adsabs/FOCAL huggingface dataset:
Unlike BH–BH mergers, binary neutron star (BNS) merger mechanisms are expected to yield an optical counterparts powered by the radioactive decay of rapid neutron capture process (r-process) elements synthesized in the merger ejecta (Li & Paczyński 1998; Metzger et al. 2010) or by the cooling of shock-heated material around the neutron star (NS; Piro & Kollmeier 2018). On 2017 August 17, the LVC detected the first and best example of a BNS merger: GW170817 (Abbott et al. 2017a). Only two seconds after its detection, a short gamma-ray burst (GRB 170817A) was detected (Connaughton et al. 2017; Goldstein et al. 2017). GRB 170817A was followed by the discovery of the optical counterpart Swope Supernova Survey 17a [SSS17a (AT 2017gfo)] ∼11 h later in the galaxy NGC 4993 (Coulter et al. 2017) and later confirmed by others teams (Abbott et al. 2017c; Arcavi et al. 2017; Covino et al. 2017; Díaz et al. 2017; Drout et al. 2017; Kilpatrick et al. 2017; Lipunov et al. 2017; McCully et al. 2017; Pian et al. 2017; Shappee et al. 2017; Smartt et al. 2017; Soares-Santos et al. 2017; Troja et al. 2017; Utsumi et al. 2017; Valenti et al. 2017). The combination of EM and GW information on the BNS can be used to constrain the mass and the radii of NS (Margalit & Metzger 2017) or their equation of state (Bauswein & Janka 2012; Annala et al. 2018). Also, observations of the optical–infrared (IR) ‘kilonova’ counterpart (Chornock et al. 2017; Cowperthwaite et al. 2017; Drout et al. 2017; Kilpatrick et al. 2017; Nicholl et al. 2017; Shappee et al. 2017; Valenti et al. 2017; Villar et al. 2017) provided the first observational confirmation that NS mergers produce the majority of the r-process elements heavier than iron (Burbidge et al. 1957; Cameron 1957; Kasen et al. 2017; Pian et al. 2017; Metzger 2019). They also permit tests of theoretical kilonova models. For example, Drout et al. (2017) showed that the temperatures cooled from 10 000 to 5100 K in between 12 and 36 h after the event, confirming model predictions. Also, Shappee et al. (2017) used spectra taken 11.76 and 12.72 h after the merger to show that the photosphere was expanding at ∼0.3c.
Citation: Connaughton et al. 2017
Citation function: Background
The formation and evolution of S0 galaxies are very important for understanding the formation and evolution of galaxies, but they are still an open question (e.g., the recent review by D’Onofrio et al. 2015). Currently, there are two possible scenarios on the origin of S0 galaxies. One is that S0 galaxies are transformed from spiral galaxies, where spirals lose their gas and star formation is rapidly quenched. The other is that S0 galaxies are intrinsically different from spiral galaxies since their formation (Kormendy & Kennicutt 2004; Barway et al. 2009; van den Bergh 2009). The transformation origin may be associated with intra-cluster medium and neighboring galaxies, via minor mergers, slow encounters, galaxy harassments (Moore et al. 1996), or tidal effects in the dense environment (Gunn & Gott 1972; Larson et al. 1980; Dressler & Sandage 1983; Mihos & Hernquist 1994; Moore et al. 1998, 1999; Neistein et al. 1999; Shioya et al. 2002). A lot of studies have discussed the environmental dependence of galaxy evolution. Generally speaking, early-type galaxies tend to be in dense environments and have low star-formation rates (SFRs; Dressler 1980; Balogh et al. 1997, 1998, 2000; Poggianti et al. 1999; Treu et al. 2003). The fraction of S0 galaxies in the field is only about 15%, while spirals are the majority (Naim et al. 1995). Within the group environment, spirals and S0s are both about 40%–45% (Postman & Geller 1984). Furthermore, S0s become dominant in dense environments, the fraction grows up to 60% in clusters (Dressler 1980; Postman & Geller 1984). Dressler et al. (1997), Fasano et al. (2000), and Desai et al. (2007) also found that the galaxy morphological distributions change abruptly in clusters at 





, about 50 ∼ 70% of spirals at high redshift (z > 0.4) are transformed into S0s, while a fraction of ellipticals, about 25%, remains nearly constant between z = 0.8 and z = 0.0. However, Wilman et al. (2009) showed that the fraction of S0s in groups is the same as in clusters but it is much higher than in the field at z = 0.4, which might suggest that S0s are formed in groups or subgroups.
Citation: Neistein et al. 1999
Citation function: Background; Background
The early studies of the cooling of hypernuclear stars did not include the high-mass astrophysical constraint on the EoS of hypernuclear matter (Haensel & Gnedin 1994; Schaab, Balberg & Schaffner-Bielich 1998; Tsuruta et al. 2009). In Paper I we started a systematic study of the cooling of hypernuclear compact stars on the basis of modern density functionals mentioned above. Several works have appeared since then in this context: Grigorian, Voskresensky & Maslov (2018) used the ‘nuclear medium cooling’ scenario (Schaab et al. 1997; Blaschke, Grigorian & Voskresensky 2004) to account for hyperons using an EoS based on the DFT (Maslov et al. 2015), pion-mediated enhanced neutrino emissivities for the modified Urca and bremsstrahlung processes, vanishing P-wave neutron pairing, and phenomenological crust–core temperature relation. The pion-mediated enhanced neutrino emissivities and vanishing P-wave neutron pairing are the key factors that distinguish their models from those in Paper I. Nevertheless, the remainder of the physical input, for example, the crust–core temperature relation, is not identical to ours. Note that Paper I does not implement the minimal cooling scenario (Page et al. 2004) as it allows for the dUrca process among nucleons and hyperons, uses pair-breaking processes with dressed (in-medium) vertices, and covers a wide range of masses up to the maximum mass of Mmax/M⊙ ∼ 2, which implies a varying cooling behaviour for the mass hierarchy (see, in particular, the conclusion Section of Paper I). The same problem was also addressed by Negreiros et al. (2018) using hyperonic EoS based on the FSU class of relativistic density functional models which feature massive objects and are tuned to available data on hypernuclei (Tolos, Centelles & Ramos2017a,b). In these models, the main cooling agents are the (Λ, p) and (n, p) dUrca processes which were regulated by variations of proton and neutron superfluid gaps. In particular the density range over which the proton S-wave gap is non-zero has been explored with the aim to obtain a satisfactory agreement with the data. The main difference with respect to Paper I is the absence of pairing in the hyperonic sector. Combined, the emerging new generation of cooling models of hypernuclear stars have the potential to further constraining the physics of dense nuclear matter, especially its composition.
Citation: Maslov et al. 2015
Citation function: Background; Background
Magnetic field fluctuations in the solar wind are highly turbulent. The measured power spectral density (PSD) of the fluctuating magnetic field always exhibits power laws k−α, where k is the wavenumber, and α is the spectral index. A single spacecraft measures the PSD as a function of f−α in the frequency domain, which can be converted to the spatial domain under the Taylor hypothesis. According to the physical processes at different scales, the PSD in the solar wind can be divided into several segments, which can be fitted with different α. The inertial range, which is dominated by magnetohydrodynamic (MHD) turbulence, follows the cascade models with spectral indices αi from around 3/2 to 5/3 (Bruno & Carbone 2013; Chen et al. 2020). The PSDs become steepened below the ion scales (ion thermal gyroradius ρi or ion inertial length di), where kinetic mechanisms should be taken into account. Sometimes a sharp transition range is observed with αt ∼ 4 (Sahraoui et al. 2010; Bowen et al. 2020a). This transition range may be caused by imbalanced turbulence (Voitenko & Keyser 2016; Meyrand et al. 2021), energy dissipation of kinetic waves (Howes et al. 2008), ion-scale coherent structures (Lion et al. 2016), or a reconnection dominated range (Mallet et al. 2017). At smaller scales, a flatter sub-ion kinetic range forms with the spectral index αk ∼ 7/3, which can be explained as the MHD Alfvénic turbulence developing into a type of kinetic wave turbulence, e.g., kinetic Alfvén waves (KAWs; Schekochihin et al. 2009) or whistler waves (Cho & Lazarian 2004). Intermittency in the kinetic range could lead to an −8/3 spectrum (Boldyrev & Perez 2012; Zhao et al. 2016). Ion-cyclotron-wave (ICW) turbulence could lead to a steeper −11/3 spectrum (Krishan & Mahajan 2004; Galtier & Buchlin 2007; Meyrand & Galtier 2012; Schekochihin et al. 2019).The kinetic range always behaves as the KAW turbulence with the slope of −2.8 in the near-Earth space (Bale et al. 2005; Chen et al. 2013; Chen 2016). The spectral indices increase again beyond the electron kinetic scales in observations, indicating the conversion of turbulence energy to electrons (Sahraoui et al. 2009; Alexandrova et al. 2012; Chen et al. 2019) or transitions to a further cascade (Schekochihin et al. 2009; Chen & Boldyrev 2017). In simulations, Meyrand & Galtier (2013) obtained a −8/3 spectrum at electron scales under the 3D electron–MHD model.
Citation: Bowen et al. 2020a
Citation function: Background
The main goal of this paper is to use the JAM method to measure the parameters MBH, (M/L)TOT, and βz of the central region of the galaxy NGC 4546. To do this, we obtained AO assisted IFU observations from the Near-Infrared Integral Field Spectrometer (NIFS), installed on the Gemini North Telescope, and also archive photometric observations performed with the Wide Field and Planetary Camera 2 (WFPC2), installed on the Hubble Space Telescope (HST). NGC 4546 (see Fig. 1) is an SB0 galaxy (de Vaucouleurs et al. 1991) with magnitude Mk = −23.30 (Skrutskie et al. 2006, K band) and is located at a distance of 14 Mpc (Tully et al. 2013; 1 arcsec = 68 pc). It hosts a type 1 LINER AGN at the centre (Ricci, Steiner & Menezes 2014b). This object was also classified as a fast rotator by Emsellem et al. (2011). Gas discs in molecular, neutral, and ionized forms are all counter rotating with respect to the stellar disc (Galletta 1987; Bettoni, Galletta & Oosterloo 1991; Sage & Galletta 1994; Sarzi et al. 2006; Ricci, Steiner & Menezes 2014a). Cappellari et al. (2013a) used JAM to model the stellar kinematics of NGC 4546 using a SAURON data cube of this galaxy from the ATLAS3D project (Cappellari et al. 2011). Although their procedure was very important to determine global parameters for the stellar dynamics of NGC 4546, such as (M/L)TOT and βz within 1 effective radius (1 Reff = 22 arcsec, Cappellari et al. 2013a), the spatial resolution of the data cubes from the ATLAS3D project is not high enough to resolve the sphere-of-influence radius of this object. Assuming $R_{\mathrm{ soi}} \, = \, GM_{\mathrm{ BH}} / \sigma _e^2$, where σe is the velocity dispersion within 1 effective radius and using the MBH × σ relation of Saglia et al. (2016), we calculate Rsoi = 18 pc (0.26 arcsec) for σe = 188 km s−1 (Cappellari et al. 2013a). We will show that Rsoi is resolved by our NIFS observations. Thus, this paper will complement the JAM results of Cappellari et al. (2013a) for NGC 4546, but for the vicinity of its SMBH, i.e. using AO assisted IFU data within a field of view (FOV) of 200 × 200 pc2.
Citation: Skrutskie et al. 2006
Citation function: Background; Background; Background
Finally, to calculate Σ(rp), we need to build an accurate model for ξhm on radial scales between 0.1$\, h^{-1}\, {\mathrm{Mpc}}$ and ${\sim}100\, h^{-1}\, {\mathrm{Mpc}}$. We adopt the two-component model of ξhm developed by Zu et al. (2014; a modified version proposed by Hayashi & White 2008):
(6)$$\begin{eqnarray*}
\xi _{\mathrm{hm}}(r) &amp;=&amp; \left\lbrace \begin{array}{ll}\xi _\mathrm{1h} &amp; \quad \mbox{if}\ \xi _\mathrm{1h} \geqslant \xi _\mathrm{2h},\\
\xi _\mathrm{2h} &amp; \quad \mbox{if}\ \xi _\mathrm{1h} \lt \xi _\mathrm{2h},\\
\end{array} \right.\nonumber \\
\xi _\mathrm{1h} &amp;=&amp; \frac{\rho _\mathrm{NFW}(r|M_h)}{\rho _\mathrm{m}} - 1, \nonumber \\
\xi _\mathrm{2h} &amp;=&amp; b \,\, \xi _{\mathrm{mm}}.
\end{eqnarray*}$$Here, ξ1h and ξ2h are the so-called ‘1-halo’ and ‘2-halo’ terms in the halo model (Cooray & Sheth 2002), ρNFW(r|Mh, c) is the NFW density profile of halo mass Mh and concentration c, b is the average halo bias, and ξmm is the non-linear matter–matter autocorrelation function predicted at Planck cosmology (Takahashi et al. 2012). Zu et al. (2014) found that this simple model provides an adequate description of the halo–matter cross-correlation measured from simulations at the level of a few per cent on scales of our concern (i.e. below the halo radius). For massive haloes, several studies found that the average density profile in the inner region of haloes deviates from the NFW shape and the Einasto profile is more accurate (Dutton & Macciò 2014; Klypin et al. 2016). However, since we are only fitting to scales above $0.1\, h^{-1}\, {\mathrm{Mpc}}$, the difference between NFW and Einasto should be negligible. We have also ignored the extra lensing effect caused by the stellar mass of the BCGs, which has negligible contribution on scales above $0.1\, h^{-1}\, {\mathrm{Mpc}}$. The uncertainties of the predicted ΔΣ can be as large as $10{{\ \rm per\ cent}}$ around the transition between the 1-halo and 2-halo scales ($2{-}4\, h^{-1}\, {\mathrm{Mpc}}$; see the fig. 5 of Zu et al. 2014), but we are only concerned with the measurement of halo mass and concentration, which are inferred primarily from scales below the transition scale, and the large-scale bias, which is estimated from scales above $10\, h^{-1}\, {\mathrm{Mpc}}$. Additional, our statistical uncertainties of weak lensing at those scales is around ${~}10{{\ \rm per\ cent}}$. Therefore, the impact of relatively large uncertainties in the transitional regime on our conclusions should be negligible. For future modelling of cluster weak lensing signals of per cent level uncertainties, the transitional behaviour can be potentially improved by using the scheme proposed by Garcia et al. (2020) and emulators developed by Salcedo et al. (2020).
Citation: Cooray & Sheth 2002
Citation function: Uses
We investigate the possibility of our clouds corotating within a ‘cold-flow disc’, in which material that is accreting on to the galaxy arises due to filamentary structure from the cosmic web to form an extended disc structure. As done previously by Kacprzak et al. (2010b), we investigate this by employing a simple halo model from Steidel et al. (2002) with the updated values for i and position angle (PA). In order to determine if all the low-, intermediate-, and high-ionization clouds can be explained by a corotating thick disc, we set the lagging scale height term to hv = 1000 kpc. This term is related to the location above the mid-plane where the lagging halo velocity component begins to matter and vlos becomes dominated by an exponentially declining term. Following previous studies (e.g. Steidel et al. 2002; Kacprzak et al. 2010a,b, 2011b, 2019; French & Wakker 2020; Nateghi et al. 2021) we are forcing the disc to be a thick, nearly rigid rotator, where material far above the disc plane is rotating along with the disc. This places our low- and high-ionization clouds high above the galactic disc at a high angle. The line-of-sight distance (Dlos) and line-of-sight velocity (vlos) can be found in Fig. 9. The figure shows that much of the low-, intermediate-, and high-ionization clouds (except Cloud #1) can be explained by a corotating thick disc, but the low-metallicity subsystem cannot be. Using such a large value for hv may be extreme, but it clearly illustrates the contrast between the low-metallicity subsystem and the other gas phases. In short, it is not at all possible that the low-metallicity subsystem is due to a cold-flow disc. Using a more plausible (smaller) value for hν will not affect the disparity between the low-metallicity subsystem and the other gas phases. We find that even by adopting a non-physically motivated scale height of hv = 10 kpc (adopted from Ho et al. 2017), we can still explain the metal-enriched clouds with a corotating thick disc (except Cloud #1 and Cloud #2), but we are still unable to explain the low-metallicity subsystem. This agrees with recent simulations in which high angular momentum gas accretes via a cold-flow disc, creating large corotating gaseous structures in the galaxy halo (Stewart et al. 2013, 2017). The Q0454−220 system also agrees with observations of galaxy/absorber pairs, which have been found to exhibit disc-like or accretion kinematics over a large range of impact parameters (e.g. Burchett et al. 2013; Bouché et al. 2016; Ho et al. 2017; Rahmani et al. 2018).
Citation: Kacprzak et al. 2010a
Citation function: Uses; Uses
The origin of the extreme magnetic fields of magnetars is still an important open question as several scenarios have been proposed, such as a convective dynamo in a fast rotating proto-neutron star (PNS) (Thompson & Duncan 1993; Raynaud et al. 2020; Raynaud, Cerdá-Durán & Guilet 2022), the Tayler-Spruit dynamo following fallback (Barrere et al. 2022), amplification in main-sequence stellar mergers (Schneider et al. 2019), or the fossil field scenario (Ferrario & Wickramasinghe 2006). To shed light on this question, it is important to assess the efficiency of each of the dynamo mechanisms in the conditions specific to a PNS. The impact of different ingredients on the MRI has been studied in recent years: the shear parameter (Masada et al. 2012), neutrino viscosity and drag (Guilet, Müller & Janka 2015), stable stratification (Guilet & Müller 2015; Reboul-Salze et al. 2021b), and spherical geometry (Reboul-Salze et al. 2021a). A remaining open question is the dependence on diffusion coefficients. Local disc simulations have shown that the efficiency of the MRI is strongly correlated to the magnetic Prandtl number Pm (the ratio of viscosity to resistivity) when Pm ∼ 0.1−16 (Fromang et al. 2007; Lesur & Longaretti 2007; Simon & Hawley 2009; Longaretti & Lesur 2010; Shi, Stone & Huang 2016; Potter & Balbus 2017) and plateaus at Pm  0.1 with an imposed external magnetic field (Meheut et al. 2015). Low and high Pm regimes are computationally challenging because a high numerical resolution is needed to resolve the small viscous (low Pm) or resistive (high Pm) scale. The regime of low Pm has attracted particular attention because it is relevant to most regions of accretion discs as well as to liquid metals (laboratory experiments and Earth core), but the opposite regime of large Pm has not been targeted specifically by previous studies of the MRI. This regime is relevant to PNSs (Thompson & Duncan 1993; Guilet et al. 2015; Lander 2021), neutron star merger remnants (Rossi, Armitage & Menou 2008), which exhibit physical conditions similar to PNSs (Guilet et al. 2017), interstellar and intergalactic media (Schekochihin et al. 2004), and to the inner parts of some accretion discs (Balbus & Henri 2008; Potter & Balbus 2014, 2017). In the latter case, the Pm dependence of the MRI may drive an instability leading to variability in the accretion rate (Potter & Balbus 2014, 2017; Kawanaka & Masada 2019).
Citation: Reboul-Salze et al. 2021b
Citation function: Background; Background
We use the Hβ λ4861 velocities to standardize the SNe II because the majority of the high redshift spectra (SDSS-II and SNLS samples) are noisy and taken at early phases where the Fe ii absorption lines are not visible. Errors on the Hβ velocities were obtained by measuring many times the minimum of the absorption changing the continuum fit. Both quantities are listed in Table 1. To find the best epoch to use the SCM we need the velocities for different epochs. As proposed by Hamuy (2001) and used in all SNe II cosmology works (Nugent et al. 2006; Poznanski et al. 2009; D’Andrea et al. 2010; Olivares et al. 2010; Poznanski et al. 2010; Rodríguez et al. 2014; de Jaeger et al. 2015) we do an interpolation/extrapolation using a power law of the form
9





where A and γ are two free parameters obtained by least-squares minimization for each individual SN and t is the epoch since the explosion. To derive the velocity error following the work done by de Jaeger et al. (2015), a Monte Carlo simulation is performed, varying randomly each velocity measurement according to the observed velocity uncertainties over more than 2000 simulations. Following Poznanski et al. (2009), we add in quadrature to the velocity uncertainty of every SN II a value of 150 km s−1 to account for unknown host-galaxy peculiar velocities. For the SNe II with one spectrum the same power law is used but this time with a fixed γ, which is derived using only the CSP-I sample for which we have many spectra per SN and a better fit can be achieved. We find a median value of γ = −0.407 ± 0.173. It is important to note that in the majority of other SN II cosmology works, the authors used the same power law for all of the SNe, whereas in our work the γ is different for all SNe II with more than two spectra. Additionally, in Section 6.4, we show the possibility of using a new relation between A and γ in order to derive the velocity when only one spectrum is acquired without assuming the same power-law exponent.
Citation: Poznanski et al. 2010
Citation function: Uses; Uses; Uses
HD 141569 is a Herbig star classified as a B9-A0 spectral type (Augereau & Papaloizou 2004), with an effective temperature of 9750 ± 250 K, an estimated age of 7.2 ± 0.02 Myr, a luminosity between 16.60 ± 1.07 L⊙ (Vioque et al. 2018) and 27.0 ± 3.6 L⊙ (Di Folco et al. 2020), a mass of 2.14 ± 0.01M⊙, and a Gaia distance of 110 ± 1 pc (Arun et al. 2019)2. It is a non-flaring disk system with little mid-IR excess classified as a group-II source (Meeus et al. 2001). It is the only known pre-main sequence star characterized by a hybrid disk (Wyatt et al. 2015; Péricaud et al. 2017; Di Folco et al. 2020), an evolutionary disk state between the protoplanetary and debris-disk regimes. Near-IR imaging spatially resolved an optically thin disk consisting of two rings located at about ~ 280 and ~ 455 au from the star Augereau et al. (1999); Weinberger et al. (1999); Biller et al. (2015). A more complex system is shown in the visible, consisting of multiple rings and outer spirals that could be explained through perturbations by two nearby (~ 7.5 arcsec) M dwarfs, or by planetary perturbations (Augereau & Papaloizou 2004; Wyatt 2005; Reche et al. 2009). Fisher et al. (2000) found a warm disk component up to 110 au at 10.8 and 18.2 μm, later confirmed by Marsh et al. (2002). The short-wavelength counterpart of this component was detected by Mawet et al. (2017) through L′ imaging and ranging between 20 and 85 au. Emission at 8.6 μm was detected by Thi et al. (2014), and was interpreted as emission from polycyclic aromatic hydrocarbons (PAHs). NOEMA and ALMA observations in the millimeter range showed continuum emission equally shared between a compact (≲ 50 au) and a smooth extended dust component (~350 au), with large millimeter grains dominating the inner regions and smaller grains in the outer ones (Di Folco et al. 2020). Finally, inner disk features were detected by SPHERE in the Y, J, H, and K bands (Perrot et al. 2016) and by Keck/NIRC2 in the L′ band (Currie et al. 2016) at physical separations of 45, 61, and 88 au. These results point out the high morphological complexity of the outer disk in the HD 141569 system.
Citation: Weinberger et al. (1999)
Citation function: Background