adsabs/FOCAL huggingface dataset:
Figure 13 shows the temporal evolution of the density, ρ, plasma flow velocity, vx, position of the wave crest, PosA, phase speed, vw, and magnetic field component in the z-direction, Bz, for the primary waves in every different case of initial amplitude, ρIA. In Figure 13(a), we observe that the amplitude of the density remains approximately constant at their initial values until the time when the shock is formed and the density amplitude of the primary wave starts decreasing (see Vršnak & Lulić 2000), i.e., at t ≈ 0.03 (blue), t ≈ 0.04 (red), and t ≈ 0.055 (green). For the case of ρIA = 1.3 (magenta), a decrease of the amplitude of the primary can hardly be observed, as expected for low-amplitude wave (Warmuth 2015). One can see that the larger the initial amplitude, ρIA, the stronger the decrease of the primary wave’s amplitude, which is consistent with observations (Warmuth & Mann 2011; Muhr et al. 2014; Warmuth 2015). The amplitudes decrease to values of ρ ≈ 1.6 (blue), ρ ≈ 1.5 (red), and ρ ≈ 1.4 (green) until the primary wave starts entering the CH. Due to the fact that the waves with larger initial amplitude enter the CH earlier than those with small initial amplitude, we can see in Figure 13(a) that the tracking of the parameters of the faster waves stops at an earlier time than the one for the slower waves. A similar behavior to the one of the density, ρ, can be observed for the plasma flow velocity, vx, in Figure 13(b) and the magnetic field component, Bz, in Figure 13(e). Here, the amplitudes decrease from vx = 0.75, Bz = 1.9 (for ρIA = 1.9, blue), vx = 0.6, Bz = 1.7 (for ρIA = 1.7, red), vx = 0.45, Bz = 1.5 (for ρIA = 1.5, green), and vx = 0.27, Bz = 1.3 (for ρIA = 1.3, magenta) to vx = 0.55, Bz = 1.6 (for ρIA = 1.9, blue), vx = 0.46, Bz = 1.5 (for ρIA = 1.7, red), vx = 0.36, Bz = 1.4 (for ρIA = 1.5, green), and vx = 0.25, Bz = 1.25 (for ρIA = 1.3, magenta). Figure 13(c) shows how the primary waves propagate in the positive x-direction. In all four cases of different initial amplitude, ρIA, the phase speed decreases slighty (consistent with observations; see Warmuth et al. 2004 and Warmuth 2015) until the waves enter the CH at different times, i.e., the values for the phase speed start at vw ≈ 2.2 (for ρIA = 1.9, blue), vw ≈ 1.9 (for ρIA = 1.7, red), vw ≈ 1.7 (for ρIA = 1.5, green), and vw ≈ 1.4 (for ρIA = 1.3, magenta) and decrease to vw ≈ 1.5 (for ρIA = 1.9, blue), vw ≈ 1.39 (for ρIA = 1.7, red), vw ≈ 1.2 (for ρIA = 1.5, green), and vw ≈ 1.13 (for ρIA = 1.3, magenta).
Citation: Warmuth 2015
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Citation: Warmuth 2015
Citation function: Similarities
Citation: Warmuth 2015
Citation function: Similarities
Another clear difference for Ark 120 with the NT thin disc predictions is the variability. From the best simultaneous fit of the 2013 and 2014 observations, we infer a significant increase in mass accretion rate through the disc from 0.03 to 0.07 Lbol/LEdd in only one year, but a standard thin disc around a SMBH cannot vary on such timescale. Indeed, the radial mass accretion rate change via viscous processes has a time-scale of torb(R)/[α(H/R)2]. For Ark,120, assuming an orbital time-scale at R = 100 Rg, a viscosity parameter of 0.1 and a H/R (H is the height scale of the disc) value of 0.1, this corresponds to ∼150 years. Moreover, the optical-UV flux significantly changes in less than a year, varying for example by a factor of 50% in the UVW2 band between 2013 and 2014 (e.g. see Fig. 1, Lobban et al. 2018). Similar rapid changes in the optical-UV flux are typically seen in other BLS1s, especially those at low Lbol/LEdd (e.g. MacLeod et al. 2010; Kozłowski 2016; Simm et al. 2016; Rakshit & Stalin 2017). These are generally assumed to be from reprocessing, where the X-ray flux illuminates the outer disc (e.g. Buisson et al. 2017), and adds to the intrinsic emission. An additional reprocessed component in the optical-UV would lead us to overestimate the value of the intrinsic Ṁ, so to an underestimate of BH spin via the efficiency argument (Kubota & Done 2018). Nonetheless, changes as large as about 50% in the UV flux are unlikely to be driven by X-ray reprocessing, as the UV flux in Ark 120 is much higher than in the X-ray band. Besides, detailed models of the expected optical-UV variability from X-ray reprocessing fail to fit the excellent long term simultaneous optical-UV-X-ray datasets, and would imply a larger disc size than expected by standard thin disc (e.g. NGC 5548: McHardy et al. 2014; Edelson et al. 2015; Gardner & Done 2017; NGC 4151: Edelson et al. 2017; Ark 120: Gliozzi et al. 2017; Fairall 9: Pal et al. 2017; NGC 4593: Cackett et al. 2018; Pal & Naik 2018; Microlensing studies: Morgan et al. 2010; Dai et al. 2010). Such large discs should be significantly brighter than observed and this discrepancy may be explained for example by a flatter temperature profile than in NT discs, from scattering of a significant part of the optical flux on larger scales, by electron scattering in the disc atmosphere (Dai et al. 2010; Morgan et al. 2010; Hall et al. 2018).
Citation: Simm et al. 2016
Citation function: Similarities
Thermonuclear explosions and AICs of WDs in AGN disks are potential sources for future joint gravitational wave (GW), EM, and neutrino multi-messenger observations. In addition to AIC of WDs, another possible formation channel of millisecond magnetars in AGN disks is mergers of BNSs. Such a merger may produce a millisecond magnetar if the NS equation of state is stiff enough (e.g., Ai et al. 2020). This would also drive a magnetar-powered explosion as discussed above. The mechanism by which a nascent millisecond magnetar might power a GRB jet has been studied in detail over the past decade (Bucciantini et al. 2008, 2012; Metzger et al. 2008, 2011; Siegel et al. 2014).8

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Some indirect evidence for the magnetar mechanism in GRBs are the presence of a plateau in the early X-ray afterglow of long-duration GRBs, which may be explained as the continuous energy injection from the spindown of a magnetar (Dai & Lu 1998; Zhang & Mészáros 2001; Zhang et al. 2006).
 A putative jet may thus be launched from the AIC of WDs and BNS mergers occurring in AGN disks. However, Zhu et al. (2021b, 2021a) and Perna et al. (2021a) recently showed that GRB jets embedded in AGN disks would be usually choked by the dense material of the accretion disks. Although it is hard to observe gamma-ray signals, Zhu et al. (2021a) suggested that these choked GRBs may effectively produce TeVPeV neutrinos that could be detected by IceCube and IceCube-Gen2. Furthermore, the dissipation of magnetic energy during the merger of BWDs is expected to accelerate cosmic rays and produce high-energy neutrinos (Xiao et al. 2016). In the GW channel, BWD mergers are promising astrophysical GW sources for space-borne GW observatories, e.g., LISA (Amaro-Seoane et al. 2017), TaiJi (Ruan et al. 2020), and TianQin (Luo et al. 2016), while the GW signals from BNS mergers are readily detected with the ground-based GW detectors such as LIGO and Virgo (Abbott et al. 2017). Future joint observations of GW, EM, and neutrino signals can reveal the existence of WD and NS populations in AGN disks.
Citation: Metzger et al. 2008
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The most common approaches to observe SPI are spectral observations of typical emission lines for stellar chromospheres and coronae. The earliest approach to observe SPI dates back to the early 2000’s and focused on non-thermal chromospheric Ca II H and K emissions (Shkolnik et al. 2003). Subsequent studies by Shkolnik et al. (2005, 2008) of the system HD 179949 saw enhanced emissions in four out of six observational epochs that appeared to be periodic with the planetary orbital period. Other authors observed different systems and likewise saw chro-mospheric excess emissions (Walker et al. 2008; Staab et al. 2017). In addition, Cauley et al. (2019) estimated possible magnetic field strengths of hot Jupiters based on chromospheric excess emissions. The derived energy fluxes from chromospheric Ca II H and K emissions are larger than typical fluxes derived from Alfvén wing models and magnetohydrostatic models (Saur et al. 2013; Lanza 2015) and the observed periodicity only appears in some epochs, which was suggested to be alternatively explained by star spots (Miller et al. 2012). Recently, Strugarek et al. (2019) investigated how modeling and observations could reconcile by including the magnetic topology of stellar coronae based on the example of Kepler-78. In the UV range, France et al. (2016, 2018) conducted large surveys with the Hubble Space Telescope, however, the authors could not identify any signals related to SPI. At coronal X-ray wavelengths, several studies have investigated the influence of planets on stellar X-ray activity, such as Kashyap et al. (2008); Scharf (2010); Poppenhaeger et al. (2010). Some studies found correlations between planets and X-ray activity in stars and some did not. Poppenhaeger & Schmitt (2011a,b) explain previously observed correlations, attributed to SPI, as a result of selection effects due to planet detection methods and the limitations in X-ray observations. The radio wavelength range is particularly promising to search for SPI. Recent observations with LOFAR showed strong hints for the existence of SPI in several systems (Vedantham et al. 2020; Callingham et al. 2021; Pérez-Torres et al. 2021). In addition, a recent campaign in the radio range provided the first tentative direct observational hints of an intrinsic magnetic field on an exoplanet (Turner et al. 2021).
Citation: Walker et al. 2008
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Active galactic nuclei (AGNs) are extragalactic sources that emit across the whole electromagnetic spectrum. Such systems are composite and each sub-structure has its own role in shaping the emerging spectrum (see Padovani et al. 2017, for a comprehensive review). It is ubiquitously accepted that the X-ray emission originates in the very inner regions of AGNs, near the central supermassive black hole (SMBH). Accretion of matter infalling on to the SMBH is responsible for the enormous amount of optical-UV photons, a fraction of which can be further energized via inverse-Compton (Sunyaev & Titarchuk 1980) off thermal electrons (the so-called hot corona: Haardt & Maraschi 1991, 1993; Madejski et al. 1995; Zdziarski et al. 1995) up to the X-rays. The maximum energy gain for these seed photons is mainly set by the hot plasma’s temperature, and, to a lower extent, by its opacity (e.g. Rybicki & Lightman 1979; Beloborodov 1999; Middei et al. 2019). In fact, the X-ray continuum in AGNs is well modelled by a power law with a high energy roll-over (e.g. Perola et al. 2002; Dadina 2007; Molina et al. 2009, 2013; Malizia et al. 2014; Fabian et al. 2015, 2017; Ricci et al. 2018; Tortosa et al. 2018). AGN X-ray spectra may show additional features due to reprocessing of the primary X-ray emission by the circumnuclear material. A fluorescence emission line from the Fe K-shell is commonly observed as the most prominent feature (e.g. Bianchi et al. 2009) and its analysis carries a wealth of information on the physics of the reflecting material. This emission line has an intrinsically narrow profile that can undergo distortions, such as broadening, due to special and general relativistic effects. In particular, the closer to the SMBH the reflectors, the more distorted (i.e. the broader) the neutral or ionized Fe line profile (e.g. Fabian et al. 1995). On the contrary, at larger distance, these effects are negligible, thus the Fe Kα shape is consistent with a narrow profile. Additionally, in the case of Compton-thick reflectors (i.e. NH ≳ 1.5 × 1024 cm−2) the X-ray spectra show a typical emission excess around 30 keV, the so-called Compton-hump (e.g. Matt, Fabian & Ross 1993). The effect of any absorbing matter crossing our line of sight can significantly attenuate the observed number of photons, especially in the soft X-rays (e.g. Cappi et al. 1999; Awaki et al. 2000; Matt 2002; Bianchi et al. 2009; Middei et al. 2021).
Citation: Dadina 2007
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The measurement of clumps at redshifts beyond the peak of cosmic star formation is notoriously difficult. The identification of clumpy galaxies in the dominant population of irregular galaxies at high redshift dates back to the first deep Hubble Space Telescope (HST) images (e.g., Williams et al. 1996). Subsequent studies at z ~ 1−3 revealed that clumpy galaxies are more numerous than in the local Universe (e.g., Abraham et al. 1996; van den Bergh et al. 1996; Giavalisco et al. 1996; Elmegreen et al. 2004, 2007, 2008, 2009; Elmegreen & Elmegreen 2006; Elmegreen et al. 2013; Kubo et al. 2013, 2016; Glazebrook 2013; Tadaki et al. 2014; Murata et al. 2014; Guo et al. 2015; Garland et al. 2015; Bournaud 2016). The fraction of clumpy galaxies was studied from z = 0 up to the most distant galaxies identified today at z ~ 10 (Ravindranath et al. 2006; Guo et al. 2012, 2015; Shibuya et al. 2016). In the sample of galaxies with 0.5 z  3 in the CANDELS survey, Guo et al. (2015) find that galaxies with low mass log (M⋆/M⊙)  9.8 have a high fraction of off-center clumps fclumpy ~ 60% constant over the observed redshift range, while for intermediate and massive galaxies fclumpy decreases from ~40% at z ~ 3 to ~15% at z ~ 0.5. Combining deep HST imaging from the 3D-HST, CANDELS, HUDF and HFF surveys Shibuya et al. (2016) claim that \hbox{$f_{\rm clumpy}^{\rm UV}$}fclumpyUV follows an evolution similar to that of the star formation rate density (e.g., Madau & Dickinson 2014), increasing from z ≃ 8 to z ≃ 1−3 and decreasing from z ≃ 1 to z ≃ 0. In comparing these observed trends with the predictions of simulations, Guo et al. (2015) conclude that VDI are likely responsible for fclumpy at high mass and that minor mergers are a viable explanation for fclumpy at intermediate mass for z  1.5, while Shibuya et al. (2016) conclude that VDI is the main origin of all clumps. Both these studies exclude major mergers as a possible contribution to the evolution of the fraction of clumps. This is somewhat surprising as the major merger fraction for star-forming galaxies is observed to increase up to ~20% at z ≃ 1−2 (e.g., Lotz et al. 2011; López-Sanjuan et al. 2011, 2013), staying high at least to z ~ 3−4 (e.g., Conselice et al. 2008; Tasca et al. 2014a), and one would therefore expect that a fraction of the clumps is related to ex situ major merging.
Citation: Ravindranath et al. 2006
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Previously it has been postulated that the SMS formation by DC is only possible in atomic cooling haloes both with intense FUV irradiation and with primordial gas composition. Here we have shown that SMSs can form also in slightly metal-enriched cases as long as FUV irradiation is intense enough. This relaxation of the condition will increase the expected number density of massive seed BHs. Several authors have estimated the number density of such seeds formed in the usual DC scenario, which requires metal-free gas composition. Using the critical intensity advocated by recent studies (Agarwal et al. 2014; Sugimura et al. 2014; Inayoshi & Tanaka 2015), this ranges from a few Gpc−3 (e.g. Dijkstra et al. 2008; Dijkstra, Ferrara & Mesinger 2014), to 10−6 to 10−4 Mpc−3 (e.g. Agarwal et al. 2012; Visbal, Haiman & Bryan 2014a; Chon et al. 2016; Habouzit et al. 2016; Valiante et al. 2016). The dynamical heating model by Wise et al. (2019) predicts the number density of 10−7 to 10−6 Mpc−3. Those seeds can account for the origin of the rare BHs in the high-z universe. If all the SMBHs down to z = 0 shares the universal origin, the seed density should be comparable to that of the SMBHs ubiquitously residing in massive galaxies, ∼0.01–0.1 Mpc−3 (Aller & Richstone 2002; Davis et al. 2014), which is much larger than the seed density predicted by the DC model. The DC in the primordial environment is terminated around z ∼ 10, as the metal enrichment proceeds (e.g. Trenti & Stiavelli 2009; Chon et al. 2016): once a Pop III star ends its life as a core-collapse SN, metallicity inside the host halo jumps up to 10−4 to 10−3 Z⊙ (e.g. Maio et al. 2010; Ritter et al. 2015; Sluder et al. 2016; Chiaki, Susa & Hirano 2018). According to our result, the SMS formation can continue even in such second-generation haloes. Since the second-generation haloes tend to be under stronger radiation field than the pristine ones (e.g. Dijkstra et al. 2014), we expect significant increase in the number density of the SMSs. For example, the first episode of star formation delays another star formation for a few hundred million years by the radiative and SNe mechanical feedback, which ejects a gas from the halo. If the haloes approach close enough to a luminous galaxy before another episode of star formation, they can be ideal sites for SMS formation. Not only the radiation sources outside the halo, but also those inside the same halo can trigger SMS formation in an irradiated massive cloud as long as the metallicity is low enough. With such new varieties of SMS formation sites, the expected seed BH number will be largely enhanced. We will pursue the validity of this scenario using samples from the cosmological simulations in the future studies.
Citation: Dijkstra et al. 2008
Citation function: Compare/Contrast; Compare/Contrast; Compare/Contrast
In the context of RI imaging techniques, the standard clean algorithm implements a non-linear, greedy approach to perform deconvolution in an iterative manner (Högbom 1974). Working pixelwise, clean implicitly assumes sparsity of the sought image, removing at each iteration, a fraction of the maximum intensity pixel convolved with the dirty beam from the computed residual image. Over the past years, many variants of this celebrated algorithm have been developed, for instance, its multiscale version (Cornwell 2008). In addition to these clean-based approaches, recent developments in the field of compressive sensing (CS) applied for astronomical imaging have given birth to many imaging algorithms, particularly for RI (Wiaux et al. 2009; Carrillo, McEwen & Wiaux 2012; Garsden et al. 2015; Onose et al. 2016). Leveraging optimization frameworks, these methods aim to solve the underlying image recovery problem by enforcing sparsity of the image of interest in some suitable domain. These techniques have shown the potential to surpass the image reconstruction quality obtained by clean -based approaches (Carrillo, McEwen & Wiaux 2014; Onose et al. 2016; Pratley & Johnston-Hollitt 2016; Onose, Dabbech & Wiaux 2017; Dabbech et al. 2018). While the aforementioned methods have been developed mainly for Stokes I imaging, as previously mentioned these can be extended to polarimetric imaging by following the same imaging approach for all the Stokes parameters. In the context of sparsity regularized approaches, one such extension for polarimetric imaging has been presented in Akiyama et al. (2017), promoting sparsity of each of the Stokes parameters using ℓ1 norm combined with total variation (TV) regularization (Rudin, Osher & Fatemi 1992). A point worth noting here is that the above-mentioned approaches, whether clean and its variants or the sparsity regularized method, solve for the Stokes parameters totally independently. However, these images are physically linked via polarization constraint, i.e. the polarized intensity cannot be greater than the total intensity. Within the CS framework, this constraint has been exploited by a recently proposed approach, namely Polarized SARA, estimating jointly the Stokes parameters (Birdi, Repetti & Wiaux 2018a, b). This approach has been shown to provide better reconstruction quality in comparison with the case when the constraint is not accounted for.
Citation: Garsden et al. 2015
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Various outbursts in GeV–TeV energies from 1ES 1959+650 have been reported by Fermi-LAT and several ground-based Cherenkov experiments like VERITAS, MAGIC, FACT (Albert et al. 2006; Aliu et al. 2014; Biland et al. 2016) including a new highest historical γ-ray activity, detected by MAGIC collaboration on 2016 June 13–14, when the VHE flux above 100 GeV reached up to 3 Crab unit (C.U) (Bhattacharyya et al. 2019). Many multiwavelength campaign from radio to TeV γ-ray energies have been conducted since its discovery (Krawczynski et al. 2004; Biland et al. 2016) and this source has exhibited a number of outbursts across all energies. One interesting VHE outburst was detected on 2002 June 4 by the Whipple telescope without any activity in contemporaneous X-ray energies, called ‘orphan flare’, during a multiwavelength campaign (Krawczynski et al. 2004). In the recent past, 1ES 1959+650 has undergone a few major outbursts in all wavelengths specially in X-ray energies since 2015 (Kapanadze et al. 2016). Unprecedented flux and spectral variability during these episodes have been observed (Kapanadze 2015; Kaur et al. 2017; Kapanadze et al. 2018; Patel et al. 2018). Such flux and spectral variations in different energy bands over time provide an opportunity to study underlying physical processes in terms of particle acceleration processes, evolution of the accelerated particles (leptons/hadrons) within the source, and their radiative cooling and escape processes. Many authors explained the broad-band emission from this source with leptonic one-zone SSC process (Krawczynski et al. 2004; Tagliaferri et al. 2008). But the lack of correlation between X-ray and VHE γ-rays several times and also during the orphan flare, mentioned earlier, challenges the leptonic one-zone SSC model of emission. Alternatively, EC, lepto-hadronic and two-zone SSC models have been also tried to explain the broad-band SED of the source (Backes et al. 2012; Aliu et al. 2014; Asano & Hayashida 2018; Patel et al. 2018). Recently, the MAGIC collaboration again introspected the VHE flaring episodes in 2016, mentioned earlier, to study the feasibility of neutrino emission from 1ES 1959+650 as well under the framework of mixed lepto-hadronic model (MAGIC Collaboration 2020a). Although they (MAGIC collaboration) have concluded that it is difficult to detect neutrinos from the same source during extreme γ-ray flares, it is customary to mention that more study with well-sampled data from multiple experiments are inevitable to understand the true physical processes inside this type of astrophysical sources in a better way.
Citation: Kapanadze et al. 2018
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