adsabs/FOCAL huggingface dataset:
In addition to some fraction of BNS mergers masquerading as long GRBs, our sample used to constrain the DTD may suffer from other issues of incompleteness. As we rely on the modeling of host galaxies when constraining the DTD, we do not consider short GRBs that do not have a confident host association. Though the properties of short GRB hosts do not seem to deviate strongly as a function of the host-association confidence (see, e.g., Figure 4 of Nugent et al. 2022), neglecting these events may have a potential impact on both the low and high ends of our inferred DTD. For example, GRBs that are highly offset from their hosts may have afterglows with much lower luminosities, making precise localization (and therefore host identification) difficult (Perna et al. 2022). Such systems may have migrated over longer timescales to reach the highly offset locations of the burst and therefore may have longer delay times than the general population. Furthermore, the P
cc method for host identification may incorrectly associate a GRB with a faint underlying host rather than a bright host at a larger offset, though Fong et al. (2022) predicted this to be an effect only at the ≲7% level. On the other hand, if such poorly associated GRBs are instead truly associated with faint galaxies that are below detection limits, we may be excluding additional systems with short delay times as these faint, low-mass galaxies are typically star-forming. Furthermore, though Swift can detect GRBs out to z ∼ 3, there is likely some fraction of short GRBs that occur beyond this horizon, when the universe was ≲2 Gyr old. Short GRBs that occur at these early stages in the history of the universe must have short delay times, and this selection effect may bias the general population in our analysis to longer delay times. This would lead to a larger inferred 



tmin

, and, due to the correlation between 



tmin

 and α, more negative values of α. However, this population of high-redshift short GRBs is likely small; assuming the SFH from Madau & Fragos (2017), 10% of stars are born beyond z = 3, and the fraction of compact object binary mergers beyond this redshift will be even smaller due to the delay time between formation and merger.
Citation: Perna et al. 2022
Citation function: Compare/Contrast
In general, hybrid-kinetic models have proved to be capable of satisfactorily catching the main kinetic physics at play for a large number of problems, ranging from fluid and kinetic instabilities (Hellinger & Matsumoto 2000; Matteini et al. 2006; Califano et al. 2008; Henri et al. 2013; Kunz et al. 2014) to collisionless shocks (Lembège et al. 2009; Caprioli & Spitkovsky 2013; Weidl et al. 2016), dynamo effects (Rincon et al. 2016; St-Onge & Kunz 2018), MR (Le et al. 2016; Palmroth et al. 2017; Cerri & Califano 2017; Franci et al. 2017; Wang et al. 2019; Califano et al. 2020), and kinetic-scale turbulence (Servidio et al. 2015; Grošelj et al. 2017; Cerri et al. 2018, 2019; Hellinger et al. 2019; Wang et al. 2019). The main goal of our project is to investigate the possibility of improving the electron description in hybrid-kinetic models. The starting point is the “hybrid Vlasov–Maxwell” (HVM) code (Mangeney et al. 2002; Valentini et al. 2007), which is already equipped with a fluid model in which electrons are described as an isotropic, isothermal fluid with finite inertia. The HVM code has recently been upgraded in order to implement a more sophisticated model for the electron fluid. This includes evolution equations for the anisotropic (gyrotropic) electron pressures, p∥, e and p⊥, e (where ∥ and ⊥ refer to the local magnetic-field direction, b = B/|B|), and a Landau-fluid (LF) closure for the parallel transport of the gyrotropic electron thermal energy along field lines (i.e., parallel heat fluxes, q∥, e and q⊥, e). Hereafter, we refer to this new model as “hybrid Vlasov–Landau-fluid” (HVLF). The idea is to include within a hybrid description the relevant electron pressure-anisotropy effects and a fluid model for the electron-kinetic response that still holds in a nonlinear regime. Therefore, the LF model implemented in the HVLF code goes beyond the early attempts to include these effects in simplified settings (e.g., Hammett & Perkins 1990; Snyder et al. 1997; Passot & Sulem 2007) and is based on the approach presented by Sulem & Passot (2015).
Citation: Valentini et al. 2007
Citation function: Uses; Background
In this paper, we report on spectroscopic CH3CN, CH3OH (methanol), and dust continuum observations with the Atacama Large Millimeter/submillimeter Array (ALMA) at 349 GHz with an angular resolution of 0′′.1. We exploit the CH3CN (19K–18K) K-ladder, with excitation energies ranging from 168 K (for K = 0) to 881 K (for K = 10), to probe, at different radii, the physical conditions in the accretion disk of an early-type young star. We targeted the star-forming region G023.01−00.41, at a trigonometric distance of 4.59

$^{+0.38}_{-0.33}$





−0.33

+0.38



 kpc from the Sun (Brunthaler et al. 2009), where we recently revealed the accretion disk around a young star of 104.6 L⊙, corresponding to a ZAMS star of 20 M⊙ (Sanna et al. 2019, their Fig. 1); the disk was imaged by means of spectroscopic ALMA observations of both CH3CN and CH3OH lines at 0′′.2 resolution inthe 230 GHz band. The disk extends up to radii of 3000 au from the central star where it warps above the midplane; here, we resolve the outer disk regions in two apparent spirals projected onto the plane of the sky. We showed that molecular gas is falling in and slowly rotating with sub-Keplerian velocities down to radii of 500 au from the central star, where we measured a mass infall rate of 6 × 10−4 M⊙ yr−1 (Sanna et al. 2019, their Fig. 5). The disk and star system drives a radio continuum jet and a molecular outflow aligned along a position angle of 57°, measured east of north (Sanna et al. 2016, their Fig. 2); their projected axis is oriented perpendicular to the disk midplane whose inclination with respect to the line-of-sight was estimated to be less than 30° (namely, the disk is seen approximately edge-on; Sanna et al. 2014, 2019). Previously, we also measured the average gas conditions over the same extent of the whole disk, by means of Submillimeter Array (SMA) observations of the CH3CN (12K–11K) emission, and we estimated a kinetic temperature of 195 K and CH3CN column density of 5.1 × 1016 cm−2 (Sanna et al. 2014, their Fig. 2 and Table 4).
Citation: Sanna et al.; 2019
Citation function: Background
Citation: Sanna et al. 2019
Citation function: Background; Background
The gravitational potential of an eccentric companion at the quadrupole order can be decomposed as a sum over circular orbits (e.g. Storch & Lai 2013; Vick, Lai & Fuller 2017):
(5)$$\begin{eqnarray*}
U\left(\boldsymbol{\mathbf {r}}, t\right) = \sum \limits _{m=-2}^2 U_{2m} \left(\boldsymbol{\mathbf {r}}, t\right) ,
\end{eqnarray*}$$(6)$$\begin{eqnarray*}
U_{2m}\left(\boldsymbol{\mathbf {r}}, t\right) &amp;=&amp; -\frac{GM_2 W_{2m} r^2}{D(t)^3} Y_{2m}(\theta , \phi) e^{-imf\!\!\!\:(t)}, \\
&amp;=&amp; -\frac{GM_2W_{2m} r^2}{a^3}Y_{2m}\left(\theta , \phi \right) \sum \limits _{N = -\infty }^\infty \!\!F_{Nm}e^{-iN\Omega t} .
\end{eqnarray*}$$Here, the coordinate system is centered on the MS star, (r, θ, ϕ) are the radial, polar, and azimuthal coordinates of $\boldsymbol{\mathbf {r}}$ respectively, $W_{2 \pm 2} = \sqrt{3\pi /10}$, W2 ± 1 = 0, $W_{20} = -\sqrt{\pi / 5}$, D(t) is the instantaneous distance to the companion, f is the true anomaly, and Ylm denote the spherical harmonics. FNm denote the Hansen coefficients for l = 2 (also denoted $X^N_{2m}$ in Murray & Dermott 1999), which are the Fourier coefficients of the perturbing function, i.e.
(7)$$\begin{eqnarray*}
\frac{a^3}{D(t)^3} e^{-imf\!\!\!\:(t)} = \sum \limits _{N = -\infty }^\infty \!\!F_{Nm} e^{-iN\Omega t}.
\end{eqnarray*}$$The FNm can be written explicitly as an integral over the eccentric anomaly (Murray & Dermott 1999; Storch & Lai 2013):
(8)$$\begin{eqnarray*}
F_{Nm} = \frac{1}{\pi }\int \limits _{0}^{\pi } \frac{\cos \left[N\left(E - e\sin E\right) - mf(E)\right]}{\left(1 - e\cos E\right)^2}\,\,\mathrm{d}E.
\end{eqnarray*}$$By considering the effect of each summand in equation (5), the total torque on the star, energy transfer in the inertial frame, and energy transfer in the star’s corotating frame (which is also the tidal heating rate) can be obtained (Storch & Lai 2013; Vick et al. 2017):
(9)$$\begin{eqnarray*}
T = \sum \limits _{N = -\infty }^\infty F_{N2}^2 T_{\rm circ}\left(N\Omega - 2\Omega _{\rm s}\right),
\end{eqnarray*}$$(10)$$\begin{eqnarray*}
\dot{E}_{\rm in} &amp;=&amp; \frac{1}{2}\sum \limits _{N = -\infty }^\infty \Bigg [ \left(\frac{W_{20}}{W_{22}}\right)^2 N\Omega F_{N0}^2 T_{\rm circ}\left(N\Omega \right) \\
&amp;&amp;+\, N\Omega F_{N2}^2 T_{\rm circ}\left(N\Omega - 2\Omega _{\rm s}\right) \Bigg ] ,
\end{eqnarray*}$$(11)$$\begin{eqnarray*}
\dot{E}_{\rm rot} = \dot{E}_{\rm in} - \Omega _{\rm s} T .
\end{eqnarray*}$$Here, dots indicate time derivatives.
Citation: Murray & Dermott 1999
Citation function: Uses; Uses
Citation: Murray & Dermott 1999
Citation function: Background
In the last few years, several experiments have reported upper limits on the power spectrum of 21-cm fluctuations during reionization (Parsons et al. 2014; Patil et al. 2017; Barry et al. 2019; Mertens et al. 2020; The HERA Collaboration 2021b) and the earlier cosmic-dawn era (Eastwood et al. 2019; Gehlot et al. 2019, 2020; Garsden et al. 2021; Yoshiura et al. 2021). Scenarios in which the bulk IGM is still colder than the cosmic microwave background (CMB) during reionization give rise to the strongest fluctuations and so will be the first models to be tested as upper limits continue to improve (e.g. Parsons et al. 2014; Pober et al. 2015; Greig, Mesinger & Pober 2016). Similarly, stronger-than-expected 21-cm signals can arise if the cosmic radio background has contributions other than the CMB (Feng & Holder 2018), e.g. synchrotron emission from accreting black holes (Ewall-Wice et al. 2018), star-forming galaxies (Mirocha & Furlanetto 2019), or from decaying particles (Fraser et al. 2018; Pospelov et al. 2018). Indeed, constraints from MWA, HERA, and LoFAR disfavour models with negligible X-ray heating at z ∼ 8–9 or very strong radio backgrounds (Ghara et al. 2020, 2021; Mondal et al. 2020; Greig et al. 2021a, b; The HERA Collaboration 2021a). Of course, the recent report of an absorption signal in the sky-averaged spectrum at z ∼ 17 from EDGES (Bowman et al. 2018) requires an even colder IGM (Barkana 2018; Boddy et al. 2018; Fialkov, Barkana & Cohen 2018; Kovetz et al. 2018; Muñoz & Loeb 2018) or a brighter background (Ewall-Wice et al. 2018; Feng & Holder 2018; Fialkov & Barkana 2019; Mirocha & Furlanetto 2019) than models in ΛCDM cosmologies generally predict. However, the most stringent power spectrum upper limits from The HERA Collaboration (2021b) are derived at sufficiently low redshifts relative to EDGES (z ≲ 10 versus z ≃ 18) that they cannot yet directly address the EDGES controversy (Hills et al. 2018; Bradley et al. 2019; Singh & Subrahmanyan 2019; Sims & Pober 2020; Tauscher, Rapetti & Burns 2020; Singh et al. 2021).
Citation: Feng & Holder 2018
Citation function: Motivation; Motivation; Motivation
Citation: Feng & Holder 2018
Citation function: Motivation
Recently, our knowledge of the MW formation and evolution has been revolutionized by the massive amount of data products from the Gaia mission and large spectroscopic surveys. Several dwarf galaxies are suggested to have been accreted by the MW since its formation (e.g., Helmi et al. 2018; Belokurov et al. 2018; Myeong et al. 2019). It is expected that GCs formed in these dwarf galaxies mix up with GCs formed in situ, i.e., in the main MW progenitor, to form the current GC system. This formation history can be revealed by precise kinematics and dynamical modeling. Massari et al. (2019) placed Pal 13 into a group of GCs related to a dwarf galaxy, named Sequoia, which was likely accreted 9 Gyr ago (Myeong et al. 2019). The existence of MPs in Pal 13 supports the statement that this phenomenon is not unique in Galactic GCs formed in situ (e.g., Li & de Grijs 2019; Milone et al. 2020). Besides its proposed accreted origin, Pal 13 is suggested to be experiencing tidal stripping (Yepez et al. 2019; Piatti & Fernández-Trincado 2020). Hamren et al. (2013) found that Pal 13 has lost a considerable amount of mass, which is related to its low present-day cluster mass. Recently, the discovery of tidal tails in Pal 13 was reported by Shipp et al. (2020). Using the RR Lyrae stars, these authors estimated the initial luminosity to be 





 L⊙, which is significantly larger than the current luminosity estimated by B11, 





 L⊙. Due to its large distance, several detailed features of Pal 13 have been discovered just recently (e.g., Piatti & Fernández-Trincado 2020; Shipp et al. 2020) and more are still waiting for further investigation, the estimation of total mass loss or initial mass of Pal 13 should be illuminative for future studies. Besides Pal 13, several remote, low mass GCs, including Whiting 1 and Eridanus GC, also show evidence of tidal tails or extra-tidal structures(Carballo-Bello et al. 2014; Myeong et al. 2017). These features agree with the predictions of dynamical evolution of low mass GCs: it is difficult to keep stars in a shallow potential well. The relatively large stellar mass lost from low mass GCs also complicates the discussion of MPs and especially of the lower mass limit needed to maintain MPs. On the other hand, the N-rich stars located in these GCs are lost to the field, contributing to the rare N-rich field stars (e.g., Martell & Grebel 2010; Fernández-Trincado et al. 2017; Tang et al. 2019; Fernández-Trincado et al. 2019; Tang et al. 2020).
Citation: Milone et al. 2020
Citation function: Similarities
SBs are complex environments. The hot rarefied plasma is spanned by primary shocks, which decay into turbulent motions and MHD waves. Collective effects such as re-acceleration processes contribute to the acceleration of particles. SBs are also delimited by a forward shock which expands into the ISM. While the size of this shock can be very large (R ≳ 100 pc, Weaver et al. 1977), it is too slow (u ∼ 30 km s−1) to accelerate PeV protons, even assuming large magnetic fields in the ISM. More generally, in the case of strong shocks, the velocity u is identified as the velocity of the shock, as it is precisely the velocity jump at the shock discontinuity which drives the acceleration of the particles via the DSA mechanism (e.g. Drury 1983). Inside SBs, there are two types of strong primary shocks: the time-dependent SNR shocks which expand after a supernova (SN) explosion, and the wind termination shocks (WTS) which surround the individual massive stars, or the entire stellar cluster if it is compact enough. SNRs and WTSs have a typical velocity of several 1000 km s−1. In the low-density SB interior, SNRs expand to a typical radius of a few tens of pc before reaching the Sedov–Taylor phase, which is generally larger than the radius of the WTS, even in the case of a WTS powered by a very massive compact cluster. The size of the latter shock depends on the mechanical power of the stellar cluster and is typically of the order of 10 pc (Weaver et al. 1977). It is believed that, due to in situ CR acceleration, efficient magnetic field amplification takes place upstream of SNR shocks, leading to a magnetic field of up to several tens of μG. On the other hand, it is less clear if the magnetic field can be as efficiently amplified in stellar winds. In both cases, the maximum energy is generally inferred to be only ∼1 PeV (e.g. Gupta et al. 2020; Morlino et al. 2021; Vieu et al. 2022b). While this is slightly larger than the maximum energy achieved at isolated SNR, atypical conditions would be required for protons to be accelerated well beyond PeV by primary shocks embedded in SBs. A promising situation combining the advantages of both WTS and SNRs might be that of an SNR expanding in a wind profile close to a compact cluster. A powerful cluster may indeed convert a substantial amount of its mechanical energy into turbulence, amplifying magnetic fields up to hundreds of μG in its vicinity, such that particles could be accelerated up to 10 PeV by a powerful SNR propagating in the wind, even in the absence of additional magnetic-field amplification. Considering a similar scenario of SNR shock – cluster wind interaction with efficient turbulence generation, Bykov et al. (2015) found that proton energies up to 40 PeV could be reached in the case of fast shocks (U = 104 km s−1).
Citation: Drury 1983
Citation function: Background
All coreshine observations were obtained with the IRAC instrument on board the Spitzer observatory and are gathered in Paladini (2014) and Lefèvre et al. (2014), from which we selected 21 starless cores in the Taurus-Perseus, Chamaeleon, Cepheus, and L183/L134 regions (see their Fig. 9 and Table 1). Lefèvre et al. (2014) summarised their results in two figures that we reproduce here: the 4.5 to 3.6 μm ratio, that they name “coreshine ratio”, as a function of the 3.6 μm intensity and the 2.2 to 3.6 μm ratio, that they name the “near-IR to mid-IR ratio”, as a function of the coreshine ratio. For the model cloud, we use the parameters of the control cloud defined in Sect. 3. After convolving our models with a 10″ FWHM Gaussian kernel to simulate the data analysis presented in Lefèvre et al. (2014) and following the Lehtinen & Mattila (1996) prescription to take into account the part of the ISRF+CM light that can be transmitted through the cloud3, we compute the synthetic photometry for each pixel along a radial cut through our model clouds. The results are shown in Figs. 10a and b, which present the coreshine ratio as a function of the 3.6 μm intensity for the 21 starless cores of Lefèvre et al. (2014) and in Fig. 10c and d, which displays the near- to mid-IR ratio as a function of the coreshine ratio for the Taurus-Perseus and L183/L134 regions. The CMM model has to be ruled out to explain coreshine since it only marginally fits the coreshine ratio and fails to reproduce the near- to mid-IR ratio (green areas in Fig. 10). On the contrary, the CMM+AMM and CMM+AMMI models can explain the coreshine observations (magenta and blue areas in Fig. 10, respectively). Based on the results presented in Sect. 3 and Fig. 7, the cloud parameters are as important as the dust model for explaining the dispersion in the observations. Varying the cloud external radius from Rout = 0.3 to 0.1 pc (Figs. 10a and c) and the central density from ρC = 104 to 5 × 105 H/cm3 (Figs. 10b and d) seem enough to explain most of the observed scatter.
Citation: Lehtinen & Mattila (1996)
Citation function: Uses; Uses
Recently, Fraija et al. (2021a) presented the afterglow light curves generated by the deceleration of sub-relativistic masses ejected from the merger of BCOs and the death of massive stars. The authors assumed that a PL velocity distribution describes the isotropic-equivalent kinetic energy of these masses and that the sub-relativistic ejected masses were decelerated, in turn, by a stratified-density environment. As a particular case, to explain the multiwavelength observations of the gravitational event GW170817/GRB 170817A at ∼900 days, they constrained the parameter space of the synchrotron light curves of a sub-relativistic mass ejected during the merger of two NSs and decelerated in a constant-density environment. The synchrotron radiation of the sub-relativistic material was consistent with the faster blue KN afterglow. Inspired by the new observations of this GW event at 3.3 yr after the initial merger (Hajela et al. 2022), in this paper, we extend the synchrotron model presented in Fraija et al. (2021a), including the continuous energy injection from the central engine (either a spinning magnetized NS or BH remnant) into the blast wave through a numerical approach and analytic arguments. In addition, we apply the current model to potential candidates of sGRB events with evidence of a KN. The paper is organized as follows: Section 2 presents the dynamical evolution of the afterglow when the central engine continuously injects energy into the blast wave. We show an analytical solution and numerical approach. In Section 3, we show a synchrotron model with energy injection from a spinning magnetized NS and BH remnants. Section 4 shows the analysis of the multiwavelength light curves using typical values of the GRB afterglow. In Section 5, we apply our model to several potential candidates, including GW170817/GRB 170817A, and finally, in Section 6, we summarize. We consider the convention 



Qx=Q10x

 in cgs units and assume for the cosmological constants a spatially flat universe Λ cold dark matter model with H
0 = 69.6 km s−1 Mpc−1, ΩM = 0.286, and ΩΛ = 0.714 (Planck Collaboration et al. 2016).
Citation: Hajela et al. 2022
Citation function: Motivation; Motivation
In order to establish the completeness and purity of the RR Lyrae stars confirmed by the SOS Cep&RRL pipeline and to estimate the number of new discoveries by Gaia, we performed a deep and careful comparison with the literature. As a first step, the catalogue of 140 784 confirmed sources was cross-matched against all major catalogues of known RR Lyrae stars that are available. We primarily used the OGLE catalogues for RR Lyrae stars (version IV of the survey, Soszyński et al. 2014, 2016), but we also used RR Lyrae stars by CTRS (Drake et al. 2013a,b, 2014, 2017; Torrealba et al. 2015), ASAS (Pojmanski 1997; Richards et al. 2012), ASAS-SN (Jayasinghe et al. 2018), ATLAS (Tonry et al. 2018), IOMC (Alfonso-Garzón et al. 2012), LINEAR (Palaversa et al. 2013), NSVS (Kinemuchi et al. 2006), Pann-Stars (PS1 Sesar et al. 2017), and from the works based on Kepler/K2 (Debosscher et al. 2011; Nemec et al. 2011; Molnár et al. 2015a,b, 2016) and on the Simbad database (Wenger et al. 2000). These cross-matches returned a list of 88 578 known RR Lyrae stars in our sample of 140 784 stars. The SOS Cep&RRL confirmed RR Lyrae stars were also cross-matched against catalogues of candidate RR Lyrae stars discovered by the VVV survey (Gran et al. 2016; Minniti et al. 2017; D. Minniti, priv. comm.) in the MW disc and bulge. This returned 319 VVV cross-identified sources in the MW disc and 222 in the MW bulge. We thus confirm these VVV candidates. For known RR Lyrae stars in GCs, the main reference was the catalogue of Clement et al. (2001), which was updated to the latest literature as described in Garofalo et al. (in prep.). For variables in dSphs, we used the following references: Kaluzny et al. (1995), Clementini et al. (2005), Kinemuchi et al. (2008), Dall’Ora et al. (2012) and Garofalo et al. (2013). These latter cross-matches returned a list of 1986 further known RR Lyrae stars. At the end of this cross-match procedure, of the 140 784 RR Lyrae stars that are confirmed by the SOS Cep&RRL pipeline, 90 564 were shown to be known previously, and 50 220 are new discoveries by Gaia.
Citation: Drake et al. 2013a
Citation function: Uses; Uses; Uses