adsabs/FOCAL huggingface dataset:
Alongside studies of their emission properties, examining the environments of FRBs on subparcsec to kiloparsec scales can be equally informative. Thus far, only a fraction of known extragalactic repeating FRBs have been localized to host galaxies (Chatterjee et al. 2017; Tendulkar et al. 2017; Marcote et al. 2020; Macquart et al. 2020; Heintz et al. 2020; Bhandari et al. 2021; Bhardwaj et al. 2021a, 2021b; Li et al. 2021).17

17
FRB 20201124A is the fifth announced repeating FRB with a host galaxy. In total, there are eight such FRB sources known as of 2021 August.
 Much closer by, FRB-like emission has been detected from a Milky Way magnetar, SGR 1935 + 2154 (CHIME/FRB Collaboration et al. 2020; Bochenek et al. 2020). All identified repeating FRB hosts have evidence for low to modest ongoing star formation rates of ∼0.06−2 M⊙ yr−1 (Gordon et al. 2004; Bhandari et al. 2020a; Heintz et al. 2020), several exhibit spiral arm morphologies (Bhardwaj et al. 2021a; Mannings et al. 2021; Tendulkar et al. 2021), and their stellar populations span a range of stellar masses, ∼108−1010.5 M⊙ (de Blok et al. 2008; Bhandari et al. 2020a; Heintz et al. 2020; Mannings et al. 2021). At face value, these characteristics, coupled with the absence of any quiescent host galaxy identifications for repeating FRBs, may indicate that FRBs are connected to host galaxies with ongoing star formation. However, studies of their more local environments reveal a rich diversity. For instance, the discovery of the repeating FRB 20200120E in an old (9.1 Gyr) globular cluster on the outskirts of the grand design spiral galaxy M81 (Bhardwaj et al. 2021a; Kirsten et al. 2021a) and the detection of the repeating FRB 20121102A embedded in a star-forming knot in its dwarf host galaxy (Bassa et al. 2017) seemingly represent polar opposite local environments. If all repeaters discovered to date originate from the same type of progenitor, then the progenitor model must accommodate the observed diversity of both local and galactic environments. Additionally, any connection in progenitors to the population of apparent nonrepeaters remains opaque.
Citation: Tendulkar et al. 2021
Citation function: Background
After checking the results of the U-net in the test area, we applied the same detection method to the full-sky SZ prediction map, with a detection threshold of pmax = 0.1 in order to recover the maximum number of Planck_z clusters. We detected 20 204 sources in the full-sky map with the U-net with pmax = 0.1. We compared the detections with the three catalogues of known galaxy clusters, Planck_z, Planck_no-z, and MCXCwP. Among the 20 204 detected sources, 98.5% of the Planck_z clusters are recovered, together with 76.4% of Planck_no-z clusters, and 20.8% of MCXCwP clusters. Moreover, 11 cluster are identified by ACT (Hasselfield et al. 2013) and 98 clusters are identified by the South Pole Telescope (SPT; Bleem et al. 2015), but are not included in the Planck PSZ2 catalogue. This means that 18 415 sources do not belong to any of the catalogues. We investigated the nature of the sources detected with the U-net. First, we cross-matched the sample of 18 415 sources with Planck point sources. Only 6.1% are matched within a cross-match radius of 5 arcmin with the positions of the Planck catalogue of galactic cold cores, and only 0.2% are matched with the positions of the Planck sources identified at 353 GHz. Second, we stacked at their positions 16 maps in different wavelengths, each of them potentially probing different galaxy cluster counterparts. Some of the 16 maps are also based on Planck data and thus are not independent, but some of the maps are independent and may show indications of galaxy cluster counterparts in other wavelengths, that is, in near-infra-red (where galaxies emit) and in X-rays (where the same gas emits as is detected with the SZ). The 16 maps are the Planck SZ MILCA map, the 6 Planck HFI frequency maps, the IRIS map at 100 μm (Miville-Deschênes & Lagache 2005), the CMB lensing map (based on Planck; Planck Collaboration VIII 2020), 4 galaxy over-density maps of all galaxies (called GAL ALL), passive galaxies (called GAL P), transitioning galaxies (called GAL T), and active galaxies (called GAL A) following Bonjean et al. (2019), star formation rate density maps (called SFR), a stellar mass density maps (called Mstar) constructed with the method from Bonjean et al. (2019), and finally, the ROSAT X-ray map (ByoPiC4 product). The Planck maps were masked from the Planck Catalogue of Compact Sources (PCCS, Planck Collaboration XXVI 2016), and the ROSAT map was masked from the point sources detected in ROSAT, Chandra, and XMM-Newton (Boller et al. 2016; Evans et al. 2010; Rosen et al. 2016, respectively). The result of the 16 stacks are shown in Fig. 5.
Citation: Bleem et al. 2015
Citation function: Uses; Differences
In previous work (Rogers & Safi-Harb 2016, hereafter RSH16), we studied a parametrized phenomenological model for describing magnetic field growth in neutron star (NS) evolution. Assuming that the external dipole field is buried by an intense process of fall-back accretion after the formation of the NS (Muslimov & Page 1995, 1996), the slow diffusion from the stellar surface exerts a time-dependent torque on the NS (Geppert, Page & Zannias 1999). This process may provide an explanation for the observed braking indices of young pulsars with n ≠ 3 (Espinoza 2012; Archibald et al. 2016), in contrast to the prediction of a rotating magnetic dipole in vacuum (n = 3; Ostriker & Gunn 1969). Magnetic field growth provides an evolutionary link between apparently disparate classes of NS, such as the extremely weak field1 subset of NSs known as central compact objects (CCOs), high magnetic field pulsars (HBPs) and the X-ray dim isolated NSs (XDINSs). On the other hand, magnetic field decay has been invoked to describe the evolution of the anomalous X-ray pulsars (AXPs) and soft gamma repeaters (SGRs). These objects are conventionally described by the magnetar model that posits the development of large magnetic fields via dynamo action from rapid rotation early in the life of the NS, necessary to support the core of the massive progenitor from collapse (Duncan & Thompson 1992; Akiyama et al. 2003; Thompson, Quataert & Burrows 2005). However, the distinction between the apparently rotation-powered HBPs, with dipole fields just above the quantum critical limit, BQED = 4.4 × 1013 G, and the magnetically powered AXPs and SGRs (B ≥ BQED) is significantly blurred after the HBPs were observed displaying magnetar-like activity (Gavriil et al. 2008; Kumar & Safi-Harb 2008; Göğüs et al. 2016), and radio emission was observed from magnetars (Camilo et al. 2007). Moreover, the SNRs associated with the AXPs and SGRs show evidence of ‘typical’ explosion energies (≤1051 erg) and not superenergetic as one would expect from a rapidly rotating proto-NS (Vink & Kuiper 2006; Kumar, Safi-Harb & Gonzalez 2012; Safi-Harb & Kumar 2013; Kumar et al. 2014). These complications are difficult to reconcile with the standard magnetar picture.
Citation: Archibald et al. 2016
Citation function: Compare/Contrast; Compare/Contrast
In the left-hand panel of Fig. 2, we compare the morphological types assigned by the STAGES collaboration for the galaxies in the whole OMEGA sample and the jellyfish candidates sample. The sample of jellyfish galaxies (JC345) is composed mainly of late-type spirals and irregulars. In the middle panel of Fig. 2, we show the distribution of SED types for both samples. Based on the SED types of the galaxies, of the 70 jellyfish galaxy candidates analysed, 66 were found to be part of the blue cloud and 4 as being dusty reds (IDs: 11633, 17155, 19108, 30604). However, contrary to what could be expected, dusty red galaxies are only a small portion of our sample of jellyfish candidates. One reason why we may not detect many dusty reds as jellyfish galaxies might be due to the fact that these galaxies, despite having relatively high SFRs (only four times lower than that in blue spirals at fixed mass, Wolf et al. 2009), have significant levels of obscuration by dust that might hamper the identification of the jellyfish signatures. Another reason for that is that we selected jellyfish galaxy candidates within a parent sample of H α-emitting galaxies that already had a low fraction of dusty red galaxies (≈15 ${{\ \rm per\ cent}}$). As these galaxies have low star formation, it is harder to perceive the morphological features of RPS. Dusty red galaxies have been previously studied in this same system (Wolf et al. 2009), and RPS was suggested to be the main mechanism acting in these galaxies (Bösch et al. 2013). While in Bösch et al. (2013) one of the main pieces of evidence suggesting the action of enhanced RPS were the existence of disturbed kinematics without disturbed morphologies, in our study we strongly base our selection on such morphological distortions. Both our jellyfish galaxy candidates and the dusty red galaxies show different characteristics that can be correlated to the effect of RPS. Nevertheless, they might be tracing different stages of the same phenomenon, where dusty red galaxies have more regular morphologies, but disturbed kinematics. Our sample of morphologically disturbed jellyfish galaxy candidates may be showing the stage where the features of RPS are the most visible and the SFRs are enhanced.
Citation: Wolf et al. 2009
Citation function: Uses
The first interstellar discovery of MF took place in 1975 from the microwave emission spectrum of Sgr B2, which agrees with the rotational constants of the syn isomer (Curl 1959; Brown et al. 1975; Nummelin et al. 2000). Subsequently, it was detected in several other sources such as comet Hale–Bopp, and protostars and hot corinos (e.g., in Orion A and the protoplanetary nebula CRL 618) (Ellder et al. 1980; Ikeda et al. 2001; Cazaux et al. 2003; Bottinelli et al. 2004a, 2004b; Remijan et al. 2005, 2006; Remijan & Hollis 2006). Later, the less stable trans rotamer was detected in both the laboratory and the interstellar medium (ISM; Ruschin & Bauer 1980; Blom & Günthard 1981; Müller et al. 1983; Neill et al. 2011, 2012). The detailed mechanistic route for the interstellar synthesis of MF is a subject of considerable debate (Horn et al. 2004; Herbst 2005; Garrod & Herbst 2006; Snyder 2006; Bennett & Kaiser 2007; Occhiogrosso et al. 2011; Lawson et al. 2012). The initially assumed gas-phase model involving an ion–molecule reaction of protonated methanol (





) with formaldehyde (H2CO) fails to reproduce the large MF column density, because this pathway demands a significant activation barrier (∼128 kJ mol−1) for one of the intermediate steps (Horn et al. 2004). Instead, gas-phase radiative association between the methyl cation (





) and formic acid (HCOOH) seems energetically promising to produce H+MF and thereafter neutral MF via dissociative recombination with electrons:
1







2





However, this reaction scheme fails to reproduce the huge MF concentration observed in hot cores (Horn et al. 2004). Alternative routes suggested more recently involve ion–molecule reactions with a low barrier or no barrier between protonated methanol and formic acid (acid-catalyzed Fischer esterification) or between protonated formic acid and methanol (





 transfer):
3







4





followed by reaction (2) or exothermic proton transfer of 





 (H+MF) to a base with a higher proton affinity than MF (Ehrenfreund & Charnley 2000; Neill et al. 2011, 2012).
Citation: Occhiogrosso et al. 2011
Citation function: Background
We compute the emission line luminosity using the following equation:
1



LCO′=3.25×107SCOΔvDL2(1+z)3νobs2Kkms−1pc2,

where ν
obs is the observed CO transition frequency, D

L
 is the luminosity distance, and S
COΔv is the line-integrated flux in units of Jy km s−1. We convert the observed CO transition luminosity into CO (1–0) luminosity (L




CO(1−0)′

) by assuming that the low-J CO transitions are thermalized and optically thick, so 



LCO4–3′=LCO3–2′=LCO1–0′

. Using the ratios (



rJ1=LCOJ→J−1′/LCO1–0′

) from Carilli & Walter (2013) with r
31=0.97 and r
41 = 0.87 did not significantly change our results. Furthermore, in 3C 298, we found that the molecular gas is consistent with being thermalized and optically thick (Vayner et al. 2017). These physical conditions are consistent with what is found for Mrk 231 (Feruglio et al. 2015). Finally, we convert the 



LCO1–0′

 line luminosity into molecular gas mass using the CO-to-H2 conversion factor: α
CO with units of (K km s−1pc2)−1. For sources where we do not detect any narrow CO emission at the systemic redshift, we place a limit on the molecular gas mass over an aperture equal to the beam size for an emission line with a velocity FWHM of 250 km s−1. The molecular gas mass limits can be linearly scaled with a different α
CO value. For sources with detected molecular gas at the quasar’s systemic redshift, we compute the radius of the molecular gas region, which allows us to measure the gas surface density. All radii are computed using a curve-of-growth method. Effective radii refer to a region that encloses 50% of the flux, while “maximum” extent refers to a size scale that encloses 90% of the flux. In all cases, narrow emission at the systemic redshift of the quasar is spatially resolved by our observations. We deconvolve the size of the beam from all radius measurements. For sources with no detected CO emission, we use the molecular gas mass limit and the beam of the observations as a proxy for the radius. Values associated with the molecular gas at the systemic redshift are summarized in Table 3.
Citation: Feruglio et al. 2015
Citation function: Similarities
Two spectroscopically confirmed X-ray AGNs in the cluster core (#607, 661 in G13) are suitable candidates for ionizing the nebula. The depth of the new Chandra observation, coupled with an optimal on-axis alignment, allowed us to perform a basic X-ray spectral analysis despite the limited photon statistics (34 and 20 net counts in the observed 0.5–7 keV band for sources #607 and 661, respectively). Source #607 is characterized by a power-law spectrum with photon index Γ = 2.0 ± 0.6; the observed 2–10 keV flux is 1.7+1.1−0.6 × 10−15 erg cm−2 s−1, corresponding to a rest-frame 2–10 keV luminosity of 





 erg s−1, typical of a luminous Seyfert galaxy. The X-ray spectrum of source #661, the point-like Lyα emitter (Figure 2), is flat: fitting the data with a power-law model provides 





, highly indicative of strong obscuration. We then included an absorption component and fixed the photon index to 1.8, as expected for the intrinsic AGN emission (e.g., Piconcelli et al. 2005). This model results in a column density of 





 cm−2, i.e., consistent with marginal Compton-thick absorption (1.5 × 1024 cm−2). The tentative detection of an iron Kα emission line at 6.4 keV (with equivalent width of ≈2.4 keV rest frame), if confirmed, would further support the heavily obscured nature of source #661. The derived 2–10 keV flux is (7.4 ± 2.2) × 10−15 erg cm−2 s−1, corresponding to a rest-frame luminosity of 





 erg s−1, placing source #661 in the quasar regime. We do not detect any bright counterpart in deep Jansky Very Large Array observations at 3 GHz down to 2.7 μJy (rms), and we thus classify source #661 as radio-quiet. From aperture photometry, we estimated a Lyα flux of 





 erg cm−2 s−1, corresponding to a luminosity of (1.9 ± 0.2) × 1042 erg s−1. The spectral energy distribution (SED) of #661 is shown in Figure 7. From SED modeling, which benefits from near-, mid-, and far-IR observations from Spitzer and Herschel, we estimated a bolometric luminosity for the AGN of (2.7 ± 1.5) × 1045 erg s−1. A similar value (3.2 ± 0.6 × 1045 erg s−1) is derived using the observed [O iii]λ5007 Å luminosity obtained from recent Subaru/MOIRCS spectroscopy of the galaxy (Valentino et al. 2015), converted into a bolometric luminosity as 





 (Heckman et al. 2004). Assuming the luminosity-dependent bolometric correction as in Lusso et al. (2012), we predict an intrinsic 2–10 keV luminosity for source #661 of 





 erg s−1. This value is consistent, within the uncertainties due to the adopted relations and measurements, with that derived from the X-ray spectral analysis reported above.
Citation: Piconcelli et al. 2005
Citation function: Uses
While the velocity field is very regular, the velocity dispersion looks more complex. Instead of one peak extended along a southeast-northwest axis (Figs. 10, and 11), we find three (Fig. 13). The central one, elongated along the main axis of the galaxy and not perpendicular to it, is accompanied by two additional regions with wider lines to the northeast and southwest, separated by a total of 5.′′ 9 ± 0.′′2. At these locations we observe the transition between rigid body rotation and steep rotation curve to a flat one (see also Chou et al. 2007). At a radius of 2.′′45 (≈45 pc) and with a rotation velocity of 140 km s−1 we obtain with Eq. (1) of Mauersberger et al. (1996; η = 1) an enclosed mass of M2.45 = 2.1 × 108 M⊙ with an estimated error of 10%, in good agreement with Cunningham & Whiteoak (2005). For comparison, Roy et al. (2010) obtain with the H92α line 3 × 107 M⊙ for the mass inside a radius of 1″ (≈19 pc). Furthermore, we note that the nuclear region is slightly lopsided: The center of the line connecting the two outer peaks of line width is located ≈1.′′4 southwest of the continuum peak (see Table 3). Systemic velocities (Vbarycentric ≈ 571 km s−1) are found about 0.′′9 (≈15–20 pc) southwest of the continuum peak and may be closer to the position of the maser disk (Greenhill et al. 1997) given in Sect. 3.1. While uncertainties with respect to the value of the systemic velocity (Sect. 2) and the relative positions of the continuum peak and the maser disk are significant, most of the star formation represented by the continuum emission (see Bendo et al. 2016) appears to originate slightly northeast of the dynamical center. The lopsidedness of the central region of NGC 4945 also explains why molecular lines with significant absorption show integrated intensity peaks slightly shifted to the southwest: With the bulk of the continuum arising from the northeast, line emission is likely more quenched by absorption at this side of the center than in the southwestern part of the inclined nuclear disk (see Sect. 4.1.2 and Fig. 7).
Citation: Chou et al. 2007
Citation function: Similarities
A key question that remains unanswered is how BBHs are formed. Viable formation channels include isolated binary evolution (e.g., Bethe & Brown 1998; Belczynski et al. 2002, 2014, 2016; Dominik et al. 2013; Mennekens & Vanbeveren 2014; Spera et al. 2015; Eldridge & Stanway 2016; Mandel & de Mink 2016; Marchant et al. 2016; Mapelli et al. 2017; Stevenson et al. 2017; Barrett et al. 2018; Giacobbo & Mapelli 2018; Kruckow et al. 2018; Mapelli & Giacobbo 2018) and dynamical encounters in dense stellar environments, such as globular clusters (e.g., Portegies Zwart & McMillan 2000; O’Leary et al. 2006; Sadowski et al. 2008; Downing et al. 2010, 2011; Rodriguez et al. 2015, 2016a, 2016b; Askar et al. 2017; Fragione & Kocsis 2018; Rodriguez & Loeb 2018; Samsing 2018; Samsing et al. 2018; Zevin et al. 2019), young star clusters (e.g., Banerjee et al. 2010; Ziosi et al. 2014; Mapelli 2016; Banerjee 2017, 2018; Di Carlo et al. 2019; Kumamoto et al. 2019), and galactic nuclei (e.g., O’Leary et al. 2009; Antonini & Perets 2012; Antonini & Rasio 2016; Petrovich & Antonini 2017; Stone et al. 2017a, 2017b; Rasskazov & Kocsis 2019). Moreover, the dynamical process known as Kozai–Lidov (KL) resonance (Kozai 1962; Lidov 1962) can trigger the merger of a BBH, even if the BBH has not been formed in a dense star cluster. In fact, if the BBH is orbited by a tertiary body (i.e., the BBH is the inner binary of a stable hierarchical triple system), the KL mechanism triggers oscillations of the BBH’s eccentricity, which might speed up the merger by gravitational-wave emission. Each channel is expected to produce black hole mergers with different mass and spin distributions (Mandel & O’Shaughnessy 2010; Abbott et al. 2016a; Rodriguez et al. 2016c; Farr et al. 2017; Abbott et al. 2019d). The limited statistics from the low number of systems detected through gravitational waves and model uncertainties so far do not allow strong constraints on the formation channels.
Citation: Rodriguez et al. 2015
Citation function: Background; Background
Ideally, however, one would want to go beyond the description of cosmic global history, and trace galaxy evolution on a galaxy-by-galaxy basis to understand the physical processes driving it. In this respect, great progress has been made by surveys at different redshifts that have established the existence of a strong dependence of galaxy histories on galaxy stellar mass. On average, more massive galaxies have formed their stars and completed their star formation activity at higher z than less massive galaxies (the so-called downsizing effect, Cowie et al. 1996; Gavazzi et al. 2006; De Lucia et al. 2007; Sánchez-Blázquez et al. 2009). The existence of relations between SFR and galaxy stellar mass (SFR–Mass) and specific star formation rate and mass (sSFR = SFR/Mass) have been established from z = 0 out to z > 2 (Brinchmann et al. 2004; Daddi et al. 2007; Noeske et al. 2007; Salim et al. 2007; Rodighiero et al. 2011; Whitaker et al. 2012; Sobral et al. 2014; Speagle et al. 2014), and many other galaxy properties have been found to be correlated with galaxy mass. Furthermore, a number of works have pointed out that galaxy properties are even more strongly correlated with a combination of galaxy mass and galaxy ‘size’, arguing for velocity dispersion (Bernardi et al. 2003; Franx et al. 2008; Smith, Lucey & Hudson 2009; Wake, van Dokkum & Franx 2012) or galaxy surface mass density (Brinchmann et al. 2004; Kauffmann et al. 2006) as principal drivers. The exact origin of these trends is still unknown, but evidence has accumulated for a dependence of galaxy stellar population ages on galaxy sizes at fixed mass (van der Wel et al. 2009; Cappellari et al. 2012; Poggianti et al. 2013), suggesting that also galaxy structure, and not just stellar mass, is relevant. In a recent paper, Omand, Balogh & Poggianti (2014) argue that the observed correlation of the quenched fraction with M/R1.5 is related to the dominance of the bulge component with respect to the disc, suggesting it might ultimately be linked with galaxy morphology (see also Driver et al. 2013). Even the sSFR–Mass relation might be due to the increase of the bulge mass fractions with galaxy stellar mass, as the ratio of SFR and stellar mass of the galaxy disc is virtually independent of total stellar mass (Abramson et al. 2014).
Citation: Brinchmann et al. 2004
Citation function: Background