A common feature of all previous models is that the relation between the central galaxy stellar mass and the halo mass reaches a maximum at halo masses ~1012 M⊙. According to Yang et al. (2012), below this threshold the mass accretion of the central galaxy is dominated by star formation. Thus, when the halo mass reaches ~1012 M⊙ a process takes place that quenches the star formation. Interestingly, this mass scale is very similar to the cold-mode to hot-mode transition scale (Birnboim & Dekel 2003; Kereš et al. 2005) in the theory of gas accretion, as derived in hydrodynamic simulations, whereas large halos primarily accrete hot gas and low mass halos cold gas. This would suggest that the quenching of central galaxies coincides with the formation of a hot gaseous halo, and thus with a lack of cold gas supply. What would be the fate of satellites? According to Simha et al. (2009), the subhalos also retain their identity for quite some time after accreting a larger halo, so satellites in subhalos less massive than ~1012 M⊙ do not immediately see the effect of the hot gas in the larger halo and accrete in cold mode. Thus, consistent with the results of Yang et al. (2012) and Béthermin et al. (2013), satellite galaxies continue to accrete gas and convert it to stars over a rather long period, which according to Simha et al. (2009) is about of 0.5−1 Gyr after the merger. The gas accretion declines steadily over this period. Since star formation follows mass accretion with a short delay, satellites should experience quenching in a similar amount of time. This scenario would be consistent with our observations. Indeed, at z ~ 1 when massive halos are just forming via merger, the SF activity in the accreted subhalos is still high. At later epochs, instead, the transition to the hot mode accretion of the satellites and the consequent progressive quenching of their SF activity would lead to the faster decline of their contribution to the CSFH with respect to lower mass halos, which evolve in a cold mode accretion phase maintaining a high SFR.


A number of ideas have been put forward to explain the formation and early evolution of the compact Kepler and radial velocity systems, which in cases such as Gliese 581 and HD 69830 appear to contain in excess of ∼30 M⊕ of solid material within a few tenths of an au (Lovis et al. 2006; Udry et al. 2007). This concentration of solids close to the star led to classical core accretion models combined with disc-driven migration being developed using population synthesis codes (Alibert et al. 2006). More recent population synthesis calculations that also include prescriptions for planet–planet interactions have also been presented (Ida & Lin 2010). N-body simulations, combined with either hydrodynamic simulations or analytic prescriptions for migration and eccentricity/inclination damping of planetary growth, have also been used to examine the origins of such systems (Cresswell & Nelson 2006, 2008; Terquem & Papaloizou 2007; McNeil & Nelson 2009, 2010; Hellary & Nelson 2012; Cossou, Raymond & Pierens 2013; Coleman & Nelson 2014; Hands, Alexander & Dehnen 2014). A common outcome of these N-body simulations is the formation of resonant convoys of planets in the presence of convergent migration, an outcome that is not reflected in the Kepler systems. Various ideas have been put forward to explain why the resonances may be unstable, including tidal eccentricity damping followed by separation of the resonance for short-period systems (Terquem & Papaloizou 2007), stochastic migration due to local turbulence (Adams, Laughlin & Bloch 2008; Rein & Papaloizou 2009; Rein 2012) – a process that is likely to only operate close to the star where the disc can be thermally ionized (Umebayashi & Nakano 1988; Desch & Turner 2015), resonance breaking due to overstable librations (Goldreich & Schlichting 2014), orbital repulsion due to non-linear spiral wave damping in planet coorbital regions (Podlewska-Gaca, Papaloizou & Szuszkiewicz 2012; Baruteau & Papaloizou 2013).


The millimetre bump in M87 as recently observed by the Atacama Large Millimeter/submillimeter Array can be naturally modelled by the synchrotron emission of the thermal electrons in the ADAF, which is different from the prediction of the jet model. Therefore, it provides an opportunity to explore the accretion process near the BH horizon. In particular, Feng etal. (2016) and Li etal. (2016) both found that the rotation measure predicted from the ADAF is roughly consistent with the observational values. It is still difficult to constrain the BH spin parameter from the modelling of the SED of M87 due to some degeneracy in model parameters (wind parameter, s, magnetic parameter ; Feng etal. 2016). The spin parameter can be better constrained from the jet model if the relativistic jet is indeed powered by the rotating BHs as suggested by MHD simulations and some observations. We find that the dimensionless BH spin parameter should be larger than 0.96 for the lower limit of jet power derived from the X-ray cavities (e.g. Rafferty et al. 2006; Russell etal. 2013b). In this work, we adopt several typical values of parameters (e.g. 0.1, 0.3 and 0.5). The larger value of  will lead to a lower accretion rate near the horizon to explain the observed millimetre bump, and the BHs need to rotate faster to reproduce the observed jet power. The peak of synchrotron emission from the thermal electrons of ADAF will move to the submillimetre waveband if  is too small, which is different from the observed millimetre bump. We adopt the equipartition case of 0.5 in our calculations, where magnetic energy will become dominant if the BH is fast spinning, considering the possible amplification of the magnetic field by the frame dragging effect. For the weaker magnetic case (e.g. 0.5), the BH needs to rotate faster to explain the observed SED and jet power. We find that our results are not sensitive to the viscosity parameter . Therefore, we suggest that the BH should be fast rotating in M87 even after considering the possible uncertainties.


Cosmic H i gas density ($\Omega _{\rm {H}\,\small {I}}$) as a function of redshift (bottom axis) and lookback time (top axis). All measurements are corrected to the same cosmological parameters. Some DLA measurements adopted a different definition of the cosmic H i density taking into account neutral gas abundance (Ωgas) including helium or contributions from Lyman-α absorbers with lower column density ($\log N({\rm H\,{\small {I}}}) < 20.3$). We have corrected all measurements to a consistent definition of $\Omega _{\rm {H}\,\small {I}}$. The large red star shows the $\Omega _{\rm {H}\,\small {I}}$ measurement from this work. The small black square and triangle at z ∼ 0 are the HIPASS and ALFALFA 21-cm emission measurements by Zwaan et al. (2005) and Martin et al. (2010), respectively. The red diamonds are from the Parkes telescope with an H i stacking technique (Delhaize et al. 2013). The open circle and square are the results from the Arecibo Ultra Deep Survey (AUDS) (Freudling et al. 2011; Hoppmann et al. 2015). Two right-pointing triangles are measured using the WSRT and H i stacking technique by Rhee et al. (2013). The pink triangle is measured by Lah et al. (2007) applying the GMRT 21-cm emission stacking. The green hexagon denotes the H i stacking measurement for the COSMOS field using the GMRT (Rhee et al. 2016). The black arrow line is the upper limit constrained using the GMRT H i stacking (Kanekar et al. 2016). The closed and open diamonds, left-pointing triangle, big circle, open triangles, open diamonds and small circles are damped Lyman-α measurements from the HST and the SDSS by Rao et al. (2006, 2017), Neeleman et al. (2016), Prochaska et al. (2005), Noterdaeme et al. (2009), Noterdaeme et al. (2012) and Bird et al. (2017), respectively. The downward triangles and big square at high redshift of z > 2 are ESO UVES and Gemini GMOS measurements of DLAs by Zafar et al. (2013), Crighton et al. (2015), respectively. The black line with a shaded area shows a weighted fit of all $\Omega _{\rm {H}\,\small {I}}$ measurements and its 95 per cent confidence interval. The solid blue line shows the semi-analytic model prediction of Kim et al. (2015), and the dashed red line shows the mufasa (Davé et al. 2017) simulation.


Five XRT observations got only low exposure (i.e., much less than 1 ks) and the X-ray source was therefore undetected in these data sets. Surprisingly, we also found that the source was undetected in a “deep” observation taken on 2015 August 19 with an exposure time of about 1.6 ks (Table 2). Within a 47″ radius circular region centered at the source position (corresponding to 90% of the encircled energy fraction of XRT at 1.5 keV; Moretti et al. 2005), only one photon (which is located near the edge of the region) was detected in this 1.6 ks observation. Even assuming that only this event is from the source, the inferred count rate is much lower than the measurements in 2010 and 2015–2017; for example, seven source counts would have been detected in a 1.6 ks observation with the count rate of 4.4 × 10−3 cts s−1 measured five days later (Table 2). Using a Bayesian approach (Kraft et al. 1991), we computed 95% upper limits for all the nondetections. As expected, the upper limits for data with 1 ks are not very much constraining (i.e., a few ×10−2 cts s−1, while the average count rate of the four individual detections is about 4 × 10−3 cts s−1). The upper limit for the 1.6 ks data is deeper (i.e., 8.7 × 10−3 cts s−1), but still insufficient to clarify whether the low-count-rate measurement is physically or statistically based. For a deeper constraint, we combined all the six XRT observations, and the X-ray source can be marginally detected in the stacked image with

 cts s−1. Although this marginal detection shows a ∼50% decrease on flux in the period from 2015 February 04 through August 19, the variability is not statistically significant (i.e., less than 2σ). To check whether this variability was seen at other frequencies, we performed a Fermi-LAT analysis with the data collected between 2015 February 04 and August 19, and the γ-ray flux (100 MeV–100 GeV) did not vary significantly. In UV, there are some Ultraviolet/Optical Telescope (UVOT) images taken simultaneously with the XRT observations. Although the UVOT magnitudes (obtained by aperture photometry using the uvotsource task in HEAsoft v6.22) significantly changed over time (Table 2), this was due to the orbital modulation (Figure 3; will be discussed in the coming sections).


There are only three exceptions from the picture discussed above. In Fig. 2 we see that the model with the smallest number of objects, $N=40\, 000$, does not show abrupt dissolution. This model is characterized with the shortest initial relaxation time (among models considered in this paper), about 1.6 Gyr, and the smallest number of retained BHs, about 30. In such a situation, according to Breen & Heggie (2013), BHs are quickly kicked out from the system and BHS cannot survive for a long time. So, the cluster evolution is not governed by strong energy generation by BHs in BHS. In Fig. 3 we can see that the model with W0 = 3 dissolves extremely fast and the model with W0 = 9 does not show the abrupt dissolution feature. The very fast dissolution of tidally filling models with low King model concentration was already extensively discussed in the literature (e.g. Weinberg 1993; Fukushige & Heggie 1995; Whitehead et al. 2013; Contenta et al. 2015). The dissolution of such clusters is controlled by very strong initial mass-loss powered by stellar/binary evolution. Relaxation process is not important at all. The situation is much different for the model with W0 = 9. As it can be seen from Fig. 5, the cluster enters the core-collapse phase, so it has to be controlled by the relaxation process. The difference between the model with W0 = 6 and the model with W0 = 9 is connected with the rate of BHS evolution. Model W0 = 9 is initially much denser, so its half-mass radius and half-mass relaxation time are shorter than for the W0 = 6 model. Nevertheless, for both models the BH mass segregation ends nearly at the same time, about 4 Gyr. At that time, the models contain about 50 BHs and 160 BHs for W0 = 9 and W0 = 6, respectively. According to Breen & Heggie (2013) the evolution of BHs is controlled by the energy flow through the cluster half-mass radius, which is proportional to Eb/Trh ≈ GM/Rh/Trh, where Eb is the cluster binding energy. Models with W0 = 9 have smaller half-mass radius and half-mass relaxation time than models with W0 = 6, so the energy demand is larger for this model and leads to a much faster burning out of BHs – larger number of dynamical interactions leading to BH removal from the system. After the time of BH mass segregation (about 4 Gyr) the model W0 = 6 contains enough BHs to form a BHS and enters the phase of balanced evolution (Breen & Heggie 2013). Contrarily, the model with W0 = 9 has too small a number of BHs and continues to remove BHs quickly to support the needed energy flow. Finally, it enters the phase when other energy sources connected with ordinary binaries take over and the cluster enters the core-collapse and then core-bounce phases. It is important to note that the tidal field plays an important role in the above picture. The less concentrated model has larger Rh and loses more mass, so it is easier for the BHS to provide the needed energy to support the cluster structure. This phase of evolution ends when the cluster suddenly loses its virial equilibrium.


Since its launch in March of 2009, the NASA Kepler mission has been monitoring ∼160,000 stars in order to detect transiting extrasolar planets with high relative photometric precision (∼20 ppm in 6.5 hr, Jenkins et al. 2010). In 2013 May, the Kepler main mission ended with the failure of a second reaction wheel; however, the first four years of Kepler data have led to a wealth of planetary discoveries with a total of 4706 announced planet candidates13

13

http://exoplanetarchive.ipac.caltech.edu/ as of 2015 November 11.
 (Borucki et al. 2010, 2011; Batalha et al. 2013; Burke et al. 2014). The confirmed and candidate exoplanets typically have orbital periods shorter than 1000 days because at least three detected transits are needed for identification by the automated Transit Planet Search algorithm. Therefore, transiting exoplanets with periods longer than ∼1000 days are easily missed. The detection of short-period planets is further favored because the transit probability decreases linearly with increasing orbital distance. For these reasons, estimates of the statistical occurrence rate of exoplanets tend to focus on orbital periods shorter than a few hundred days (e.g., Dong & Zhu 2013; Fressin et al. 2013; Petigura et al. 2013). Radial velocity (RV) techniques also favor the detection of shorter period orbits. While gas giant planets have been discovered with orbital periods longer than a decade, their smaller reflex velocity restricts detection of sub-Neptune mass planets to orbital radii less than ∼1 AU (Lovis et al. 2011). In principle, astrometric observations favor longer period orbits; however, high precision needs to be maintained over the correspondingly longer time baselines. For shorter periods, the planets need to be massive enough to introduce a detectable astrometric wobble in the star and Gaia should begin to contribute here (Perryman et al. 2001). Microlensing offers sensitivity to planets in wider orbits and has contributed to our statistical knowledge about occurrence rates of longer period planets (e.g., Gaudi 2010; Cassan et al. 2012) and direct imaging of planets in wide orbits is also beginning to contribute important information (Oppenheimer & Hinkley 2009).


Even though most of the energy released by a gamma-ray burst (GRB) is emitted during the prompt emission phase, the emission mechanism is still not understood. In order to determine the emission process typically, the low-energy spectral index, α, of the GRB spectrum is analysed. This analysis has not given a conclusive answer since the α-distribution is broad and has not been uniquely explained by a single emission process. For instance, the peak of the distribution has been used as an argument for synchrotron emission since it is close to the expected value. However, a large fraction (${\sim } 28{{\ \rm per\ cent}}$) was found to be inconsistent with the theoretical limit of −2/3 (‘line of death’; Preece et al. 1998; Ghirlanda, Celotti & Ghisellini 2002; Guiriec et al. 2015; Goldstein et al. 2016; Yu et al. 2016); and only specific physical scenarios remain plausible (large emitting radii and Lorentz factors of the flow, Beniamini & Piran 2013; Iyyani et al. 2016; Beniamini, Barniol Duran & Giannios 2018; Burgess et al. 2018). Alternatively, the spectral width or sharpness angle has been used as a tool to characterize the spectrum, which also take into account the high-energy spectral slope (Axelsson & Borgonovo 2015; Yu et al. 2015b). Around 80 per cent of the bursts were found to have fitted Band spectra that are narrower than what was expected for synchrotron emission. However, direct fitting with a synchrotron model decreases this fraction (Burgess 2017). Therefore, a firm conclusion based on the spectral width or sharpness angle alone cannot either be reached. Yet another approach has been to study the correlation between spectral parameters. A well-studied and strong correlation is the Golenetskii correlation (Golenetskii et al. 1983) which relates the instantaneous flux, F, and peak of the spectrum Epk (Kargatis et al. 1994; Borgonovo & Ryde 2001; Lu et al. 2012). Another prominent correlation is between Epk and α which is valid in a fraction of bursts (Crider et al. 1997; Kaneko et al. 2006). Both of these correlations show a variety of behaviours. In some individual pulses, the spectral parameters track each other, while in others the correlation is different during the rise and decay phases of the pulses. The variety of behaviours have complicated the search for physical explanations.


Six out of the eight sources (75 per cent of the sample) show evidence for high obscuration by cold gas (column density NH≳ 1023 cm−2). The number of AGN with NH≳ 1024 cm−2 is in the range 2–4 when 90 per cent confidence errors on the column density are considered (see Table 4). Although our sample is definitely too small to make any strong conclusion in terms of statistical incidence of obscured AGN in dual systems, we note that the fraction of CT AGN in our sample of dual AGN is 25–50 per cent, i.e. higher than the fraction of hard X-ray selected CT AGN in isolated systems in a similar luminosity interval (BAT: 27 ± 4 per cent, Ricci et al. 2015, NuSTAR: ∼4 per cent, Marchesi et al. 2018). Recently, it has been suggested (Ricci et al. 2017) that AGN in multiple systems in a late stage of merging (with a projected distance d ≤ 10 kpc) are often type 2 AGN due to the huge amount of gas and dust obscuration channelled towards the nuclear region by the close galaxy encounter. Furthermore, in a sample of 44 UltraLuminous IR Galaxies (ULIRG) from the GOALS survey (Sanders et al. 2003) in the local universe, it has been demonstrated that the fraction of CT AGN in late mergers is higher (65$^{+12}_{-13}$ per cent) than in local isolated hard X-ray selected AGN (27$^{+4}_{-4}$ per cent, Ricci et al. 2017), while it is marginally higher than isolated AGN in early mergers (35$^{+13}_{-12}$ per cent). The increasing of the obscuration with the disturbed or interacting morphology of the galaxy host has been also observed up to higher redshift (up to z = 1.5) using X-ray Chandra data from deep surveys (Kocevski et al. 2015). These authors have analysed 154 heavily obscured AGN and compared their morphologies with control samples composed of moderately obscured and unobscured AGN. They found that the fraction of galaxies undergoing mergers is higher in AGN with NH above 3 × 1023 cm−2 with respect to unobscured AGN with NH  1022 cm−2 (7.8$^{+1.9}_{-1.3}$ versus 21.5$^{+4.2}_{-3.3}$ per cent). X-rays (Kocevski et al. 2015; Ricci et al. 2017) as well as more recent mid-IR colour selection (Satyapal et al. 2017) strongly suggest that AGN in late state of merging (10 kpc) are highly absorbed (but see also Ellison et al. 2017), with NH above 1024 cm−2. All these observational results suggest that CT AGN in dual systems correspond to a phase where the supermassive black hole is still accreting gas during an early stage of an interaction or merging.


The multiwavength UV to mid-IR SED of J1224+5555 supports the scenario of a highly obscured AGN, based on the empirically derived templates of Assef et al. (2010). These templates consist of a set of three galaxy templates and one AGN template. The galaxy templates, E, Sbc, and Im, are based on templates from Coleman et al. (1980) with wavelength ranges extended using stellar models from Bruzual & Charlot (2003) and added dust and polycyclic aromatic hydrocarbon components from Devriendt et al. (1999). The AGN template is based on the average type 1 AGN template from Richards et al. (2006) with efforts to remove host galaxy contamination and create a template for an unobscured AGN. This method has since been used to successfully uncover highly obscured and even Compton-thick AGNs in several other works thus far (Chung et al. 2014; Hainline et al. 2014; Lansbury et al. 2014; Stern et al. 2014; Assef et al. 2015; Chen et al. 2015; Lansbury et al. 2015). Figure 7 shows the best-fit linear combination of the three galaxy templates to the observed SDSS, Two Micron All Sky Survey (2MASS), and WISE photometry of the target galaxy, where the best fit is obtained through a

 minimization. The multiwavelength fluxes were obtained from the NASA/IPAC Extragalactic Database (NED).17

17

https://ned.ipac.caltech.edu/

 The maximum wavelength of the Assef et al. (2010) templates is 30 μm; therefore, we cannot use the IRAS  60 μm flux in our fits. As discussed in Section 5, the 60 μm flux for J1224+5555 is likely overestimated, and its use is in any case questionable. While the SDSS

, 2MASS

, and WISE W1 bands are well fit by a combination of the Sbc and Im templates, the combined model flux is much lower than the observed fluxes in the W2–W4 bands. In fact, Table 2 of Assef et al. (2010) shows that at z = 0 and z = 0.1 no galaxy template has a W1–W2 color greater than 0.2 (Vega). However, when we add in the AGN template with a variable level of obscuration (red line) in Figure 7, we see that the observed multiwavelength photometry of the galaxy is better fit with a combination of a young and old stellar population and highly obscured AGNs (red line) dominating the WISE bands, with a

 value almost a factor of 7 times lower than the best fit obtained using the galaxy templates alone. As an additional test, we investigated the effect of adding variable additional extinction to the galaxy templates alone to determine if a comparable fit can be obtained without invoking the presence of an obscured AGN contribution. However, as can be seen from Figure 7, the best-fit model cannot adequately fit the SED in the WISE bands and results in a

 value that is significantly worse than the model that includes an AGN (Figure 7). We note that the improvement in fit when the AGN template is added cannot be attributed simply to having additional model parameters in the fit. In the first model without an AGN, we achieve a

 of 1323.29, with three model parameters. While the addition of an AGN template more than doubles the number of model parameters to 8,

 is reduced by a factor of

, indicating that a model that includes an AGN template produces a significant improvement to the fit of the data.


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).


Table 2

Catalogue description.

Column
 Description

1
Source number

2−3
X-ray coordinates, right ascension and declination (epoch 2000.0)

4
Uncertainty of X-ray position [′′]. For XMM-Newton positions taken from Sturm et al. (2013c) the 1σ  error includes a systematic uncertainty of 0.5′′.

5
Origin of the X-ray coordinate (A: ASCA, C: Chandra, E: Einstein, I: Integral, N: XMM-Newton, R: ROSAT, S: Swift, X: RXTE). When no reliable position could be determined from the non-imaging RXTE collimator-instruments, a radius of 30′ for the position error indicates the size of the field of view.

6
Reference for source discovery.

7
Identification of optical counterpart with emission-line star from Meyssonnier & Azzopardi (1993). The negative number indicates a star found in the catalogue of Murphy & Bessell (2000).

8−15
Flags indicating different source properties. For their description see Table 3.

16
Confidence class (values 1−6, see Table 4).

17−18
Optical coordinates, right ascension and declination (epoch 2000.0) for the identified counterpart from Zaritsky et al. (2002), or – when not available there – from Massey (2002).

19−26
The Magellanic Clouds Photometric Survey (MCPS): U, error(U), B, error(B), V, error(V), I, error(I) [mag] from Zaritsky et al. (2002).

27−32
Colour indices U−B, error(U−B), B−V, error(B−V), V−I, error(V−I) [mag] derived from MCPS photometry.

33−34
Reddening-free Q-value (Q = U−B−0.72 × (B−V)) and error(Q) [mag].

35
Near-IR counterpart to the optical star from the Two Micron All Sky Survey (2MASS, Skrutskie et al. 2006).

36−41
Near-IR magnitudes with corresponding errors: J, error(J), H, error(H), K, error(K) [mag].

42−45
Near-IR colour indices, J−H, error(J−H), H−K, error(H−K) [mag].

46−53
Spitzer IRAC fluxes at 3.6, 4.5, 5.8 and 8.0 μm [mag] with respective errors (from the SAGE project, for a description see Meixner et al. 2006).

54
Angular distance between X-ray and optical position [′′].

55
Angular distance between optical and near-IR position [′′].

56
Neutron star spin period [s] inferred from X-rays.

57
Orbital period [days] (see flags for origin).

58−59
Maximum and minimum X-ray flux [erg cm-2 s-1] when available in the 0.2−10 keV band. Fluxes in the SMC XMM-Newton catalogue of Sturm et al. (2013c) are given for the 0.2 to 4.5 keV band. To convert them into the 0.2−10 keV band, we multiplied them by a factor of 2.6 assuming a standard power law with photon index 0.9 (Haberl et al. 2008) and a column density of 1021 cm-2 (solar abundance). For Swift XRT count rates we used a flux conversion factor of 1.1 ×  10-10 erg cm-2 cts-1

60
Flag for minimum flux: 1 for a non-detection with an upper limit; −1 when unknown; 0 for detection.

61−62
References for maximum and minimum X-ray flux.

63
X-ray variability factor (ratio of maximum to minimum flux).

64
Equivalent width of the Hα line [Å] (minimum value if more than one measurement is available).

65
Maximum equivalent width of the Hα line [Å].

66
References for the Hα measurements.

67
Comments with key references.


When energy flows from the interior of the Sun outward into the solar atmosphere, why is the Sun’s outer atmosphere, the corona, much hotter than the inner atmosphere, the underlying chromosphere and photosphere? This is the long-standing problem of the coronal heating, which is one of the eight key mysteries in modern astronomy (Kerr 2012). For about 80 yr since the discovery of the extremely hot corona around the late 1930s (Grotian 1939; Edlen 1945), people have worked hard on addressing this issue, and great advances have been made in observation and theoretical studies (Parnell & De Moortel 2012; Amari et al. 2015; Arregui 2015; Cargill et al. 2015; De Moortel & Browning 2015; Jess et al. 2015; Klimchuk 2015; Longcope & Tarr 2015; Peter 2015; Schmelz & Winebarger 2015; Velli et al. 2015; Wilmot-Smith 2015). Especially during recent decades, high-resolution observations of solar super-fine structures indicate that small spicules, minor hot jets along small-scale magnetic channels from the low atmosphere upwards to the corona, petty tornados and cyclones, and small explosive phenomena such as mini-filament eruptions and micro- and nano-flares—all of these small-scale magnetic activities contribute greatly to coronal heating (De Pontieu et al. 2011; 2018; Zhang & Liu 2011; Parnell & De Moortel 2012; Klimchuk 2015; Peter 2015; Schmelz & Winebarger 2015; Henriques et al. 2016; Li et al. 2018a). Additionally, contributions of MHD waves to heating the corona have been observationally illustrated (van Ballegooijen et al. 2011; Jess et al. 2015; Kubo et al. 2016; Morton et al. 2016; Soler et al. 2017; Morgan & Hutton 2018). Meanwhile, with the progress of observational studies, two groups of theoretical models, magnetic reconnection models and magnetohydrodynamic wave models, have traditionally attempted to explain coronal heating, but so far no models can address the key mystery perfectly (van Ballegooijen et al. 2011; Arregui 2015; Cargill et al. 2015; Peter 2015; Velli et al. 2015; Wilmot-Smith 2015). Maybe we do not need to intentionally take to heart such the classical dichotomy, because waves and reconnections may interact with each other (De Moortel & Browning 2015; Velli et al. 2015). Additionally, statistical studies may look at coronal heating from a comprehensive perspective. Li et al. (2018b) found that the long-term variation of the heated corona, which is represented by coronal spectral irradiances, and that of small-scale magnetic activity are in lockstep, indicating that the corona should statistically be effectively heated by small-scale magnetic activity. Observational and theoretical model studies through heating channels and modes, and statistical studies by means of heating effect (performance of the heated corona), both suggest that coronal heating originates from small-scale magnetic activity.


Figure 6 shows the orbital eccentricities as a function of [M/H] for the metal-rich disc sample. The solid lines correspond to the required eccentricity (see Eq. (2)) for different values of ISM radial metallicity gradients: −0.10 dex kpc−1 (black), −0.07 dex kpc−1 (our measured gradient for young stars in Table 1; see also Minchev et al. 2018, red), −0.04 dex kpc−1 (orange), and −0.06 dex kpc−1 (Cepheids analysis from Genovali et al. 2014, green). For the three first cases, we assumed ISM[M/H](R⊙) = 0.0 to estimate Rbirth from the stellar metallicity. However, Genovali et al. (2014) have their own zero point, defined as: [Fe/H] = −0.06 * Rg + 0.57, with a clear shift in the relation compared to the other ones assumed in this work. The impact of the ISM gradient value and the zero-point assumption on the derived Rbirth, and therefore on the required eccentricity to reach the solar vicinity without the need for churning, is clearly observed. As described in Hayden et al. (2020), given the measured [M/H] and eccentricity, stars lying to the left are able to reach the solar neighbourhood through blurring, while the stars to the right of the line are possible candidates to have migrated through churning. This is the case for most of the SMR stars (70% of the SMR stars lie below the line that corresponds to the Cepheids analysis); they are therefore likely to have been brought to the solar neighbourhood by churning, which is in close agreement with previous studies (e.g., Kordopatis et al. 2015a; Wojno et al. 2016). However, it is worth noting that the observed metallicity distribution function in Fig. 2 peaks around 0.2 dex, which is higher than previous reported solar vicinity MDFs (see e.g., Fuhrmann et al. 2017). A possible ignored bias towards more metal-rich objects in the sample selection could be pulling the percentage of possible migrators to higher values. Among the entire distribution, our churned candidates with [M/H] >  + 0.1 comprise around 17% of the sample. If we constrain the number of migrators to only stars with [M/H] >  + 0.25, the global percentage decreases to 8% of the sample.


Our basic model for acceleration is shown in Figure 1. The left panel shows a cartoon of the model for acceleration, transport, and radiation at the X-reconnection site in the corona. Acceleration can be by the second order Fermi (or stochastic) acceleration by turbulence (see, e.g., Petrosian & Liu 2004), by a standing shock produced by the downflow from the X-reconnection site (see, e.g., Guo & Giacalone 2012), or in the merging of islands shown to arise in PIC simulations during reconnection (Drake et al. 2006, 2013; Le et al. 2012; Oka et al. 2010). The latter model has been invoked as a possible mechanism in the downstream of the CME shock (Le Roux et al. 2015; Zank et al. 2015). The right panel shows a similar cartoon for acceleration in the environment of the CME. The (red) rectangles show the cross section of the box within which particles are accelerated. For our purposes here the relevant parameters of the acceleration and transport of particles are the momentum (p) and pitch angle cosine (μ) diffusion rates

 and

, and direct energy gain

 and loss

 rates, where E is the particle energy.1

1
Here and in what follows we neglect the effects of the third diffusion coefficient

, which are generally small (see Schlickeiser 1989; Petrosian & Liu 2004).
 For an isotropic pitch angle distribution the evolution of the volume-integrated energy spectrum,

, is governed by the so-called leaky-box model kinetic equation (see Petrosian 2012)
1

where

 is a source term,

, and the direct acceleration coefficients for stochastic and shock accelerations are given as
2

Here

 is the shock velocity,

 is the coefficient of spatial diffusion (along the field lines),

,

 is the Lorentz factor, and ζ depends on the shock compression ratio and other factors (see, e.g., Steinacker et al. 1988). The energy-loss rate at low energies (for both electrons and ions) is dominated by Coulomb collisions with background particles (mainly electrons).2

2
Coulomb collisions also cause pitch angle scattering, and therefore spatial diffusion along the field lines, with a rate that is comparable to the energy-loss rate at nonrelativistic energies but decreases as

. They also cause energy diffusion, which is negligible in the cold target case (

), but can be comparable to the energy-loss rate for a hot target as E approaches kT (see Petrosian & Kang 2015), where T is the background plasma temperature.
 At higher energies (not relevant for the discussion here) inelastic interactions (synchrotron, inverse Compton, and bremsstrahlung) for electrons and (nuclear line excitation, neutron and pion productions) for ions become important.


In order to explode as an ECSN, several ingredients need to be in place. Nomoto (1984) argued that stars with helium cores between 2.0 and

 (which corresponds roughly to initial masses between 8 and

) would explode as an ECSN. His models, however, did not develop a second dredge-up, which can significantly reduce the mass of the helium core and diminishes the predictive power of this criterion (see Podsiadlowski et al. 2004; Poelarends et al. 2008). Since then, several authors (Siess 2007; Poelarends et al. 2008; Doherty et al. 2010, 2015; Jones et al. 2013) have established precise initial mass ranges for ECSNe to occur, although these mass ranges are highly sensitive to the adopted convection criteria, overshooting, and mass-loss prescriptions (Poelarends et al. 2008; Doherty et al. 2010; Langer 2012). Jones et al. (2013) produced several detailed models, computed all the way to electron captures on 24Mg and 20Ne and found that CO cores with masses over

 are able to reach densities high enough for this to occur (

 g cm−3). If the CO core is massive enough, neon will ignite off-center (Jones et al. 2013; Schwab et al. 2016), but Jones et al. (2014) also found that the upper boundary for ECSNe is affected by uncertainties regarding the progression or stalling of the neon flame. There seems to be consensus, however, that the mass of the CO core is a reliable indicator for the final fate of stars in this mass range. How this translates into the initial mass of the star depends on the adopted convection criterion, with the Schwarzschild criterion producing more massive cores than the Ledoux criterion for the same initial mass. Inclusion of overshooting will also lead to larger cores (see Siess 2007). However, although the initial mass range for ECSNe is therefore quite sensitive to the adopted convection criteria, this does not seem to be the case for the final MCO, as most authors find similar values for MCO at which neon ignites, even though they treat convection differently. In the context of binary evolution, this provides an additional reason to adopt MCO as an indicator for whether the star explodes as an ECSN or not, as the CO core is generally not eroded by mass transfer, in contrast to the He core.


Local dwarf galaxies were the focus of large Herschel and Spitzer surveys (e.g., The Dwarf Galaxy Survey, DGS; Madden et al. 2006). Studies on both resolved and integrated-galaxy scales have highlighted some distinctively unique observational signatures of star-forming low-metallicity dwarf galaxies. A non-linear relation of the dust-to-gas mass (D/G) with metallicity is observed, with extremely low dust masses observed for the lowest metallicity galaxies (Z ≤ 0.1 Z⊙; Herrera-Camus et al. 2012; Fisher et al. 2014; Rémy-Ruyer et al. 2015; Galliano et al. 2018, 2021; Cigan et al. 2021). Furthermore, the hard radiation fields in star-forming dwarf galaxies, along with their lower dust abundance, result in extended ionized gas regions prominent on global galaxy scales (Hunter et al. 2011; Cormier et al. 2015, 2019). The consequence is often a largely photodissociated molecular phase, existing in clumps which are difficult to observe with the usual molecular gas tracer, CO (1-0) (Cormier et al. 2014; Hunt et al. 2015; Accurso et al. 2017b), beckoning the presence of a CO-dark molecular gas phase (Grenier et al. 2005; Röllig et al. 2006; Wolfire et al. 2010; Glover & Clark 2012; Bolatto et al. 2013; Accurso et al. 2017a; Madden et al. 2020). Other emission lines, however, such as the far-infrared [C ii]λl58 µm line, are strikingly enhanced on global scales in dwarf galaxies (e.g., Cormier et al. 2015, 2019; Cigan et al. 2016; Lebouteiller et al. 2017; Jameson et al. 2018), making the [C ii]λl58 µm line a potential tool for tracing star formation activity (Malhotra et al. 2001; Papadopoulos et al. 2007; Pineda et al. 2014; De Looze et al. 2014; Olsen et al. 2015; Herrera-Camus et al. 2015, Herrera-Camus et al. 2018; Carniani et al. 2018; Matthee et al. 2019; Izumi et al. 2021; Fujimoto et al. 2021) and a tracer of the total H2 in galaxies, near and far (Poglitsch et al. 1995; Wolfire et al. 2010; Pineda et al. 2013; Nordon & Sternberg 2016; Fahrion et al. 2017;
Accurso et al. 2017b; Zanella et al. 2018; Madden et al. 2020; Schaerer et al. 2020; Tacconi et al. 2020).


The first study, to our knowledge, to notice the net increase in the abundance of slow (s) neutron capture elements in young stellar populations is D’Orazi et al. (2009), in which the abundance of barium in young star clusters was seen to be higher than in the older ones. Maiorca et al. (2011, 2012) added a few more elements with important s-process contributions (yttrium, zirconium, lanthanum, and cerium), confirming the increasing trend towards younger ages. Subsequently, a number of works have attempted to both clarify the origin of this increase (see, e.g., Bisterzo et al. 2014; Mishenina et al. 2015; Trippella et al. 2016; Magrini et al. 2018; Spina et al. 2018; Busso et al. 2021) and to use their abundances to estimate the ages of stars, often using neutron capture s-process elements in combination with other elements with opposite behaviours, such as α elements – that we indicate as chemical clocks – and thus maximising the dependence of the relationship with age (see, e.g., Tucci Maia et al. 2016; Nissen 2016; Feltzing et al. 2017; Fuhrmann et al. 2017; Slumstrup et al. 2017; Titarenko et al. 2019). Once the existence of a relationship between age and chemical clocks was established (see, e.g., Spina et al. 2016; Delgado Mena et al. 2019; Jofré et al. 2020), the next steps were the following: (i) to clarify the applicability of these relationships with luminosity class (dwarf or giant) (see, e.g., Tucci Maia et al. 2016; Slumstrup et al. 2017; Casamiquela et al. 2021), metallicity (see, e.g., Feltzing et al. 2017; Delgado Mena et al. 2019; Casali et al. 2020), and population type (thin disc, thick disc, halo) (see, e.g., Titarenko et al. 2019; Nissen et al. 2020; Tautvaišienė et al. 2021), or even in dwarf galaxies (Skúladóttir et al. 2019; ii) to calibrate them with a sample of stars with reliable age determination, which are usually open star clusters (OCs), solar twins, or targets with asteroseismic observations. Finally, it is essential to understand whether these relationships are valid throughout the Galactic disc, or whether they are necessarily position-dependent. For the first time, Casali et al. (2020) applied the relations derived from a large sample of solar-like stars located in the solar neighbourhood and noted that they fail to reproduce the ages of star clusters in the inner disc. They concluded that the relationship between age and chemical clocks is not universal and that it varies with galactocentric position. Later, Magrini et al. (2021b) suggested that the differences in the relationships between age and chemical clocks in different parts of the Galactic disc are due to the strong dependence on the metallicity of the yields of low-mass stars, which produce s-process elements during the final stages of their evolution. Casamiquela et al. (2021) used red clump stars in open clusters to investigate the age dependence of several abundance ratios, including those that contain s-process and α elements. They found that the relationship between [Y/Mg] and ages outlined by open clusters is similar to the one found using solar twins in the solar neighbourhood. They also found that the abundance ratios involving Ba are those with the highest correlation with age. However, they also note that as one moves away from the solar neighbourhood, the dispersion increases and is in agreement with the findings of Casali et al. (2020), which attributed this to the spatial variation of the star formation history along the galactocentric radius.


Within the same scenario, Smith et al. (2004, 2005) have theoretically predicted that, when observed in polarized light, NLS1s should have a low polarization fraction, the polarization angle should change monotonically from one wing of the BELs to the other, and the BELs should be significantly broader than those observed in direct light (Smith et al. 2004, 2005). We assume that polarized BELs are scattered into the line of sight by material that is close to coplanar with the BLR (i.e. the torus). This prediction has been searched for in RQ-NLS1s, but many objects have not been detected in polarized light, and none of the detected show any significant broadening of the BELs (e.g. Goodrich 1989; Breeveld & Puchnarewicz 1998; Kay et al. 1999). On the other hand, only one RL-NLS1 (PKS 2004-447) has been the object of a spectropolarimetry study (Baldi et al. 2016), and the three predicted features have been observed that have a FWHM in the polarized Hα of ≈ 9000 km s−1. We propose for the first time a simple scenario that can resolve such apparent tension between the two classes of NLS1s. We speculate that in those NLS1s that are intrinsically without a jet (a fraction of the RQ-NLS1) a significant amount of polarized light could be due to polar scattering of the BELs, by material above the BLR, as observed for type II AGN (Antonucci 1983; Miller & Antonucci 1983; Antonucci & Miller 1985). A polar-scattered component could easily overwhelm the planar-scattering contribution, predicted to result in low polarization fractions for face-on objects because of the axi-symmetry of the scattering material. In RL-NLS1s, on the contrary, the jet could evacuate the intervening material, resulting in a lower degree of polarization in which only the broadened component of the BEL is left. The evacuation of the polar material by a jet closely aligned withthe line of sight has been already proposed by Ghisellini & Sbarrato (2016) to account for the lack of high-z blazar parent population. This simple speculative scenario can be tested through spectropolarimetry of a larger sample of RL-NLS1s.


Fig. 6 displays the MLR for the investigated CCs whose intrinsic stellar parameters were derived from the best-fitting models listed in Table 1. These data are compared with the predicted canonical (no overshooting, no mass-loss) MLR (Bono et al. 2000; the solid lines) and with the relations obtained by increasing the zero-point of the canonical MLR by 0.25 dex (the dashed lines) and 0.5 dex (the dotted lines) to reproduce the effect of mild and full overshooting,5 respectively (see Chiosi et al. 1993; Bono et al. 1999, for details). Inclusion of mass-loss and/or rotation would produce a similar increase in the Cepheid luminosity level at fixed mass (see Neilson et al. 2012, for details). As the light curves of $OGLE\_CEP\_LMC\_2019$ are best reproduced adopting a different value of the helium content (see above), in Fig. 6 we also show the MLR for Y = 0.30, Z = 0.008 (the green lines). Note that this relation is slightly more luminous than those calculated for the standard LMC elemental composition (Y = 0.25, Z = 0.008). According to the location of the variables in the ML plane, the canonical MLR is not strictly satisfied, as the points are spread between the canonical and full overshooting predictions. Even if at this stage we cannot disentangle the role of overshooting, mass-loss, and rotation in producing the quoted excess luminosity, at fixed mass, the detected dispersion might indicate a combination of these different non-canonical phenomena. Indeed, if only overshooting were efficient, one would in principle expect the same amount of excess luminosity for all stellar masses (within small uncertainties). Rotation produces similar effects as overshooting because it implies a larger He burning core and a brighter luminosity at fixed mass (see e.g. Anderson et al. 2016) On the other hand, if the mass-loss process were efficient, this could be inferred from the predicted deviation of the best-fitting stellar mass from the value corresponding to the canonical MLR. Such a deviation is represented in Fig. 7 as a function of the canonical mass (top) and of the pulsation period (bottom) for the CCs in our sample. We note that the expected mass differences range from $0{{\ \rm per\ cent}}$ to almost $\sim 50{{\ \rm per\ cent}}$ and are not clearly correlated with the pulsation period or the stellar mass.


Leitzinger et al. (2014) briefly compared the mass estimates of stellar mass ejection events and the estimated X-ray energies of their associated flares with the solar scaling. We extend their comparison here. Specifically, we consider the events observed as blueshifted extra emissions in Balmer lines on the active M dwarfs AD Leo (Houdebine et al. 1990), AT Mic (Gunn et al. 1994), V374 Peg (Vida et al. 2016) and the pre-main-sequence star DZ Cha (Guenther & Emerson 1997). Published estimates of their minimum masses are 7.7 × 1017, ∼1015, ∼1016 and 1.4 × 1018…7.8 × 1019 g, respectively. Since no simultaneous X-ray observations of the associated flares are available, we estimate the corresponding X-ray flare energies based on the assumption that the published U-band energies are of comparable magnitude (cf. Hawley & Pettersen 1991). Doing so, we find EX ≈ EU ∼ 2 × 1032 erg for AD Leo (Rodono et al. 1985; Hawley & Pettersen 1991), 3 × 1031 erg for AT Mic (Gunn et al. 1994) and 1.2…2.5 × 1035 erg for DZ Cha (Guenther & Emerson 1997), respectively. The recently published event at V374 Peg has an Hγ flare energy of ∼4 × 1030 erg, which corresponds to ∼1032 erg in X-rays (Butler, Rodono & Foing 1988). We add two other events, which are, however, more uncertain in their interpretation as real CME events. Doyle et al. (1988) observed a strong increase in neutral hydrogen column density during a flare on the active M dwarf YZ CMi, which may be interpreted as a rising filament obscuring parts of the flare region. The estimated mass of this event is 3 × 1017 g and the estimated X-ray flare energy EX ≈ EU ∼ 8 × 1030 erg, since the observed X-ray flare emission was likely partly absorbed by the neutral material (Doyle et al. 1988). The second one is a long-decay flare on the young, active M dwarf AU Mic, which was interpreted as an eruptive event with a CME mass in the order of 1020 g by Cully et al. (1994). However, there is a different model of this strong flare (∼3 × 1035 erg) including post-eruptive energy release, which does not involve a CME (Katsova, Drake & Livshits 1999). Other studies that observed stellar CME events (cf. Section 1) did not provide mass estimates and had to be excluded.


In the expectation of a relationship between the observed IRX and the rest-frame UV slope there is the assumption that the stars and dust are well mixed, which leads to the coupling of any observed reddening in the UV to the FIR emission detected (e.g. Meurer et al. 1999; Charlot & Fall 2000; Calzetti 2001). If instead the galaxy consists of regions of significantly different obscuration, then the relationship will break down for the galaxy as a whole. Indeed, geometric effects have been put forward as an explanation of the discrepant results at z > 5 (Faisst et al. 2017; Popping, Somerville & Galametz 2017). In local starburst galaxies, the existence of an IRX–β relation and a clear morphological similarity, indicates that the rest-frame UV emission from young stars is being attenuated from dust that is tracing broadly the same star-forming regions of the galaxy (e.g. in the spiral arms; Kennicutt et al. 2003; Gil de Paz et al. 2007). In the high-redshift Universe, however, where galaxies become more turbulent and irregular (e.g. Förster Schreiber et al. 2011; Buitrago et al. 2013; Guo et al. 2015) the expected morphology of the dust relative to the observed UV emission is not clear. Evidence for offset dust continuum emission relative to the rest-UV has been identified in several high-redshift Lyman-break galaxies (LBGs; Koprowski et al. 2016; Faisst et al. 2017; Laporte et al. 2017; Bowler et al. 2018), and similar trends have been found when comparing the [C ii] FIR line and the rest-UV continuum (e.g. Maiolino et al. 2015; Carniani et al. 2017). Although some of these offsets have been attributed to astrometric systematics (e.g. Dunlop et al. 2017) there is a growing consensus that FIR continuum and line emission are frequently physically offset as compared to the observed rest-UV emission (see Carniani et al. 2018). Whether high-redshift galaxies show large and distinct regions of obscured and unobscured star formation has implications for the use of the IRX–β relation in deriving the cosmic SFR density (e.g. Bouwens et al. 2016b), as the assumed energy balance will break down (Buat et al. 2019) and the global β measurement will not be representative of the full source.


Observations by a single coronagraph allow us to obtain CME 2D images projected on the plane of the sky. Using the image sequences of CME expansions, we can determine CME 2D parameters such as speed and angular width. Such observations do have some limitations, mostly the projection effect. Kwon et al. (2015) showed that the apparent width of CMEs does not represent an accurate measure of the CME ejecta size. In previous studies, geometrical models based on the assumption of different cone shapes were developed to obtain CME 3D parameters from single spacecraft observations (Xie et al. 2004; Xue et al. 2005; Michalek 2006). Na et al. (2013) compared the 3D parameters of halo CMEs obtained from the cone shapes (elliptical-cone, ice-cream-cone, and asymmetric-cone models) and found that their 3D speeds are comparable to one another but their angular widths and inclination angles are quite different. We note that stereoscopic observations by twin spacecraft STEREO make it possible to obtain the CME 3D parameters using geometrical triangulation methods (Mierla et al. 2008, 2009; Liewer et al. 2009; Temmer et al. 2009; Liu et al. 2010a, 2010b) and a flux rope method (Chen et al. 2000; Thernisien et al. 2009). Recently, the STEREO CME Analysis Tool (StereoCAT), based on a triangulation method from multiple spacecraft observations, was developed by the Community Coordinated Modeling Center (CCMC) for real-time CME analysis carried out by the CCMC/SWRC forecasting team. CME 3D parameters can be determined by selecting two coronagraphs from SOHO/LASCO C3 and STEREO/SECCHI COR2 images using the model. The triangulation algorithm and the performance of the CME measurements from StereoCAT are described in Sections 3.1 and 3.2 of Mays et al. (2015). Lee et al. (2015) showed, using 44 halo CMEs from 2010 to 2011, that the CME 3D parameters (speed, angular width, and source location) obtained from StereoCAT are similar to those obtained from the graduated cylindrical shape flux rope model (Thernisien et al. 2009; Thernisien 2011). Jang et al. (2016) showed, using 306 LASCO front-side halo CMEs from 2009 to 2013, that around 20% of the CME 2D speeds in the LASCO CME catalog are underestimated compared with the 3D ones obtained from StereoCAT.


A realistic EoS that is able to reproduce the properties of compact astrophysical objects has to fulfill several requirements. The possibility of including many particle species, which is known as multicomponent character, is of crucial importance for modeling the NS interiors, which in even the simplest treatment include neutrons, protons, and electrons, while more advanced descriptions have to account for the presence of hyperons (Schaffner-Bielich et al. 2002). Therefore, the grand canonical ensemble is the natural choice for the formulation of such an EoS. Another element of the realistic phenomenological hadronic EoS corresponds to the short-range repulsive interaction of the hard core nature between particles  (Andronic et al. 2006; Bugaev et al. 2016). Analysis of the particle yields produced in relativistic A+A collisions within statistical (thermal) models, i.e., the Hadron Resonance Gas (HRG) model (Bugaev et al. 2016), shows the importance of the particle hard core repulsion. In this approach every particle species is defined as a rigid sphere with a fixed radius estimated from experimental data analysis. These radii do not exceed 0.5 fm  (Andronic et al. 2017; Sagun et al. 2017a). Note that the hard core of hadrons in phenomenological EoSs is important in order to suppress thermal excitations of the hadronic spectrum and provide deconfinement of the color degrees of freedom expected at high temperatures/densities  (Satz 2012). Another requirement to the phenomenological EoS is related to its causal behavior when the speed of sound cannot exceed the speed of light. At sufficiently high densities this condition is violated by the hard core repulsion. As was shown by Sagun et al. (2017a), introducing the induced surface tension (IST) of particles to the model with the hard core repulsion between an arbitrary number of hadron species makes the EoS significantly softer and extends its causality range up to 7.5 normal nuclear densities, where formation of the quark-gluon plasma is expected. The IST is the key element of this approach (Sagun et al. 2014), as it allows us to account for the hard core repulsion between constituents in the most accurate way, and to properly reproduce the virial expansion of the multicomponent EoS. Recently, the IST EoS was used to describe the experimental data of hadron multiplicities measured at AGS, SPS, RHIC, and LHC energies of nuclear collisions (Sagun et al. 2017b), as well as the nuclear matter properties (Sagun et al. 2014). In this work the focus is on the application of IST EoS to the study of NS properties.


Here we report the X-ray spectral and timing analysis of the joint XMM–Newton and NuSTAR observations of an IRAS 04416+1215, a nearby (z = 0.0889; Boller et al. 1992) hyper-Eddington AGN. The source is part of a XMM–Newton/NuSTAR campaign that aims to constrain the broad-band X-ray properties of eight super-Eddington AGN from the best sample of bona fide super-Eddington sources available, i.e. super-Eddington accreting massive black holes (SEAMBHs; Du et al. 2014, 2015; Wang et al. 2014) that contains exclusively objects with black hole masses estimated from reverberation mapping. In this campaign we are carrying out to study the broad-band X-ray properties of super-Eddington AGN, all the sources have new NuSTAR observations performed simultaneously with XMM–Newton or Swift-X-ray Telescope (XRT). IRAS 04416+1215 has bolometric luminosity $\log (L_{\rm bol}/\rm erg\, s^{-1})=47.55$, according to Castelló-Mor, Netzer & Kaspi (2016), and $\log (L_{\rm bol}/\rm erg\, s^{-1})=45.52$, according to Liu et al. (2021). The former estimate is computed using, for the SED fitting procedure, the Slone & Netzer (2012) code, including the comparison of the observed SED with various combinations of disc SEDs covering the range of mass, accretion rate, spins, and taking into account the correction for intrinsic reddening and host galaxy contribution. In the latter estimate, the SED fitting is done using the more semplicistic templates from Krawczyk et al. (2013). The dimensionless accretion rate (Du et al. 2014) and black hole mass of the source are $\log (\dot{\mathscr {M}})$ = $2.63^{+0.16}_{-0.67}$ and log (MBH/M⊙) = $6.78^{+0.31}_{-0.06}$ with the reverberation mapping technique (Du et al. 2015), respectively, where $\dot{\mathscr {M}}\equiv \dot{M}_{\bullet }c^2/L_{\rm Edd}$, $\dot{M}_{\bullet }$ is mass accretion rates, c is speed of light, and LEdd is the Eddington luminosity. The dimensionless accretion rate is estimated by $\dot{\mathscr {M}}=20.1\, \ell _{44}^{3/2}M_7^{-2}$ from the Shakura–Sunyaev disc model (Du et al. 2015), where ℓ44 is the 5100 Å luminosity in units of $10^{44}\, {\rm erg\, s^{-1}}$ and $M_7=M_{\bullet }/10^7\, \mathrm{M}_{\odot }$. This approximation is valid for $\dot{\mathscr {M}}\lesssim 10^3$. To compute the Eddington ratio we assumed the bolometric luminosity value from Castelló-Mor et al. (2016), which is a better and more trustable estimate of the bolometric luminosity of the source, obtaining λEdd ∼ 472. This value is in perfect agreement with the dimensionless accretion rate from Du et al. (2014). However even assuming the luminosity from Liu et al. (2021), with which the value of the accretion rate would be λEdd ∼ 4.40, the source would remain a super-Eddington accreting AGN. IRAS 04416+1215 turned out to be the most peculiar of our sample, it is classified as NLS1 galaxy, showing narrow Hβ line [full width at half-maximum (FWHM) = $1670 \, \rm km \, \rm s^{-1}$; Moran, Halpern & Helfand 1996] and very broad [O iii] (FWHM = $1150 \, \rm km \, \rm s^{-1}$; Véron-Cetty, Véron & Gonçalves 2001) lines, which is typically found in sources accreting at such high Eddington accretion rates (Greene & Ho 2005; Ho 2009). The source shows a photon index in the Roentgen Satellite (ROSAT) (0.1–2.4 keV) energy band, of Γ = 2.96 ± 0.50 (Boller et al. 1992) and of $\Gamma =2.46^{+0.27}_{-0.26}$ for the rest-frame >2 keV spectrum, according to Liu et al. (2021).


In Fig. B.1 we plot the soft X-ray light curves for our candidate highly variable AGN using available X-ray data taken by the satellite missions Einstein, ROSAT, XMM, Suzaku and Swift. The count rates were obtained from different archives including HEASARC, the XMM Science Archive, the Swift UKSSDC and from our own Swift XRT data analysis, and for upper limits the 1SXPS catalogue (Evans et al. 2014) and the XMM upper limit server2 were queried. The count rates of the different satellites were converted to fluxes between 0.2–2.0 keV using PIMMS3 assuming a power law with a photon index of 1.7 as a spectral shape taking into account Galactic extinction as given by Willingale et al. (2013). Sobolewska & Papadakis (2009) found a positive correlation between flux and spectral slope for a sample of bright RXTE AGN in the 2–10 keV band. This could affect the relative fluxes seen in our sample which are plotted in Fig. B.1. We have attempted to quantify this for the different detectors used in the creation of our light curves. The sample of Sobolewska & Papadakis (2009) showed spectral changes with observed power-law slope varying between 1.0 and 2.0 (see their Fig. 7). For a typical Galactic absorption of 3 × 1020 cm-2 the change from slope of 1.0 to 2.0 would alter our estimated fluxes by −14% (ROSAT), −13% (XMM-Newton), +7% (Swift-XRT), +76% (Suzaku), +25% (Einstein-IPC). The change is large for Suzaku observations since in this case we use the count rate between 2.0–10.0 keV and extrapolate it to the soft band. All other satellites are sensitive in the soft band and hence the fluxes are less dependent upon the assumed spectral index. Six sources within our sample (XMMSL1 J024916.6-041244, J034555.1-355959, J045740.0-503053, J051935.5-323928, J070841.3-493305, and J193439.3+490922) display variation in flux of a factor of ten or greater between at least one pair of XMM and Swift observations, on timescales of months to years. The ratio between the soft X-ray flux observed with Swift and that observed with XMM for the remaining sources is typically a factor of a few. We observed the two TDE candidates with XRT, and found that both had faded significantly, following expectations from previous and later fluxes and upper limits (Figs. 1p and h).


Pair-instability supernovae (PI SNe) are the most energetic thermonuclear explosions known and can be detected near the edge of the observable universe. They have now been studied by several groups for their potential to probe the properties of the first stars and galaxies (Greif et al. 2008, 2010, 2011, 2012; Johnson et al. 2009, 2014; Turk et al. 2009; Stacy et al. 2010, 2012; Clark et al. 2011; Hosokawa et al. 2011; Smith et al. 2011; Glover 2013; Jeon et al. 2012; Susa 2013; Pawlik et al. 2011, 2013; Wise et al. 2012; Whalen 2013; Hirano et al. 2014). They can also shed light on the origins of supermassive black holes and early cosmological reionization and chemical enrichment (Mackey et al. 2003; Whalen et al. 2004, 2008a, 2010; Abel et al. 2007; Smith & Sigurdsson 2007; Wise & Abel 2008; Alvarez et al. 2009; Tanaka & Haiman 2009; Smith et al. 2009; Park & Ricotti 2011, 2012, 2013; Volonteri 2012; Agarwal et al. 2012; Ritter et al. 2012; Whalen & Fryer 2012; Choi et al. 2013; Latif et al. 2013a, 2013b; Reisswig et al. 2013; Schleicher et al. 2013; Johnson et al. 2014; Chiaki et al. 2013; Safranek-Shrader et al. 2014). For example, detections of both PI and core-collapse (CC) SNe at high redshift could be roughly binned by mass, thereby building up a simple Pop III IMF over time as enough events are discovered (de Souza et al. 2013, 2014). If the Pop III initial mass function (IMF) proves to be top heavy, this could account for the origins of SMBHs because enough 200–300

  seed black holes might be formed at

 20 for a few to reach 109

  by

 7. If not, alternatives for SMBH seeds must be found, such as BHs forming by direct collapse in atomically cooled halos at slightly later epochs, z ∼ 10–15 (Johnson et al. 2012, 2013b). PI SN candidates such as SN 2007bi (Gal-Yam et al. 2009; Kozyreva et al. 2014) and SN 2213–1745 (Cooke et al. 2012) have now been discovered at z = 0.126 and 2.05, respectively.


In the standard AGN unification theory, the anisotropic radio emission produced by relativistic jets points to a description where properties like radio morphology and radio spectral index depend on orientation. Two widely used orientation indicators in the radio regime are the (a) core-to-lobe flux density ratio (R), which is the ratio of the flux densities of the core and lobe in radio wavelengths, and (b) radio-to-optical ratio of the quasar core (RI), which is the ratio of flux densities of the core in radio and optical wavelengths (See equation 5 and 6 for the exact definition of R and RI parameters). A high R-value suggests a small viewing angle to the jet axis, while a low R-value indicates more of an edge-on view to the quasar (Kapahi & Saikia 1982; Orr & Browne 1982; Morisawa & Takahara 1987; Morganti et al. 1997). Studies have shown that R correlates with core radio luminosity (Hardcastle & Worrall 2000) as well with the core optical luminosity (Kharb & Shastri 2004). Factors like age and environment are believed to introduce the scatter in the correlation involving the R parameter suggesting that R might not be the best measure for orientation. The R parameter depends on both the core’s radio flux as well as the radio flux of the lobes arising from the quasar jets. These jets are suggested to be affected by environmental factors. Wills & Brotherton (1995) and Kharb, Lister & Cooper (2010) showed that the radio-to-optical ratio (RI) of the quasar core might be a better measure of studying the orientation. The RI parameter measures the core-boosting factor by normalizing the core’s radio flux by the core’s optical flux. Although with some scatter in the correlation, Kimball et al. (2011) demonstrated a strong correlation between R and RI parameter suggesting that the two parameters are good indicators of quasar orientation. The scatter in the correlation supports the idea that other factors like age, size, environment, and luminosity also influence these two measurements.


The linear relationship that we observe between breakpoint energy and core temperature is in line with previous measurements (e.g. McComas et al. 1992; Štverák et al. 2009), for both the halo and strahl. According to Scudder & Olbert (1979), a linear trend in the halo relation also follows under the assumption that binary Coulomb collisions dominate electron dynamics in the solar wind. However, in order to align with available experimental data, Scudder & Olbert (1979) set a scaling factor of Ebp/kBTc = 7, which differs from our scaling factor of Ebp/kBTc = 5.5 ± 0.1. With a scaling factor of Ebp/kBTc = 7, Scudder & Olbert (1979) predict that a transformation of thermal electrons into the suprathermal population occurs as the solar wind flows out from the Sun. Findings by Štverák et al. (2009), on the other hand, show that the (nh + ns)/nc ratio remains roughly constant with heliocentric distance in the slow wind, suggesting a lack of interchange between the thermal and suprathermal populations. However Štverák et al. (2009) observes some variability in the (nh + ns)/nc ratio in the fast wind, which they attribute to either statistical effects due to a lack of samples or a possible “interplay” between thermal and suprathermal electrons. Scudder & Olbert (1979) also predict that the halo Ebp/kBTc ratio remains constant with heliocentric distance, whereas Štverák et al. (2009) find that the halo Ebp/kBTc ratio decreases with heliocentric distance. These findings by Štverák et al. (2009), along with the discrepancy between our calculated ratio of Ebp/kBTc = 5.5 ± 0.1 and the prediction of Ebp/kBTc = 7, suggest that the model of Scudder & Olbert (1979) requires a minor update to either the theory or to the input parameters. The discrepancy, however, may also be indicative of other processes, such as wave-particle scattering (e.g. Gary et al. 1994), that possibly modifies the ratio between breakpoint energy and core temperature while preserving its linear relationship.


Armstrong et al. (2020) recently announced the discovery of TOI-849b, a planet having a size comparable to the one of Neptune (

R
pl

=
3
.

45

−
0.12

+
0.16

R
⊕

$ R_{\mathrm{pl}} = 3.45^{+0.16}_{-0.12}\,R_{\oplus} $

), but an anomalously larger mass (

M
pl

=
40
.

8

−
2.5

+
2.4

M
⊕

$ M_{\mathrm{pl}} = 40.8^{+2.4}_{-2.5}\,M_{\oplus} $

) and a density similar to the one of the Earth (ρ = 5.5 ± 0.8 g cm−3). This is the densest Neptune-sized planet discovered to date. Interior structure models suggest that for this kind of planet, any H/He envelope would consist of no more than 3.9% of the total mass (Armstrong et al. 2020). TOI-849b orbits around a late G-type star, with an orbital period of P = 18.4 h. Its equilibrium temperature is Teq = 1800 K (for a Bond Albedo 0.3). With such properties, TOI-849b represents one of the few planets populating the hot Neptune desert, a region on the radius-orbital distance plane characterised by a surprising deficit of Neptune-sized planets on very short orbits (e.g. Lecavelier Des Etangs 2007; Beaugé & Nesvorný 2013; Mazeh et al. 2016). Growing evidence suggests that the evaporation of hot Neptunes due to stellar irradiation represents a major process in shaping the desert (Owen 2019). Some of the planets within the desert or at its lower-radius border could thus be the remnant cores of larger gas-rich progenitors that lost most of their atmosphere (e.g. Lecavelier des Etangs et al. 2004). Alternatively, orbital migration may have also played a part in sculpting the desert, with different classes of planets forming differently (e.g. Batygin et al. 2016), or following different dynamical tracks (e.g. Matsakos & Königl 2016). Even so, the origin of this key feature in the demographics of close-in planets remains debated (Zahnle & Catling 2017; Owen & Lai 2018), and investigating the past history of planets such as TOI-849b contributes to the global effort towards disentangling the important mechanisms at the root of the desert. Since the mass of the planet is larger than the threshold value for runaway gas accretion (roughly 10 − 20 M⊕, Mizuno et al. 1978; Rafikov 2006; Piso et al. 2015), TOI-849b might have been a gas giant before undergoing extreme mass loss via thermal self-disruption, collisions with other giant planets, or it was prevented from accreting gas because of gap openings in the protoplanetary disc, or because of late formation (Armstrong et al. 2020). In their work, Armstrong et al. (2020) found that their estimated photoevaporation rates cannot provide the mass-loss rates needed to remove a roughly 280 M⊕ envelope from a Jupiter-like gas giant. Nevertheless, in their estimations the planetary atmosphere is not self-consistently monitored, and neither is the luminosity emitted by the host star relative to its rotational history. This is critical, considering how recent works emphasise the interplay between all these elements in driving the evolution of close-in worlds (e.g. Owen & Alvarez 2016; Kubyshkina et al. 2018; King & Wheatley 2021; Pezzotti et al. 2021).


Carbon-enhanced metal-poor (CEMP) stars form an important class of metal-poor giants, sub-giants, and dwarfs, with a large fraction of them showing enhanced abundances of heavy elements (see Beers & Christlieb 2005; Frebel 2018 for a general review). Among the different types of CEMP stars, the CEMP-s stars are enriched with products of s-process nucleosynthesis, the CEMP-r stars are enriched with the products of r-process nucleosynthesis, and CEMP-r/s stars are enriched with products of i-process nucleosynthesis. Understanding the diverse abundance patterns exhibited by different groups of CEMP stars that are believed to be associated with different formation mechanisms has been a challenge. In Goswami et al. (2021), we present a detailed analysis and discussion on the classification criteria of CEMP stars, as well as the formation scenarios of CEMP stars put forward by different authors (Cowan & Rose 1977; Hill et al. 2000; Qian & Wasserburg 2003; Cohen et al. 2003; Jonsell et al. 2006; Campbell & Lattanzio 2008; Campbell et al. 2010; Stancliffe et al. 2011; Herwig et al. 2011; Doherty et al. 2015; Abate et al. 2016; Jones et al. 2016; Banerjee et al. 2018; Clarkson et al. 2018; Denissenkov et al. 2017; Côté et al. 2018). In this paper, we report an extremely metal-poor carbon-enhanced star, HE 1005–1439, whose surface chemical composition is found to be enriched with both s-process and i-process nucleosynthesis that forms a new class of object with a distinct abundance pattern. The peculiar abundance pattern, observed for the first time in a CEMP star, was investigated based on a parametric-model-based analysis that revealed almost equal contributions from both the s-process and the i-process to its surface chemical composition. We examined various production mechanisms and formation scenarios for this object. A formation scenario involving effective proton ingestion episodes (PIEs) triggering i-process nucleosynthesis followed by s-process asymptotic giant branch (AGB) nucleosynthesis with limited third-dredge-up (TDU) episodes seems to be most promising for this type of object.


Among the GRB empirical correlations, the Amati correlation is a very popular one, which connects the spectral peak energy in the GRB cosmological rest frame and the isotropic equivalent radiated energy (E

p
 − E
iso; Amati et al. 2002; Amati 2006a, 2006b; Amati et al. 2008, 2009; Amati & Della Valle 2013). Recently, we proposed an improved Amati correlation (Liu et al. 2022) by using the Gaussian copula, which is a powerful statistical tool capable of describing the dependence structures between multivariate random variables, and has been applied to various fields by the astronomical community (Benabed et al. 2009; Jiang et al. 2009; Koen 2009; Scherrer et al. 2010; Takeuchi 2010; Yuan et al. 2018; Qin et al. 2020; Takeuchi & Kono 2020). In Liu et al. (2022), by choosing the spatially flat ΛCDM model with Ωm0 = 0.30 and H
0 = 70 km s−1 Mpc−1 as the fiducial model, we utilize the low-redshift (z  1.4) GRB data to calibrate the standard and improved Amati correlations, and then extrapolate the results to the high-redshift GRB data to achieve the GRB Hubble diagram, where Ωm0 is the present dimensionless matter density parameter. Using these calibrated GRBs to constrain the flat ΛCDM model, we found that the improved Amati correlation can give results well consistent with the fiducial model, while the standard one cannot. Thus, in Liu et al. (2022), the reliability of the improved Amati correlation was ascertained with a fiducial model, but its cosmological application was not carried out. In this work, we will fill this gap. In order to obtain the Hubble diagram of the latest A220 GRB samples (Khadka et al. 2021) model-independently, we use the Pantheon SN Ia data (Scolnic et al. 2018) to calibrate the standard and improved Amati correlations, and then use these calibrated GRB data to constrain the ΛCDM and wCDM models. Besides the GRB data, the H(z) data set is also added in our analysis to obtain a tight constraint on model parameters. The rest of the paper is organized as follows: Section 2 introduces the improved Amati correlation briefly, and standardizes the GRB samples by using the method of the low-redshift calibration. Section 3 studies the constraints on the ΛCDM and wCDM models from the GRB data and the GRB + H(z) data. Our conclusions are summarized in Section 4.


Resonant dynamical friction may have applications beyond the relaxation of IMBHs examined in this paper. It may affect all objects in stellar clusters much more massive than the individual constituents of the disk, if present, including massive stars, stellar mass black holes (BHs), or the center of mass of massive binaries. Furthermore, it is also expected to operate in any type of disk with a high number of particles, including active galactic nucleus (AGN) accretion disks. Previously, it has been argued that stars and BHs crossing the disk on low-inclination orbits get captured by Chandrasekhar dynamical friction into the disk (Bartos et al. 2017; Panamarev et al. 2018; Tagawa et al. 2020). An interesting implication is that, if BHs settle into the disk, they interact dynamically and form BH–BH binaries efficiently, and frequent dynamical interactions and gas effects drive the BHs to merger, producing gravitational waves (GWs) detectable by LIGO, VIRGO, and KAGRA (McKernan et al. 2014, 2018; Bartos et al. 2017; Leigh et al. 2018; Yang et al. 2019; Tagawa et al. 2020, 2021; Samsing et al. 2020). Mergers are also facilitated by Lidov–Kozai oscillations in anisotropic systems (Heisler & Tremaine 1986; Petrovich & Antonini 2017; Hamilton & Rafikov 2019). The results in this paper show that resonant dynamical friction may accelerate the capture of objects in the accretion disks by a factor proportional to the SMBH mass over the local disk mass for large orbital inclinations. Pressure and viscosity in a gaseous disk do not inhibit the orbit-averaged torque from the IMBH, which leads to realignment and the warping of the disk (Bregman & Alexander 2012). Thus, RDF may efficiently catalyze the alignment of the orbital planes of BHs even in low-luminosity AGN or Seyfert galaxies with relatively small disk masses, which may not be possible for Chandrasekhar dynamical friction. In fact, this mechanism extends the scope of the “AGN merger channel” for GW source populations even beyond low-luminosity AGN and Seyfert galaxies, as it may organize BHs into disks also in nonactive galaxies with nuclear stellar disks.


The applicability of mid-IR and far-IR FSL as diagnostic tools of the ISM has received a boost thanks to the publication of samples of nearby galaxies observed with the infrared space observatory (ISO) and Herschel (e.g. Brauher et al. 2008; Farrah et al. 2013; Sargsyan et al. 2014; Kamenetzky et al. 2014; Cormier et al. 2015; Cigan et al. 2016; Herrera-Camus et al. 2016; Fernández-Ontiveros et al. 2016; Zhao et al. 2016; Díaz-Santos et al. 2017; Zhang et al. 2018). At high redshift (z ≳ 1), the far-IR FSL conveniently shift into the (sub)millimetre atmospheric windows. The most popular line is clearly [CII] 158 μm, followed by the [CI] 370, 609 μm lines (e.g. Walter et al. 2011; Bothwell et al. 2017). Both single dish submillimetre telescopes and interferometers have also detected the [NII] 122 μm and 205 μm lines at high redshift (Ferkinhoff et al. 2011, 2015; Combes et al. 2012; Nagao et al. 2012; Decarli et al. 2012, 2014; Béthermin et al. 2016; Pavesi et al. 2016; Lu et al. 2017; Tadaki et al. 2019; Novak et al. 2019). After a slow start, the [OIII] 88 μm line is quickly becoming a popular line to confirm redshifts of galaxies in the epoch of reionization (z ≳ 6), where it shifts into the submillimetre atmospheric windows below 500 GHz (Ferkinhoff et al. 2010; Inoue et al. 2016; Carniani et al. 2017; Marrone et al. 2018; Vishwas et al. 2018; Hashimoto et al. 2018, 2019; Walter et al. 2018; Tamura et al. 2019; Tadaki et al. 2019; Novak et al. 2019). Deep Herschel/SPIRE spectroscopy has also revealed a number of FSL detections, either in individual objects (Valtchanov et al. 2011; Coppin et al. 2012; George et al. 2013; Uzgil et al. 2016; Rigopoulou et al. 2018; Zhang et al. 2018), or in stacked spectra (Wardlow et al. 2017; Wilson et al. 2017; Zhang et al. 2018). These include the only detections of the [OI] 63 μm line at high redshift reported thus far. This [OI]3P2−3P1 line is arguably the best tracer for the star-forming gas, as it traces the very dense (

n
crit
H

$ n^{\mathrm{H}}_{\mathrm{crit}} $

 = 5×105 cm−3) neutral gas (one caveat being the frequeny presence of self-absorption observed in local ultra-luminous IR galaxies, Rosenberg et al. 2015). Like the [CII] line, the [OI] 63 μm line also shows a “deficit” in the most luminous far-IR sources, though with a higher scatter (Graciá-Carpio et al. 2011; Cormier et al. 2015; Díaz-Santos et al. 2017). Surprisingly, this bright FSL has not been frequently observed with ALMA, probably because it is only observable in the highest frequency bands. At least as surprising is that the fainter, but more accessible [OI]3P1−3P0 transition at λrest = 145 μm (

n
crit
H

$ n^{H}_{\mathrm{crit}} $

 = 9.5×104 cm−3) has only recently been detected at high redshifts (Novak et al. 2019 report a tentative detection in a z = 7.5 quasar). Also in nearby galaxies, this [OI] 145 μm line has not been observed very frequently as in most cases, it is fainter than the nearby [CII] 158 μm line. After initial detections with ISO (Malhotra et al. 2001; Brauher et al. 2008), Herschel has now detected [OI] 145 μm in significant samples of nearby galaxies (Spinoglio et al. 2015; Cormier et al. 2015; Fernández-Ontiveros et al. 2016; Herrera-Camus et al. 2018), and recently in a z = 6.5 lensed quasar (Yang et al. 2019).


If the variation in 10Be/9Be ratios of CAIs reflects those episodic accretion events, 10Be/9Be ratios of CH–CB CAIs observed in this study would give important constraints on the evolution of the solar protoplanetary disk. Astronomical observations suggest that FUori-type outbursts are confined to the first few hundreds of thousands of years, which correspond to the class I stage of the protoplanetary disk evolution (e.g., Schulz 2012). We propose that the high and variable 10Be/9Be ratios recorded in CH–CB CAIs reflect episodic cosmic-ray fluxes caused by FUori-type outbursts. On the other hand, relatively low and less variable 10Be/9Be ratios recorded in CV CAIs may reflect less intensive episodic accretion events, possibly the EXori-type outbursts, which are confined to the evolutional stage of a few million years after the formation of the protoplanetary disk (=class II). Note that CH–CB CAIs studied here show no (or very low) signs of 26Al-derived 26Mg excesses, while most CV CAIs show clear evidence for the past presence of 26Al. If 26Al was introduced into the solar system at the earliest stage of the disk evolution (e.g., Sahijpal & Goswami 1998), differences in Be–B and Al–Mg systematics between CH–CB and CV CAIs imply that the injection of 26Al have occurred between the evolutionary stages class I and class II of the solar protoplanetary disk. This scenario is in agreement with arguments by other authors that the 26Al-free CAIs formed prior to injection and homogenization of 26Al in the early solar system (Sahijpal & Goswami 1998; Sahijpal et al. 2000; Krot et al. 2008a see more discussion in Krot et al. 2012a). Importantly, as mentioned in the introduction, CH–CB chondrites may have accreted a significant amount of outer solar system materials (Murty et al. 2007; Ivanova et al. 2008; Briani et al. 2009; Bonal et al. 2010; Olsen et al. 2016; Van Kooten et al. 2016), suggesting that CH–CB chondrites formed at outer parts of the solar protoplanetary disk relative to CV chondrites. In this case, our new Be–B and Al–Mg data set implies that the earliest formed CAIs tend to be transported into the outer part of the solar protoplanetary disk, where the parent bodies of CH–CB chondrites likely accreted. Yang & Ciesla (2012) modeled the evolution of the protoplanetary disk and material transport in the protoplanetary disk. Interestingly, Yang & Ciesla (2012) showed that outward radial transport in class I would have been greater than that of later stages of YSO evolution, suggesting that the earliest formed CAIs could be preserved in primitive bodies that accreted in the outer part of the disk. This model is consistent with our interpretation for the Be–B and Al–Mg systematics on CH–CB CAIs. It should be noted, however, that it is possible that 26Al were heterogeneously distributed in the CAI-forming regions at the earliest stage of the solar system evolution (e.g., Krot et al. 2008a; Holst et al. 2013; Park et al. 2017 and reference therein). Because no Pb–Pb ages of CH–CB CAIs are available at present, we cannot discard that possibility. Very recently, Kööp et al. (2018) found helium and neon excesses in the 26Al-free hibonite-rich CAIs, which can be attributed to in situ irradiation by energetic particles. Because 26Al-rich CAIs in CV chondrites lack comparable noble gas irradiation records (Vogel et al. 2004), Kööp et al. (2018) concluded that 26Al-free hibonite-rich CAIs experienced intense energetic particle irradiation at the earliest stage of solar protoplanetary disk evolution. This conclusion seems to be consistent with our above scenario for 26Al-free CH–CB CAIs. Note, however, that 10Be/9Be ratios of 26Al-free hibonite-rich CAIs in CM chondrites tend to be in the range of those for 26Al-rich CV CAIs (Liu et al. 2009, 2010), which is inconsistent with the above scenario. Therefore, the relationship between 10Be and 26Al in the early solar system would be more complicated than we thought.


First detected as a radio source during the Vermilion River Observatory Sky Survey (Dickel et al. 1967) and the Ohio Sky Survey (Kraus 1977), OJ 287 has been studied extensively in the radio regime. Its relativistic jet is pointing at us with an average viewing angle of ∼2° (Jorstad et al. 2005; Agudo et al. 2012) and shows remarkable short-time-scale variability interpreted as a turbulent injection process and/or a clumpy accretion disc structure (Agudo et al. 2012) or as a binary-induced wobble (Valtonen & Pihajoki 2013; Dey et al. 2021). The inner jet is the source of highly variable γ-rays (e.g. Agudo et al. 2011; Hodgson et al. 2017) and displays strong and variable radio polarization (e.g. Aller et al. 2014; Cohen et al. 2018; Myserlis et al. 2018). Though only faintly detected in the very-high-energy (VHE) band (>100 GeV; Mukherjee et al. 2017; O’Brien 2017), OJ 287 is a well-known X-ray emitter and has been detected with most major X-ray observatories, including Einstein (Madejski & Schwartz 1988), EXOSAT (Sambruna et al. 1994), ROSAT (Comastri, Molendi & Ghisellini 1995), BeppoSAX (Massaro et al. 2003), ASCA (Idesawa et al. 1997), the Neil Gehrels Swift observatory (Swift hereafter; Massaro et al. 2008), Suzaku (Seta et al. 2009), XMM–Newton (Ciprini et al. 2007), and most recently with NuSTAR (Komossa et al. 2020a). These observations established OJ 287 as a bright and variable X-ray source, and allowed single-component X-ray spectral fits. Its (0.5–10) keV X-ray spectrum was interpreted as a mix of synchrotron and inverse-Compton (IC) emission, with the former more variable (Urry et al. 1996). OJ 287 is found most of the time in a rather flat spectral state with a photon index Γx ≈ 1.5–1.9, though Γx is as steep as 2.6 during one ROSAT observation (Urry et al. 1996) in the band (0.1–2.4) keV. Its steepest state, with an equivalent Γx = 2.8, (better fit with a logarithmic parabolic model) was detected with XMM–Newton in the band (0.3–10) keV, which caught the source right at the peak of one of the brightest X-ray outbursts measured so far (Komossa et al. 2020a). Imaging with the Chandra X-ray observatory has revealed a long, curved X-ray jet consisting of multiple knots and extending out to 20 arcsec or a de-projected scale of >1 Mpc, and bright central emission (Marscher & Jorstad 2011).


The bolometric luminosity of 3C 84 is about 0.4% of the Eddington luminosity. Thus, the accretion flow of 3C 84 is likely to be a radiatively inefficient accretion flow (RIAF: Narayan & Yi 1995) rather than a standard disk (Shakura & Sunyaev 1973). However, we note that 3C 84 has a cold (

T

e

∼

10

4

 K) disk-like accretion flow, as identified by FFA of the emission from the counter jet in the parsec scale (Walker et al. 2000) and inhomogeneous gas distribution around the black hole (Fujita et al. 2016). A number of theoretical studies predicted that the accretion flow components of hot geometrically thick (RIAF-like) and cold geometrically thin can coexist in either horizontal or vertical stratification (e.g., Miller & Stone 2000; Merloni & Fabian 2002; Liu et al. 2007; Ho 2008; Liu & Taam 2013). The measured Faraday rotation can be caused by such an RIAF-like component. We thus estimate the accretion rate of the RIAF-like component using the measured RM. For a simplicity, we assume that the RIAF-like component is a quasi-spherical Bondi accretion flow with a power-law density profile. We can calculate the accretion rate, following the formulation as follows (Quataert & Gruzinov 2000; Marrone et al. 2006; Kuo et al. 2014).

M

˙

=

1.3
×

10

−
10

[

1
−

(

r

out

/

r

in

)

−
(
3
β
−
1
)

/

2

]

−
2

/

3

×

M

BH

8.0
×

10

8

M

⊙

4

/

3

2

3
β
−
1

−
2

/

3

r

in

7

/

6

RM

rad

m

−
2

2

/

3

.

For an inner effective radius

r

in

 of 1 pc (

1.3
×

10

4

R

s

), where the hotspot is located, the observed RM implies an accretion rate of

∼
4.3
×

10

−
2

M

⊙

 yr−1 and

∼
8.6
×

10

−
2

M

⊙

 yr−1 for

β
=
0.5

 and

β
=
1.5

, which are corresponding to convection-dominated accretion flow (CDAF: Narayan et al. 2000; Quataert & Gruzinov 2000) and advection-dominated accretion flow (ADAF: Ichimaru 1977; Narayan & Yi 1995), respectively. Here we assumed the outer effective radius

r

out

 of

10

5

R

s

 (∼8 pc), which is approximately the same with the Bondi radius of 8.6 pc (Fujita et al. 2016). The derived accretion rate is roughly consistent with that estimated from the bolometric luminosity with a black hole mass of

8
×

10

8

M

⊙

 and a radiative efficiency of 10% (

M

˙

∼

L

bol

/

(
0.1

c

2

)
≃
7.1

×

10

−
2

M

⊙

 yr−1).


Two-point statistics are central to many of the leading cosmological analyses of the large-scale structure searching for deviations from ΛCDM (Simpson et al. 2013; Song et al. 2015; Amon et al. 2018; Abbott et al. 2019a; Chudaykin, Dolgikh & Ivanov 2021; Lee et al. 2022; Muir et al. 2021; Tröster et al. 2021; Vazsonyi et al. 2021), and a great deal of effort has gone into accurately modelling the non-linear matter power spectrum in modified gravity and dark energy cosmologies – a theoretical ingredient essential to extract the cosmological information locked in small scales (e.g. Koyama, Taruya & Hiramatsu 2009; Brax & Valageas 2012; Takahashi et al. 2012; Heitmann et al. 2014; Zhao 2014; Casarini et al. 2016; Mead et al. 2016; Cusin, Lewandowski & Vernizzi 2018; Cataneo et al. 2019; Winther et al. 2019; Euclid Collaboration 2021; Ramachandra et al. 2021). However, non-linear gravitational clustering converts the nearly Gaussian initial density field (Planck Collaboration VI 2020) to a late-time density field with significant non-Gaussian features that these standard analyses are unable to access (Bernardeau et al. 2002). Non-Gaussian statistics, such as the bispectrum (Brax & Valageas 2012; Munshi 2017; Yamauchi, Yokoyama & Tashiro 2017; Crisostomi, Lewandowski & Vernizzi 2020; Bose et al. 2020b), higher order weak lensing spectra (Munshi & McEwen 2020), the halo mass function (Lam & Li 2012; Cataneo et al. 2016; Hagstotz et al. 2019; McClintock et al. 2019; Bocquet et al. 2020), the void size function (Perico et al. 2019; Verza et al. 2019; Contarini et al. 2021) and Minkowski functionals (Kratochvil et al. 2012; Fang, Li & Zhao 2017), respond strongly to modified gravity and dark energy through the induced changes in the higher moments of the cosmic density field, and their remarkable complementarity to traditional two-point functions leads to tighter joint constraints on the extra non-standard parameters (Shirasaki et al. 2017; Peel et al. 2018; Sahlén 2019; Liu et al. 2021).


There have been many efforts in the past that have used the optical emission lines to understand the AGN and the host galaxy properties (e.g. Baldwin, Wampler & Burbidge 1981; Tadhunter et al. 1998). While broad and high ionization lines tend to arise from the immediate vicinity of the central massive black hole, low ionization lines arise further out (e.g. Richardson et al. 2014), tracing the host-galaxy dynamics, chemical composition and gas distribution. It is known that low ionization gases trace the H i distribution in the galactic discs (e.g. Lin & Zou 2001). Indeed, Mg ii is known as a good tracer of high column density H i gas in damped Lyman α absorbers (e.g. Bergeron & Stasińska 1986; Péroux et al. 2004). [O ii] λ3727 is widely used to trace star formation in galaxy surveys, particularly for redshifts $z$ ≳ 0.4 (e.g. Lilly et al. 1996; Hippelein et al. 2003). In optically selected quasars, [O ii] emission in excess of the base level that is expected from non-stellar photoionization indicates star formation accompanying quasar activity (e.g. Ho 2005). These lines will hence allow a better understanding of the distribution and kinematical properties of the atomic gas. Recently, Shen & Ménard (2012) and Khare et al. (2014) studied the properties of [O ii] λ3727 emission from a large sample of associated Mg ii absorbers, in order to understand the origin of associated Mg ii absorption systems relative to the AGN. Khare et al. (2014) find a high excess [O ii] emission in the composite quasar spectrum (constructed in the quasar rest frame) for the sample of quasars that show outflows, i.e. the systems for which the Mg ii absorption is blueshifted relative to the AGN. Further, the authors find that the presence of an associated Mg ii absorption enhances [O ii] emission in the composite quasar spectrum. However the excess emission flux in the [O ii] line does not depend on the strength of the Mg ii absorption line (Khare et al. 2014). The excess [O ii] line flux could be originating either from the host galaxy, or the parent AGN; in the former case, the excess [O ii] emission could indicate higher star formation in the host galaxy. The association of high [O ii] line emission with the occurrence of outflows suggests either strong stellar outbursts in the host galaxy, or strong AGN jet–gas interactions.


We stress that the age values discussed here should be taken with care as they are based on simulations which rely on a number of assumptions. One of these is the magnetic field which is typically computed using the often unrealistic equipartition conditions. For example, a factor two difference in the magnetic field value translates directly into about a factor two age difference. Another important parameter is the assumed total source age distribution which is assumed to be different in the various simulations. For example, in the case of the modelling presented in Brienza et al. (2017) and Godfrey et al. (2017), the distribution of ton was taken to be a truncated log-normal distribution between 20 and 200 Myr, with a median of 30 Myr. Hardcastle (2018) on the other hand explored two types of models, with ages being either uniformly distributed between 0 and 1000 Myr, or linearly distributed in log space between 1 and 1000 Myr. These latter authors found that the uniform-age models efficiently explain the size and luminosity statistics of bright sources observed by LOFAR, while the log-uniform models were a better representation of the faint population. Similarly, Shabala et al. (2020) showed that power-law age-distribution models (i.e. models where the radio-source population is dominated by short-lived sources) can explain the observed properties of both active and remnant and/or restarted LOFAR sources. These findings are consistent with the modelling assumptions of Brienza et al. (2017) and Godfrey et al. (2017), as the majority of short-lived sources will simply never grow large or luminous enough to make it into the observable LOFAR sample1. Modelling by both Shabala et al. (2020) and Hardcastle (2018) confirms the picture in which the remnant lobes fade quickly below the LOFAR detection limit. This fading is more rapid for older sources. Future studies, including modelling of the spectral index and follow-up observations, will help us to put tighter constraints on the age of the candidate restarted radio galaxies and the time that passed between the two bursts of activity.

Differences Compare/Contrast Similarities Similarities