adsabs/FOCAL huggingface dataset:
In order to understand this discrepancy, we consider the BL Lac object nature of NGC 1275 (Veron 1978). The BL Lac object properties of 3C 84 have been noted in connection with the synchrotron optical emission that hides the accretion generated continuum. The BL Lac object aspect can be very pronounced with optical polarization that changes dramatically in amplitude and position angle, with the largest optical polarization reaching 6% (Angel & Stockman 1980). Such extreme blazar-like properties suggest a small LOS to the jet axis (Lind & Blandford 1985). We qualify this statement by noting that the evidence does not support extreme BL Lac object behavior. The range of polarization is 1%–6% in Angel & Stockman (1980) and more extensive polarization data9

9
From http://www.bu.edu/blazars/VLBAproject.html.
 covering the time period 2011–2018 indicates an optical polarization that is usually below 2% and extremely rare (2) instances of polarization >3% were reported. This suggests a slightly off angle BL Lac object. The Hβ BEL is considered to be rotating gas in a flat “pancake-like” region in which the normal to the BEL disk is parallel to the jet axis (Wills & Browne 1986). Thus, the FWHM that appears in Equation (4) depends on the LOS. For polar lines of sight, Equation (4) will under estimate Mbh. In the virial mass estimation,
5





where G is the gravitational constant, RBLR is the orbital radius of the BLR (broad line region) and vBLR is the velocity of the BEL gas. In order to relate vBLR to an observed quantity, one defines the de-projection factor, f,
6





Equations (5) and (6) indicate that 





. The de-projection factor has been estimated for various classes of objects that are believed to be differentiated by an LOS (Antonucci 1993; DeCarli et al. 2011). The method of DeCarli et al. (2011) was to estimate Mbh from the bulge luminosity of the host galaxy. Using this estimate to set the value of Mbh in the virial formula, they were able to estimate f for various classes of objects,
7





Equation (4) is derived based on assuming an isotropic distribution of BEL gas velocity. Thus, we adopt a correction factor for the estimate in Equation (4) of 





 for BL Lac object orientations such as the one that exists in 3C 84. Taking the nominal value of DeCarli et al. (2011) in Equation (7), we expect a correction factor of 64. The orientation corrected central black hole mass estimate based on the data in Table 1 yields 





. Alternatively, if we use the blazar correction of 42 associated with the nominal value of fblazars in Equation (7) instead, we get 





.
Citation: Wills & Browne 1986
Citation function: Uses; Uses
Strong gravitational lensing of high-z star-forming galaxies offers a unique way to examine highly magnified molecular gas. The method for selecting strongly lensed dusty galaxy candidates, at z > 1, is primarily based on unusually bright (sub)millimeter fluxes compared to the expected steep drop-off in (sub)millimeter number counts (e.g., Negrello et al. 2007, 2010). This method has since identified a large number across the extragalactic sky, i.e., more than 100 lensed candidates at z > 1 (Ivison et al. 2010; Vieira et al. 2010, 2013; Bussmann et al. 2013, 2015; Wardlow et al. 2013; Weiß et al 2013; Cañameras et al. 2015; Harrington et al. 2016; Strandet et al. 2016; Díaz-Sánchez et al. 2017; Negrello et al. 2017; Bakx et al. 2018). The lensed population of dusty star-forming galaxies selected by the South Pole Telescope (SPT), Herschel Space Observatory, and Planck has now been detected in more than two CO transitions (e.g., Spilker et al. 2016; Strandet et al. 2017; Yang et al. 2017; Bakx et al. 2020; this work). The Herschel-selected, strongly lensed galaxy sample (Bussmann et al. 2013) offered the first systematic approach to producing a statistically significant sample of CO/[C i] lines (Yang et al. 2017), followed by a compilation in 11 Planck- and Herschel-selected lensed galaxies (Cañameras et al. 2018b), including four galaxies with both [C i] lines detected (Nesvadba et al. 2019). The IR-to-CO luminosity relations of local starbursts and high-z star-forming galaxies explored by Greve et al. (2014) indicate that the ISM radiation field is an important component to consider when understanding CO line excitation, yet this investigation was limited to 23 unlensed and 21 lensed dusty star-forming systems—all with more than three frequency measurements of the dust continuum and usually a single CO line detection. Most previous studies used only single- and/or double-component gas-emitting regions to reproduce the observed CO emission, excluding the simultaneous modeling of the available [C i] emission, but also ignoring the role of the dust continuum emission as a heating source of the gas.
Citation: Harrington et al. 2016
Citation function: Background
Studies of molecules play a prominent role in explaining the physical, chemical, and kinematic properties of the interstellar medium (ISM) in galaxies (Omont 2007; Tielens 2013). One such molecule is H2O, the third most abundant molecule in the warm dense ISM after H2 and CO (Neufeld et al. 1995). As an asymmetric rotor with a large electric dipole moment, H2O has a rich and complex spectrum giving rise to emission and absorption lines mainly in the submillimeter and far-infrared (FIR) regimes of the electromagnetic spectrum. Observations from local galaxies (van der Werf et al. 2010; Weiß et al. 2010; Rangwala et al. 2011; Yang et al. 2013), high-redshift ultra-luminous infrared galaxies (ULIRGs; Omont et al. 2013; Yang et al. 2016), and active galactic nuclei (AGNs; van der Werf et al. 2011) have shown H2O emission to be ubiquitous with intensities as bright as CO lines. Modeling has shown that, in addition to infrared pumping where H2O is excited by FIR photons, collisions also contribute to the intensities of low-excitation transitions (e.g., González-Alfonso et al. 2010, 2012). This is best represented in Figure 3 from Liu et al. (2017), which shows prominent H2O lines in different ISM components. The low-excitation lines become weaker or completely disappear in the warm and hot regions (>40 K) where infrared pumping dominates over collisions. The higher excitation transitions that require strong far-infrared radiation density are mainly found in the hotter regions (100–200 K) of the galaxy. The cascading emission lines, 





 (Eup = 100.8 K,νrest = 987.927 GHz), 





 (Eup = 137 K, νrest = 752.033 GHz), and p-H2O (22,0 − 21,1) (Eup = 196 K, νrest = 1228.789 GHz) are pumped by 101 μm photons from the base 11,1 level and are primarily excited in the warm regions of the galaxy. The collisional excitation of the low-lying levels (11,1 and 20,2) in optically thin or high-density hot regions might also contribute to the emission of the 





 line. Hence, H2O transitions probe the infrared radiation field density and physical properties of the ISM such as gas density and kinetic temperature (e.g., Weiß et al. 2010; González-Alfonso et al. 2014; Liu et al. 2017).
Citation: Yang et al. 2013
Citation function: Background; Background
Recent theoretical work has suggested that the presence, or lack thereof, of long-period giant planets could affect the formation of such systems. Batygin & Laughlin (2015) argued that the migration of Jupiter within our own solar system might have disrupted a massive primordial inner protoplanetary disk that could have formed multiple short-period super-Earths; they predicted that systems like the Kepler short-period multiple systems should typically lack long-period giant planets. A related question is, how common are planetary systems broadly similar in architecture to our solar system, with small close-in planets and more distant giant planets? We can begin to answer these questions in the near future through the combination of searches for short-period super-Earths and data from the long-term RV programs that have been monitoring many bright FGK stars for well over a decade. Super-Earths can be found with either high-precision RV observations or space-based transit searches. Such high-precision RV surveys include those being undertaken currently with HARPS (e.g., Díaz et al. 2016), HARPS-N (M15), APF (Vogt et al. 2014), and CHIRON (Tokovinin et al. 2013), and in the near future with MINERVA (Swift et al. 2015), CARMENES (Quirrenbach et al. 2014), ESPRESSO (Mégevand et al. 2014), and SPIRou (Artigau et al. 2014). The major upcoming space-based transit survey is that of TESS (Ricker et al. 2015). Long-term RV programs include the McDonald Observatory Planet Search (e.g., Endl et al. 2016), the Anglo-Australian Planet Search (e.g., Jones et al. 2010), the Lick-Carnegie Exoplanet Survey (e.g., Rowan et al. 2016), the CORALIE planet search (Marmier et al. 2013), and the planet search at ESO (e.g., Zechmeister et al. 2013). Long-period giant planets will also be found by Gaia, which will produce a huge sample of astrometrically detected planets (Perryman et al. 2014). While most of the Kepler sample is too faint to have been observed previously by long-term RV surveys (e.g., Coughlin et al. 2015), Gaia will be able to astrometrically detect long-period planets around many of these stars. Our own McDonald Observatory Planet Search program now has a baseline of 12–15 years for ∼200 FGKM stars, and a handful of stars also have lower precision observations dating back more than 25 years. HD 219134 is one of these stars, and here we present an analysis of our RV observations of this star, as well as our data on the stellar activity.
Citation: Díaz et al. 2016
Citation function: Background
Citation: Tokovinin et al. 2013
Citation function: Background; Background
For instance, due to its dissipative nature, gas can be very efficient in absorbing and transporting outwards the angular momentum of the pair, likely leading to a rapid evolution and eventual coalescence. From observations, as well as numerical simulations, it has been established that in gas-rich galaxy mergers there is a large inflow of gaseous material to the central kiloparsec of the galactic remnant, often resulting in a massive circumnuclear disc (Barnes & Hernquist 1992; Sanders & Mirabel 1996; Mayer et al. 2007). Driven by dynamical friction and global torques from this disc, the pair of MBHs decays very efficiently down to separations of the order of ∼1–0.1 pc, where it forms a gravitationally bound binary (Escala et al. 2004, 2005; Mayer et al. 2007; Fiacconi et al. 2013; del Valle et al. 2015; Roškar et al. 2015). At these sub-parsec scales, most theoretical and numerical studies have focused on the evolution of binaries surrounded by a gaseous circumbinary disc, often either co-rotating (see e.g. Ivanov, Papaloizou & Polnarev 1999; Armitage & Natarajan 2005; Cuadra et al. 2009; Haiman, Kocsis & Menou 2009; Lodato et al. 2009; Nixon et al. 2011; Roedig et al. 2011, 2012; Kocsis, Haiman & Loeb 2012; Amaro-Seoane, Brem & Cuadra 2013; D’Orazio, Haiman & MacFadyen 2013; Muñoz & Lai 2016; Miranda, Muñoz & Lai 2017; Tang, MacFadyen & Haiman 2017) or counter-rotating (see e.g. Roedig & Sesana 2014; Nixon & Lubow 2015) with respect to the binary’s orbital motion. These discs are generally assumed to be well-defined, smooth, and relaxed, with no attempt to link their presence to the gaseous environment around the binary, nor to the fuelling mechanisms that bring gas to the nucleus. Furthermore, all these idealized scenarios are subject to the disc consumption problem, namely, if the disc dissolves through some process (e.g. star formation, active galactic nucleus (AGN)/supernovae feedback, Lupi et al. 2015), the evolution of the binary orbit stops. In fact, the evolution of MBHBs in gas-rich environments is intimately related to the unsolved problem of gas supply to the centre of galactic nuclei.
Citation: Escala et al.; 2005
Citation function: Background
The stellar photospheric parameters we collected from literature for the benchmark stars are summarized in Table A.1. Although most benchmark stars have v sini  2 km s−1 (Reiners et al. 2018), there are two stars with larger values: J07558+833 (12.1 km s−1) and J13005+056 (16.4 km s−1). These stars are useful to investigate the performance of the algorithms when dealing with higher rotational velocities. The literature values were derived with different methods. These methods include: interferometry to estimate the stellar radius and Teff (Boyajian et al. 2012; Ségransan et al. 2003; von Braun et al. 2014; Berger et al. 2006; Newton et al. 2015), synthetic model fitting using BT-Settl models to determine Teff (Gaidos et al. 2014; Lépine et al. 2013; Gaidos & Mann 2014; Mann et al. 2015) and log g (Lépine et al. 2013), empirical relations to derive stellar mass in the form of mass-luminosity relations (Mann et al. 2015; Khata et al. 2020; Boyajian et al. 2012; Berger et al. 2006; Ségransan et al. 2003), along with the mass-magnitude relations (Maldonado et al. 2015), mass-radius relations (von Braun et al. 2014), mass–Teff relations (Gaidos & Mann 2014; Gaidos et al. 2014), empirical relations to derive the stellar radius in the form of mass-radius relations (Maldonado et al. 2015) and Teff–radius relations (Gaidos & Mann 2014; Gaidos et al. 2014; Houdebine et al. 2019), pEW measurements to determine Teff (Maldonado et al. 2015; Neves et al. 2014; Newton et al. 2015) and [Fe/H] (Maldonado et al. 2015; Neves et al. 2014; Gaidos et al. 2014; Mann et al. 2015), the definition of spectral indices such as the H2O-K2 index to estimate Teff (Rojas-Ayala et al. 2012), as well as the combination of the H2O-K2 index with pEWs to derive [Fe/H] (Rojas-Ayala et al. 2012; Khata et al. 2020), the stellar radius and Teff (Khata et al. 2020), and spectral curvature indices for the determination of Teff (Gaidos & Mann 2014). Additionally, [Fe/H] was derived by using color-magnitude metallicity relations (Dittmann et al. 2016), atomic line strength relations (Gaidos & Mann 2014), and spectral feature relations (Terrien et al. 2015). Terrien et al. (2015) used K-band magnitudes and the Dartmouth Stellar Evolution Program (Dotter et al. 2008) to derive the stellar radius, whereas Mann et al. (2015) employed the Boltzmann equation with Teff determined from synthetic model fits. Last, but not least, Houdebine et al. (2019) derived Teff from photometric colors. For more details on the individual methods, we refer to the descriptions in the corresponding works.
Citation: Khata et al. 2020
Citation function: Background; Background
Citation: Khata et al. 2020
Citation function: Background
Citation: Khata et al. 2020
Citation function: Background
The EWs were measured with the methods described in Vásquez et al. (2015). As in Paper I, we used the sum of the EWs of the two strongest CaT lines (λ8542, λ8662) as a metallicity estimator, following the Ca II triplet method of Armandroff & Da Costa (1991). Different functions have been tested in the literature to measure the EWs of the CaT lines, depending on the metallicity regime. For metal-poor stars ([Fe/H] ≲ −0.7 dex) a Gaussian function was used with excellent results (Armandroff & Da Costa 1991), while a more general function (such as a Moffat function or the sum of a Gaussian and Lorentzian, G + L) is needed to fit the strong wings of the CaT lines observed in metal-rich stars (Rutledge et al. 1997; Cole et al. 2004). Following our previous work (Gullieuszik et al. 2009; Saviane et al. 2012) we have adopted here a G+L profile fit. To measure the EWs, each spectrum was normalised with a low-order polynomial using the IRAF continuum task, and set to the rest frame by correcting for the observed radial velocity. The two strongest CaT lines were fitted by a G+L profile using a non-linear least squares fitting routine part of the R programming language. Five clusters from the sample of Saviane et al. (2012) covering a wide metallicity range were re-reduced and analysed with our code to ensure that our EWs measurements are on the same scale as the template clusters used to define the metallicity calibration. Figure 5 shows the comparison between our EWs measurements and the line strengths measured by Saviane et al. (2012) (in both cases the sum of the two strongest lines) for the five calibration clusters. The observed scatter is consistent with the internal errors of the EW measurements, computed as in Vásquez et al. (2015). The measurements show a small deviation from the unity relation, which is more evident at low metallicity. A linear fit to this trend gives the relation: ΣEW(S12) = 0.97 ΣEW(this work) + 0.21, with an rms about the fit of 0.13 Å. This fit is shown in Fig. 5 as a dashed black line. For internal consistency, all EWs in this work have been adjusted to the measurement scale of Saviane et al. (2012) by using this relation. In Table 3 we provide the coordinates, radial velocities, and the sum of the equivalent widths for the cluster member stars, both measured (“m”) and corrected (“c”) to the system of Saviane et al. 2012.
Citation: Saviane et al. 2012
Citation function: Uses
Citation: Saviane et al. 2012
Citation function: Uses
Citation: Saviane et al. (2012)
Citation function: Uses
Citation: Saviane et al. (2012)
Citation function: Compare/Contrast; Compare/Contrast; Similarities
Citation: Saviane et al. (2012)
Citation function: Uses; Uses
Citation: Saviane et al. (2012)
Citation function: Uses; Uses
Citation: Saviane et al. 2012
Citation function: Uses; Uses
The Orion Bar is an ideal source for probing the excitation and chemistry of molecules in PDRs, thanks to its close distance of 414 pc (Menten et al. 2007) and its well-known structure and geometry. The Orion Bar is located between the Orion molecular cloud and an Hii region illuminated by the Trapezium cluster. The FUV radiation field of the Trapezium cluster at the location of the Orion Bar is equivalent to (1−4) × 104χ0 in Draine (1978) units. Its orientation changes from face-on to nearly edge-on where the molecular emissions peak. The observations presented in this paper also correspond to the nearly edge-on orientation part of the Orion Bar. The geometrical enhancement of the column densities toward the nearly edge-on part of the Orion Bar was derived by multiple studies and is in the range between 4 and 20. The tilt angle compared to a completely edge-on orientation was suggested to be 3° in the model of Hogerheijde et al. (1995) and Jansen et al. (1995). A tilt angle of 3° is equivalent to an enhancement factor of 20 for the measured column densities. Based on Oi 1.317 μm emission, Walmsley et al. (2000) find a model that requires a geometrical enhancement factor of 5 to convert the observed column densities into face-on values. Neufeld et al. (2006) find the geometrical enhancement factor to be 4 based on measured C+ column densities. Using a clumpy 3D PDR model, Andree-Labsch et al. (2014) successfully reproduced the Orion Bar stratification using a clumpy edge-on cavity wall, and they claim that a model of a convex filament fails to describe the structure of the Orion Bar. The average kinetic temperature was estimated to be 85 K (Hogerheijde et al. 1995). Closer to the ionization front, higher temperatures are also measured; for example, OH transitions observed with Herschel/PACS are consistent with 160−220 K gas (Goicoechea et al. 2011) and CH+ observations with temperatures around 500 K (Nagy et al. 2013). Part of the molecular line emission measured toward the Orion Bar corresponds to an “interclump medium” with densities between a few 104 and 2 × 105 cm-3 (Simon et al. 1997). It has been suggested that other molecular lines originate in clumps with densities in the range between 1.5 × 106 and 6 × 106 cm-3 (Lis & Schilke 2003).
Citation: Jansen et al. (1995)
Citation function: Uses; Uses
Remnant BHs of Population III stars with ${\sim }10^{2}\, \mathrm{M}_{\odot }$ are one of the candidates for such seeds (e.g. Madau & Rees 2001). They can grow SMBHs in the available time if either with continuous accretion at the Eddington limit or with short episode of super-Eddington growth. In reality, however, the BH’s growth can easily be hindered by its own radiative feedback (Milosavljević, Couch & Bromm 2009; Park & Ricotti 2011; Orofino, Ferrara & Gallerani 2018; Sugimura et al. 2018). An alternative and attractive pathway for seed BH formation is via the so-called direct collapse (e.g. Bromm & Loeb 2003), where a supermassive star (SMS) with ${\sim } 10^{5}\, \mathrm{M}_{\odot }$ collapses into a BH with a similar mass by general relativistic instability (Umeda et al. 2016; Woods et al. 2017; Haemmerlé et al. 2018). Here, SMSs are supposed to form from a primordial gas in some peculiar sites. Unlike in ordinary first star formation, which is driven by H2 cooling (see e.g. Bromm & Larson 2004; see also Glover 2013 for a review), H2 is dissociated by strong external far-ultraviolet (UV) radiation from nearby galaxies and the contraction of clouds is solely caused by H atomic cooling in the most intensively studied channel for SMS formation (Omukai 2001; Omukai, Schneider & Haiman 2008; Shang, Bryan & Haiman 2010; Regan, Johansson & Wise 2014; Sugimura, Omukai & Inoue 2014). Such clouds contract almost isothermally at ∼104 K without experiencing vigorous fragmentation. The protostar formed at the centre subsequently accretes the gas at a high rate of $10^{-1}\, \mathrm{M}_{\odot }\, \mathrm{yr}^{-1}$ due to high temperature (Latif et al. 2013; Inayoshi, Omukai & Tasker 2014; Becerra et al. 2015; Chon, Hosokawa & Yoshida 2018). Such rapidly accreting protostar inflates greatly in radius with effective temperature of several 1000 K and grows supermassive with ${\gtrsim } 10^{5}\, \mathrm{M}_{\odot }$ avoiding ionization radiation feedback on the accretion flow (Hosokawa, Omukai & Yorke 2012; Hosokawa et al. 2013), before collapsing by general relativistic instability.
Citation: Umeda et al. 2016
Citation function: Compare/Contrast; Compare/Contrast
SN2006jc, PS15dpn and other narrow-line SNe. Two out of the three SNe we considered above are super-luminous, however the final collapse of a PPI+CC progenitor or PISNe does not need to be superluminous (Woosley 2017). The PPI is just one possible mechanism to create CSM, which can produce extreme luminosities by generating radiation from the kinetic energy of the ejecta and/or narrow emission lines (even if the luminosity does not reach extreme values). The detection of narrow H lines determines the classification of a SN as a type IIn, while the detection of narrow He emission lines determines the classification as type Ibn. Both kinds of event are too common to be entirely explained with PPI+CC progenitors, and it is likely that both classes contain events with a diversity of physical mechanisms (e.g., Pastorello et al. 2008 but see also Hosseinzadeh et al. 2017). Nevertheless, it is possible that at least some of these events might correspond to the observational counterpart of the death of PPI+CC progenitors. In particular, our simulations can produce several solar masses of H-free CSM moving at a few thousand km s−1, which correspond to the width of the He lines detected in some SN Ibn without any fine-tuning required. Even if the detection of narrow lines alone is not sufficient to associate a specific SN to a PPI event, combining evidences from previous coincident transients, large ejecta masses or long lightcurve durations, large 56Ni yields, an extremely young surrounding stellar population, and/or nucleosynthetic signatures might strengthen the case for associating specific event with this scenario. Possible examples of SN Ibn that might correspond to PPI+CC are SN2006jc and PS15dpn. The former showed relatively narrow He lines possibly hinting to asphericity of the CSM (Foley et al. 2007) and was spatially coincident with an unexplained outburst two years earlier (e.g., Pastorello et al. 2007; Foley et al. 2007). For the latter, Wang & Li (2019) proposed to fit the light curve by combining CSM interaction and radioactive decay, and inferred CSM and 56Ni masses of ∼0.8 M⊙ and ∼0.1 M⊙, respectively, in good agreement with our models.
Citation: Hosseinzadeh et al. 2017
Citation function: Compare/Contrast