adsabs/FOCAL huggingface dataset:
With the advent of the LOw Frequency ARray (LOFAR; Röttgering et al. 2011; van Haarlem et al. 2013) which combines a large field of view with high sensitivity on both small and large angular scales, we can now study the FIRC at lower frequencies where the contribution from thermal free-free emission is even less important than at 1.4 GHz. Operating between 30 and 230 MHz, LOFAR offers complementary information to the wealth of data collected at higher frequencies. Using deep LOFAR 150 MHz observations in the 7 deg2 Boötes field (Williams et al. 2016), Calistro Rivera et al. (2017) studied the FIRC at 150 MHz from z ∼ 0.05 out to z ∼ 2.5. They found fairly mild redshift evolution in the logarithmic IR to radio luminosity ratio in the form of qIR ∼ (1 + z)−0.22 ± 0.05. However, if the FIRC is non-linear (i.e. the logarithmic slope is different from one), then it implies that the qIR parameter would depend on luminosity. Therefore the reported redshift dependence of qIR may simply be a consequence of the non-linearity of the FIRC (Basu et al. 2015) as the mean SFR of galaxies is generally larger at higher redshifts (e.g., Hopkins & Beacom 2006; Madau & Dickinson 2014; Pearson et al. 2018; Liu et al. 2018; Wang et al. 2019). Based on LOFAR observations of the Herschel Astrophysical Terahertz Large Area Survey (H-ATLAS; Eales et al. 2010) 142 deg2 North Galactic Pole (NGP) field (Hardcastle et al. 2016), Gürkan et al. (2018) found that a broken power-law (with a break around SFR ∼1 M⊙ yr−1) compared to a single power law is a better calibrator for the relationship between RC luminosity and SFR, possibly implying additional mechanisms for generating cosmic rays and/or magnetic fields. Also using LOFAR data in the NGP field, Read et al. (2018) found evidence for redshift evolution of the FIRC at 150 MHz. Heesen et al. (2019) studied the relation between radio emission and SFR surface density using spatially resolved LOFAR data of a few nearby spiral galaxies. They found a sublinear relation between the resolved RC emission and the SFR surface densities based on GALEX UV and Spitzer 24 μm data.
Citation: Calistro Rivera et al. (2017)
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The recent report by Buisson et al. (2020b) of the detection of Type-I bursts in Sw J1858 shows the presence of a neutron star primary. Therefore, here, we first consider whether Eddington-limited accretion in such a neutron star LMXB might explain the observed radio behaviour in Sw J1858. The neutron star LMXBs with the highest mass accretion rates are the Z-sources, which are thought to accrete near or at the Eddington luminosity, tracing out Z-shaped tracks in their X-ray colour–colour diagrams (Hasinger & van der Klis 1989; Homan et al. 2010). Z-sources can show strong changes in radio brightness, related to the branch they are positioned on in their colour–colour diagram track; they are radio brighter and more variable in the Horizontal Branch than in the Flaring and Normal Branches (Penninx et al. 1988; Hjellming et al. 1990a,b; Spencer et al. 2013; Motta & Fender 2019). Time-resolved radio studies of the different branches in Sco X-1 by Hjellming et al. (1990b) and Cyg X-2 by Hjellming et al. (1990a) show that these sources have similar levels of radio variability and luminosity to Sw J1858 during their radio-faint Flaring and lower Normal Branches. However, in the X-ray – radio luminosity plane, these sources are located to the right of Sw J1858, around the neutron star Eddington limit of 2 × 1038 erg s−1. For Sw J1858 to be similar to these sources, it would therefore have to be viewed a high inclination, reducing its observed X-ray flux and masking the Z-source variability properties through obscuration. However, a clear radio difference between Sw J1858 and the Z-sources is the radio spectral index: contrary to Sw J1858, Z-sources typically show steep spectra and are indeed associated with the launch of (resolved) discrete ejecta (Motta & Fender 2019). Also, in this scenario, Sw J1858 should not have resided in the much more radio bright Horizontal Branch during any of the observations, which is unlikely given the time-scales and commonness of transitions between the branches (Homan et al. 2010).
Citation: Homan et al. 2010; Homan et al. 2010
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In general, the Hall time-scale for magnetic field evolution depends on the strength of the magnetic field, as seen in equation (4). For young NSs with fields below 1014 G, this time-scale may be longer than the observed SNR age. Therefore, magnetic field growth does not have a dramatic effect on these young NSs. However, many of these systems are observed to have a braking index n  3 (Espinoza 2012; Archibald et al. 2016). A possible explanation for these low braking indices may be through the emission of a relativistic particle wind (Thompson & Blaes 1998; Harding, Contopoulos & Kazanas 1999; Tong et al. 2013). However, the conclusive detection of wind nebulae around magnetars in particular is challenging due to the presence of dust-scattering haloes that accompany these X-ray bright, heavily absorbed objects (Esposito et al. 2013; Safi-Harb 2013). Only a handful of such nebulae have been proposed to be associated with highly magnetized NSs. For example, a wind nebula has been proposed to surround the magnetar Swift J1834.9–0846 in W41 (Younes et al. 2016), and the luminosity of a particle wind was estimated for SGR 1806–20 based on the X-ray and radio observations of the wind-powered nebula G10.0–0.3 (Thompson & Duncan 1996; Marsden, Rothschild & Lingenfelter 1999; Gaensler et al. 2005). In the pulsar wind model, relativistic particles load the magnetosphere with charge and distort the dipole field at large scales outside of the light cylinder. Besides affecting the NS spin-down, the emission of a relativistic wind can also offer an explanation for the significant timing noise that generally affects magnetar observations (Tsang & Konstantinos 2013). The HBPs J1119–6127 and J1846–0258 are clearly associated with pulsar wind nebulae (Gavriil et al. 2008; Kumar & Safi-Harb 2008; Ng et al. 2008; Safi-Harb & Kumar 2008; Safi-Harb 2013), suggesting that particle wind emission should play an important role in their evolution. We also expect that AXP and SGR evolution may be affected by wind emission due to candidate wind nebulae, but do not consider these models for the CCOs that do not show any evidence of PWN. However, we note that braking exclusively due to a steady particle wind produces a torque with a braking index n = 1, too low for the NSs with secure SNR associations in Table 1.
Citation: Tong et al. 2013
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The B/T ratio of galaxies has been traditionally measured by modelling their surface brightness profile. Such modelling allows one to find the light-profiles of individual photometric components of galaxies and then measure their properties such as luminosity, shape, size etc. The earliest such works aimed at photometric decomposition of galaxies involved fitting for the azimuthally averaged 1D surface brightness profile (Kent 1985). However, this method leads to systematic errors as creation of 1D light profile by either averaging over the galaxy image or by taking a single cut along the major axis do not properly take into account non-axisymmetric features like a bar or isophotal twist (for more details see Gao & Ho 2017). Due to these shortcomings, techniques to fit the whole 2D image of the galaxy, pixel-by-pixel, were developed (Wadadekar, Robbason & Kembhavi 1999). There are now a number of implementations of 2D bulge disc decomposition available – gim2d: Simard (2010), galfit: Peng et al. (2002), imfit: Erwin (2015), profit: Robotham et al. (2016), budda: de Souza, Gadotti & dos Anjos (2004), gasp2d: Méndez-Abreu et al. (2008) – which fit the 2D images of galaxies. They usually employ different algorithms and methods to find the best-fitting model for a galaxy image. In the last two decades, some pipeline codes have been developed (pymorph: Vikram et al. (2010), galapagos: Barden et al. (2012)) to carry out bulge disc decomposition for large galaxy samples, in a mostly automated fashion. This 2D bulge-disc decomposition approach, while being powerful, is time consuming. The employed fitting algorithm and code carries out a search in parameter space defined by various model parameters and then finds the best fit. Most of the time, one also needs to carry out additional steps (masking other objects in the field and providing PSF images) before the actual fitting of galaxies can commence. Thus estimating the B/T ratio for a single galaxy involves a significant amount of time and computational cost. More importantly, the fitting procedure scales only linearly with the size of the galaxy sample. This limitation will be a major problem in the era of next generation sky surveys (e.g. LSST, Euclid sky surveys) which will observe several billion galaxies.
Citation: Gao & Ho 2017
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The X-ray transient MAXI J1820+070 was first detected on 2018 March 11 (Kawamuro et al. 2018) by the Monitor of All-sky X-ray Image (MAXI, Matsuoka et al. 2009) and was associated with the optical transient ASASSN-18ey (Denisenko 2018; Tucker et al. 2018). In the X-rays, the source flux exceeded 3 Crabs (Bozzo et al. 2018; Mereminskiy et al. 2018), and in the optical, the source reached a magnitude of mV = 12 − 13 (Littlefield 2018; Russell et al. 2018). The parallax of the source π = 0.3 ± 0.1 mas was presented in the Gaia DR2 catalogue (Gaia Collaboration 2018). This corresponds to a distance of 



3
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9

−
1.3


+
3.3




$ 3.9^{+3.3}_{-1.3} $


 kpc (Gandhi et al. 2018a). This unusually bright event allows a detailed investigation of multi-wavelength spectral and timing properties. The course of the outburst was monitored in radio (Trushkin et al. 2018; Polisensky et al. 2018), sub-millimeter (Tetarenko et al. 2018a), optical (Baglio et al. 2018), X-rays (Uttley et al. 2018), and γ-rays (Bozzo et al. 2018; Kuulkers et al. 2018). Because the object was sufficiently bright even for small telescopes, the target was almost continuously monitored, and a rich variety of phenomena was observed. Fast variability and powerful flares in the optical and infrared (Littlefield 2018; Sako et al. 2018; Gandhi et al. 2018b; Casella et al. 2018), optical, and X-ray quasi-periodic oscillations (Mereminskiy et al. 2018; Yu et al. 2018a,b; Buisson et al. 2018; Zampieri et al. 2018) as well as low linear polarisation (Berdyugin et al. 2018) were detected in the source. A 17 h photometric period was recently reported (Patterson et al. 2018) and was tentatively associated with the orbital or superhump period (previously, the source showed a 3.4 h periodicity, Richmond 2018). The X-ray spectral and timing properties as well as the optical-to-X-ray flux ratio suggests that the source is a black hole binary (Baglio et al. 2018; Mereminskiy et al. 2018).
Citation: Tetarenko et al. 2018a
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It is generally believed that reionization started first in high-density regions, where the first luminous objects formed first. In the ‘bubble model’ of reionization (Furlanetto, Zaldarriaga & Hernquist 2004), the amount of star formation and the resulting ionizing photons are estimated using the excursion set model. Based on this idea, the so-called ‘seminumerical simulations’ have been developed. For example, Mesinger, Furlanetto & Cen (2011) developed the 21cmFAST,1 to simulate the evolution of the 3D density, ionization, and 21cm brightness temperature fields efficiently. The ‘bubble model’ considers spherical regions of increasingly smaller scales and identifies ionized bubbles by comparing the cumulative number of ionizing photons produced within the region with the number consumed in reionization process. It has been demonstrated that the statistical predictions of the ‘bubble model’ and the 21cmFAST agree fairly well with radiative–hydrodynamic simulations (Zahn et al. 2007; Mesinger et al. 2011; Zahn et al. 2011), at least when recombination, feedback, etc. are ignored. However, during most epochs of reionization, the topology of the ionization field is much more complicated than the isolated bubbles configuration. The bubbles start to connect with each other as early as when the global ionized fraction is just about 10 per cent, and the Universe starts the percolation process when the ionized fraction gets ∼30 per cent (see e.g. Furlanetto & Oh 2016; Chen et al. 2018). Inspired by the bubble model, the 21cmFAST determines the ionization state of each point by comparing the expected ionizing photon production in the surrounding region with the required number, but it allows non-spherical geometry for the ionized regions. In order to give a better description of the evolution of neutral regions after percolation, Xu et al. (2014) developed the so-called ‘island model’, assuming isolated neutral islands topology. The island model also takes into account an ionizing background that is inevitable during the late EoR (Furlanetto & Oh 2005; McQuinn, Oh & Faucher-Giguère 2011; Emberson, Thomas & Alvarez 2013). Based on the ‘island model’, a seminumerical code named islandFAST was developed to mimic the islands evolution during the last stage of reionization (Xu et al. 2017). In the islandFAST, the effect of small-scale absorbers is taken into account empirically by adopting a fitting formula for the evolution of mean free path (MFP) of ionizing photons (Songaila & Cowie 2010), based on the observed number density of Lyman limit systems up to redshift 6. Before the completion of reionization, the MFP is limited by both the underdense islands and the overdense absorbers. The evolution of the ionization field and the intensity of the ionizing background are derived self-consistently by an iterative procedure to ensure convergence in the total effective MFP of the ionizing photons.
Citation: Zahn et al. 2011
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Over the last decade, the radial and vertical dependences of the metallicity-alpha-element distribution have been studied in more and more detail with increasingly larger samples (e.g., Bensby et al. 2011; Anders et al. 2014; Nidever et al. 2014; Hayden et al. 2015; Queiroz et al. 2020). Figure 6 is mostly consistent with similar plots shown in the above papers. In the inner 10 kpc, it displays two over-densities, a high alpha-element (here [Mg/Fe]), and a low one. Between Rg = 6 and 10 kpc, the two over-densities define two different sequences. In Appendix E, we note that when the sample is restricted to a ±500 pc layer around the Galactic plane, two close but separated sequences are observed in the Rg ∈ [4, 6] kpc interval. Because of their scale height (Bovy et al. 2012), kinematics (Bensby et al. 2003), and age properties (Haywood et al. 2013), these two sequences are associated with the thick disc (high-alpha) and thin disc (low-alpha), respectively. Moving inward of Rg = 4 − 6 kpc, Fig. 6 shows that the two over-densities connect through a zone of lower density to form a single sequence. This is in agreement with the observations of Hayden et al. (2015), Bensby et al. (2017), Zasowski et al. (2019), Bovy et al. (2019), and Lian et al. (2020a,b), who also report a single sequence in the inner disc and/or in the bulge/bar area. Conversely, Rojas-Arriagada et al. (2019) and Queiroz et al. (2020) observe two sequences in the inner regions. In Appendix F, we compare the distributions of different APOGEE DR16 alpha elements in the ([Fe/H], [α/Fe]) plane (restricting the sample to the stars contained in the Rg ∈ [0, 2] kpc interval). The different elements produce different patterns: the global alpha-element abundance5 and oxygen show a double sequence, while magnesium, silicon, and calcium present a single sequence. This could explain, at least partly, why Queiroz et al. (2020), who use a combined α-element abundance, observe a double sequence, while we see a single one with magnesium. However, this does not explain the discrepancy with Rojas-Arriagada et al. (2019), who also used magnesium. Beyond Rg = 10 kpc, the high-alpha sequence gradually vanishes. This is in agreement with the finding that the thick disc has a shorter scale length than the thin disc (Bensby et al. 2011; Cheng et al. 2012; Bovy et al. 2012). It should be emphasised that in this paragraph the term ‘sequence’ is used in the geometrical sense. It does not presuppose the number of chemical tracks that form the sequence or sequences. In particular, based on Fig. 6, it can not be excluded that the single geometrical sequence observed in the inner disc be made of two chemical tracks, with the low-alpha one restricted to a narrow metallicity range. We discuss and propose an interpretation of the inner disc sequence in Sect. 5.
Citation: Rojas-Arriagada et al. (2019)
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Citation: Rojas-Arriagada et al. (2019)
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Citation: Bovy et al. 2012; Bovy et al. 2012
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It is a pressing question as to how the radiation that is observed in relativistic jets in active galactic nuclei (AGNs) is generated (e.g., Blandford & Königl 1979; Marscher 1980; Zensus 1997; Laing & Bridle 2002; Honda 2010; Levinson & Rieger 2011; Mościbrodzka et al. 2011; Ito et al. 2013; Mason et al. 2013; Potter & Cotter 2013; Hovatta et al. 2014; Scott & Stewart 2014; Shih & Stockton 2014; Wang et al. 2014; Turner & Shabala 2015; Asada et al. 2016; Hirotani et al. 2016; Koay et al. 2016; Khabibullin et al. 2016; Prieto et al. 2016). Although there is a common consensus that the emitters are energetic particles, how these particles are accelerated to such high energies, how they dissipate their energy, and how they are transported with the jets themselves are still the subject of feverish investigation (e.g., Blandford & Eichler 1987). It is argued that relativistic outflows from black holes are associated with accretion flows (Blandford 1976; Fender et al. 2004; Meier 2005; Ferreira et al. 2006; Trump et al. 2011; Pu et al. 2012; Wu et al. 2013; Ishibashi et al. 2014; Sbarrato et al. 2014). In the case of collimated relativistic jets, magnetic fields must play an important role (Camenzind 1986a, 1986b, 1987; Fendt & Greiner 2001; Vlahakis & Königl 2004; Komissarov et al. 2007; Lyubarsky 2009; Nakamura & Asada 2013; Homan et al. 2015), and it is argued that jets are powered at the expense of the black hole, wherein energy is extracted from a reservoir of rotational energy from the black hole itself, either by electromagnetic means (Blandford & Znajek 1977; Komissarov 2004, 2005; Toma & Takahara 2014) or through magnetohydrodynamical processes (Phinney 1983; Takahashi et al. 1990; Koide et al. 2002; McKinney & Gammie 2004; Hawley & Krolik 2006). Models have been proposed for both of these cases, and they both in principle possess certain testable predictions. In particular, for the latter, numerical GRMHD simulations (e.g., McKinney 2006) and analytical GRMHD studies (e.g., Takahashi et al. 1990; Pu et al. 2016) consistently show the presence of a stagnation or separation surface (a separatrix). This surface separates the (inner) inflow region from the (outer) outflow region, both of which follow the same global, black-hole-threading magnetic field lines. The relatively slow radial velocities near the stagnation surface imply a high concentration of fluid particles. If energetic particles are injected in the vicinity of the stagnation surface or near the black hole event horizon, they must accumulate in high concentrations near the stagnation surface, provided that the cooling timescale is not significantly shorter than the dynamical timescale of the jet fluid flow. This surface is a unique feature of relativistic GRMHD jets and in contrast to an ideal force-free magnetic jet (e.g., McKinney & Narayan 2007; Tchekhovskoy et al. 2008; Broderick & Loeb 2009).
Citation: Camenzind 1986a
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Citation: Lyubarsky 2009
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The first transit of HD 21749b was identified by both the MIT Quick Look Pipeline (which searches for planet candidates in the 30 minutes Full Frame Images) and the Science Processing Operations Center (SPOC) pipeline based at the NASA Ames Research Center (Jenkins et al. 2016). No other matching transits were found in the publicly released data from sectors 1 and 2. After TOI 186.01 was alerted, we searched for archival spectroscopy of this very bright star and found 59 High-accuracy Radial-velocity Planet Searcher (HARPS) radial velocities (RVs) in the European Southern Observatory (ESO) archive (see Section 2.4). A periodogram of these RVs showed a clear signal at 35.57 days, but the TESS photometry and the RHK index (Boro Saikia et al. 2018) indicate a stellar rotation period of around 35 days, calling for caution. If the strongest period in the RVs did correspond to the planet, then we expected to see additional transits in sectors 3 and 4. Once the sector 3 data were released, we discovered that a momentum dump29

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“Momentum dumps” consist of resetting the momentum wheel speed every 2.5–3 days and are used to mitigate the noisier-than-expected measurements of the spacecraft momentum wheel speeds at higher wheel speeds (see https://archive.stsci.edu/files/live/sites/mast/files/home/missions-and-data/active-missions/tess/_documents/TESS_Instrument_Handbook_v0.1.pdf for details). Momentum dumps require brief interruptions to Fine Pointing mode, during which an increase in the flux dispersion is noticeable in the science data, so data acquired during these intervals are excluded from our analysis.
 occurred approximately 35.6 days after the sector 1 transit (see Figure 1). We did not let this unexpected turn of events foil our search efforts, and upon close inspection of the light curve we succeeded in recovering a partial transit (including egress) immediately following the momentum dump. Finally, we observed a third transit in sector 4, thus allowing for a robust ephemeris determination (see Section 3.3). Serendipitously, when applied to the first three sectors of the HD 21749 light curve, the SPOC pipeline yielded an additional planet candidate with a period of 7.9 days (TOI 186.02).
Citation: Boro Saikia et al. 2018
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