adsabs/FOCAL huggingface dataset:
For normal single AGNs, the bolometric luminosity, 





, can usually be estimated from either optical or X-ray luminosities. However, for our targets, the hard X-ray luminosities are much lower than for single AGNs with the same [O iii] luminosities (Paper I). Three possible reasons are: (1) the observed X-rays are low due to high gas absorption (e.g., Bassani et al. 1999); (2) the X-rays are intrinsically weak due to a high accretion state that radiates little coronal emission (e.g., Desroches et al. 2009; Dong et al. 2012); (3) there are significant [O iii] luminosity excesses due to shocks in our galaxy mergers (e.g., Dopita & Sutherland 1995). The emission-line diagnostic ratios featuring [O iii]/Hβ and [N ii]/Hα show that our targets are well in the regime of Seyferts, with only one displaying relatively large uncertainty (Liu et al. 2010a). It is unlikely that the emission lines of our AGNs are mainly powered by shocks, as they do not lie in the region of LINERs on the diagnostic diagrams (Ho et al. 1993). It is more likely that our targets are highly obscured rather than in a high accretion state, since the estimated Eddington ratios (Table 4; see discussion below) are not close to unity. Additionally, our targets are all detected by the Wide-field Infrared Survey Explorer (Wright et al. 2010) with λLλ(12 μm) ∼ 1044 erg s−1. The X-ray–MIR ratios of our target binary AGNs (≲0.01) are much smaller than those of normal AGNs (∼0.3; Horst et al. 2008), indicating that the X-ray emission of the AGNs is absorbed and reprocessed into IR emission. As shown by Heckman et al. (2005) at low redshift, selection by narrow optical emission lines will recover most AGNs selected by hard X-rays (with the exception of BL Lac objects). On the other hand, selection by hard X-rays misses a significant fraction of the local AGN population with strong emission lines. In view of the uncertainties associated with X-ray absorption in gas-rich mergers, we estimate the bolometric luminosities of the binary AGNs from their extinction-corrected [O iii] luminosity (Paper I),
3





with intrinsic scatter 0.4 dex (see Trump et al. 2015 and references therein). The results are listed in Table 4. We adopt extinction-corrected rather than observed [O iii] luminosity (e.g., Stern & Laor 2012) to estimate the bolometric luminosity, since the former provides a smaller scatter.
Citation: Ho et al. 1993
Citation function: Uses
The Stokes V profile of maser A+B from our 2009 observations (lower panel of Figures 2 and 3) reveals an S-shaped structure that is usually taken to be a detection of the Zeeman effect (also see the discussion in Section 4 below); so does the Stokes V profile from our 2017 observations (lower panel of Figures 4 and 5). No other compact maser spots in this field show any such feature. Whenever an S-shaped feature is observed in Stokes V, the magnetic field strength is usually determined by fitting a numerical frequency derivative of the Stokes I spectrum to the Stokes V spectrum; details are in, e.g., Momjian & Sarma (2017). The Stokes V profile is usually fit simultaneously to the derivative of the I profile and a scaled replica of the I profile itself via the equation (Troland & Heiles 1982; Sault et al. 1990)
1





The scaled replica of the I spectrum is included in the fit to account for small calibration errors in RCP versus LCP; for all results reported in this paper, a ≲ 10−3. The magnetic field values are contained in the fit parameter b, which is equal to z 





, where z is the Zeeman splitting factor, and 





 is the line-of-sight (LOS) magnetic field strength (assuming, of course, that the signature in Stokes V is due to the magnetic field in the region; see Section 4.2). We used the AIPS task ZEMAN (Greisen 2015) to carry out the fit in Equation (1). This task allows multiple Gaussian components in I to be fitted simultaneously to V, with each Gaussian component fitted for a different b, and hence a different LOS magnetic field strength. For the three Gaussian components of maser A+B from our 2009 observations (listed in Table 2, and shown in the upper panel of Figure 2), the derivative profiles scaled by the respective values fitted for b = z 





 are each shown in the lower panel of Figure 2; the composite profile obtained by adding together these three scaled derivative profiles is shown in the lower panel of Figure 3. Of these three components for maser A+B from our 2009 observations, only component 2 showed a significant fit, with the fitted value given by z 





 = 152 ± 12 Hz. Following customary practice in the field of Zeeman observations, we consider fits to be significant only if the ratio of fitted value to the fitted error is at the 3σ level or greater. Meanwhile, for the three Gaussian components of maser B from our 2017 observations (listed in Table 2, and shown in the upper panel of Figure 4), the derivative profiles scaled by the respective fitted values for 





 are each shown in the lower panel of Figure 4, and the composite profile is shown in the lower panel of Figure 5. Only component 2 of maser B from our 2017 observations showed a significant fit, with the fitted value given by z 





 = 149 ± 19 Hz. Since component 2 of maser A+B from our 2009 observations and component 2 of maser B from our 2017 observations are very likely the same maser spot (see Section 4.1), this is a remarkable coincidence in 





 over observations taken eight years apart.
Citation: Greisen 2015
Citation function: Uses; Uses; Background
A number of semi-analytic models (SAMs) have attempted to reproduce submm number counts (e.g. Granato et al. 2000; Fontanot et al. 2007; Somerville et al. 2012). One such model is the galform (SAM), which has been tuned to successfully reproduce the number counts of 850 μm and $\mathrm{1.1 \, mm}$ selected galaxies.1 However, in order to achieve this good agreement galform invokes a top-heavy initial mass function (IMF). Early versions of the model used a flat IMF above $1 \, \mathrm{M_{\odot }}$, in sub-L* mergers (Baugh et al. 2005; Swinbank et al. 2008). This is required to produce sufficiently bright submm emission during frequent low-mass merger events. Later versions of the model used a more moderately top-heavy IMF in starbursts, triggered by disc instabilities rather than mergers, and found similarly good agreement with the number counts (Cowley et al. 2015, 2019; Lacey et al. 2016; Park et al. 2016). However, such IMF variability is still controversial, particularly extreme forms and any dependence on merger state (Bastian, Covey & Meyer 2010; Hopkins 2013; Krumholz 2014), and is inconsistent with the constraints on the IMF in massive star-forming galaxies that are significantly less extreme (e.g. Tacconi et al. 2008), though there is tentative evidence of a bottom-light/top-heavy IMF in both local star-forming region analogues (Motte et al. 2018; Schneider et al. 2018) and some gravitationally lensed high-redshift starbursts (Zhang et al. 2018). Safarzadeh, Lu & Hayward (2017) showed that a variable IMF is degenerate with a number of other modelling processes in SAMs, such as the form of stellar feedback. They highlight that taking in to account dust mass allows for a good fit to the number counts without resorting to a variable IMF. Most recently, the shark SAM (Lagos et al. 2018) is able to broadly reproduce the 850 μm counts (whilst slightly overestimating the bright-end counts compared to S2CLS; Geach et al. 2017) using a fixed Chabrier (2003) IMF (Lagos et al. 2019). They attribute the good agreement to their use of physically motivated attenuation curves obtained from a self-consistent galaxy evolution model (eagle; Trayford et al. 2020).
Citation: Hopkins 2013
Citation function: Compare/Contrast
The fireball model remains the most popular scenario for the gamma-ray burst (GRB) phenomenon (Cavallo & Rees 1978; Goodman 1986; Paczynski 1986; Shemi & Piran 1990; Rees & Mészáros 1992, 1994; Mészáros & Rees 1993). In this model, the GRB central engine is a stellar-mass black hole or a rapidly spinning and highly magnetized neutron star formed by either the collapse of a supermassive star (collapsar; Woosley 1993; MacFadyen & Woosley 1999; Woosley & Heger 2006) or the merger of two compact objects (Paczynski 1986; Fryer et al. 1999; Rosswog 2003). In both cases, the original explosion creates a bipolar collimated jet composed mainly of photons, electrons, positrons, and a small fraction of baryons. The relativistic explosion ejecta within the jet are not homogeneous—they form multiple high density layers, which propagate at various velocities. When the fastest layers catch up with the slowest, the charged particles contained in the layers are accelerated through mildly relativistic collisionless shocks (internal shocks; Rees & Mészáros 1994; Kobayashi et al. 1997; Daigne & Mochkovitch 1998). The particles subsequently cool via emission processes such as synchrotron, Synchrotron Self Compton (SSC), and Inverse Compton (IC). The internal shock phase is usually associated with the so-called GRB prompt emission,17

17
See Pe’er (2015) for a recent review of GRB prompt emission.
 mainly observed in the keV−MeV energy range (see, e.g., the spectral catalogs by Gruber et al. 2014; von Kienlin et al. 2014) and usually lasting from a few ms up to several tens to hundreds of seconds. As the ejecta interact with the interstellar medium they slow down via relativistic collisionless shocks (external shocks; Rees & Mészáros 1992; Mészáros & Rees 1993), accelerating charged particles, which then emit non-thermal synchrotron photons. This external shock phase is usually associated with the so-called GRB afterglow emission observed at radio wavelengths up to X-rays and in some cases even up to the GeV regime hours after the prompt phase, and days and even years for the lowest frequencies. The detailed origin of the gamma-ray emission, however, is not fully understood and many theoretical difficulties remain, such as the composition of the jet, the energy dissipation mechanisms, as well as the radiation mechanisms (e.g., Zhang 2011).
Citation: Rees & Mészáros; 1994
Citation function: Background
The GALEX NUV observations span a wide wavelength range, from 1771 to 2831 Å. While these data provide no spectral information within that bandpass, we rely on solar and stellar flare studies to inform the likely contributors to the flare flux. The NUV spectral region has not had as many observational constraints as the far-UV region in flare studies. Few NUV stellar flare spectra exist at all, and the few that do were obtained either for solar flares or on nearby M dwarfs. Flares observed in the UV are often associated with the more impulsive phases of solar flares, starting with early observations showing a close temporal association between UV and hard X-ray emission (Cheng et al. 1981). Welsh et al. (2006) reported on high time-resolution NUV+FUV flares seen with GALEX on nearby M dwarfs, and Hawley et al. (2007) presented high spectral-resolution NUV flare measurements of an M-dwarf flare. From these two studies, the contribution of emission lines relative to continuum emission could be determined; the main emission lines in the flare NUV spectrum were Mg ii, Fe ii, Al iii, and C iii. However, the main emission component overall was a continuum component. Recent results from solar flares observed from space (Heinzel & Kleint 2014; Kleint et al. 2016) demonstrate an NUV spectrum originating largely from Hydrogen Balmer continuum emission. The formation of the NUV emission appears to originate from an impulsive thermal and nonthermal ionization caused by the precipitation of electron beams through the chromosphere. This explains the temporal correlation with solar flare hard X-ray emission observed previously. More recently, Kowalski et al. (2019) presented accurately calibrated NUV flare spectra at high time cadence on an M dwarf, again finding a large flux enhancement due to continuum radiation. They commented that the oft-used 9000 K blackbody used to describe blue-optical stellar flare emission (Hawley et al. 2003) under-predicts the NUV continuum flare flux by a factor of two. Based on general similarities in radiative properties between solar and stellar flares studied thus far (Osten 2016), it is likely that a combination of line and continuum emission enhancements are present in the NUV flare flux from the flares being considered, but we cannot speculate about the relative contribution of one versus the other. These sources originate from different layers of the stellar atmosphere: singly and doubly ionized emission lines likely originate in the chromosphere, and the Balmer continuum emission also originates from the chromosphere. Some lines, such as Mg ii, exhibit absorption components and self-reversals, indicating optical depth effects in the atmosphere, while other lines such as Fe ii appear to be optically thin. Any hot blackbody emission might originate further in the photosphere.
Citation: Heinzel & Kleint 2014
Citation function: Compare/Contrast; Compare/Contrast
Galaxies seem to retain a memory of their spin orientation with respect to the cosmic web filaments and walls, as suggested by the results from large-scale cosmological hydrodynamical simulations (Dubois et al. 2014; Codis et al. 2018; Wang et al. 2018; Ganeshaiah Veena et al. 2019; Kraljic, Davé & Pichon 2020b). The mass dependence of the spin alignment signal is however debated. While some works confirmed the existence of a galaxy spin transition from parallel to perpendicular with respect to the filament’s direction (Dubois et al. 2014; Codis et al. 2018; Kraljic et al. 2020b), and analogously with respect to walls (Codis et al. 2018; Kraljic et al. 2020b), others (Ganeshaiah Veena et al. 2019; Krolewski et al. 2019) found preferential perpendicular orientation with respect to filaments at all masses with no sign of a spin transition. A possible interpretation of this lack of detection of a clear transition is the nature of the filaments, with galaxies in thinner filaments having their spins more likely perpendicular to the filament’s axis, compared to galaxies of similar mass in thicker filaments (Ganeshaiah Veena et al. 2019). This can be in turn understood recalling the multiscale nature of the problem and the conditional TTT (Codis et al. 2015) predicting larger transition mass for denser, thus thicker, filaments. Further support for this interpretation was provided by the findings of stronger impact of large-scale tides on the galaxy spin orientation in denser filaments (Kraljic et al. 2020b, using filament density as a proxy for the thickness of filaments). In addition to the stellar mass, the spin-filament alignment was shown to depend on other internal properties of galaxies. Blue or rotation-supported galaxies were found to dominate the alignment signal at low stellar mass, while red or dispersion-dominated galaxies tend to show a preferential perpendicular alignment (Codis et al. 2018; Wang et al. 2018; Kraljic et al. 2020b).
Citation: Kraljic et al. 2020b
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Citation: Kraljic et al. 2020b
Citation function: Compare/Contrast
It is well known that giant molecular clouds (GMCs) are the major gas reservoirs for star formation (SF) and the sites where essentially all stars are born. Understanding the properties of GMCs is thus key to unravelling the interplay between gas and stars within galaxies. Early GMC studies were restricted to our own Milky Way (MW) and the late-type galaxies (LTGs) in our Galactic neighbourhood (e.g. Engargiola et al. 2003; Rosolowsky 2005, 2007; Rosolowsky et al. 2007; Gratier et al. 2012; Colombo et al. 2014; Wu, Sakamoto & Pan 2017; Faesi, Lada & Forbrich 2018), where GMCs have relatively uniform properties and generally follow the so-called Larson relations (between size, velocity dispersion, and luminosity; e.g. Blitz et al. 2007; Bolatto et al. 2008). However, more recent studies of other local galaxies have raised doubts on the universality of cloud properties. The cloud properties in some LTGs (such as M51 and NGC 253) vary with galactic environment and do not universally obey the usual scaling relations (e.g. Hughes et al. 2013; Leroy et al. 2015; Schruba, Kruijssen & Leroy 2019). The first study of individual GMCs in an early-type galaxy (ETG; NGC 4526) has also clearly shown that the clouds in that galaxy do not follow the usual size–linewidth correlation and tend to be more luminous, denser and to have larger velocity dispersions than the GMCs in the MW and other Local Group galaxies (Utomo et al. 2015). The differences in NGC 4526 may be due to a higher interstellar radiation field (and/or cloud extinctions), a different external pressure relative to each cloud’s self-gravity, and/or different galactic dynamics. GMCs in ETGs seem to have shorter orbital periods and be subjected to stronger shear/tidal forces, analogous to the highly dynamic environment in the MW central molecular zone (CMZ; e.g. Dale, Kruijssen & Longmore 2019; Henshaw et al. 2019; Kruijssen et al. 2019). Although we are entering an era of large surveys of GMC populations (e.g. Sun et al. 2018), current samples of ETGs are still very limited. More studies of GMCs in varied LTGs and ETGs are thus required to provide a comprehensive census of GMC properties across different galaxy environments.
Citation: Leroy et al. 2015
Citation function: Background
The CO photodesorption yield reaches its highest value when this ice is deposited at low temperatures (down to 7 K, the lowest temperature studied experimentally) and decreases gradually at higher deposition temperatures (Öberg et al. 2007; Öberg et al. 2009; Muñoz Caro et al. 2010, 2016; Sie et al. 2022). The explanation for this phenomenon motivated further research. It was found that the columnar structure of CO ice samples, grown at incidence angles larger than 45°, increases the effective ice surface exposed to UV photons and therefore the photodesorption efficiency (González Díaz et al. 2019), but ice surface effects cannot account for the large variations observed in the photodesorption of CO ice samples deposited at different temperatures (Muñoz Caro et al. 2016). Absorption band shifts of CO ice in the UV and IR ranges only occurred at deposition temperatures above 20 K (Lasne et al. 2015; Muñoz Caro et al. 2016), suggesting that CO ice grown at lower temperatures is amorphous below 20 K in our experiments, and therefore, the decreasing photodesorption yield is not related to a transition from amorphous to crystalline ice, instead it might be associated to a different degree of molecular disorder in CO ice samples, depending on their deposition temperature. Photon energy transfer via Wannier-Mott excitons between the first photoexcited molecule in the ice and a molecule on the ice surface capable to desorb was proposed (Chen et al. 2017; McCoustra & Thrower 2018). Molecular disorder seems to enhance this energy transfer between neighbour molecules. The colour temperature variations measured at different deposition temperatures could also be the result of molecular disorder (Carrascosa et al. 2021). Urso et al. (2016), Cazaux et al. (2017), and Carrascosa et al. (2021) did not find significant changes in the desorption behaviour or the colour temperature of pure CO ice during controlled warm-up, which points to a low value of the diffusion in the ice. Finally, Sie et al. (2022) investigated the CO photodesorption yield dependence on ice thickness.
Citation: Muñoz Caro et al. 2016
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Sagittarius A* (Sgr A*) is a supermassive black hole system that allows one to observationally test the aforementioned GRMHD models of accretion flows (Goddi et al. 2017). Millimeter-Very Long Baseline Interferometry (mm-VLBI) is capable of resolving the shadow of the event horizon (Falcke et al. 2000), making this an ideal laboratory not only to tests Einstein’s General Theory of Relativity (GR) but also to investigate electron acceleration in the vicinity of a black hole. Most of the radiative models for Sgr A*, which are based on post-processing GRMHD simulations, assume that electrons have a thermal, relativistic (Maxwell–Jüttner) distribution function, and that the proton-to-electron temperature ratio is constant across the simulation domain (Goldston et al. 2005; Noble et al. 2007; Mościbrodzka et al. 2009; Dexter et al. 2010, 2012a; Shcherbakov et al. 2012). When the proton-to-electron temperature is constant, the disk dominates the images and spectra since most of the matter resides there. We have recently extended these radiative models by making the temperature ratios a function of the plasma β parameter, where 

$\beta = \frac{P_{\textrm{gas}}}{P_{\textrm{B}}}$


β=


P

gas




P
B





 is the ratio of gas to magnetic pressures. In these extended models, the electrons are hotter in the more magnetized plasma, which is usually outflowing from the system. The reason for this is that the previously mentioned models do not recover the flat radio spectra. The β parameterization enforces that the disk emission is suppressed by significantly decreasing the temperature of the electrons in those regions. As a consequence of this, the jet will be the dominant source of emission. These modifications to the electron temperature model allowed us to recover some basic observational characteristics of Sgr A* (a roughly flat radio spectral slope and a size vs. wavelength relationship that is in agreement with observations) (Mościbrodzka & Falcke 2013; Mościbrodzka et al. 2014; Chan et al. 2015b,a; Gold et al. 2017). Our model for the electron temperatures as a function of the β plasma parameter is now roughly confirmed with extended-GRMHD simulations that self-consistently take into account the evolution of the electron temperatures (Ressler et al. 2015, 2017). Moreover, GRMHD simulations with the new electron temperatures can naturally explain the symbiosis of disks and jets observed in many accreting black hole systems (Falcke & Biermann 1995; Mościbrodzka et al. 2016a).
Citation: Goldston et al. 2005
Citation function: Uses
Irregular refraction (scattering) of radio waves due to density turbulence in the solar corona broadens the angular size of the background “radio Sun” and the compact radio sources there at any given observing frequency. But the extent of broadening and the interferometer baselines required to image the solar atmosphere in detail at radio frequencies remain a puzzle, particularly at frequencies 100 MHz (Erickson 1964; Aubier et al. 1971; Riddle 1974; McMullin & Helfer 1977; Robinson 1983; Melrose & Dulk 1988; Subramanian & Sastry 1988; Thejappa & Kundu 1992, 1994; Sastry 1994; Schmahl et al. 1994; Ramesh & Sastry 2000; Ramesh 2000; Ramesh et al. 2001; Kathiravan et al. 2002; Bastian 2004; Subramanian 2004; Thejappa & MacDowall 2008). Past/present solar-dedicated radio imaging instruments in the above frequency range had/have limited angular resolution (>3′; Wild 1967; Labrum 1972; Kundu et al. 1983; Ramesh et al. 1998). Interestingly, observations carried out during different solar eclipses (similar to the lunar occulation technique) and the bandwidth, duration of some of the solar radio transients indicate that radiowave emitting sources with angular sizes ≲1′ are present/observable in the solar atmosphere from where radio emission at the above frequencies originate (Letfus et al. 1967; Melrose 1980; Benz & Wentzel 1981; McConnell 1983; Ramesh et al. 1999; Ramesh & Ebenezer 2001; Kathiravan et al. 2011; Ramesh et al. 2012). These are consistent with the theoretical predictions that solar radio sources of angular sizes ≲10″ and ≲30″ are observable at the typical frequencies of 327 MHz and 80 MHz, respectively (Riddle 1974; Subramanian & Cairns 2011). Encouraged by the above results and the availability of logistics support, we conducted interferometeric observations of the solar corona at frequencies 100 MHz on a ≈200 km long interferometer baseline for the first time, the description of which is the subject matter of this article.
Citation: Melrose & Dulk 1988
Citation function: Motivation