The growth of supermassive BHs is also altered when considering non-Gaussianities. After deriving the merger history of the most massive haloes at z = 6.5 in both the Gaussian and non-Gaussian simulations, we study the evolution of BHs in massive haloes down to z = 6.5. To investigate the cumulative effect over cosmic times on the BHs assembly, we model the growth of BHs in three different ways. Different probabilities for a halo of hosting a seed BH, and different accretion models (either each BH accretes at the Eddington limit for a dynamical time after a major merger or using an accretion rate based on a distribution probability derived from a large-scale hydrodynamical simulation) are adopted. We have not included in our models the effects of ‘kicks’ caused by asymmetric emission of gravitational waves, which have been proposed to be possibly responsible for ejecting BHs from haloes with shallow potential wells, thus halting or reducing the growth of high-redshift BHs hosted in small haloes (e.g. Yoo & Miralda-Escudé 2004; Volonteri & Rees 2006; Tanaka & Haiman 2009). This effect, however, seems to affect less than 10 per cent of binaries and it becomes negligible for BH mergers at z  10 (Volonteri & Rees 2006). We find that non-Gaussianities imply a larger number of massive BHs and also an increase in the mean BH mass (up to 0.36 in the most favourable experiment). A population of supermassive BHs will then grow faster and to higher masses in a universe with scale-dependent non-Gaussian primordial density fluctuations. If the seed masses are similar to those of Pop III star remnants, BHs will not be able to grow above few × 105 M⊙ by z = 6. However, our simulations do not resolve mini-haloes, and we may underestimate the growth of seeds at earlier times. We argue that, in a simulation resolving mini-haloes, BHs would have formed earlier through the Pop III remnant scenario, leading to a longer time for them to grow in mass. If we assumed that Pop III remnant seeds with mass 100 M⊙ form at z ∼ 30 in haloes unresolved in our simulations, they would have grown, assuming, optimistically, constant growth at the Eddington rate (but see Johnson & Bromm 2007; Alvarez, Wise & Abel 2009; Milosavljević, Couch & Bromm 2009; Park & Ricotti 2011) to ∼103 M⊙ by z = 18, where we start our analysis. The final BH mass at z = 6 would then be ∼ one order of magnitude larger, a few × 106 M⊙, still short of the ∼109 M⊙ required. The very limited growth obtained for the Pop III remnant case suggests that large seeds or super-Eddington accretion (see Volonteri, Silk & Dubus 2015, and references therein) may be necessary for successful BH growth. We have done the same experiments on BH growth starting with initial 105 M⊙ BH masses (not shown in the paper, but see Section 4). In this case we found that it is much easier for BHs to grow to higher BH masses, but still only to several 108 M⊙. This is not unexpected, because our simulation box does not contain the very rare and biased DM haloes with masses ∼1013 M⊙ believed to be hosting these extreme BHs.


The results presented in this work answer a number of questions regarding the character of the GG Tau A system, while raising others and leaving others untouched. First among the questions raised by our results are the questions of whether the detailed morphology of features in the circumbinary torus as seen in our simulations are actually present in the GG Tau A system. The current best resolution observations of the torus provide tantalizing hints that such features do exist, but a definitive statement that such features are present must await higher resolution observations of quality similar to those described by the ALMA Partnership (2015), for the HL Tau circumstellar disk. Another important question concerns planet formation in multiple systems. Whether or not any theoretical mechanism presently exists to explain the formation of giant planets in binary systems, the fact remains that at least a few binary systems, such as γ Cephei (Neuhauser et al. 2007), GI86 (Queloz et al. 2000), and HD 41004 (Zucker et al. 2004), do harbor planets. Therefore, some formation mechanism does in fact exist. While we find that accretion into the circumstellar disks occurs rapidly enough so that they can survive for the comparatively long timescales needed to form planets, other conditions, such as temperatures, remain quite unfavorable. What are the mechanisms still missing from our models that permit such objects to form? Finally, our simulations model the evolution of the “A” component of the full GG Tau system over a time span extending over only a tiny fraction of its formation timescale and in only two spatial dimensions. We neglected the full dynamical effects expected to be present in the system, insofar as our results include neither the distant binary “B” component of the GG Tau system, nor the newly discovered tight binary nature of the GG Tau Ab component. Even so, we find large scale morphological changes even over this short time span, and restricted dimensionality and physical system. Given the vigor of the activity over such a short timescale, we would expect activity of similar scale to occur over longer time spans as well, with correspondingly large consequences on the system morphology. To what extent will 3D effects also play a role in the evolution? What will be the end state configuration of the GG Tau system as a whole? Will the components eventually break apart? Merge? Future investigations extending the work presented here will be required in order to answer these questions.


The Orion Bar is probably the best studied PDR in our Galaxy. It is located between the Orion Molecular Cloud 1 and the H ii region excited by the Trapezium cluster, and is exposed to an FUV field a few 104 times the mean interstellar radiation field. Owing to its proximity (417 pc, Menten et al. 2007) and nearly edge-on orientation, the Bar provides an ideal laboratory for testing PDR models (e.g., Jansen et al. 1995; Gorti & Hollenbach 2002; Andree-Labsch et al. 2017) and a primary target for observational studies of physical and chemical structures of PDRs (e.g., Tielens et al. 1993; Walmsley et al. 2000; van der Wiel et al. 2009; Arab et al. 2012; Peng et al. 2012; Goicoechea et al. 2016; Nagy et al. 2017). Observations of various molecular spectral lines have shown that the emissions could be better interpreted with an inhomogeneous density structure containing an extended and relatively low density (

 cm−3) medium and a compact and high-density (

 cm−3) component (e.g., Hogerheijde et al. 1995; Young Owl et al. 2000; Leurini et al. 2006, 2010; Goicoechea et al. 2016). However, due to the scarcity of high-resolution observations capable of spatially resolving the density structure, the nature of the high-density clumps or condensations is still not well understood. Lis & Schilke (2003) mapped the Bar in H13CN (1–0) with the Plateau de Bure Interferometer (PdBI) at an angular resolution of about 5″, and detected 10 dense clumps. They proposed that the H13CN clumps are in virial equilibrium and may be collapsing to form stars. Goicoechea et al. (2016) performed Atacama Large Millimeter/submillimeter Array HCO+ (4–3) observations of the Bar and detected over-dense substructures close to the cloud edge, and found that the substructures have masses much lower than the mass needed to make them gravitationally unstable. These two interferometric observations both target molecular spectral lines. A high-resolution map of the dust continuum emission of the Bar, which is highly desirable in constraining the mass and density of the dense condensations, is still lacking. Here we report our Submillimeter Array (SMA) observations of the dust continuum and molecular spectral line observations of the Bar.


Before leaving this section, in Fig. 8 we display the spatial dependence of the self-consistent neutron and proton density profiles for the optimal solutions in spherical WS cells with average baryon densities nb = 0.0475 fm-3, 0.065 fm-3, and 0.076 fm-3. It is observed that in denser matter the size of the WS cell decreases, as we discussed previously, and that the amount of free neutrons in the gas increases, as expected. It can be seen that the nuclear surface is progressively washed out with increasing average baryon density as the nucleon distributions become more uniform. At high nb the density profile inside the WS cell extends towards the edge of the cell, pointing out that the WS approximation may be close to its limits of validity (Negele & Vautherin 1973; Chamel et al. 2007; Baldo et al. 2007; Pastore et al. 2011; Gögelein & Müther 2007; Newton & Stone 2009). Although the proton number Z is similar for the three average baryon densities of Fig. 8, the local distribution of the protons is very different in the three cases. In Fig. 8c the proton density profile extends more than 3 fm farther from the origin than in Fig. 8a, while the central value of the proton density has decreased by more than a factor 2, hinting at the fact that the neutrons have a strong drag effect on the protons. Figure 9 presents the nucleon density profiles obtained for cylindrical and planar geometries at the same average density nb = 0.076 fm-3 as in Fig. 8c. From Figs. 8c (droplets),  9a (rods), and 9b (slabs) we see that the size of the WS cells decreases with decreasing dimensionality, i.e. Rc,droplet>Rc,rod>Rc,slab. At high average densities near the crust-core transition, nucleons inside the WS cell can arrange themselves in such a way that the region of higher density is concentrated at the edge of the cell, leaving the uniform region of lower density in the inner part of the cell. This distribution of nucleons corresponds to the cylindrical tube and spherical bubble configurations. In Figs. 9c and d, we plot the neutron and proton density profiles of the optimal solution for tubes and bubbles at nb = 0.076 fm-3. At equal average density, the size of the cells containing tubes and bubbles is larger than the size of the cells accommodating rods and droplets, respectively, as can be appreciated by comparing Fig. 9a for rods with Fig. 9c for tubes, and Fig. 8c for droplets with Fig. 9d for bubbles. As a consequence of this fact and of the effectively larger value of the integration factors 2πr and 4πr2 when the densities are accumulated near the edge of the cell, the total number of nucleons and the atomic number in the tube and bubble cells is about 1.5−2 times larger than in their rod and droplet counterparts. The proton fraction xp = Z/A is, however, practically the same for all geometries.


It has also become evident that the normalization of the MS rapidly increases with redshift: distant galaxies form stars at higher paces than in the local Universe, at fixed stellar mass (e.g., Daddi et al. 2007; Elbaz et al. 2007; Whitaker et al. 2012; Speagle et al. 2014; Schreiber et al. 2015). This trend could be explained by the availability of copious molecular gas at high redshift (Daddi et al. 2010a; Tacconi et al. 2010, 2018; Scoville et al. 2017a; Riechers et al. 2019; Decarli et al. 2019; Liu et al. 2019a), ultimately regulated by the larger accretion rates from the cosmic web (Kereš et al. 2005; Dekel et al. 2009a). Moreover, higher SFRs could be induced by an increased efficiency of star formation due to the enhanced fragmentation in gas-rich, turbulent, and gravitationally unstable high-redshift disks (Bournaud et al. 2007, 2010; Dekel et al. 2009b; Ceverino et al. 2010; Dekel & Burkert 2014), reflected on their clumpy morphologies (Elmegreen et al. 2007; Förster Schreiber et al. 2011; Genzel et al. 2011; Guo et al. 2012, 2015; Zanella et al. 2019). IR-bright galaxies with prodigious SFRs well above the level of the MS are observed also in the distant Universe, but their main physical driver is a matter of debate. While a star formation efficiency (SFE = SFR/Mgas) monotonically increasing with the distance from the main sequence (ΔMS = SFR/SFRMS, Genzel et al. 2010, 2015; Magdis et al. 2012; Tacconi et al. 2018, 2020) could naturally explain the existence of these outliers, recent works suggest the concomitant increase of gas masses as the main driver of the starbursting events (Scoville et al. 2016; Elbaz et al. 2018). In addition, if many bright starbursting (sub)millimeter galaxies (SMGs, Smail et al. 1997) are indeed merging systems as in the local Universe (Gómez-Guijarro et al. 2018, and references therein), there are several well documented cases of SMGs hosting orderly rotating disks at high redshift (e.g., Hodge et al. 2016, 2019; Drew et al. 2020), disputing the pure merger scenario. The same definition of starbursts as galaxies deviating from the main sequence has been recently questioned with the advent of high spatial resolution measurements of their dust and gas emission. Compact galaxies with short depletion timescales typical of SBs are now routinely found on the MS, being possibly on their way to leave the sequence (Barro et al. 2017a; Popping et al. 2017; Elbaz et al. 2018; Gómez-Guijarro et al. 2019; Puglisi et al. 2019; Jiménez-Andrade et al. 2019); or galaxies moving within the MS scatter, due to mergers unable to efficiently boost the star formation (Fensch et al. 2017) or owing to gravitational instabilities and gas radial redistribution (Tacchella et al. 2016).


The Pleiades open cluster does not only offer us a beautiful spectacle during the fall, it is one of the best-studied stellar associations and one of the cornerstones to understand stellar properties and evolution. In fact, the literature includes more than one thousand refereed papers dealing with the Pleiades, most of which use the Pleiades as a reference, just in the last ten years. In spite of this, the Pleiades cluster still has many secret and basic parameters, such as its distance and age, that are not clearly established. Even the census of this cluster is incomplete, although the recent works by Stauffer et al. (2007), Lodieu et al. (2012), and Bouy et al. (2013) have considerably improved the membership list, In fact, Bouy et al. (2015) has increased the number of known members by 50% with public archival data, very accurate proper motions, and multiwavelength photometry (see additional details in Sarro et al. 2014). Regarding its distance, there are currently two different methodologies based on parallaxes from Hipparcos (Perryman et al. 1997) and isochrone fitting, respectively. Pre-Hipparcos distances for the Pleiades range between 125 and 130 pc (see, for instance, Soderblom et al. 1993d), whereas the initial distance derived by Hipparcos is much closer, about 119 pc (van Leeuwen 1999). This last value is significantly different from the distance derived by Pinsonneault et al. (1998), who used color-magnitude diagrams and fitting isochrones and obtained 133.5 ± 1.2 pc. More recently, van Leeuwen (2009), by reanalyzing Hipparcos data, derived a distance of 120.2 ± 1.9 pc. These values should correspond to the distance to the cluster center, whose core radius should be around 3 degrees, which corresponds to 5–6 pc. Another trigonometric parallax, based on Hubble Space Telescope data and three members, was found by Soderblom et al. (2005), who obtained 134.6 ± 3.1 pc. Recently, Melis et al. (2014) derived a distance of 136.2 ± 1.2 pc based on an accurate parallax for four bona fide members obtained with the VLBI. This value also agrees with that derived by Galli et al. (in prep.) using accurate proper motions and the convergence point method (137.7 ± 2.5 pc).


The driving mechanisms and chronology of the buildup of bulges in late-type galaxies (LTGs) is an issue of key relevance to our understanding of galaxy evolution. According to our current knowledge on bulge demographics in the local universe, a large fraction of LTGs host pseudo-bulges (PBs; e.g., Gadotti 2009; Fisher & Drory 2011; Fernández Lorenzo et al. 2014) that substantially differ from classical bulges (CBs) in their spectrophotometric and kinematical characteristics. The latter resemble in many respects “old and dead” elliptical galaxies, lacking ongoing star-formation (SF), exhibit a spheroidal shape with inwardly steeply increasing surface brightness profiles (SBPs) being well approximated by the Sérsic (1963) fitting law with a high (≳3) exponent η, show stellar kinematics dominated by velocity dispersion (σ⋆) and obey the Kormendy (1977) scaling relations for normal elliptical galaxies (Fisher & Drory 2010). It is observationally established that CBs contain a super-massive black hole (SMBH) with a mass M∙ tightly correlating with their stellar mass

${\cal M}_{\star,\textrm{B}}, \sigma_{*}$M⋆,B,σ*
 and optical luminosity (Ho 2008; Kormendy & Ho 2013; see also Ferrarese & Merritt 2000). Traditionally, bulges were thought to invariably form early-on via violent quasi-monolithic gas collapse (Larson 1974) or mergers (Bender et al. 1992; Aguerri et al. 2001; Keselman & Nusser 2012) associated with vigorous nuclear starbursts (Okamoto 2012), with the disk gradually building up around them. Whereas this inside-out galaxy formation scenario appears consistent with important integral characteristics of CBs (e.g., their red colors), it does not offer a plausible explanation for the presence of PBs in present-day LTGs. These generally show ongoing SF, a significant degree of rotational support (Kormendy & Kennicutt 2004, for a review) and flatter/ellipsoidal shapes with nearly exponential SBPs (η≲2; e.g., Drory & Fisher 2007; Fisher & Drory 2010). Even though there is observational evidence that PBs also contain a SMBH (Kormendy et al. 2011; Kormendy & Ho 2013), in some cases revealing itself as an active galactic nucleus (AGN; e.g., Kotilainen et al. 2016; see Kormendy & Ho 2013 for a review), these do not follow the M∙ –σ* correlation for CBs, which appears to be consistent with a different formation route. Indeed, the prevailing concept on PB formation is that these entities emerge gradually out of galactic disks through gentle gas inflow spawning quasi-continuous SF and the emergence of a central bulge-like luminosity excess at their centers (e.g., Courteau et al. 1996; Carollo et al. 2001; Kormendy & Kennicutt 2004). Besides bar-driven gas inflow (e.g., Springel & Hernquist 2005), various other mechanisms, such as inward stellar migration, minor mergers with low-mass satellites, or a purely dynamical re-arrangement of the disk (Scannapieco et al. 2010; Guedes et al. 2013; Bird et al. 2012; Roskar et al. 2012; Grand et al. 2014; Halle et al. 2015) have been proposed as further contributors to PB growth along the Gyr-long secular evolution of LTGs.

Background Background Background Compare/Contrast


The majority of the particles collected by Stardust are olivine and pyroxene silicates with solar isotopic compositions, which suggests an origin in our solar system rather than an interstellar provenance. These polymineralic particles dominate those made of a single mineral even down to sizes smaller than 100 nm, indicating that the dust composition is surprisingly consistent at different scales and that the smallest subunits of the dust may be as small as tens of nanometers (Hörz et al. 2006; Zolensky et al. 2006). The sizes of these smallest single mineral impactors are similar to those of the nanocrystals determined by Rietmeijer (1993). As discussed above, they might also be existing in MIDAS dust particles and might be fused into the 100 nm features. Price et al. (2010) and Wozniakiewicz et al. (2012) investigated the sizes of particles smaller than 10 μm that impacted the aluminum foils of the Stardust probe. The distribution peaks at about 175 nm, but if we assume that the particles areagglomerates of smaller subunits, as indicated by their common polymineralic nature, then the subunit size distribution would peak at sizes below 100 nm (Price et al. 2010). A study of over 450 particles that do not seem to be agglomerates, that is, those that show single mineral impactors of silicate or sulfide, found geometric mean sizes of

$532^{741}_{-310}$

532

−310

+741

 nm for the silicate particles and

$406^{491}_{-222}$

406

−222

+491

 nm for the sulfides (Wozniakiewicz et al. 2013). These sizes are notably larger than the 175 nm (or less) found for the whole dataset. This large spread of subunit sizes could indicate a size distribution with a large width. No fits of these sizedistributions are available, but the figures in Wozniakiewicz et al. (2012) and Price et al. (2010) indicate that the differential sizes may follow a log-normal distribution. When we assume that the smallest subunit sizes are possibly between tens and hundreds of nanometers, the subunit size range found for MIDAS smallest features would be encompassed. The determination of the size distributions for the small Stardust particles and a detailed comparison to the distributions obtained for comet 67P could be the work of an interesting future project.


In particular, stellar wind P-Cygni profiles and photospheric absorption lines are detected with high significance (green and yellow dashed lines in Fig. 2). The detection of photospheric lines in J0121+0025 indicates unambiguously that the UV luminosity is dominated by stellar emission, rather than an AGN. We identify more than ten photospheric features. Some of them are resolved and detected with high significance (e.g. C ii 1324 Å, O iv 1343 Å, and S v 1501 Å; see Fig. 3). We use these to determine the systemic redshift zsys = 3.244 ± 0.001 of J0121+0025. Others are seen in blends from multiple transitions (e.g. Si iii 1417 Å, C iii 1427 Å, and Fe v 1430 Å at λ0 ≃ 1415–1435 Å). These stellar absorption lines are intrinsically weak in star-forming galaxies, with EW0 typically well bellow 1 Å (e.g. Shapley et al. 2003; Steidel et al. 2016; Rigby et al. 2018). As they are formed in the photospheres of hot stars and are seen in absorption, the background radiation should be dominated by the starlight, otherwise they would not be detected. Even a small contribution of an AGN to the UV continuum (≲ 25 per cent), that is featureless in these spectral regions, would make these lines disappear at the SNR of our spectrum. In addition, the observed P-Cygni profiles in N v 1240 Å and C iv 1550 Å can be also well explained/modelled by stellar models with a very young age (≃3 Myr burst; see Fig. 3 and Section 3.2 for details), similar to those seen in other very young starbursts (e.g. Rivera-Thorsen et al. 2019; Vanzella et al. ), some of them also very/extremely luminous (Vanzella et al. 2018; Marques-Chaves et al. 2020b). While some rare AGNs, such as broad or narrow absorption line QSOs (BAL/NAL QSOs), can show N v and C iv profiles mimic those of stellar P-Cygni, from the combination of a broad emission and a redshifted absorption (see for example Bentz, Osmer & Weinberg 2004; Appenzeller et al. 2005), photospheric lines are not present in the spectra of AGNs.


NLFFF extrapolation provides the reconstructed coronal magnetic field for AR 11158 from 2011 February 13 − 2011 February 15 (Figures 1(d)–(f)). The field lines (yellow lines) within the core of the AR have arcade-like structure with a relatively strong twist mainly near the PIL. These figures show that the magnetic field evolved during this period. Although we did not quantitatively compare the field lines with the observation, in general, the reconstructed coronal field morphologies match with the observations in Figures 1(a)–(c). The general morphologies and the locations of the high-twist fields are also in agreement with many previous studies (Jing et al. 2012; Sun et al. 2012; Dalmasse et al. 2013; Inoue et al. 2013, 2014a; Liu et al. 2013; Wang et al. 2013; Aschwanden et al. 2014; Malanushenko et al. 2014; Zhao et al. 2014). Unlike Zhao et al. (2014) who could identify the twisted flux rope from the topology of the reconstructed coronal field, we could not find an obvious topological signature of a flux rope existing in our NLFFF during our analysis time window. It might be due to the fact that there was little magnetic flux with twist higher than one turn in our NLFFF and it is difficult to topologically define it as a flux rope. However, our result is consistent with other NLFFF results (Jing et al. 2012; Sun et al. 2012; Liu et al. 2013; Wang et al. 2013; Inoue et al. 2014a; Malanushenko et al. 2014). The high-twist region in our result is also in agreement with the region with high helicity flux (Dalmasse et al. 2013) and the location of the flare ribbons (Bamba et al. 2013; Liu et al. 2013), as well as the high current density region (Janvier et al. 2014). Figures 1(g)–(i) show the evolution of the twist distribution map, with the magnetic twist of the field lines plotted at the footpoints of field lines according to a color scale. This shows that the high-twist (strongly right-handed twist corresponding to Tw > 0.5) areas are concentrated in only a limited part of the AR. The high-twist area grew and became even more twisted just before the X2.2 flare (Figure 1(i)). Most parts of the AR have twist values less than 0.25, but near the PIL the twist can reach more than 0.5, even up to about a full turn. This is consistent with the results of Sun et al. (2012) and Inoue et al. (2014a). A high-twist (strong negative/left-handed twist) area also developed in the eastern part of the AR, which did not exist initially on February 13. Both of these high-twist areas produced several flares. However, here we focus on the flares that resulted from the high-twist core region near the center of the AR, where the M6.6 and X2.2 flares occurred.


While previous studies mainly focused on the optical-UV properties of a MAD in RLAGNs, for the first time we try to investigate their X-ray properties in this work. The origin of X-ray emission in RLAGNs is still under debate, which may come from a corona, jet, or both. In observations, there is a big difference between the X-ray properties of radio-quiet AGNs (RQAGNs) and RLAGNs. Firstly, the average X-ray flux in RLAGNs is found to be 2–3 times higher than that in RQAGNs (e.g., Zamorani et al. 1981; Wilkes & Elvis 1987; Li & Gu 2021). Secondly, Laor et al. (1997) reported that RLAGNs have harder 2–10 kev X-ray spectra than RQAGNs by compiling a sample of 23 quasars observed with ROSAT, which was subsequently confirmed by Shang et al. (2011) with a larger sample. Comparing the X-ray spectrum of 3CRR quasars and that of radio-quiet quasars, Zhou & Gu (2020) also gave a similar result. In addition, the X-ray reflection features of RLAGNs are weaker than those of RQAGNs (Wozniak et al. 1998). All of these results seem to indicate that the contribution of a jet to X-ray spectra cannot be neglected. However, several recent works suggested a totally different result. First, the slope of LUV − LX is found to be consistent for RLAGNs and RQAGNs (Zhu et al. 2020, 2021; Li & Gu 2021). Second, Gupta et al. (2018, 2020) discovered that the distributions of X-ray photon spectral indices between RLAGNs and their radio-quiet counterpart are very similar (see Zhu et al. 2021 either). This opposite conclusion may be due to the effect of sample selection. The sample of Gupta et al. (2018, 2020) was X-ray selected (and optically selected for the sample of Zhu et al. 2021), which may lead to the radio jet power being very feeble compared to the bolometric luminosity in most of the RLAGNs. These weakly jetted RLAGNs can therefore have different X-ray photon indices compared to the strong jetted RLAGNs, such as the 3CRR quasars of Zhou & Gu (2020). P15 also indicated that the weakly jetted RLAGNs have similar αEUV as RQAGNs. However, interestingly, Markoff et al. (2005) demonstrated that both the corona model and the jet model can fit the X-ray data of some Galactic X-ray binaries well and that the jet base may be subsumed to corona in some ways. The 3CRR quasars are low frequency radio selected and have a strong jet on a large scale. However, it is still unclear whether all the objects with a strong jet harbor a MAD, or just containing MAD when jet is firstly launching millions of years ago. We focus on the RLAGNs with an EUV deficit in this work, which should possess a MAD in the inner disk region as suggested by P15. The presence of a MAD surrounding the black hole may bring a remarkable difference to the X-ray emissions since the structure of disk-corona greatly changes in the case of MAD (e.g., Tchekhovskoy et al. 2011; McKinney et al. 2012; White et al. 2019). In theory, it has been suggested that X-ray emission increases when an advection-dominated accretion flow (ADAF) becomes a MAD in its inner region (Xie & Zdziarski 2019). Nevertheless, how MAD affects the disk-corona corresponding to the X-ray emission of quasars is still an open issue. This work can constrain a future theoretical model for MAD in RLAGNs.


A comparison of the He and O abundances in RWT 152 with evolutionary models of stellar yields by Marigo (2001) suggests a progenitor star with an initial mass of ∼1.3 M⊙ and very low metallicity Z = 0.004. The stellar mass is compatible with that expected for sdOs (Heber 2009) and the low metallicity indicates that the progenitor was formed in a poor-metal environment. In contrast, similar models of initial mass ∼1.25 M⊙ and Z = 0.004 by Karakas (2010) predict abundances substantially different from those found in RWT 152. Even for the lowest metallicity in the Karakas models (Z = 0.0001), we do not recover the chemical abundances of RWT 152. Therefore, it is clear that drawing conclusions about the progenitor star of RWT 152 from the chemical abundances of the nebula may be a bit misleading. For this reason, we have obtained information about the progenitor from the current status of its CS. Fig. 6 shows the position of RWT 152 in the Hertzsprung-Russell diagram log g–Teff with the post-AGB tracks by Bloecker (1995) and Schoenberner (1983). The location of the star (Teff ≃ 45 000 K, log g ≃ 4.5; Ebbets & Savage 1982) is consistent with a current mass of M∼0.55 M⊙ which implies an initial mass in the main sequence of ∼1 M⊙. For such a low-mass star, the ejected mass during the AGB evolution is expected to be small. In fact, with the electron density derived above and the size of the nebula, we obtain values of 1.3 × 10−2–1.6 × 10−1 M⊙ for the ionized nebular mass (assuming 2.4 and 6.5 kpc, respectively, and a filling factor of 0.6). These values are much smaller than ionized masses usually obtained for PNe (see e.g. Hua & Kwok 1999), further supporting a low-mass progenitor for RWT 152. Moreover, taking into account that 0.1–0.3 M⊙ are lost in the red giant branch phase of a low-mass star (Dorman, Rood & O'Connell 1993), the current mass of the CS and the obtained ionized mass, we recover a progenitor star with a ∼0.8–1.0 M⊙, in agreement with the value obtained from the position of the CS in the log g–Teff diagram. Such a low ionized mass, combined with a relatively high kinematical age, may explain the low surface brightness of RWT 152. If a low-mass progenitor is involved in the evolution of other PNe+sdO systems, it is not surprising that these PNe are very faint and, in some cases, may have faded beyond detection.


As done by Le Fèvre et al. (2020), we visually inspect the ancillary data, the intensity maps, the velocity and velocity dispersion fields presented in Sect. 3.1 to search for the presence of multiple components or disturbed morphology near the position of the targets. The channel maps, the spectra and the PVDs are checked together searching for consistent emission features. By taking into account the results of the initial qualitative classification by Le Fèvre et al. (2020) and of the more recent quantitative analysis of a subsample of the ALPINE targets by Jones et al. (2021), we proceed with a more in-depth characterization of the [CII]-detected galaxies aimed at obtaining a robust merger fraction at z ∼ 5. Adopting the same criteria described in Sect. 2 to differentiate the targets and considering the S/N of the minor merger component as described in Sect. 3.1, we find a slightly lower fraction (∼31%, 23 out of 75 [CII]-detected sources) of mergers7 if compared to the 40% found by Le Fèvre et al. (2020), with 12, 20 and 7% of the sample made by rotating, extended and compact dispersion dominated sources, respectively. To be more conservative in the classification of the galaxies (especially for obtaining a more robust merger statistics), we define the remaining 30% of the sample as ‘uncertain’, a new category that includes the weak galaxies (as described in Le Fèvre et al. 2020) and also objects that, by visual inspection, present features that are intermediate to those of various classes. This category is similar to the ‘uncertain’ (UNC) class introduced in Jones et al. (2021) that, based on the results of the 3DBarolo fits, contains sources they are unable to classify because of the low S/N and/or spectral resolution, or contrasting evidence in their classification criteria. Although the uncertain category is populated by a significantly larger fraction of sources with respect to the weak class (∼16%) in Le Fèvre et al. (2020), we recover the same qualitative morpho-kinematic distribution of the previous analysis, confirming the high fraction of rotators and mergers at these early epochs.


Altogether, the approximately thirty molecules recently detected have confirmed the chemical complexity in the nebula, and generated our interest for the present study. Of these species, we will focus our attention on the seventeen species listed by molecular families in Table 1. As can be seen in this table, the WHISPER survey allowed the detection of some organic molecules in the Horsehead nebula, such as formaldehyde (H2CO) and methanol (CH3OH), which constitute key species in the likely synthesis of more complex organic molecules such as some prebiotic molecules (Bernstein et al. 2002; Muñoz Caro et al. 2002; Garrod et al. 2008). Because they are detected in a wide variety of interstellar sources – in hot cores (Sutton et al. 1995; Ceccarelli et al. 2000), dark clouds (Bergman et al. 2011), shocked regions (e.g. Sakai et al. 2012; Codella et al. 2012; Tafalla et al. 2010) and even in comets (Mumma & Charnley 2011; Cordiner et al. 2015) – it is of prime importance to understand well how these precursor molecules form. H2CO is commonly thought to form both in the gas-phase and on grain surfaces, while CH3OH is believed to be only formed on grain surfaces (Garrod et al. 2006; Geppert et al. 2006). Guzman et al. (2013) reported the observations of these two molecules toward the Horsehead nebula in both the PDR and Core positions. Unable to reproduce the observed abundances of either H2CO or CH3OH at the PDR position with only pure gas-phase models, they concluded that, for this region, both species are formed on grain surfaces and then photodesorbed into the gas phase. On the other hand, at the Core position, a pure gas-phase model can reproduce the observed H2CO abundance, while photodesorption of ices is still needed to explain the observed abundance of CH3OH. Other organic molecules were reported in the Horsehead nebula as first detections in a PDR environment, including HCOOH (formic acid), CH2CO (ketene), CH3CHO (acetaldehyde), and CH3CCH (propyne) (Guzman et al. 2014). Their abundances were found to be higher at the PDR position than at the Core, revealing that complex organic chemistry is also occurring in UV-illuminated neutral gas (Guzman et al. 2014). Of these molecules, some – HCOOH, CH2CO, and CH3CHO – have now also been detected in the Orion bar PDR (Cuadrado et al. 2016, 2017).


The duration of the manually scheduled follow-up observations were also increased to improve the likelihood of a radio detection. The recent investigation of the entire sample of radio-detected GRBs before 2011 April by Chandra & Frail (2012) demonstrated that the majority of GRBs detected in the radio band at 8.5 GHz had a peak flux between 0.1 and 0.2 mJy beam−1 at 5–10 d post-burst (see fig. 4 of Chandra & Frail 2012). A 4 h AMI observation is therefore required to reach an rms noise of ∼0.03–0.04 mJy beam−1 that will allow the reliable detection of >0.1–0.2 mJy beam−1 sources. However, it is worth noting that since GRB relativistic blast waves generate synchrotron radiation as they expand into the circumstellar (wind generated) medium (Granot & Sari 2002), we expect the forward-shock of the afterglow to peak more brightly at 15.7 GHz and at earlier times than the peaks recorded by Chandra & Frail (2012). We therefore require a higher monitoring cadence at early times (within 5 d post-burst) to detect similar radio peaks. As the range of radio peaks observed by Chandra & Frail (2012) will be brighter at 15.7 GHz, the rms achieved by a 4 h AMI observation will be sufficient for detecting events similar to those seen in their sample. The follow-up observations are manually scheduled to occur near transit approximately 24 h, 3, 7, and 10 d post-burst, with this temporal spacing designed to catch the peak of the forward- or reverse-shock at 15.7 GHz at a range of redshifts (z ≲ 5; e.g. see figs 22 and 23 of Chandra & Frail 2012). In the event that a GRB radio counterpart was detected, the AMI observing cadence was increased to a 4 h observation every 1 or 2 d. As part of the AMI GRB observing programme, we also obtained manually scheduled observations of GRBs that were detected with the Fermi Large Area Telescope (LAT; Atwood et al. 2009), the Fermi Gamma-ray Burst Monitor (GBM; Meegan et al. 2009) and the International Gamma-Ray Astrophysics Laboratory (INTEGRAL; Winkler et al. 2003), whose positions had been more precisely localized through the identification of X-ray and/or optical counterparts, usually by the Swift-XRT, Swift-UVOT, or one of the ground-based GRB follow-up programmes.


The high-twist region in our result is also in agreement with


The ESPRESSO DFS concept (Di Marcantonio et al. 2018) was conceived during its preliminary design phases with the goal of maximizing operational efficiency, flexibility, and scientific output while complying with the standard Paranal Observatory operational scheme. The main challenge derives from the requirement to operate ESPRESSO in a seamless way with any of the UT’s or with all four UT’s simultaneously. This must be possible not only with a predetermined schedule, but also “on the fly”. The flexibility in ESPRESSO’s operations has been tackled by adopting a new DFS deployment plan described in Di Marcantonio et al. (2018) that is exceptional under various aspects because it has to cope with various telescope and instrument configurations while remaining operationally simple. Figure 8 shows the main ESPRESSO DFS elements and their final deployment. Besides the software packages already described, part of the software for the control of the CT devices has been incorporated into the VLT telescope CS to allow CT operations even when ESPRESSO is offline (thus avoiding conflicts, e.g., with instruments of the VLT Interferometer operations). In addition to the standard DFS software packages, ESPRESSO is the first instrument to also provide a data analysis package that is able to extract relevant astronomical observables from the reduced data. The following DFS subsystems are specific to ESPRESSO: (1) the ETC hosted on the ESO web page9; (2) the CS with the full suite of acquisition, observation, and calibration templates that are able to control all vital parts of the instrument and the CT (Calderone et al. 2018); (3) the data reduction software (DRS) package (or “pipeline”) capable of providing “science-ready” reduced data only minutes after the end of the individual observation; (4) the data analysis software (DAS) package that produces higher-level astronomical observables with no or limited supervision; (5) the DRS and DAS are distributed to the community10.


Eigenbrod et al. (2006) first modelled this system with HST imaging and suspected that the second set of bluer arcs in F814W band (see Fig. 1) inside and outside the area delimited by the red arcs in F160W band could be either a second source at a different redshift or a star-forming region in the source galaxy. We examine the possibility of a second source plane existing at a lower redshift than the source (z = 1.52) due to the bluer colour of the arc and find that the scenario is very unlikely, as the macro model determined by the red arc cannot reproduce a reasonable source for the blue arcs given a possible range of the source redshift from z = 0.5 to z  1.52. In contrast, we do find that a star-forming region can be reconstructed at the same source redshift. Faure et al. (2011) modelled the lens with high-resolution H and Ks imaging obtained using the European Southern Obseratory (ESO) Very Large Telescope (VLT) with AO and the laser guide star system. They identified a luminous object, located ∼0.3 arcsec to the north of the lens galaxy, but showed that it cannot be responsible for the anomalous flux ratios. Many studies (e.g. Metcalf & Madau 2001; Bradač et al. 2002; Dalal & Kochanek 2002; Pooley et al. 2012; Schechter et al. 2014; Glikman et al. 2018; Badole et al. 2020) have shown that the macro model cannot explain the flux ratios in this lens, which suggested the presence of microlensing or dark matter substructures. Thus, to avoid possible biases caused by flux ratios, we only use the lensed quasar positions and the extended arc to constrain the mass model, which is also the standard procedure for H0 measurements in TDCOSMO Collaboration. Ignoring the fluxes of the lensed quasar images will not affect the constraining power of the imaging data, since the lensed arc emission is much more constraining than the lensed quasar fluxes. In addition, For the error budget of H0 contributed from the uncertainties of the lensed quasar positions, Chen et al. (2021a) have showed that with high-resolution AO imaging it is a subdominate term given configuration of the J0924+0219 lens system. Gilman et al. (2020) also show that the presence of substructures do not bias H0 above the per cent level. We use glee, a strong lens modelling code to model (Suyu & Halkola 2010; Suyu et al. 2012) the three HST bands and one Keck AO band simultaneously. We describe the models in the following for fitting the high-resolution imaging data. We show the imaging, models, normalized residuals, and reconstructed sources in Fig. 4. Note that since the source in F555W band has more clumpy star-forming regions, the reconstructed source is less regular, with small-scale structures and more noise.5 In addition, the noise-overfitting problem is due to the fact that the outer region of the source plane is under-regularized, but this effect will not affect the uncertainty because the uncertainty will be dominated by the time delay and velocity dispersion measurements. Besides, we model the imaging with different source resolutions and marginalize over them to control the systematics.


The origin of the GC VHE emission remains undetermined, due in part to source confusion and the limitations of current instruments. The source of VER J1745–290 may be Sgr A* (Atoyan & Dermer 2004; Aharonian & Neronov 2005; Ballantyne et al. 2011; Chernyakova et al. 2011; Fatuzzo & Melia 2012; Kusunose & Takahara 2012; Fujita et al. 2017; Rodríguez-Ramírez et al. 2019) or PWN G359.95-0.04 (Wang et al. 2006; Hinton & Aharonian 2007), with which it is spatially coincident (Acero et al. 2010). Other possible origins include the annihilation of dark matter particles (Bergström et al. 2005a, 2005b; Horns 2005; Profumo 2005; Aharonian et al. 2006c; Belikov et al. 2012; Cembranos et al. 2012, 2013; Gammaldi et al. 2016) or a population of millisecond pulsars (Bednarek & Sobczak 2013; Bartels et al. 2016; Guépin et al. 2018). The mechanism of gamma-ray emission may be predominantly due to hadronic processes, where relativistic protons interact with gas and subsequently produce gamma rays through neutral pion decay (Aharonian & Neronov 2005; Ballantyne et al. 2011; Chernyakova et al. 2011; Fatuzzo & Melia 2012; Linden et al. 2012; Guépin et al. 2018), leptonic processes where gamma rays are produced when electrons and positrons undergo inverse Compton scattering off a radiation field (Atoyan & Dermer 2004; Hinton & Aharonian 2007; Kusunose & Takahara 2012; Lacroix et al. 2016), or a combination of processes (hybrid scenario), where leptons produce high-energy, but not VHE, gamma rays (Guo et al. 2013). Both the correlation of VHE emission with the CMZ and the lack of a cutoff in the diffuse spectrum support a hadronic scenario, capable of explaining both VER J1745–290 and the diffuse emission (Aharonian et al. 2006b; Linden et al. 2012; Abramowski et al. 2016). Measurement of the diffuse spectrum by H.E.S.S. up to energies of tens of TeV with no evidence of a cutoff has also been interpreted as evidence for the existence of PeV protons within the central 10 parsecs of the GC, accelerated by Sgr A* (Abramowski et al. 2016). While cosmic rays are known to extend up to PeV energies (e.g., Hörandel 2003), few, if any, accelerators of PeV cosmic rays, or “PeVatrons,” have been clearly established (e.g., Abramowski et al. 2016; Abeysekara et al. 2020). Discovering the nature of PeVatrons in our Galaxy is thus a particularly important step in understanding the origins of cosmic rays.


In the last two decades, it has been demonstrated that massive galaxies harbour supermassive black holes (SMBHs) in the centres of their bulge components (Kormendy & Ho 2013, and references therein). Also, at redshifts higher than 6, quasars are found that possess SMBHs with the mass higher than 109 M⊙ (Fan et al. 2001; Kurk et al. 2007). The formation history of these SMBHs is among the most significant unsolved issues in astrophysics (Volonteri & Bellovary 2012; Haiman 2013). Two recently discovered high-redshift quasars, ULAS J112010+641 with the mass of MBH = 2 × 109 M⊙ at redshift z = 7.085 (Mortlock et al. 2011) and SDSS J01001+2802 with MBH = 1.2 × 1010 M⊙ at z = 6.30 (Wu et al. 2015) have raised a serious problem for the formation of SMBHs. Possible building blocks for such high-redshift SMBHs are the remnants of first stars. The initial mass function of first stars is thought to be more or less top-heavy (Abel, Bryan & Norman 2000; Nakamura & Umemura 2001; Bromm, Coppi & Larson 2002; Yoshida et al. 2006; Greif et al. 2011; Hirano et al. 2014; Susa, Hasegawa & Tominaga 2014). First stars of several tens M⊙ undergo supernovae, leaving black holes (BHs) of few tens M⊙ (Heger & Woosley 2002). For SMBHs to grow from such first star remnants through mass accretion at z ≳ 6, a super-Eddington accretion rate is requisite. If SMBHs grow continuously by mass accretion from BH remnants of ∼20 M⊙, the Eddington ratio (λ) is required to be λ = 1.4 for ULAS J112010+641, or λ = 1.3 for SDSS J01001+2802. However, the continuous accretion is unlikely to be sustained due to feedbacks, and thus the average mass accretion rates should be lower than the Eddington rate (Alvarez, Wise & Abel 2009; Milosavljevic, Couch & Bromm 2009). On the other hand, seed BHs may stem from supermassive stars of 104–6 M⊙ as a result of the direct collapse of primordial density fluctuations (Umemura, Loeb & Turner 1993; Bromm & Loeb 2003; Inayoshi & Omukai 2012). These BHs are thought to be incorporated into a primordial galaxy of ∼108–109 M⊙ (Greene 2012). If an SMBH grows via gas accretion from such a massive BH (MBH), the constraint on the accretion rate can be alleviated.


So far, the discussion has emphasized particle energization by local plasma regions of contracting and fast reconnecting (merging) small-scale flux ropes generated in the vicinity of large-scale primary current sheets through turbulent current sheet reconnection. However, one can also approach this topic from the perspective of MHD turbulence theory, simulations, and related solar wind observations. Theoretical considerations and simulations of MHD turbulence in the presence of a significant background/guide magnetic field suggest that solar wind turbulence can to lowest order be modeled in terms of a combination of a dominant quasi-2D turbulence component of coherent structures (small-scale magnetic islands) perpendicular to the background/guide field and a minor parallel-propagating Alfvén wave turbulence component (Shebalin et al. 1983; Zank & Matthaeus 1992, 1993; Dmitruk et al. 2004; Zank et al. 2017), a view that is consistent with analysis of solar wind observations (Matthaeus et al. 1990; Bieber et al. 1996) and with the finding that quasi-2D turbulence alone is not sufficient to explain observed solar wind turbulence (Turner et al. 2012). A recent analysis of Wind data to identify inertial-scale flux ropes indicates that these structures are much more commonplace in the solar wind near 1 au than previously thought. Zheng (2017), Zheng et al. (2017), and Zheng & Hu (2018) identified an unprecedented number of small-scale flux ropes at 1 au with scales in the inertial range using the Grad–Shafranov reconstruction approach (∼3500 per year on average) with a clear solar cycle dependence, a number that is expected to grow when the data analysis shifts to shorter timescales. Furthermore, an axial (out-of-plane) current density distribution constructed from the Grad–Shafranov-based data analysis yielded a non-Gaussian probability density function (pdf) entirely consistent with the out-of-plane current density pdf produced from compressible 2D MHD turbulence simulations with a strong out-of-plane guide field, in which merging magnetic island structures are a common occurrence (Greco et al. 2009). This result, combined with the sheer number of small-scale flux ropes being identified, suggests that the common occurrence of small-scale flux ropes in the low-latitude solar wind near 1 au is a natural consequence of local MHD turbulence evolution in a highly conductive plasma with a strong guide field, independent of additional flux-rope production at primary current sheets. Furthermore, observational evidence of merging (magnetic reconnection) of neighboring small-scale flux ropes at Earth (Khabarova et al. 2015, 2016), including evidence on the basis of Grad–Shafranov reconstruction of small-scale flux ropes (Zheng & Hu 2016; Zheng 2017; Zheng et al. 2017), is consistent with the concept of quasi-2D turbulence theory of an inverse cascade of magnetic island energy to smaller wavenumbers.


The realistic structures of the pulsar magnetosphere still remain uncertain. Knowledge about the pulsar magnetosphere structures can be used to identify the potential sites of particle acceleration and gamma-ray emission. A vacuum dipole field is generally adopted in the early study of pulsar emission, because it has an exact analytical solution given by Deutsch (1955). Based on this field structure, different theoretical models have been developed to explain the observed pulsar emission. In these models, it is widely believed that particles are accelerated in the gap region where an accelerating electric field is created because of the deficit of charges. Gamma-ray emission is produced by the curvature or inverse-Compton radiation from high-energy particles accelerated in these gaps. Due to different emission zone locations, standard pulsar radiation models include the polar cap (PC; e.g., Ruderman & Sutherland 1975; Daugherty & Harding 1982), the slot-gap (SG) (e.g., Dyks & Rudak 2003; Dyks et al. 2004; Muslimov & Harding 2004), and the outer gap (OG; e.g., Cheng et al. 1986, 2000; Zhang & Cheng 1997, 2001; Zhang et al. 2004) models. These gap models have achieved great successes in explaining pulsar high-energy emissions and light curves (e.g., Watters et al. 2009; Romani & Watters 2010). However, the vacuum solution has no plasma; it is not able to reproduce any pulsar phenomena. It is well known that the pulsar magnetosphere should be filled with plasma (Goldreich & Julian 1969). In the presence of abundant plasma, all accelerating electric fields can be efficiently screened to form a force-free (FF) magnetosphere. The FF solution for an aligned rotator was first obtained by Contopoulos et al. (1999). The CKF solution consists of a closed field line region extending to the light cylinder (LC), an open field line region, and an equatorial current sheet beyond the LC. Moreover, the time-dependent simulations for the FF axisymmetric rotator also confirmed the closed–open CKF solution (e.g., Komissarov 2006; McKinney 2006; Timokhin 2006; Yu 2011; Parfrey et al. 2012; Cao et al. 2016a; Etienne et al. 2017).


Simon & Leighton (1964) described supergranulation as a system of atmospheric currents in the photosphere. The currents form a cellular network, which is visible in Doppler maps after the reduction of other larger scale flows (e.g. differential rotation and convective blueshift). Inside the cells, the flow is radially directed from the cell centre to its boundary. The diameters of supergranules are in the range of 10 Mm up to 45 Mm (Simon & Leighton 1964; Roudier et al. 2014; Orozco Suárez et al. 2012). Depending on the method, different values are measured (Hirzberger et al. 2008). A mean horizontal velocity of 0.4 km s−1 in supergranules was measured by Simon & Leighton (1964). Orozco Suárez et al. (2012) investigated supergranular convective flows using Fourier local correlation tracking and intergranular magnetic elements. The flow velocity in supergranules increases outwards for larger distances from the centre. After reaching a maximum of 0.35 km s−1, the velocity decreases monotonically. By averaging over 222 976 supergranular cells and applying time-distant helioseismic inversions, Švanda et al. (2014) found a symmetrical flow in supergranules, which is directed radially away from the centre of the cells to its periphery, with horizontal velocities in the range of 0.3–0.6 km s−1. Their velocity profile is found to increase with increasing distances to the supergranule centre up to a maximum value from where on a continuous decrease for larger distances starts. According to Simon & Weiss (1968), the non-stationary cells only survive their turnover time, then the flow lapses into disorder. Supergranular cells have a mean lifetime of 1.5 days, which can extend up to 4 days (Roudier et al. 2014; Hirzberger et al. 2008). De Rosa & Toomre (2004) observed the creation of supergranules due to fragmentation or merging of older cells and stated that each supergranular cell takes part in a minimum of one merging or splitting event during its lifetime. This interaction seems to be the preferred mode of evolution of the observed supergranular pattern. The mechanism of advecting magnetic field elements to the network and the similar appearance suggests a relation between the supergranular pattern and the quiet-Sun magnetic network, as was recognised by Simon & Leighton (1964). De Rosa & Toomre (2004) observed a strengthening or weakening of the network lanes due to splitting or merging of supergranules. Although there is no definite evidence of a one-to-one relation (see review by Rieutord & Rincon 2010 and references therein), in the large-scale picture, the supergranular pattern can be equated with the magnetic network. The horizontal, radially outward directed gas motion of the moat flow, which is visible in Doppler maps, resembles the characteristics of supergranular flows. In addition, magnetic lanes can form around the moat cell due to magnetic features (MMFs), which cross the moat and conglomerate, resembling the magnetic network. This relation has been described by Simon & Leighton (1964), Sheeley (1972), and Vrabec et al. (1974), for example, who proposed the sunspot to be sitting in the centre of a supergranule. Meyer et al. (1974) concluded from the observations that sunspots are related to supergranular convection, but these two flows should be distinguished because of their difference in size and occurrence by taking into account the unique relation of the moat flow to its sunspot. Švanda et al. (2014) found the moat flow to be asymmetric, while flows in supergranules are symmetric, but the investigated cells approximately show the same size. In addition, the authors described the moat cell as a downflow region, while the motion in supergranules shows an upflow−downflow behaviour. Various investigations have been carried out to study the differences and similarities between the moat flow and supergranules (see e.g. Sobotka & Roudier 2007 and the reviews by Solanki 2003 and Rieutord & Rincon 2010 and references therein). An overview of the characteristics of the moat flow and supergranules is given in Table 1.


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.


The above results, when considered as relating to different aspects of the same physical process, suggest the probable existence of correlations between periods of significant Alfvénic turbulence, higher time-integrated AE and HILDCAA periods, and longer storm recovery phases. However, a statistical demonstration of the coupling between the solar wind Alfvénicity and the ring current dynamics, aimed at providing evidence in support of this possible scenario and distinguishing between a temporal coincidence and a causal linkage, is still missing in the literature and motivates the present paper, which deals with the correlation between the presence of Alfvénic streams and long-living periods of geomagnetic activity at low latitudes, after the main phases of storms induced by recurrent or non-recurrent solar transients. Specifically, by exploiting a 16 yr survey of interplanetary and geomagnetic data, an attempt is here sought to provide an empirical law that quantitatively links the duration of the recovery phase (as inferred from the SYM-H index-related magnetospheric output at low latitudes) with the Alfvénic content of the solar wind fluctuations. The goal is twofold. First, this would allow a robust establishment of the extent to which Alfvénic intervals contribute to slow recovery phases. Second, interesting predictions of the duration of the total geomagnetic storm (main plus recovery phase) would be available once the empirical law between the two periods was established. This is crucial since, although space weather forecasting often focuses on estimating the onset and intensity of geomagnetic storms (Joselyn 1995; Lundstedt et al. 2002; Abunina et al. 2013), the impacts on ground infrastructure are also related to the whole storm duration (closely linked to the storm intensity; e.g., Haines et al. 2019) through time-integrated effects (Balan et al. 2016; Lockwood et al. 2016; Mourenas et al. 2018). Indeed, the investigation of the different storm indices for low, intermediate, and high latitudes has shown that infrastructures on Earth, namely, the equipment of electric power transmission networks, are affected by the geomagnetically induced currents. In particular, transformers and some electrical substations have been observed to be very sensitive to prolonged periods of substorms, with delayed increases in anomalies (Švanda et al. 2020). Moreover, highly energetic charged particles in near-Earth space represent a risk for satellite operations. A strong relationship exists between relativistic electrons and HILDCAA events: the longer the HILDCAAs last, the higher the energy reached by relativistic electrons (Thorne et al. 2013). Hajra et al. (2015), taking into account solar cycle 23, found a flux enhancement of relativistic electrons in the outer radiation belt during HILDCAAs, corroborating the results of Meredith et al. (2002, 2003). Furthermore, there, large events are always followed by high flux peaks of 2 MeV electrons (Mourenas et al. 2019). It is worth pointing out that the investigation of the particle flux for electron energy larger than 2 MeV is extremely important for the space weather forecasting as a key indicator of serious hazard for damage of spacecraft in low, medium, and geosynchronous Earth orbits (Forsyth et al. 2020). Therefore, prolonged periods of strong Alfvénic turbulence in the solar wind could be precursors of many adverse effects in both ground installations and space satellites.


In Fig. 6, we plot the [C/O] ratios dependence on [O/H] for the different stellar populations. This figure serves to evaluate the balance between the two different elements directly with the evolution of one of them, which in this case has a well-known, single production site. Here, we chose to only show the ratios with the oxygen abundances from the 6158 Å indicator as it was shown to present less dispersion and be moretrustworthy (Bertran de Lis et al. 2015). Nevertheless, for completeness, we present the figure with the forbidden oxygen line in the appendix (Fig. A.1). Previous works in the literature have shown an increase of [C/O] ratios as [O/H] increases (e.g. Bensby & Feltzing 2006; Nissen et al. 2014; Amarsi et al. 2019), but our ratios present a quite large dispersion that together with the shorter range in [O/H] prevents us from seeing a clear behaviour. If we focus on the results with the 6158 Å line, the [C/O] ratios present a general flat trend (see the upper panel of Fig. 6 with the running mean for each population), as [O/H] increases to then clearly decrease at [O/H] ≳ 0 dex, both for thin-disk and hαmr stars. To better evaluate the significance of this apparent trend, we applied a weighted least squares fit to the [C/O] values of thin disk and hαmr stars at [O/H] ≳ 0 dex. We then used the values of the slopes and the associated uncertainties to assess the significance of the fits. The p-values come from the F-statistics that tests the null hypothesis that the data can be modelled accurately by setting the regression coefficients to zero. The resulting p-values are 3.3 e−2 and 1.6 e−2 for the thin disk and hαmr stars, meaning that the correlations are significant. This trend is in agreement with the previously mentioned turning point of [C/O] at [O/H] ~ 0.0 dex (Carigi et al. 2005). However, this is in contrast with the flattening at super-solar [O/H] presented by Nissen et al. (2014) or the increasing trend found by Franchini et al. (2021) which might be caused by the use of different oxygen indicators. On the other hand, the [C/O] ratios for thick-disk stars present no trend but show a clear offset in its running mean with respect to the thin disk as also reported by Amarsi et al. (2019). This separation between the thick and thin disk, also observed for other elements, supports the different formation episodes of both populations (Chiappini et al. 2001; Amarsi et al. 2019).


HATLAS12-00 had already been identified as a candidate gravitationally lensed galaxy as a result of its high submm flux (i.e. F500 > 100 mJy), red Herschel colours and the lack of a bright optical or radio counterpart (see e.g. Negrello et al. 2010 for a discussion of the selection of lens candidates in H-ATLAS and other Herschel surveys). This source was therefore observed spectroscopically in the submm. A CO spectroscopic redshift of 3.26 was first suggested by Z-spec (Bradford et al. 2004) observations, then subsequently confirmed by observations by the CARMA interferometer (Leeuw et al., in preparation) and the Zpectrometer instrument (Harris et al. 2007) on the Greenbank Telescope (Harris et al. 2012; see also Fu et al. 2012). Additional followup observations in the optical, near-IR, submm and other wavelengths were targeted at the lensed z = 3.26 source and the foreground objects responsible for the lensing, resulting in detailed analyses of this lensing system by Fu et al. (2012) and Bussmann et al. (2013). Their conclusions are that the z = 3.26 source HATLAS12-00 is subject to gravitational lensing, with a magnification of 9.6 ± 0.5 in both the submm continuum and CO, and 16.7 ± 0.8 in the K′ band, by two foreground galaxies, one at a spectroscopically determined redshift of 1.22, and another with photometry suggesting that it lies at a similar redshift. The submm photometry of HATLAS12-00 at 890 μm acquired with the Submillimeter Array (SMA) as part of this programme (Fu et al. 2012; Bussmann et al. 2013) is fully consistent with the 870 μm and 850 μm fluxes derived for this source from the LABOCA and SCUBA2 data to be presented here. The spectral energy distribution (SED) of the lensed source, after correcting for the lensing amplification, is well matched by the optically thick SED model for Arp220 from Rangwala et al. (2011), with a lensing-corrected far-IR luminosity of 1.2 ± 0.2 × 1013 L⊙, and an implied star formation rate of 1400 ± 300  M⊙ yr−1. In many ways the unlensed properties of this object match those of the broader population of bright submm selected galaxies first discovered by the SCUBA submm imager (see e.g. Chapman et al. 2005; Clements et al. 2008; Michałowski, Hjorth & Watson 2010). The unlensed 870 μm flux of this object would be ∼7.7 mJy.


Star formation is an inefficient process, as evidenced by observed gas depletion times,1 which are two orders of magnitude above the dynamical time, both in galaxies (e.g. Leroy et al. 2017; Utomo et al. 2018), and in individual giant molecular clouds (GMCs) (e.g. Krumholz & Tan 2007; Evans, Heiderman & Vutisalchavakul 2014; Heyer et al. 2016; Pokhrel et al. 2020; Hu et al. 2022). Theoretical models explain this inefficiency through a combination of mechanisms that provide support against gravitational collapse, including turbulence, magnetic fields, stellar feedback, and dynamical stabilization (Krumholz & McKee 2005; Ostriker, McKee & Leroy 2010; Federrath & Klessen 2012; Krumholz, Klein & McKee 2012b; Federrath 2013b; Padoan et al. 2014; Federrath 2015; Burkhart 2018; Meidt et al. 2018; Krumholz & Federrath 2019; Evans, Kim & Ostriker 2022). Recent progress in both theory and observations have highlighted the pivotal role that feedback, especially due to massive (main-sequence) stars, plays in star/star-cluster formation (Krumholz et al. 2014; Krumholz, McKee & Bland-Hawthorn 2019), and the lifecycle of GMCs (see Chevance et al. 2020, 2022a for reviews). This massive-star feedback has been suggested to be largely responsible for limiting the integrated star formation efficiency (ϵ*) to low values in typical environments, where ϵ* is given by
(1)$$\begin{eqnarray}
\epsilon _* = \frac{M_{*}}{M_{\mathrm{gas}}},
\end{eqnarray}$$which quantifies the net efficiency of star formation over the lifetime of a GMC, i.e. the ratio of the final stellar mass M* and the available gas mass in the parent molecular cloud Mgas. Feedback achieves this by (i) disrupting GMCs in order ∼ unity dynamical time-scales, through the momentum and energy carried by feedback processes (e.g. Grudić et al. 2018), and (ii) driving turbulent motions that could further provide support against collapse (e.g. Mac Low & Klessen 2004; Krumholz, Matzner & McKee 2006; Elmegreen 2009; Gritschneder et al. 2009; Federrath et al. 2010; Wibking, Thompson & Krumholz 2018; Gallegos-Garcia et al. 2020; Menon, Federrath & Kuiper 2020; Menon et al. 2021).


With our joint analysis of stellar mass fraction and source size, we find a larger stellar mass fraction than earlier statistical studies. In FigureÂ 2, we compare our determination of the stellar surface density fraction to a simple theoretical model and to the best fit of a sample of lens galaxies by Oguri etÂ al. (2014). The simple theoretical model is the early-type galaxy equivalent of a maximal disk model for spirals. We follow the rotation curve of a de Vaucouleurs component for the stars outward in radius until it reaches its maximum and then simply extend it as a flat rotation curve to become a singular isothermal sphere (SIS) at large radius (see details in the Appendix). The ratio of the surface mass density of the de Vaucouleurs component to the total surface mass density is shown as a dashed curve in FigureÂ 2. We also show as a gray band the best fit for the stellar fraction in the form of stars determined by Oguri et al (2014) in a study of a large sample of lens galaxies using strong lensing and photometry, as well as the best model using a Hernquist component for the stars and an NFW halo for the dark matter with and without adiabatic contraction, also from Oguri etÂ al. (2014). We have used the average and dispersion estimates for the Einstein and effective radii available for 13 of the objects in our sample from Oguri etÂ al. (2014), Sluse etÂ al. (2012), Fadely etÂ al. (2010), and LehÃ¡r etÂ al. (2000; see TableÂ 1) as an estimate of RE/Reff in FigureÂ 2. The average value and dispersion of the sample is RE/Reff = 1.8 Â± 0.8. This also averages over the different radii of the lensed images. The agreement of our estimates with the expectations of the simple theoretical model and with estimates from other studies (Oguri etÂ al. 2014) is quite good. For comparison, the estimate of Pooley etÂ al. (2012; using the Einstein and effective radii estimates for 10 out of 14 of their objects from Schechter etÂ al. 2014) seems somewhat lower than expected at those radii. The range of stellar mass fractions from MED09 for source sizes in the range 0.3â15.6 light days is also shown in FigureÂ 2. In this case, the discrepancy between our estimate and their reported value of Î± = 0.05 is completely due to the effect of the source size. Although accretion disk sizes are known to be smaller in X-rays, recent estimates are in the range of 0.1â1 light-days, depending on the mass of the black hole (see Mosquera etÂ al. 2013), and these finite sizes will increase the stellar surface densities implied by the X-ray data. Another possible origin for this discrepancy is that Pooley etÂ al. (2012) use the macro model as an unmicrolensed baseline for their analysis. It is well known that simple macro models are good at reproducing the positions of images, but have difficulty reproducing the flux ratios of images due to a range of effects beyond microlensing. Recently, Schechter etÂ al. (2014) found that the fundamental plane stellar mass densities have to be scaled up by a factor 1.23 in order to be compatible with microlensing in X-rays in a sample of lenses with a large overlap with that analyzed by Pooley etÂ al. (2012). It is unclear how this need for more mass in stars at the position of the images found by Schechter etÂ al. (2014) can be reconciled with the apparently low estimate of mass in stars at those radii by Pooley etÂ al. (2012). Our estimate of the stellar mass fraction agrees better with the results of microlensing studies of individual lenses (Keeton etÂ al. 2006; Kochanek etÂ al. 2006; Morgan etÂ al. 2008, 2012; Chartas etÂ al. 2009; Pooley etÂ al. 2009; Dai etÂ al. 2010) that reported values in the range 8%â25%, and with the estimates from strong lensing studies (see for example Jiang & Kochanek 2007; Gavazzi etÂ al. 2007; Treu 2010; Auger etÂ al. 2010; Treu etÂ al. 2010; Leier etÂ al. 2011; Oguri etÂ al. 2014) which produced stellar mass fractions in the range 30%â70% integrated inside the Einstein radius of the lenses.


We note in passing that recent studies address improved satellite modellimg that ameliorates many of these issues, including the core–cusp issue via non-sphericity of the stellar velocity distribution (Hayashi, Chiba & Ishiyama 2020) and the detectability of MWG satellites (Nadler et al. 2020). Other proposed solutions include those invoking baryonic physics, ranging from inclusion of baryon-contraction-induced diversity (Lazar et al. 2020), through dynamical friction-mediated coupling with baryonic clumps (El-Zant, Shlosman & Hoffman 2001; El-Zant et al. 2004; Tonini, Lapi & Salucci 2006; Romano-Díaz et al. 2008; Goerdt et al. 2010; Cole, Dehnen & Wilkinson 2011; Del Popolo et al. 2014; Nipoti & Binney 2015), or through dynamical feedback driven by starbursts or active galactic nuclei (AGNs; Read & Gilmore 2005; Mashchenko, Couchman & Wadsley 2006; Mashchenko, Wadsley & Couchman 2008; Peirani, Kay & Silk 2008; Governato et al. 2012; Pontzen & Governato 2012; Zolotov et al. 2012; Martizzi, Teyssier & Moore 2013; Teyssier et al. 2013; Madau, Shen & Governato 2014; Ogiya & Mori 2014; Pontzen & Governato 2014; El-Zant, Freundlich & Combes 2016; Silk 2017; Freundlich et al. 2020). Alternatively, modifications to the particle physics model of the dark matter have been proposed. Such proposals include ‘pre-heated’ warm dark matter (e.g. Colín, Avila-Reese & Valenzuela 2000; Bode, Ostriker & Turok 2001; Macciò et al. 2012; Schneider et al. 2012; Shao et al. 2013; Lovell et al. 2014; El-Zant, Khalil & Sil 2015) and self-interacting dark matter, whereby energy flows into the central cores of haloes through conduction (e.g. Burkert 2000; Kochanek & White 2000; Spergel & Steinhardt 2000; Miralda-Escudé 2002; Peter et al. 2013; Zavala, Vogelsberger & Walker 2013; Elbert et al. 2015). Ultralight axions, with boson mass ∼10−22 eV, have also been considered as dark matter candidates in connection with these same small (sub)galactic scale problems (e.g. Peebles 2000; Hu, Barkana & Gruzinov 2000; Peebles 2000; Marsh & Silk 2014; Schive et al. 2014b; Hui et al. 2017; Mocz et al. 2019; Nori et al. 2019; see Niemeyer 2019 for recent review). Here the zero-point momentum associated with a long de Broglie wavelength corresponding to the small mass comes along with ‘fuzziness’ in particle positions. This in turn leads to a hotter halo core with non-diverging central density and a cut-off in halo mass. Such axion fields can also be relevant for inflationary scenarios or late dark energy models. The non-thermal production implies that the axions are present with the required abundance for dark matter; they behave as cold dark matter on larger scales despite the tiny masses (Marsh 2016, 2017).


The radio versus optical DCF peaks at non-zero lags (τdelay) with the optical emissions leading the radio emissions, implying that the radio and optical emission regions are not cospatial with the optical/IR emission region being closer to the base of the jet. This is consistent with the observation that optical emission regions are smaller than radio emission regions (Table 8). These observations are in accordance with standard shock-in-jet models where higher frequencies are emitted closer to the shock front while lower frequencies are produced from larger volumes that extend farther away from the shock (e.g., Marscher & Gear 1985; Marscher et al. 2008). This could also be understood as a manifestation of the position offset of optically thick features that can be interpreted as a frequency-dependent shift of the self-absorbed core of the jet (e.g., Lobanov 1998; Pushkarev et al. 2012). The linear separation of the V and 37 GHz emission region can be estimated using the relation (Pushkarev et al. 2010; Lisakov et al. 2017):
10

where βapp is the apparent jet speed and θ is the viewing angle. Using a range of 15–100 days for τdelay from Table 6, the maximum linear separation between the emission regions (

) is estimated to be in the range of 4.06 (for Segment 6) to 27.04 pc (for Segment 1) and the corresponding projected separation varies from 0.18 to 1.20 pc using a viewing angle of θ = 13. The resulting angular separation is 0.022–0.151 mas. The simplest jet geometry is that of a conical jet. It cannot, however, explain the change in the separation between emission regions. In a conical jet geometry, the distance of an emission region from the central engine can be calculated using

 (Abdo et al. 2011) assuming the emission region fills the cross-section of the jet and the opening angle

. The obtained dce is given in Table 8 and the separation comes out to be 1.4 pc. Thus, the conical jet model also severely underestimates the separation between the emission regions, as it does not take into account the jet collimation. An alternative model is that of an inhomogeneous curved jet, where synchrotron radiation of decreasing frequency is produced in an outer and wider jet region that changes orientation with time. It is possible that the long-term variability behavior of 3C 454.3 during our extended observation is dominated by geometrical effects that also lead to temporal delays between the radio and optical bands.


In local galaxies, the ISM conditions are often described by some physical quantities such as ionization parameter (q), gaseous metallicity (Z) and electron density (ne). At high redshift, the ionization parameter is raised by a large flux of ionizing photons in ISM originated from hot O, B stars due to intensive star formation in relatively small galaxies. Previous studies suggest a high ionization parameter of SF galaxies at z > 2 compared to that of local galaxies (Erb et al. 2010; Nakajima et al. 2013; Masters et al. 2014; Nakajima & Ouchi 2014). Secondly, the chemical abundance of SF galaxies at z ∼ 2 is lower by 0.1–0.3 dex for a given stellar mass compared to those at low-z (Erb et al. 2006a; Sanders et al. 2014; Steidel et al. 2014). This leads to more compact and hotter O, B stars due to lower opacity (Ezer & Cameron 1971; Maeder 1987), and thus UV radiation becomes harder and produces more ionizing photons. Thirdly, the strength of collisionally excited emission lines (e.g. [O iii], [N ii]) strongly depends on the electron density. It is closely related to the number of electrons to collide since the excitation potential of this line is ∼1 eV, which is nearly the same as the energy of electrons at the virial temperature (∼104 K) (Dyson & Williams 1980). Due to low excitation potential, the transition of collisionally excited line is reliant on electron density compared to gaseous metallicity. Recent observations have suggested a high electron density (ne > 100 cm−3) in SF galaxies at z ∼ 2 (Newman et al. 2012; Masters et al. 2014; Shirazi, Brinchmann & Rahmati 2014; Wuyts et al. 2014). This value is larger than that of normal SF galaxies at low-z by an order of magnitude, and close to that of interacting galaxies which are seen as (ultra)luminous infrared (IR) galaxies in the present-day Universe (Krabbe et al. 2014). Such a large electron density contributes to the offset of galaxy distributions on the BPT diagram together with other physical parameters (Brinchmann, Pettini & Charlot 2008a). In this way, the cosmic dependence of the BPT diagram can be attributed to such physical parameters which determine ISM conditions.


A more equitable view of the catalogue is given by considering the distributions of band values across all haloes. In Fig. 14, we quantify the relative probability for a given integrated ($\mathcal {T}_\rm{IGM}^\rm{\,int}$, bottom panels) and maximum ($\mathcal {T}_\rm{IGM}^\rm{\,max}$, top panels) band transmission at each redshift. For reference, we also show the cumulative distribution functions as dashed curves and the median and 1σ summary statistics in the middle panels. It is extremely unlikely to have non-negligible transmission spikes in the blue bands at z ≳ 8, although this is certainly allowed (int) or even common (max) at z ≲ 6. When also considering the central band in rare cases it is plausible to witness multiple-peaked or otherwise complex spectral line profiles (e.g. see the discussions by Byrohl & Gronke 2020; Mason & Gronke 2020; Gronke et al. 2021; Park et al. 2021). This also stresses the need for systemic tracers beyond  Lyα to distinguish IGM signatures and pinpoint the origins of various spectral features. Perhaps more significant is the broad distribution in the red band. The transmission is relatively high after the midpoint of reionization z ≲ 7.67, and roughly follows the global ionized fraction (see Fig. 13). However, as will be shown later the exponential sensitivity on optical depth ($\mathcal {T}_\rm{IGM} = e^{-\tau _\rm{IGM}}$) allows the same galaxy to have both large and small values, i.e. the bimodality is not due to halo mass or environment alone. In fact, there will always be sightlines with very low transmission due to filaments and other self-shielding structures common at high-z. The location and broad nature of the higher transmission peak is redshift dependent, which is an important consideration when using LAEs as probes of reionization. We emphasize that the sharp cutoffs below $\mathcal {T}_\rm{IGM} \approx 1$ are not physical but are due to the global treatment of the long-range damping-wing absorption. For example, at z = 7 the red band values cannot exceed $\mathcal {T}_\rm{IGM} \approx 0.95$ as every sightline on average experiences at least $5{{\ \rm per\ cent}}$ absorption via cosmological integration throughout the remainder of the EoR. Finally, we note that the ultra red band is less susceptible to resonant scattering and halo proximity effects and is therefore a cleaner probe of the global state of the IGM if this can be robustly measured with deep spectroscopy for large numbers of high-z galaxies.


Douchin & Haensel (2001; DH) formulated a unified EoS for NS on the basis of the SLy4 Skyrme nuclear effective force (Chabanat et al. 1998), where some parameters of the Skyrme interaction were adjusted to reproduce the Wiringa et al. calculation of neutron matter (Wiringa et al. 1988) above saturation density. Hence, the DH EoS contains certain microscopic input. In the DH model the inner crust was treated in the CLDM approach. More recently, unified EoSs for NS have been derived by the Brussels-Montreal group (Chamel et al. 2011; Pearson et al. 2012; Fantina et al. 2013; Potekhin et al. 2013). They are based on the BSk family of Skyrme nuclear effective forces (Goriely et al. 2010). Each force is fitted to the known masses of nuclei and adjusted among other constraints to reproduce a different microscopic EoS of neutron matter with different stiffness at high density. The inner crust is treated in the extended Thomas-Fermi approach with trial nucleon density profiles including perturbatively shell corrections for protons via the Strutinsky integral method. Analytical fits of these neutron-star EoSs have been constructed in order to facilitate their inclusion in astrophysical simulations (Potekhin et al. 2013). Quantal Hartree calculations for the NS crust have been systematically performed by (Shen et al. 2011b,a). This approach uses a virial expansion at low density and a RMF effective interaction at intermediate and high densities, and the EoS of the whole NS has been tabulated for different RMF parameter sets. Also recently, a complete EoS for supernova matter has been developed within the statistical model (Hempel & Schaffner-Bielich 2010). We shall adopt here the EoS of the BSk21 model (Chamel et al. 2011; Pearson et al. 2012; Fantina et al. 2013; Potekhin et al. 2013; Goriely et al. 2010) as a representative example of contemporary EoS for the complete NS structure, and a comparison with the other EoSs of the BSk family (Chamel et al. 2011; Pearson et al. 2012; Fantina et al. 2013; Potekhin et al. 2013) and the RMF family (Shen et al. 2011b,a) shall be left for future study.


There is now renewed interest in the literature concerning these transitions, because some of the O iv and S iv intercombination lines, together with the Si iv resonance lines, are routinely observed with the Interface Region Imaging Spectrograph (IRIS; De Pontieu et al. 2014) at much higher spectral, spatial and temporal resolution than previously. For example, Peter et al. (2014) used the intensities of the O iv vs. Si iv lines to propose that very high densities, on the order of 1013 cm-3 or higher, are present in the so-called IRIS plasma “bombs”. Line ratios involving an O iv forbidden transition and a Si iv allowed transition have been used in the past to provide electron densities during solar flares and transient brightenings (e.g. Cheng et al. 1981; Hanssen 1981). However, the validity of using O iv to Si iv ratios has been hotly debated because these ratios gave very high densities compared to the more reliable ones obtained from the O iv ratios alone (see, e.g. Hayes & Shine 1987). In addition, Judge (2015) recalled several issues that should be taken into account when considering the Si iv/O iv density diagnostic. The main ones were: (1) O iv and Si iv ions are formed at quite different temperatures in equilibrium and hence a change in the O iv/Si iv ratio could imply a change in the temperature rather than in the plasma density (2) the chemical abundances of O and Si are not known with any great accuracy and could be varying during the observed events (3) density effects on the ion populations could increase the Si iv/O iv relative intensities by a factor of roughly three to four. Judge (2015) has also mentioned the well known problem of the “anomalous ions”, that is, the observed high intensities of the Li- and Na-like (as Si iv) ions (see also Del Zanna et al. 2002). Another important aspect to take into account is the effect of non-equilibrium conditions on the observed plasma diagnostics. It is well known that strong variations in the line intensities are obtained when non-equilibrium ionisation is included in the numerical calculations (see, e.g. Shen et al. 2013; Raymond & Dupree 1978; Mewe & Schrijver 1980; Bradshaw et al. 2004). In particular, Doyle et al. (2013) and Olluri et al. (2013) investigated the consequences of time-dependent ionization on the formation of the O iv and Si iv transition region lines observed by IRIS. In addition, Dudík et al. (2014) showed that non-Maxwellian electron distributions in the plasma can substantially affect the formation temperatures and intensity ratios of the IRIS Si iv and O iv lines. These authors also suggested that the observing window used by IRIS should be extended to include S iv. Recent IRIS observation sequences have indeed included the S iv line near 1406 Å. The S iv line ratios have a higher limit for density sensitivity than the O iv line ratios and are thus particularly useful for diagnosing high densities which might occur in flares. Previous flare studies have in fact reported line ratios involving O ions which lay above the density sensitivity range, indicating an electron density in the excess of 1012 cm-3 (e.g. Cook et al. 1995; Polito et al. 2016).


The search for low-mass BHs (MBH ≲ 106 M⊙) in dwarf galaxies is mostly based on the detection of X-ray emission (e.g. Greene & Ho 2007a; Desroches, Greene & Ho 2009; Reines et al. 2011; Dong et al. 2012; Schramm et al. 2013; Baldassare et al. 2015, 2017; Lemons et al. 2015; Secrest et al. 2015; Pardo et al. 2016; Chen et al. 2017), in some cases spatially coincident with jet radio emission (e.g. Reines et al. 2014; Nucita et al. 2017), or the use of standard virial techniques to estimate the BH mass (e.g. Barth et al. 2004; Greene & Ho 2004, 2007b; Peterson et al. 2005; Reines, Greene & Geha 2013; La Franca et al. 2015; Bentz et al. 2016; Onori et al. 2017; Liu et al. 2018; Chilingarian et al. 2018; see Mezcua 2017 for a review). Additional searches in the infrared (IR) regime have yielded a few more candidates (e.g. Satyapal et al. 2007, 2008, 2009, 2014; Sartori et al. 2015; Marleau et al. 2017). Most of these samples are however incomplete, very local (z  0.3), skewed towards high Eddington ratios, or skewed towards type-1 active galactic nucleus (AGN) in the case of optical searches (e.g. Greene & Ho 2004, 2007b; Reines et al. 2013) which can hamper the detection of BHs lighter than 105 M⊙ if the size of the broad-line region is controlled by BH mass (e.g. Chakravorty, Elvis & Ferland 2014). Baldassare et al. (2015) found an AGN with MBH ∼ 5 × 104 M⊙ estimated using the virial technique, Yuan et al. (2014) performed a study of four low Eddington ratio sources, and Pardo et al. (2016) searched for AGN in dwarf galaxies out to z  1. Yet, these studies include very few sources. To circumvent the biases mentioned above, in Mezcua et al. (2016) we performed an X-ray stacking analysis of ∼50 000 dwarf galaxies selected in the COSMOS field making use of the recently completed Chandra COSMOS-Legacy survey (Civano et al. 2016). We found that a population of IMBHs with X-ray luminosities ∼1039−1040 erg s−1 does exist in dwarf galaxies out to z = 1.5 and that their detection beyond the local Universe is most likely hampered by their low luminosity and mild obscuration unless deep surveys like the Chandra COSMOS-Legacy are used.


The question is, which of these proposed merger channels dominate the merger rate? Are several channels operating with a possible dependence on redshift? Or are the majority of GW sources formed through a still unknown mechanism? Several studies show that one can distinguish at least classes of channels, such as isolated binaries and dynamically induced mergers, by considering the observed distribution of merger masses (Zevin et al. 2017) or the relative spin orientation of the merging objects (Rodriguez et al. 2016c), as well as the orbital eccentricity at some reference GW frequency (Gültekin et al. 2006; Samsing et al. 2014; Samsing & Ramirez-Ruiz 2017; Samsing & Ilan 2018; Samsing et al. 2018b; Samsing 2018; Samsing et al. 2018a; Samsing & D’Orazio 2018; Rodriguez et al. 2018; Zevin et al. 2019; Samsing et al. 2019a, 2020). Other “indirect” probes have also been suggested, such as stellar tidal disruptions (e.g., Samsing et al. 2019b; Lopez et al. 2019; Kremer et al. 2019b). In this picture, it is now largely believed that dynamically assembled mergers are likely to have mass rations near one (e.g., Rodriguez et al. 2018), random relative spin orientations (e.g., Rodriguez et al. 2016c), and a nonnegligible fraction of mergers with measurable eccentricity in LISA (Samsing & D’Orazio 2018; Kremer et al. 2019c), DECIGO/Tian-Qin (e.g., Chen & Amaro-Seoane 2017; Samsing et al. 2020), and LIGO (Samsing 2018). This is in contrast to isolated binary mergers, which likely have correlated spins (e.g., Kalogera 2000), a bimodal distribution for the effective spin parameter (Zaldarriaga et al. 2018; Hotokezaka & Piran 2017; Piran & Piran 2020), larger mass ratios, and which merge on orbits with eccentricities indistinguishable from ~0 near LISA and LIGO. This picture is rather clean when comparing mergers forming in highly dynamical systems, such as globular clusters (GCs) and GNs, to completely isolated field binary mergers; however, it becomes less clean when considering, e.g., the proposed subpopulation of field binaries that undergo secular interactions with nearby single or binary objects (e.g., Naoz et al. 2013; Naoz 2016; Toonen et al. 2016; Antonini et al. 2017; Silsbee & Tremaine 2017; Liu & Lai 2018; Rodriguez & Antonini 2018; Randall & Xianyu 2018a; Antonini et al. 2018; Liu & Lai 2019; Fragione & Loeb 2019; Fragione & Kocsis 2019; Hamers & Thompson 2019; Safarzadeh et al. 2020). In this case, secular exchanges of especially angular momentum can drive the binary to merge with random spin orientations (e.g., Liu & Lai 2017) and notable eccentricities (e.g., Randall & Xianyu 2018b; Liu et al. 2019; Fragione & Kocsis 2020), which makes it more challenging to disentangle cluster mergers from field binary mergers.


In this paper, we presented our analysis of the spectral and timing behavior of GRS 1915+105 using the TCAF paradigm. For this purpose, the θ class data of the source as obtained by the LAXPC instrument of AstroSat satellite were used. To the best of our knowledge, this is the first time that the θ class data of AstroSat are analyzed with the TCAF paradigm. In this paradigm, different spectral states as well as the QPOs resulting from resonance oscillation of the Compton cloud arise out of the interplay between two types of accretion rates, namely, the Keplerian disk rate (

m

d

) and the sub-Keplerian halo rate (

m

h

). In the sub-Keplerian flow, shocks are formed when the Rankine–Hugoniot conditions are satisfied (Chakrabarti 1989, 1996; Chakrabarti & Das 2004). If the Keplerian disk rate increases, it increases the soft-seed photons, and the post-shock region (CENBOL, which acts as a “Compton cloud”) is cooled down rapidly, causing the shock to proceed toward the black hole. This has indeed been obtained in our spectral analysis. In the process, both the compressional heating timescale and the cooling timescale of the CENBOL decrease and become comparable (Chakrabarti et al. 2015), triggering shock oscillations that manifest as QPOs in the light curve. We first obtained the state of the system by estimating the photon index using the phabs*(diskbb+power-law) model. Throughout our analysis, the photon index (Γ) was found to be above 2.4, indicating that the source was either in the soft-intermediate state (SIMS) or in the soft state (SS). In the case of both orbits, Γ monotonically increases with the increase in photon flux, implying the transition from SS to SIMS. In the first orbit, Γ made an excursion from 2.5 to 3.0 during the 400 s span of spectral analysis, indicating the transition of the source from SIMS to SS. The gradual enhancement of the disk accretion rate and the consequent decline of the shock location were also obtained from the TCAF model fitting. The shock location moves from 39rs to 17rs, while the disk accretion rate increases from 0.77 to 0.84

M

Edd

. A similar trend was followed in the second orbit as well (Table 2). In the two orbits (02345 and 02346) we have analyzed, within a span of 400–500 s, the total flux changes by a factor of 3–5. The corresponding shock velocity is ∼2200 and 900 m s−1, respectively. This is significantly higher than the shock velocity obtained earlier in the case of transient sources, which lie within 10–20 m s−1 (Chakrabarti et al. 2008; Nandi et al. 2012; Debnath et al. 2013). This suggests the possibility that the local modulation of the accretion rate is due to feedback from the outflow (Chakrabarti & Manickam 2000; Chakrabarti & Nandi 2000). All these results indicate that GRS 1915+105 went through repeated microflares when it was in θ class.


One of the most astonishing characteristics of Liller 1 is the extremely large value of the collision rate parameter. Verbunt & Hut (1987) showed that Liller 1 has the second-highest value of stellar encounter rate (after Terzan 5; see also Lanzoni et al. 2010) among all star clusters in the Galaxy, thus suggesting that it represents an ideal environment where exotic objects, generated by collisions, can form. In fact, it is commonly believed that dynamical interactions in GCs facilitate the formation of close binary systems and exotic objects such as cataclysmic variables (CVs), low-mass X-ray binaries (LMXBs), millisecond pulsars (MSPs), and blue straggler stars (BSSs) (e.g., Bailyn 1992; Paresce et al. 1992; Ferraro et al. 2001, 2009a, 2012; Ransom et al. 2005; Pooley & Hut 2006). Moreover, Hui et al. (2010) found that clusters with large collisional parameters and high metallicity (see also Bellazzini et al. 1995) usually host more MSPs. Indeed, Terzan 5 hosts the largest population of MSPs among all Galactic GCs (Ransom et al. 2005).
6

6

Note that Terzan 5 is suspected to not be a genuine GC, because it harbors at least three stellar populations with different iron abundances (Ferraro et al. 2009a; Origlia et al. 2011, 2013; Massari et al. 2014).
 A strong γ-ray emission has been recently detected in the direction of Liller 1 by the Large Area Telescope (LAT) on board Fermi (Tam et al. 2011). This is the most intense emission detected so far from a Galactic GC, again suggesting the presence of a large number of MSPs. However, no direct radio detection of these objects has been obtained so far in this system (Ransom et al. 2005). The only exotic object identified in the cluster is the rapid burster MXB 1730-335, an LMXB observed to emit radio waves and type I and type II X-ray bursts (Hoffman et al. 1978). It seems to be located in the central region of Liller 1, but no optical/IR counterpart of this object has been found so far (Homer et al. 2001).