adsabs/FOCAL huggingface dataset:
However, two recent pieces of evidence may suggest that standard disk accretion is not the only process at work in growing a BH to the (super)massive regime. The first is the discovery of an increasing number of active BHs with masses M
• ≳ 109
M
⊙ at very high redshifts, z ≳ 7 (e.g., Mortlock et al. 2011; Wu et al. 2015; Venemans et al. 2017, 2018; Reed et al. 2019; Banados et al. 2018, 2021; Wang et al. 2019, 2021), when the age of the universe was shorter than 0.8 Gyr. The second is the robust measurements of extreme BH masses M
• ≳109−10
M
⊙ at the center of early-type galaxies with stellar mass M
⋆ ≳ 1011
M
⊙ (e.g., McConnell et al. 2011; Ferre-Mateu et al. 2015; Thomas et al. 2016; Mehrgan et al. 2019; Dullo et al. 2021) that formed most of their old stellar component during a star formation episode lasting some 108 yr at z ≳ 1, as demonstrated by astroarcheological measurements of their stellar ages and α-enhanced metal content (e.g., Thomas et al. 2005, 2010; Gallazzi et al. 2006, 2014; Johansson et al. 2012; Maiolino & Mannucci 2019; Morishita et al. 2019; Saracco et al.2020). These observations concur to raise the issue of how billion-solar-mass BHs may have grown in less than a gigayear. In fact, this is somewhat challenging if standard disk accretion starts from a light seed of ∼102
M
⊙ of stellar origin and proceeds with the typical Eddington ratios λ ≲ 1 as estimated out to z ∼ 6 in active BHs (see Vestergaard & Osmer 2009; Nobuta et al. 2012; Kelly & Shen 2013; Dai et al. 2014; Kim & Im 2019; Duras et al. 2020; Ananna et al. 2022), which would require an overall time ≳0.8/λ Gyr to attain ∼109
M
⊙. Solutions may invoke mechanisms able to rapidly produce heavy BH seeds of 103−5
M
⊙, thus reducing the time required to attain the final masses by standard disk accretion (see Natarajan 2014, Mayer & Bonoli 2019, Inayoshi et al. 2020, and Volonteri et al. 2021 for recent and exhaustive reviews). Viable possibilities comprise direct collapse of gas clouds within a (proto)galaxy, possibly induced by galaxy mergers or enhanced matter inflow along cosmic filaments (e.g., Lodato & Natarajan 2007; Mayer et al. 2010, 2015; Di Matteo et al. 2012, 2017); merging of stars inside globular or nuclear star clusters (e.g., Portegies Zwart et al. 2004; Devecchi et al. 2012; Latif & Ferrara 2016; Kroupa et al. 2020); and migration of stellar BHs toward the nuclear galaxy regions via dynamical friction against the dense gas-rich environment in strongly star-forming progenitors of local massive galaxies (e.g., Boco et al. 2020, 2021).
11


11
When referring to the dynamical friction mechanism, the term “seeds” is used in a broader sense with respect to the classic meaning in the literature. A seed is usually referred to as the first compact object on which subsequent disk accretion occurs, eventually leading to the formation of a supermassive BH. The heavy seeds formed with the dynamical friction mechanism are by-products of multiple mergers of already-existing stellar mass BHs (that in turn could be referred to as light seeds) forming across a wide redshift range.
Citation: Morishita et al. 2019
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Torques caused by a misalignment between the spin axis of a black hole and the angular momentum of the accretion disc can also induce jet precession. In an overall misaligned system the inner part of the accretion disc is aligned with the black hole spin axis by the combined effect of Lense–Thirring precession and internal viscosity of the disc. This process is known as the Bardeen–Petterson effect (Bardeen & Petterson 1975). The outer disc will maintain its angular momentum direction. Through the accretion of material of the outer disc on to the black hole, the black hole spin axis will precess and align with the outer disc. The precession time-scale is of the same order as the alignment time-scale (Scheuer & Feiler 1996). Two aspects of the disc-induced precession, namely the relatively slow precession rate and the single possible precession cycle in the alignment lifetime (Scheuer & Feiler 1996; Lodato & Pringle 2006), make this process less likely to be a cause of jet precession for the given radio sources. For example, following Lodato & Pringle (2006), adopting their numerical values for accretion efficiency ε and disc viscosity parameters α1, 2 and assuming black hole spin a ≈ 1, we obtain a precession period of
(3)
\begin{eqnarray*}
\begin{split} t_\mathrm{prec} &amp;\approx t_\mathrm{align} \approx 7\, \, \mathrm{Myr} \\
&amp;\times \, a^{11/16} \left(\frac{\epsilon }{0.1}\right)^{7/8} \left(\frac{L}{0.1 L_\mathrm{Edd}}\right)^{-7/8} M_9^{-1/16} \\
&amp;\times \left(\frac{\alpha _1}{0.03}\right)^{15/16} \left(\frac{\alpha _2}{0.3}\right)^{-11/16} \, .\\
\end{split}
\end{eqnarray*}
This is larger than the precession period we estimated for Cygnus A. Simulations show that if the precession time-scale becomes comparable to the source age, very asymmetric lobe morphologies result (Donohoe & Smith 2016). As discussed in Section 3.3, the sources in our sample do not show such asymmetries, but often show smooth, ring-like hotspot features that suggest a local impact time-scale short compared to the source ages and thus many precession cycles. This is not consistent with precession induced by a misaligned accretion disc. A process that produces one precession cycle only, where the period of the cycle can be much shorter than the source age (equation 3), would have only a limited probability to be detected at the time a given source is observed. However, we also find precession frequently in comparatively old sources, e.g. 3C 33.1, 3C 303, and 3C 436. These are three of the seven sources with modelled ages of log(age/yr) ≥ 7.9. This makes disc-induced precession less likely for this subpopulation.
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We end by studying the BH mass versus stellar velocity dispersion, σ, in Fig. 7. This is also a benchmark scaling, and it is considered to be one of the best predictors for BH masses, i.e. to have a small intrinsic scatter (0.3 dex). This is perhaps the most problematic for Horizon-AGN. If we fit a single logarithmic slope to the BH mass versus stellar velocity dispersion, we find a value of 4.02 for galaxies with σ > 60 km s−1, compared to 4.38 found by Kormendy & Ho on a culled subsample of the galaxies shown in Fig. 7. The main reason is that velocity dispersions in the simulation are underestimated at Mbulge  1011 M⊙ (σ  200 km s−1), and overestimated above, when comparing to the observed Faber–Jackson relation (Bernardi et al. 2003). The dearth of BHs with MBH  5 × 107  M⊙ in intermediate-mass galaxies highlighted in the discussion of Fig. 5 also contributes, but given the good agreement in the MBH versus Mbulge relation, it is a minor contribution, for this specific comparison. One contribution to the discrepancy is that galaxy sizes in the simulation are larger than for observed galaxies with mass ≲ 1011 M⊙. In fact, the galaxy sizes are in good agreement with observed low-mass late-type galaxies, but are larger, by a factor of ∼ 3–4, than observed low-mass early-type galaxies (Dubois et al., in preparation). Since the complete sample is a mix of late- and early-type galaxies, there appears to be an overall mismatch at σ  200 km s−1 of a factor of ${\sim } \sqrt{2}$ in the velocity dispersion (roughly corresponding to an average overestimate of the galaxy size of a factor of 2). The velocity dispersion is also sensitive to the DM halo central profile, which depends not only on DM halo mass but also on a subtle interplay with the baryons and AGN feedback. Indeed, the central DM densities in Horizon-AGN are lower than for the Horizon-noAGN simulation, which differs only in that AGN are not included in the latter. We defer a more detailed study quantifying the exact contribution of this effect as a function of galaxy mass and assembly history to another paper (Peirani et al., in preparation).
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However, our understanding of the scaling relations in lower-mass galaxies is lacking. This includes galaxies near the break of the galaxy stellar-mass function 





 (e.g., Baldry et al. 2012) like the Milky Way (MW), as well as galaxies at lower masses. Measurements of black hole (BH18

18
In this article, we use the abbreviations SMBH and BH interchangeably.
) masses are lacking in these galaxies. Despite the MW’s precisely measured central BH mass of 





 





 (Boehle et al. 2016; Gillessen et al. 2017), we actually know very little about the demographics of BHs in galaxies of this mass. In particular, there are only 11 published 





 in spiral galaxies within ∼1 dex in stellar mass of the MW, but there are strong indications that the relations between 





 and galaxy properties that hold at higher mass (Kormendy & Ho 2013; McConnell et al. 2013; Saglia et al. 2016; van den Bosch et al. 2016) break down for M⋆ spiral galaxies (Greene et al. 2010, 2016; Läsker et al. 2016). For instance, Andromeda is only ≲2 times more massive than the MW, but its SMBH is two orders of magnitude higher, 





 ∼1.4 × 108M⊙ (Bender et al. 2005). At even lower galaxy masses, the changes in the inferred 





 are twofold: not only does the scatter increase enormously at low mass (e.g., Reines et al. 2013), but it also seems that the 





 are lower at fixed stellar mass or σ of the lower-mass systems (e.g., Kormendy & Ho 2013; Scott et al. 2013; Greene et al. 2016; Saglia et al. 2016; Nguyen et al. 2017, 2018, 2019, hereafter N17, N18, N19). Despite the incomplete sample of 





 at lower masses, there are already two intriguing findings from the limited data available: (1) a factor of 2 difference in normalization seen in the 





–σ correlations for early-type (ETGs) and late-type (LTGs) galaxies (McConnell et al. 2013) and (2) a large scatter up to two orders of magnitude of MBH seen around the global scaling relations for low-mass systems (Greene et al. 2010; Scott et al. 2013; Graham & Scott 2015; Greene et al. 2016; Läsker et al. 2016; Chilingarian et al. 2018).
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The 21-cm line is the transition line of atomic hydrogen between two hyperfine states that are split from the ground state due to the spin coupling of the nucleus and electron. The 21-cm brightness temperature from a bulk of hydrogen gas depends on the matter density, the neutral fraction, the velocity gradient, and the spin temperature. During the EoR, it is reasonable to assume that X-ray heating and Ly α-pumping effects are so efficient that the H i spin temperature is much greater than the CMB temperature. In this limit, the power spectrum of 21-cm brightness temperature fluctuations can be written as (Datta et al. 2012; Mao et al. 2013)
(3)
\begin{eqnarray*}
P_{\Delta T}({\boldsymbol k},{\boldsymbol n},z) = \widetilde{\delta T}_b^2 \bar{x}_{\rm H\,{\small I}}^2\, \left[ b_{\rho _{\rm H\,{\small I}}}(z) + \mu _{\boldsymbol k}^2 \right]^2\, P_{\zeta }(k,{\boldsymbol n},z),
\end{eqnarray*}
at large scales, with
(4)
\begin{eqnarray*}
\widetilde{\delta T}_b(z) = (23.88\, {\rm mK}) (\frac{\Omega _b h^2}{0.02}) \sqrt{\frac{0.15}{\Omega _M h^2} \frac{1+z}{10}}.
\end{eqnarray*}
Note that $\bar{x}_{\rm H\,{\small I}}$ is the global neutral fraction that implicitly depends on the redshift, and μk ≡ k · n/|k| (cosine of angle between line of sight n and wave vector k of a given Fourier mode). We can relate the neutral hydrogen density fluctuations to the total matter density fluctuations with a simple bias parameter, the neutral density bias, $b_{\rho _{\rm H\,{\small I}}}(k,z) \equiv \tilde{\delta }_{\rho _{\rm H\,{\small I}}}({\boldsymbol k},z)/\tilde{\delta }_{\rho }({\boldsymbol k},z)$. At large scales, the bias approaches to a scale-independent value, $b_{\rho _{\rm H\,{\small I}}}(z)$. Here, we assume that the baryonic distribution traces the cold dark matter distribution at large scales, so $\delta _{\rho _{\rm H}} = \delta _{\rho }$. Also, we assume that no additional hemispherical power asymmetry was generated during the evolution of the Universe, so the dipolar modulation of the curvature perturbation is transferred to the 21-cm power spectrum with the same form.
Citation: Mao et al. 2013
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The sensitivity of dwarfs to galaxy formation physics, like the gas density at which stars form (e.g. Kravtsov 2003; Saitoh et al. 2008) or the details of galactic outflows (e.g. Read et al. 2006b), makes them a natural ‘rosetta stone’ for constraining galaxy formation models. Early work simulating dwarf galaxies focused on high resolution small box simulations, stopping at high redshift (z ∼ 5–10) to avoid gravitational collapse on the scale of the box (e.g. Abel et al. 1998; Gnedin & Kravtsov 2006; Read et al. 2006b; Mashchenko, Wadsley & Couchman 2008). These simulations demonstrated that stellar winds, SN feedback, and ionizing radiation combine to prevent star formation in halo masses below ∼107–8 M⊙ (Read et al. 2006b; Bland-Hawthorn, Sutherland & Webster 2015). Furthermore, once cooling below 104 K is permitted and gas is allowed to reach high densities $n_{\rm max} {\gtrsim}10\,{\rm cm}^{-3}$ (requiring a spatial and mass resolution better than $\Delta x {\lesssim}100$ pc and mmin  103 M⊙; Pontzen & Governato 2012; Benitez-Llambay et al. 2019; Bose et al. 2019; Dutton et al. 2019), star formation becomes much more stochastic and violent (Mashchenko et al. 2008; Pontzen & Governato 2012; Dutton et al. 2019). The repeated action of gas cooling and blow-out due to feedback expels DM from the galaxy centre, transforming an initially dense DM cusp to a core (e.g. Read & Gilmore 2005; Pontzen & Governato 2012). Finally, independent of any internal sources of feedback energy, cosmic reionization can halt star formation in low-mass galaxies (e.g. Gnedin & Kravtsov 2006). Dwarfs that have not reached a mass of $M_{200} {\gtrsim}10^8$ M⊙ (see Section 2.4 for a definition of M200) by the redshift that reionization begins (z ∼ 8–10; Gnedin & Kaurov 2014; Ocvirk et al. 2018) are gradually starved of fresh cold gas, causing their star formation to shut down by a redshift of z ∼ 4 (Oñorbe et al. 2015). This is similar to the age of nearby ‘ultra-faint’ dwarf galaxies (UFDs) that have $M_* {\lesssim}10^5$ M⊙, suggesting that at least some of these are likely to be relics from reionization, inhabiting pre-infall halo masses in the range M200 ∼ 108–9 M⊙ (Gnedin & Kravtsov 2006; Bovill & Ricotti 2009, 2011; Brown et al. 2014; Weisz et al. 2014; Jethwa, Erkal & Belokurov 2018; Read & Erkal 2019).
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Using a single-fluid approach Chian & Oliveira (1994, 1996), and Oliveira & Chian (1996) studied the effects of parametric instabilities on standing waves with applications to the magnetosphere. These authors developed a theory of MHD parametric instabilities driven by a left- or right-hand circularly polarised standing Alfvén wave in a low plasma-β under very specific initial conditions. Israelevich & Ofman (2004) and Israelevich & Ofman (2011) using both the MHD approach and hybrid simulations investigated numerically the generation of an electric field parallel to the magnetic field, due to non-linear induced ion-acoustic waves produced by the presence of a standing Alfvénic pump. In a different context, Mottez (2015) also explored the generation of a parallel electric field using PIC (particle in cell) simulations when two counter-propagating Alfvén waves are excited. However, the mechanism proposed by this author is not related to the ponderomotive force that it is usually operating under such conditions. A ponderomotive force (see the review of Lundin & Guglielmi 2006) generates density enhancements periodically in time and space which are normally called cavities and depend on gradients in density or electric fields. The periodicity depends on the value of gas pressure of the system and the amplitude of these enhancements is related quadratically with the amplitude of the initial transverse excitation (see also Rankin et al. 1994; Tikhonchuk et al. 1995). Interestingly, ponderomotive forces can lead to separation of species, although this topic is not addressed in the present work since the quasi-neutrality condition is imposed for the plasma composed of kinetic protons and the treatment of electrons is that of a fluid. Recently, Martínez-Gómez et al. (2018) used the multi-fluid approach to model high-frequency waves in partially ionised plasmas similar to those in quiescent solar prominences. These authors show that non-linear Alfvén waves generate density and pressure perturbations and that the friction due to collisions dissipates a fraction of the wave energy, which is transformed into heat that raises the temperature of the plasma.
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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
.

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: Gandhi et al. 2018a
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We study the impact of secondary outgassing associated with partial melting and volcanism following accretion and a possible magma ocean phase. Raymond et al. (2007) showed that the accretion of planets around M-dwarf stars is faster than around more massive stars, which has also been calculated by Lissauer (2007) and Ida & Lin (2005). We therefore assumed that the accretion of a planet around an M dwarf is very fast (1 Myr), while for the other stellar cases we assumed an accretion time of 20 Myr. This corresponds to the lower limit in accretion times of Raymond et al. (2007) for planets around stars with masses greater than 0.8 MSun. After accretion, the planet may reside in a magma ocean stage, which we assume to last 1 Myr, following the results by Spohn & Schubert (1991) and Lebrun et al. (2013), among others. For high stellar irradiation, the planet may be trapped in a long-term magma ocean stage, which may last from several millions of years to billions of years (see e.g. Hamano et al. 2013, Nikolaou et al. 2019). For planets in a long-term magma ocean stage, Hamano et al. (2013) suggested very low water concentrations in the mantle at the end of the magma ocean stage because the planet can only exit this stage via severe atmospheric water loss because thick, water-rich atmospheres show a thermal blanketing effect suppressing the cooling of the planet via the Planck feedback. Their estimates of interior water reservoirs for such cases correspond to our lower limit of the initial mantle water concentration of 34 ppm. We nevertheless assumed in our calculations that the magma ocean has a short duration. This allows the high-luminosity pre-main sequence phase of the M dwarfs to have a larger impact on the HZ evolution. Any CO2 and H2O outgassed during the magma ocean phase is neglected in our calculations, i.e. we started our outgassing calculations with a 1 bar N2 atmosphere. Neglecting any atmospheric H2O and CO2 from the magma ocean phase results in a very narrow HZ at the beginning of our calculations, which is then only extended by secondary outgassing. This facilitates the evaluation of the potential of secondary outgassing from a stagnant-lid planet to form a HZ.
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