the maximal eccentriciy with additional general relativistic precession in Eq. (\ref{eq:jmin_e0}) and the conditions for orbital flip in Eq. (\ref{eq:flipgr}). We compare to N-body integrations which include 2.5PN effects and find that our formulae underestimate the maximal eccentricity in some cases, but overall it is in good agreement. We manifest our results by finding the merger time for black-hole binaries in the Galactic Centre and argue that the rate or black-hole mergers due to emission of gravitational waves should be larger when accounting for non-secular effects. In addition, we found a regime where the maximal eccentricity is unconstrained, direct collisions and/or eccentric mergers of binary black holes are possible, similar to direct collisions of white-dwarfs found by \citet{katz_sne}. \item We apply our results to hot-Jupiter formation rates. We include tidal effects and find the maximal eccentricity in this case in Eq. (\ref{eq:emaxtide}). We incorporate our new maximal eccentricity in a recent analytical model, and find that the total migration rate and the disruption rate are increased, but the rate of Hot-Jupiter formation is unchanged. Nevertheless, the rate for Warm-Jupiter migration can increase and the underlying observed distributions of migrating warm Jupiters and their properties are altered, namely Warm-Jupiters could migrate from further out separations and achieve larger eccentricities. \end{enumerate} | 18 | 8 | 1808.02030 |
1808 | 1808.00729_arXiv.txt | {% Radio detection of inclined air showers is currently receiving great attention. To exploit the potential, a suitable event reconstruction needs to be developed. A crucial step in this direction is the development of a model for the lateral distribution of the radio signals, which in the case of inclined air showers exhibits asymmetries due to ``early-late'' effects in addition to the usual asymmetries from the superposition of charge-excess and geomagnetic emission. We present a model which corrects for all asymmetries and successfully describes the lateral distribution of the energy fluence with a rotationally symmetric function. This gives access to the radiation energy as a measure of the energy of the cosmic-ray primary, and is also sensitive to the depth of the shower maximum. } | \label{intro} Due to the superposition of geomagnetic and charge-excess emission, the distribution of the energy fluence of the radio emission from extensive air showers is asymmetric on the ground \cite{HuegePLREP}. Lateral distribution functions (LDFs) have to take into account this asymmetry by either a two-dimensional description \cite{NellesLDF}, by correcting for the asymmetry induced by the charge-excess contribution \cite{KostuninLDF}, or by treating the two contributions individually \cite{Glaser:2018byo}. In case of inclined air showers, additional asymmetries arise from ``early-late effects'', i.e., the fact that the emission above the shower axis propagates longer through the atmosphere than the emission below the shower axis. In this article, we first present methods to correct for the early-late asymmetry as well as the charge-excess-induced asymmetry in the energy-fluence footprints of inclined air showers. We then propose a rotationally symmetric LDF that, when fit to the symmetrized energy fluences, allows precise determination of the cosmic-ray energy for inclined air showers. | 18 | 8 | 1808.00729 |
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1808 | 1808.06613_arXiv.txt | ALMA has found multiple dust gaps and rings in a number of protoplanetary disks in continuum emission at millimeter wavelengths. The origin of such structures is in debate. Recently, we documented how one super-Earth planet can open multiple (up to five) dust gaps in a disk with low viscosity ($\alpha\lesssim10^{-4}$). In this paper, we examine how the positions, depths, and total number of gaps opened by one planet depend on input parameters, and apply our results to real systems. Gap locations (equivalently, spacings) are the easiest metric to use when making comparisons between theory and observations, as positions can be robustly measured. We fit the locations of gaps empirically as functions of planet mass and disk aspect ratio. We find that the locations of the double gaps in HL Tau and TW Hya, and of all three gaps in HD 163296, are consistent with being opened by a sub-Saturn mass planet. This scenario predicts the locations of other gaps in HL Tau and TW Hya, some of which appear consistent with current observations. We also show how the Rossby wave instability may develop at the edges of several gaps and result in multiple dusty vortices, all caused by one planet. A planet as low in mass as Mars may produce multiple dust gaps in the terrestrial planet forming region. | \label{sec:intro} The Atacama Large Millimeter Array (ALMA) has discovered multiple gaps and rings in a number of protoplanetary disks in dust continuum emission at millimeter (mm) wavelengths at $\sim$10--100 AU. Examples include the disk around HL Tau \citep{brogan15}, TW Hya \citep{andrews16, tsukagoshi16}, HD 163296 \citep{isella16hd163296}, AS 209 \citep{fedele18}, AA Tau \citep{loomis17}, Elias 24 \citep{cieza17, cox17, dipierro18}, GY 91 \citep{sheehan18}, and V1094 Sco \citep{ansdell18, vanterwisga18}. The origins of these gaps are being debated. Several hypotheses have been put forward, including disk-planet interaction \citep[e.g.,][]{dong15gap, dipierro15hltau, pinilla15twoplanets, jin16}, zonal flows \citep{pinilla12dusttrapping}, secular gravitational instability \citep[e.g.,][]{takahashi14}, self-induced dust pile-ups \citep{gonzalez15}, radially variable magnetic disk winds \citep{suriano17, suriano18}, and dust sintering and evolution at snowlines \citep{zhang15, okuzumi16}. In nearly all mechanisms, perturbations in gas are followed by radial drift and concentration of mm-sized dust \citep[e.g.,][]{weidenschilling77, birnstiel10, zhu12}. In planet-disk interaction scenarios, even super-Earth planets may open dust gaps \citep[e.g.,][]{zhu14votices,rosotti16, dipierro16, dipierro17}. At first glance, multiple gaps seem to imply multiple planets, with one planet embedded within each gap. Recently, however, it has been appreciated that a single planet can generate multiple gaps. Using gas+dust simulations, \citet{dong17doublegap} showed that one super-Earth planet (with mass between Earth and Neptune) can produce up to five dust gaps in a nearly inviscid disk, and that these gaps would be detectable by ALMA (see also \citealt{muto10, duffell12, zhu13, bae17, chen18, ricci18}). In particular, a pair of closely spaced narrow gaps (a ``double gap'') sandwiches the planet's orbit. This structure is created by the launching, shocking, and dissipation of two primary density waves, as shown in \citet[see also \citealt{rafikov02, rafikov02migration}]{goodman01}. Additional gaps interior to the planet's orbit may also be present, possibly opened by the dissipation of secondary and tertiary density waves excited by the planet \citep{bae18theory,bae18simulation}. Motivated by the boom of multi-gap structures discovered by ALMA, and the mounting evidence of extremely low levels of turbulence at the disk midplane (from direct gas line observations, e.g., \citealt{flaherty15, flaherty17, teague18twhya}, and from evidence for dust settling, e.g., \citealt{pinte16,stephens17}), we investigate here what properties of the planet and disk can be inferred from real-life gap observations. We study systematically how gap properties depend on planet and disk parameters, and propose generic guidelines connecting ALMA observations with disk-planet interaction models. We introduce our models in \S\ref{sec:simulation}, and present our main parameter survey of gap properties in \S\ref{sec:results}. These results are applied to three actual disks in \S\ref{sec:applications}, discussed in \S\ref{sec:discussions}, and summarized in \S\ref{sec:summary}. | \label{sec:discussions} \subsection{The Observed Double Gap}\label{sec:doublegap} To explain the double gap in the HL Tau and TW Hya disks, two mechanisms have been proposed: dust sintering and evolution at snowlines \citep{zhang15, okuzumi16}, and planet-disk interactions along the lines of our models. The former needs the snowlines of two relevant volatiles to be rather close (separation/radius ratio $\sim17\%$), while the latter posits a sub-Saturn planet at dozens of AU. While accurate assessments of snowline locations are needed to test the former hypothesis, the slow variation of disk temperature with radius may be difficult to make work. Recently, an eccentric double gap with $e\sim0.1$ was tentatively discovered in the MWC 758 disk \citep{dong18mwc758}. A non-zero eccentricity may be compatible with a planet on an eccentric orbit, but is not natural to the snowline mechanism, which has no explicit azimuthal dependence. \subsection{Low Viscosity and the Rossby Wave Instability}\label{sec:lowalpha} A low viscosity ($\alpha\lesssim10^{-4}$) at the disk midplane (where $\sim$mm-sized dust particles are expected to settle) is required for a planet with $\mplanet\lesssim\mth$ to open multiple gaps. Figure~\ref{fig:alpha} compares two runs of Model H003MP02 with $\alpha=5\times10^{-5}$ (left; standard $\alpha$) and $3\times10^{-3}$ (right). No dust gaps are opened in the latter, while five gaps are opened in the former. \begin{figure*} \begin{center} \includegraphics[trim=0 0 0 0, clip,width=0.9\textwidth,angle=0]{plot_sigma_alpha.pdf} \end{center} \figcaption{Two runs of Model H004MP02 at 1600 orbits with $\alpha=5\times10^{-5}$ (left) and $\alpha=3\times10^{-3}$ (right). A low viscosity is required for a sub-thermal mass planet to open multiple dust gaps. See \S\ref{sec:lowalpha} for details. \label{fig:alpha}} \end{figure*} Evidence for little-to-no turbulence at the disk midplane at tens of AUs is accumulating both theoretically \citep[e.g.,][]{perezbecker11td, bai13ad, bai15, schlaufman18} and observationally \citep[e.g.,][see also the discussion in \citealt{fung17}, final section]{pinte16, flaherty17}. If $\alpha\lesssim10^{-4}$ is common, narrow dust gaps should mostly appear in pairs or multiples. In low viscosity environments, gas gaps sufficiently deep are prone to the Rossby wave instability \citep[RWI; e.g.,][]{li01, li05} at their edges. Gaps in this paper are generally too shallow to develop the RWI. Experiments (not shown) suggest that $\sigmag$ needs to be depleted by $\gtrsim 50$\% to trigger the RWI when $\alpha=5\times10^{-5}$. The RWI may develop at the edges of more than one gap, forming multiple dust-trapping vortices. Figure~\ref{fig:rwi} shows an example Model RWI, in which the RWI is triggered at both the inner edge of IG1 and the outer edge of OG1, forming two vortices. Dust are also collected at the triangular Lagrange points L$_4$ and L$_5$. In total, four dust clumps are produced. \begin{figure} \begin{center} \includegraphics[trim=0 0 0 0, clip,width=0.45\textwidth,angle=0]{plot_sigma_rwi.pdf} \end{center} \figcaption{The dust surface density map for Model RWI at 1200 orbits in log scale with an aggressive color stretch. The green dot and the green plus symbol mark the locations of the planet and the star, respectively. Two vortices triggered by the RWI form at the outer edge of OG1 and the inner edge of IG1. Two dust clumps at the triangular Lagrange points L$_4$ and L$_5$ are also present. In total there are 4 dust clumps in this disk with one planet. See \S\ref{sec:lowalpha} for details. \label{fig:rwi}} \end{figure} \subsection{Implications for Planet Formation}\label{sec:implication} An emerging trend from recent ALMA high resolution disk surveys is that multi-gap structures at 10--100 AU are common in disks around a variety of host stars (S. Andrews, F. Long, private comm.). This demands a robust gap opening mechanism insensitive to disk and host star properties. We argue that such structures can be produced by one or several planets ranging in mass from 0.1 to 10s of Earth masses, located at 1--100 AU. Figure \ref{fig:mars} shows a ``Model Mars'' where a 0.1$\me$ ($3\times10^{-7}\msun$) planet is seen to generate multiple dust gaps over $10^5$ orbits. The low $h/r=0.02$ of this model may characterize disk regions within a few AU from the star, where terrestrial planets in the Solar System reside. Mars-mass planets are capable of being formed by the streaming instability + pebble accretion \citep{johansen07youdin, ormel17, johansen17, lin18}. \begin{figure} \begin{center} \includegraphics[trim=0 0 0 0, clip,width=0.45\textwidth,angle=0]{plot_sigma_mars.pdf} \end{center} \figcaption{The dust surface density map for Model Mars at 20,000 orbits. A Mars mass planet ($0.1\me = 3\times10^{-7}\msun$; green dot) can significantly perturb the dust distribution. See \S\ref{sec:implication} for details. \label{fig:mars}} \end{figure} If most multi-gap structures at 10--100 AU are produced by planet-disk interactions, the ubiquity of disk structures implies a ubiquity of planets analogous to the ice giants (Uranus and Neptune) in our Solar System, and a short time for the formation of their cores (e.g., HL Tau and GY 91 are believed to be $\lesssim$1 Myr old; \citealt{brogan15, sheehan18}). Microlensing surveys might be able to uncover such a population of ``cold Neptunes'' \citep[e.g.,][]{poleski14}. Analyzing microlensing survey data available at the time, \citet[Fig. 15; see also \citealt{zhu17ulensing}]{suzuki16} concluded that cold Neptunes are common, with an order unity occurrence rate around main-sequence stars. We carry out two-dimensional, two-fluid hydrodynamical simulations of protoplanetary disks with low viscosities ($\alpha\lesssim10^{-4}$) to study the spacings, depths, and total number of dust gaps opened by one sub-thermal mass planet ($\mplanet<\mth=M_\star(h/r)^3$). Our results provide basic guidelines for interpreting multi-gap structures seen in mm continuum emission. Our main findings are: \begin{enumerate} \item Among the observable properties of gaps and rings, gap location is the least affected by observing conditions, finite optical depth, and time evolution (\S\ref{sec:comments}; \S\ref{sec:time}; Figure~\ref{fig:time}). \item Gap spacings increase with increasing $h/r$ and decreasing $\mplanet$ (\S\ref{sec:spacing}; Figure~\ref{fig:rp}). We provide empirical fitting functions for the double gap OG1+IG1 (defined in Figure~\ref{fig:example}; Eqn.~\ref{eq:fit:og1ig1}) and IG2 (Eqn.~\ref{eq:fit:ig1ig2}; Figure~\ref{fig:spacing}). \item Gaps are opened faster by planets with higher masses, and in disks with higher $h/r$ (\S\ref{sec:depth}). \item More gaps are opened within a given radius range in disks with lower $h/r$. The number of gaps is less sensitive to $\mplanet$. \item The spacings of the double gap in HL Tau \citep{brogan15} and TW Hya \citep{andrews16}, and all three gaps in HD 163296 \citep{isella16hd163296} match those of modelled gaps produced by a sub-Saturn mass planet (\S\ref{sec:applications}; Figures~\ref{fig:realdisks} and \ref{fig:sidebyside}). \item A low midplane viscosity ($\alpha\lesssim10^{-4}$) is needed for a sub-thermal mass planet to open multiple gaps (Figure~\ref{fig:alpha}). The Rossby wave instability may develop at the edges of gaps, producing multiple dusty vortices (Figure~\ref{fig:rwi}). A planet as low mass as Mars may significantly perturb the dust disk in terrestrial planet forming regions (Figure~\ref{fig:mars}). \end{enumerate} | 18 | 8 | 1808.06613 |
1808 | 1808.03200_arXiv.txt | Perturbations of Friedmann-Lema\^{i}tre (FL) cosmologies play an essential role in confronting theoretical models with observations of the anisotropy of the cosmic microwave background (CMB) and the inhomogeneity of the large scale structure (LSS) of the Universe. Initially linear perturbations were adequate but now the increasing accuracy of the observations necessitates the use of second order (nonlinear) perturbations to analyze, for example, the presence of non-Gaussianity in the CMB and the LSS.\footnote {See, for example, Bartolo \emph{et al} (2010)~\cite{baretal10} and Tram \emph{et al} (2016)~\cite{traetal16}.} In this paper we consider first and second order scalar perturbations of FL universes subject to the following assumptions: \begin{itemize} \item[i)] the spatial background is flat; \item[ii)] the stress-energy tensor can be written in the form $T^a\!_b = \left(\rho + p\right)\!u^a u_b + p\delta^a\!_b$, thereby describing perfect fluids and scalar fields; \item[iii)] the linear perturbation is purely scalar. \end{itemize} The dynamics of perturbations of FL universes are governed by the perturbed Einstein equations and the perturbed matter equations. For \emph{scalar perturbations} the perturbed Einstein equations give four equations (linear combinations of the components of the perturbed Einstein tensor) which include evolution equations for the metric perturbations. The perturbed conservation equations provide evolution equations for two primary matter perturbations, the density perturbation and the scalar velocity perturbation. Only four of these six equations are needed to fully describe the perturbations, but in order to obtain a well-defined system the gauge freedom has to be eliminated by fixing the gauge. Since 2004 much work aimed at confronting theoretical models with observations has been done using second order perturbations. In one respect second order perturbations are analogous to first order perturbations: the \emph{leading order terms} in the equations have exactly the same form. The greater complexity at second order arises from the fact that each equation is augmented by so-called \emph{source terms} that depend quadratically on the first order perturbations. There are various ways of formulating the governing equations for second order perturbations, depending on the choice of variables and gauge. These choices are influenced by various factors such as the problem to be investigated, for example, long wavelength perturbations or perturbations of the $\Lambda CDM$ universe, or in the case of numerical work, by the availability of numerical packages. A number of detailed formulations of the governing equations have been given,\footnote {See for example, Noh and Hwang (2004)~\cite{nohhwa04}, and Nakamura (2007)~\cite{nak07}.} but mainly due to the complexity of the source terms no standard systems have emerged: it is as though the necessary technical infrastructure for analyzing second order perturbations has not been sufficiently well developed. With this as motivation, our goal in this paper is to present five systems of equations that are suitable for analyzing the dynamics of both first and second order scalar perturbations of FL universes. To accomplish this we begin by imposing the so-called C-gauge of Hwang and Noh~\cite{nohhwa04} up to second order which fixes the spatial gauge, but we initially keep an arbitrary temporal gauge. Within this framework we construct a set of leading order and quadratic source terms for the perturbed Einstein field equations (a set of four scalar equations) and the perturbed energy-momentum conservation equations (a set of two scalar equations). Finally we construct five specific systems of gauge invariant equations by also fixing the temporal gauge. Firstly, specializing the perturbed Einstein field equations to the Poisson (longitudinal, zero shear) gauge and the uniform (flat) curvature gauge yields two systems of governing equations. Secondly, specializing three of the perturbed Einstein field equations together with the perturbed momentum conservation equation to the Poisson gauge and the total matter gauge results in two more systems. Finally we create a fifth system by using the perturbed energy-momentum equations to describe the evolution of the density perturbation in the total matter gauge, and the velocity perturbation in the Poisson gauge, with two of the perturbed Einstein equations acting as constraints to determine the metric perturbations. We regard these five systems of equations as \emph{ready-to-use} since they are gauge invariant, contain no redundant equations or variables, and do not require that any further simplifications be made before use. It is important to note that in general the above systems of equations are not closed (not fully determined) since the non-adiabatic pressure perturbation has to be specified. However, for a barotropic perfect fluid and a minimally coupled scalar field, the systems are fully determined, once an equation of state and a scalar field potential, respectively, has been given. Moreover, we present the systems in a manner that makes it possible to apply them to more general matter models such as multi-fluids and multiple scalar fields. The present paper is the second of four related papers by the authors. The first paper~\cite{uggwai19a}, hereafter referred to as UW1, gives a unified and simplified formulation of gauge change formulas at second order, while the third paper~\cite{uggwai19b}, called UW3, uses the present paper, in conjunction with UW1, to give new conserved quantities and derive the general explicit solution at second order for adiabatic perturbations in the long wavelength limit, results that are subsequently adapted to inflationary universes with a single scalar field in~\cite{uggwai19c}, which we refer to as UW4. The outline of the paper is as follows. In section~\ref{variables} we introduce the metric and matter perturbation variables. In section~\ref{pert.einst} we present leading order and quadratic source terms for the perturbed Einstein field equations, and in section~\ref{pert_cons} we present the leading order and quadratic source terms for the perturbed conservation equations. In both cases the details of the source terms are deferred to an appendix. The central goal of the paper is reached in section~\ref{gov.eq} where we derive the five ready-to-use systems of governing equations. Finally in section~\ref{discussion} we comment on specific applications of the five systems and on their relative merits. | } In this paper we have given five ready-to-use systems of governing equations for second order scalar perturbations, subject to the assumption that at first order the perturbations are purely scalar. Here we summarize their identifying features and give their active dynamical variables: \begin{itemize} \item[i)] Equations~\eqref{p_gov2} using the Poisson gauge (variable $\psi_{\mathrm p}$), \item[ii)] Equations~\eqref{evol2_psi,V} using the Poisson gauge (variables $\psi_{\mathrm p}, V_{\mathrm p}$), \item[iii)] Equations~\eqref{ucg_gov2} using the uniform curvature gauge (variables $\phi_\mathrm{c}, B_{\mathrm c}$), \item[iv)] Equations~\eqref{totmat_2} using the total matter gauge (variables $\psi_{\mathrm v}, B_{\mathrm v}$), \item[v)] Equations~\eqref{delta_v,V_p_evol2} using the conservation equations (variables $\bdelta_{\mathrm v}, {\bf D}^2 V_{\mathrm p}$). \end{itemize} Other systems of equations that are more general then ours, as regards matter content and gauge choices, have been developed in the extensive series of papers by Hwang and Noh (see for example~\cite{nohhwa04,hwanoh07a}) and Nakamura~\cite{nak07,nak10}. We regard our less general but more focussed framework, which comprises the above five ready-to-use systems of equations, as complementing the more general systems in the above references. Each of our systems is minimal in the sense that there are no redundant equations or variables, and the matter content is restricted so that the systems are closed once the non-adiabatic pressure perturbation $\Gamma$ is specified. Although we are primarily motivated by the needs of second order perturbation theory we note that our framework can be specialized to linear perturbations by simply dropping the source terms. Because of this we hope that our framework will form a useful reference for both linear and second order perturbations. We now make some remarks concerning the utility of the five systems of governing equations as regards applications. The unified nature of our formulation of these systems of equations enables one to easily compare their relative merits as regards a chosen application. We begin by noting that in cosmological perturbation theory the evolution of the perturbations is described in general by partial differential equations. Usually, in order to obtain explicit, approximate or numerical solutions in a particular physical context, one applies the Fourier transform to the partial differential equations which converts them to ordinary differential equations for the Fourier coefficients of the perturbation variables, with the wave number $k$ as a parameter, together with algebraic constraints relating the Fourier coefficients. For first order perturbations the spatial derivatives appear only via the spatial Laplacian ${\bf D}^2$, and one can implement the transition by simply making the replacement ${\bf D}^2\rightarrow - k^2$. At second order, however, the process is more complicated since one has to use the Convolution Theorem to take the Fourier transform of products of the first order perturbations that appear in the source terms.\footnote{See for example, Tram \emph{et al} (2016)~\cite{traetal16}, equations (1.1)-(1.3) and Vretblad (2005)~\cite{vre05} for details.} There are, however, two important applications of cosmological perturbation theory, namely, adiabatic perturbations in the super-horizon regime (the long wavelength limit) and perturbations of $\Lambda$CDM universes, in which it is not necessary to make the transition to Fourier space since the evolution equations automatically simplify to ordinary differential equations. We regard these applications as elementary but important benchmark problems in cosmological perturbation theory. First we note that long wavelength adiabatic perturbations are defined by the requirement that terms of order 2 in the scaled dimensionless spatial differential operator ${\cal H}^{-1}{\bf D}^i$ can be neglected, and that the non-adiabatic pressure perturbation is negligible (${}^{(r)}\!\Gamma\approx 0,\,r=1,2$). However, the background matter scalars $w$ and $c_s^2$ are unrestricted. Second, as shown in appendix~\ref{frac.density.pert}, when the background model is the $\Lambda CDM$ universe we have $w_m=0$ and hence the background matter scalars are given by \begin{equation} c_s^2=0, \qquad 1+w=\Omega_m, \end{equation} which implies that the perturbations are adiabatic (${}^{(r)}\!\Gamma= 0,\,r=1,2$). In both cases the term $c_s^2{\bf D}^2$, which appears in the leading order terms in the evolution equations, is negligible, and it is this property that reduces the evolution equations to ordinary differential equations. It turns out that in these two benchmark problems it is possible to explicitly solve the ordinary differential equations and obtain the general time dependence of the perturbations at first and second order, including both growing and decaying modes. The spatial dependence is described by arbitrary spatial functions that arise as constants of integration. In order to achieve this goal it is necessary to make an appropriate choice from among the five ready-to-use systems. Observe that in both systems iii) (uniform curvature gauge) and iv) (total matter gauge) the two evolution equations decouple, and can thus be solved successively for the two active variables, first at linear order and then, after using the linear solution to calculate the source terms, at second order. However, on evaluating the source terms one finds that \emph{system iv), using the total matter gauge, provides the simplest method of solution for the two benchmark problems.} Details concerning the derivation of the solution in the case of long wavelength perturbations are given in UW3~\cite{uggwai19b}. We conclude with some brief remarks on the relative merits of the five systems for problems other than the two benchmark problems. An immediate conclusion is that system iv) is no longer the simplest system since the presence of a non-zero $c_s^2$ complicates the evolution equations considerably, since the term $c_s^2\bdelta_{\mathrm v}$, which depends on $\psi_{\mathrm v} -{\cal H}B_{\mathrm v}$, appears on the right side of both evolution equations in~\eqref{totmat_1} and~\eqref{totmat_2}. In addition the presence of this term makes the source terms more complicated. Instead it appears that \emph{system iii), based on the uniform curvature gauge, is the simplest system}, which makes it a natural choice for numerical experiments or qualitative analysis using dynamical systems methods. This system has not been given before.\footnote{We mention, however, that Malik and co-workers have used the uniform curvature gauge to study second order perturbations of inflationary universes with single and multiple scalar fields (see for example, Malik (2007)~\cite{mal07}, Huston and Malik (2009)~\cite{husmal09} and Christopherson \emph{et al} (2015)~\cite{chretal15}.) The structure of the governing equations in these references is specifically adapted to the scalar fields and as a result they do not have much in common with our governing equations.} Our analysis suggests that it is worthy of further study. \begin{appendix} | 18 | 8 | 1808.03200 |
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