The λ Boötis stars are a group of chemically peculiar objects on the upper main-sequence, showing underabundances (~1–2 dex) of iron-peak elements and near-solar abundances of C, N, O, and S (e.g., Kamp et al. 2001; Heiter 2002; Andrievsky et al. 2002). The class was discovered by Morgan et al. (1943) and named following the bright prototype λ Boötis, which is one extreme member of the class. This group of refractory-poor objects comprises about 2% of early B through early F stars (Gray & Corbally 1998; Paunzen 2001). However, among the pre-main-sequence Herbig Ae/Be stars, that is to say the progenitors of A-type stars, the λ Boötis-like fraction is about 33% (Folsom et al. 2012). The origin of the peculiarity still remains a puzzle, as we can read in the recent discussion of Murphy & Paunzen (2017). Unlike common chemical peculiarities seen in Am and Ap stars, λ Boötis stars are not constrained to slow rotation (Abt & Morrell 1995; Murphy et al. 2015). Cowley et al. (1982) first suggested that λ Boötis could possibly originate from the interstellar medium (ISM) with a nonsolar composition, or from the separation of grains and gas. Then, Venn & Lambert (1990) proposed that λ Boötis stars likely occur when circumstellar gas is separated from grains and then accreted to the stars due to the similarity ofits abundance pattern with the ISM. Other proposed mechanisms include the interaction of a star with a diffuse interstellar cloud (Kamp & Paunzen 2002; Martinez-Galarza et al. 2009), where the underabundances are produced by different amounts of accreted material. Turcotte & Charbonneau (1993) estimated that once the accretion stops, the photospheric mixing and meridional circulation would erase this peculiar signature on a ~1 Myr timescale. By studying thedistribution of λ Boötis stars on the HR diagram, Murphy & Paunzen (2017) conclude that multiple mechanisms could result in a λ Boötis spectra, depending on the age and environment of the star.


Feedback from massive stars plays a critical role in the star formation process and evolution of molecular clouds. In particular, expanding H ii regions may have a positive effect on star formation, i.e., they can trigger a new generation of star formation in molecular clouds either by sweeping ambient clouds into dense shells or by compressing nearby dense clouds into bound clumps/cores (for details, see Deharveng et al. 2010). In both cases, dense material eventually fragments to form new stars. Since the PG108.3 cloud consists of three evolved H ii regions, one can thus speculate that the formation of young sources in the region might have been triggered by the expanding bubble. However, our observations largely do not favor such a process because in such a scenario one would expect the distribution of young YSOs or star-forming cores at the outskirts of the H ii regions (see, e.g., Zavagno et al. 2006; Jose et al. 2013; Panwar et al. 2014; Liu et al. 2016; Samal et al. 2018). In contrast, we find that no young Class I sources around S147 and most of Class I sources around S148 are cospatial with the Class II sources located near the center of the ionized gas. Since young clusters often harbor Class III to Class I sources as a part of the cluster formation process (e.g., Jose et al. 2017; Panwar et al. 2017, 2018), the Class I YSOs of the S148 are thus most likely part of the central cluster responsible for the ionization of S148. Given the fact that the H ii region S149 is located near the boundary of S148 and smaller in size, one can argue that S149 might have been influenced by S148. Unfortunately, from the current data we have no way to estimate precisely either the age of the H ii regions or the embedded YSOs in S148. Nonetheless, we searched for signatures of the early phases of high-mass star formation such as masers and outflows in the vicinity of S148, and our search resulted in no such observations. The small size of the S149 H ii regions could be due to the fact that it is ionized by a less massive star compared to S148, and as a result, S149 is expanding at a slow rate compared to S148, assuming that both of them formed in a similar environment. Considering the fact that S149 is optically visible, its dynamical age is close to the age of S148, and the ionized gas pressures of both regions are within a factor of two, we regard the probability that the formation of S149 is solely due to expansion of S148 to be rather low, although we cannot ignore the possibility that the expanding ionization front of S148 might have accelerated star formation in S149 after the initial clump formation and fragmentation. A way to prove such a hypothesis is to compare the molecular gas pressure of S149 with external gas pressure generated by S148 (e.g., Thompson et al. 2004; Kim et al. 2017; Liu et al. 2017) in the early stage of star formation; however, in S149, neither do we find significant cold gas nor do Azimlu & Fich (2011) find any molecular clumps to obtain some clue in this direction.


CCSNe are also of interest for GW astronomy as targets in their own right. As the sensitivity of GW detectors increases, they will begin to detect not only binary mergers but also other lower amplitude sources of GWs such as CCSNe. Accurate knowledge of the GW emission from CCSNe will be essential for detection and parameter estimation. The GW signal from rotational core bounce has already been well covered in the literature (e.g. Dimmelmeier et al. 2008; Abdikamalov et al. 2014; Fuller et al. 2015; Richers et al. 2017). In the non-rotating case, the GW emission from the post-bounce phase has been studied using self-consistent 3D simulations by many groups (Kuroda, Kotake & Takiwaki 2016; Andresen et al. 2017, 2019; Kuroda et al. 2017, 2018; Powell & Müller 2019, 2020; Radice et al. 2019; Andresen, Glas & Janka 2020; Mezzacappa et al. 2020; Pan et al. 2020). The structure of the GW emission has shown common features in different simulations from recent years. The dominant emission feature in the GW emission is due to the quadrupolar surface f/g mode 1 of the proto-neutron star (PNS), which produces GW frequencies rising in time from a few hundred Hz up to a few kHz (Müller, Janka & Wongwathanarat 2012; Sotani et al. 2017; Kuroda et al. 2018; Morozova et al. 2018; Torres-Forné et al. 2018, 2019). In addition, some models (Kuroda et al. 2016, 2017; Andresen et al. 2017; Mezzacappa et al. 2020; Powell & Müller 2020) exhibit low-frequency GW emission due to the standing accretion shock instability (SASI; Blondin, Mezzacappa & DeMarino 2003; Blondin & Mezzacappa 2006; Foglizzo et al. 2007). In rapidly rotating models, very strong GW emission can also occur during the post-bounce phase due to a corotation instability (Takiwaki & Kotake 2018). The emerging understanding of the GW emission features has led to the formulation of universal relations for the GW emission (Torres-Forné et al. 2019) and paved the way for phenomenological modelling for CCSN signals (Astone et al. 2018). Further work is still needed, however, to extend these models to fully explore CCSN GW signals from across the progenitor parameter space. The majority of 3D simulations that include GW emission are for progenitor stars below $30\, \mathrm{M}_{\odot }$. In this paper, we perform simulations of high-mass Population III (Pop-III) stars in the pulsational pair instability regime to expand the parameter space coverage of 3D simulations and to provide further insights into the massive and very massive star remnant BH population.


The X-ray luminosity of CG X-1 is variable by a factor of ≈10 (Bianchi et al. 2002; Weisskopf et al. 2004). The highest flux reported in the literature (5.2 × 10−12 erg cm−2 s−1 for a power-law fit or 5 × 10−12 erg cm−2 s−1 for an MCD fit, in the 0.5–8 keV band; Weisskopf et al. 2004) would imply, for a distance of 4.2 Mpc, a 0.5–10 keV luminosity of LX = (1.5–2) × 1040 erg s−1. If the system is Eddington-limited, the lower limit on the mass of the accreting BH is MBH ≳ 75 M⊙ for an He or C/O donor. For the system to shine in X-rays, the velocity of the WR star wind has to be slow enough to allow the formation of an accretion disc. This condition corresponds for CG X-1 to the requirement MBH ≳ 1.5 vw, 10004δ2 M⊙, where MBH is the BH mass, vw, 1000 is the wind velocity in units of 1000 km s−1, and δ ≈ 1 is a dimensionless parameter (adapted from Carpano et al. 2007b, see also Illarionov & Sunyaev 1975). In the simplest wind-accretion case (e.g. Edgar 2004), the luminosity can be estimated as

(1)

\begin{equation}
L_{\rm{X}}\approx \eta \frac{\dot{M}_{\rm{w}} c^2G^2M_{\rm{BH}}^2}{a^2 (v_{\rm{orb}}^2+v_{\rm{w}}^2)^2},
\end{equation}

where η is the efficiency, $\dot{M}_{\rm{w}}$ is the wind mass-loss rate, a is the orbital separation, vorb is the orbital velocity, and vw is the wind velocity at the BH orbit. Assuming $\dot{M}_{\rm{w}} = 10^{-5}$ M⊙ yr−1 and vw = 1000 km s−1 for the WR star (e.g. Crowther 2007), a = 5.8 × 1011 cm (for a 10 M⊙ companion), MBH = 75 M⊙, and the formation of a disc with η = 0.1, the corresponding luminosity is LX ≃ 2 × 1040 erg s−1. More in general, for MBH > 10 M⊙ and all the other things being equal, one finds LX ≳ 3 × 1039 erg s−1. In case of Roche lobe overflow, even higher X-ray luminosity could be achieved. However, we note that if CG X-1 is indeed a WR–BH binary, the WR star is probably not filling its Roche lobe [unless it is very massive; see for example the discussion of the case of Cyg X-3, where the orbital period is much shorter, in Szostek & Zdziarski (2008)]. An X-ray luminosity of ∼2 × 1040 erg s−1 can be therefore accounted for. We finally notice that, although we do not regard the question as crucial, the problem of the lifetime of the system discussed by Weisskopf et al. (2004) would be significantly attenuated, since the WR phase of a massive O-type star is thought to last a few ×105 yr (Meynet & Maeder 2005).


Solar vortex tubes can be spontaneously generated by turbulent convection. In simulations of quiet Sun regions, vortices are found along intergranular lanes (Shelyag et al. 2011a; Kitiashvili et al. 2012; Moll et al. 2012; Silva et al. 2020). These structures have an average lifetime of around 80 s (Silva et al. 2021) and a radius between 40 and 80 km (Shelyag et al. 2013; Silva et al. 2020). Solar kinetic vortex tubes (Silva et al. 2021) act as a sink for magnetic field, creating magnetic flux tubes that expand with height (Kitiashvili et al. 2012; Moll et al. 2012; Silva et al. 2020). The concentration of magnetic flux leads to a high magnetic field tension, which can prevent the magnetic field lines from being twisted by the rotational motion (Shelyag et al. 2011b; Moll et al. 2012; Nelson et al. 2013; Silva et al. 2021). In some cases, twisted magnetic flux tubes appear close enough to flow vortices, leading to magnetic and kinetic vortex structures closely coexisting in regions with high plasma-β (Wedemeyer & Steiner 2014; Rappazzo et al. 2019; Silva et al. 2021). The vortical motions can still trigger perturbations along magnetic lines that could lead to wave excitation, e.g., Battaglia et al. (2021). The vorticity evolution in the magnetized solar atmosphere is mainly ruled by the magnetic field, which also influences the general shape of vortices (Shelyag et al. 2011a). Based on the analysis of swirling strength, the part of the vorticity only linked to swirling motion (Shelyag et al. 2011b; Canivete Cuissa & Steiner 2020) showed that the magnetic terms in the swirling equation evolution tend to cancel the hydrodynamic terms close to the solar surface, whereas the magnetic terms dominate alone the production of swirling motion in the chromosphere. The magnetic field also tends to play an important role in the plasma dynamics along the whole vortex tube, as the Lorentz force has a magnitude comparable to the pressure gradient (Silva et al. 2020; Kitiashvili et al. 2013). High-speed flow jets have also been linked to simulated vortex tubes, driven by high-pressure gradients close to the photosphere and by Lorentz force in the weakly magnetized upper solar photosphere (Kitiashvili et al. 2013). In general, the averaged radial profile of magnetic field, angular velocity, pressure gradient inside of the vortex tube at the lower chromosphere and photosphere levels show similar behavior (Silva et al. 2020).


Since the discovery of the radio burst, there have been extensive follow-up observations of SGR J1935+2154 across the electromagnetic spectrum. The lack of another radio pulse coincident with an X-ray flare puts interesting constraints on the emission mechanism and begs the question of whether we should be able to see such radio bursts in other active Galactic magnetars. In this context, connecting FRBs with extragalactic magnetars is tantalizing. Current theories that propose FRB emission from a magnetar can be broadly divided into two categories: (1) far-away models, where the FRB is generated by a maser away from the neutron star, and (2) close-in models where the FRB is produced in the magnetospehere of the star. The maser emission model runs into difficulties when explaining all the observed radio and X-ray properties of the contemporaneous radio/X-ray burst seen from SGR J1935+2154 (see Lu, Kumar & Zhang 2020, for more details). Younes et al. (2020a) have shown that the X-ray burst contemporaneous with the radio burst was spectrally unique compared to all other burst in the activity period and it also supports it having a polar cap origin. Hence, if we assume that FRBs produced by magnetars are created in the magnetosphere close to the polar cap, we can expect them to be significantly beamed (Lu, Kumar & Zhang 2020). This of course also means that the source must still be exhibiting bursts infrequently in the radio and that more of them might be associated with an X-ray burst than we observe. Observational evidence so far does suggest a connection between the X-ray and the radio emission mechanisms prevalent in neutron stars and any changes in one of them affects the other (Archibald et al. 2017). It is believed that while the X-ray bursts and radio pulsations in these sources come from different regions in the magnetosphere, the pair plasma causing the X-ray flares can affect the acceleration of radio-emitting particles. This was shown in PSR J1119−6127, a high B-field radio pulsar where a series of X-ray bursts from the source quenched the radio pulsations with the radio emission returning a few minutes after the last X-ray burst (Archibald et al. 2017).


Einstein’s theory has been tested successfully mainly in the regime of weak gravity through Solar system tests and laboratory experiments. But the validity of this highly accepted theory still faces stringent constraint in the regime of strong gravity, viz., the region near to a black hole, ultradense compact stars, and expanding universe. The recent discovery of peculiar highly overluminous SNeIa, e.g. SN 2003fg, SN 2006gz, SN 2007if, and SN 2009dc (Howel et al. 2006; Scalzo et al. 2010) indicates a huge Ni-mass and confirms the highly super-Chandrasekhar white dwarfs, having mass 2.1–2.8 M⊙, as a suitable progenitors (Howel et al. 2006; Hicken et al. 2007; Yamanaka et al. 2009; Scalzo et al. 2010; Silverman et al. 2011; Taubenberger et al. 2011). Recently, Linares, Shahbaz & Casares (2018) have discovered a highly massive pulsar of mass $2.27^{+0.17}_{-0.15}~\mathrm{M}_{\odot }$ in their observation of compact binaries PSR J2215+5135. Clearly, these observations are not only questioning the standard Chandrasekhar limit for the compact stellar objects but also invoking the necessity of modification of GR in the strong gravity regime. Interestingly, our study reveals that due to the effect of $f\left(R,\mathcal {T}\right)$ gravity theory, the maximal mass limits rise higher than their standard values in GR for the chosen parametric values of χ. Hence, the stellar systems in the framework of $f\left(R,\mathcal {T}\right)$ gravity theory may also explain the observed massive stellar systems, viz., massive pulsars, super-Chandrasekhar stars, and magnetars, etc., which GR hardly can explain suitably so far. In support of the achieved result in the present investigation, it is worth mentioning that the important study by De Laurentis (2018) also reveals that the application of the Noether Symmetry Approach can explain the extreme massive stars which are supposed to be unstable in the framework of GR. However, in the limit χ = 0, one may retrieve the solutions of the standard Einstein gravity.


Solar cycle predictions are needed to plan long-term space missions and are of high importance for space weather applications. Currently, precursor methods are the most favored models for the prediction of solar cycle strength (Conway 1998; Svalgaard et al. 2005; Kane 2008; Hathaway 2009). These precursor techniques often use geomagnetic activity levels near or before the time of solar cycle minimum (Ohl & Ohl 1979; Feynman 1982; Gonzalez & Schatten 1988; Thompson 1993; Wilson et al. 1998). Predicting the amplitude of a solar cycle can be done using solar polar magnetic fields from the previous cycle as “precursors” of the next cycle (Schatten & Sofia 1987; Schatten 2005; Svalgaard et al. 2005; Wang & Sheeley 2009; Muñoz-Jaramillo et al. 2012). The other class of precursor techniques that do not need an a priori physical understanding of the causal relations (i.e., that do not require any knowledge of the physics involved) is based on finding particular sunspot number characteristics that serve as indicators of the next cycle strength (Ramaswamy 1977; Lantos 2006; Cameron & Schüssler 2008; Brajša et al. 2009). A number of techniques are used to predict the amplitude of a cycle during the time near and before sunspot minimum (Hathaway 2010). The two precursor types that have received the most attention are polar field precursors and geomagnetic precursors. The strength of the solar polar magnetic fields at solar minimum is a very accurate indicator of the maximum amplitude of the following solar cycle. Forecasts using the polar field method have proven to be consistently in the right range for cycles 21, 22, and 23 (Schatten & Sofia 1987; Schatten et al. 1996). The polar fields reach their maximal amplitude near the minima of the sunspot cycle. However, the maxima of the polar field curves are often rather flat, so approximate forecasts are feasible several years before the actual minimum (Hathaway 2010; Petrovay 2010). Using the rather flat and low maximum in polar field strength, Svalgaard et al. (2005) have been able to predict a relatively weak cycle 24. Such an early prediction is not always possible: early polar field predictions of cycles 22 and 23 had to be corrected later and only forecasts made shortly before the actual minimum finally converged (Hathaway 2010; Petrovay 2010).


To confirm the importance of VO, water and an inversion layer (Evans et al. 2018; Mikal-Evans et al. 2019, 2020) obtained repeated HST observations of the transmission spectrum and the secondary eclipse using the STIS and WFC 3 instruments. The optical transmission spectrum displays rich variation, with multiple features consistent with VO absorption that Evans et al. (2018) could reproduce by assuming an isothermal T-P profile at 1500 K and a metallicity equivalent to 10× to 30× solar. Absorption bands of TiO appeared to be muted in the transmission spectrum, which was explained by Evans et al. (2018) as evidence of condensation of Ti-bearing species, which commences at higher temperatures than condensation of V-bearing species, producing for example, calcium titanates (Lodders 2002) while VO remains in the gas phase. Mikal-Evans et al. (2019) observed the day side emission spectrum with the G102 grism of WFC3 (0.8–1.1 μm), augmenting their earlier observations with the G141 grism. The G102 spectrum does not show the VO bands expected to be present there, and this led Mikal-Evans et al. (2019) to question the interpretation that the 1.2 μm feature is caused by VO emission. The secondary eclipse was observed at 2 μm (Kovács & Kovács 2019) and at optical wavelengths with the TESS instrument. These were analysed together with the preceding Hubble, Spitzer, and ground-based observations to yield tighter constraints on the atmospheric structure, composition and overall system parameters (Bourrier et al. 2020a; Daylan et al. 2019). These studies found that the hottest point on the day side exceeds a temperature of 3000 K, that the atmosphere is inverted on the day side, and a metallicity that is consistent with solar (Bourrier et al. 2020a) or slightly elevated (Daylan et al. 2019). Although the chemical retrievals follow different strategies (equilibriumversus free-chemistry), both indicate that a depletion of TiO relative to VO is needed to explain the observed emission spectrum, supporting the earlier findings by Mikal-Evans et al. (2019). Recently, Mikal-Evans et al. (2020) obtained new secondary-eclipse observations using the G141 grism of WFC3. Although confirming the presence of emission by H2O, a joint analysis with their previous WFC3 observations did not reproduce the emission feature at 1.2 μm, prompting the authors to entirely discard their previous interpretation of emission caused by VO.


Galaxies seem to retain a memory of their spin orientation with respect to the cosmic web filaments and walls, as suggested by the results from large-scale cosmological hydrodynamical simulations (Dubois et al. 2014; Codis et al. 2018; Wang et al. 2018; Ganeshaiah Veena et al. 2019; Kraljic, Davé & Pichon 2020b). The mass dependence of the spin alignment signal is however debated. While some works confirmed the existence of a galaxy spin transition from parallel to perpendicular with respect to the filament’s direction (Dubois et al. 2014; Codis et al. 2018; Kraljic et al. 2020b), and analogously with respect to walls (Codis et al. 2018; Kraljic et al. 2020b), others (Ganeshaiah Veena et al. 2019; Krolewski et al. 2019) found preferential perpendicular orientation with respect to filaments at all masses with no sign of a spin transition. A possible interpretation of this lack of detection of a clear transition is the nature of the filaments, with galaxies in thinner filaments having their spins more likely perpendicular to the filament’s axis, compared to galaxies of similar mass in thicker filaments (Ganeshaiah Veena et al. 2019). This can be in turn understood recalling the multiscale nature of the problem and the conditional TTT (Codis et al. 2015) predicting larger transition mass for denser, thus thicker, filaments. Further support for this interpretation was provided by the findings of stronger impact of large-scale tides on the galaxy spin orientation in denser filaments (Kraljic et al. 2020b, using filament density as a proxy for the thickness of filaments). In addition to the stellar mass, the spin-filament alignment was shown to depend on other internal properties of galaxies. Blue or rotation-supported galaxies were found to dominate the alignment signal at low stellar mass, while red or dispersion-dominated galaxies tend to show a preferential perpendicular alignment (Codis et al. 2018; Wang et al. 2018; Kraljic et al. 2020b).


In the southern sky, the large background of atmospheric muons reduces the efficiency to select through-going tracks induced by neutrinos below the PeV regime. A very large fraction of the aforementioned background can be rejected by imposing an active veto at the detector boundary, as, for example, used in Aartsen et al. (2013a). This reduces the detector volume to a smaller fraction of the instrumented volume sacrificing statistics for signal purity. Furthermore, the more clearly an event is identified as a starting track, the more probable it is to be an astrophysical rather than atmospheric background. Down-going atmospheric neutrino events at high energy are likely to be accompanied by muons produced in the same cosmic-ray shower that triggers the veto and reduces the atmospheric neutrino background (Schönert et al. 2009; Gaisser et al. 2014). In analyses using veto techniques (Aartsen et al. 2013a, 2015b), the selection is usually more efficient for cascade-like events than tracks, and high astrophysical neutrino purity demands neglecting energies below 60 TeV, where backgrounds are more abundant. In searches for point-like sources of astrophysical neutrinos, track-like events are of great importance given their good angular resolution compared to cascade-like events. Furthermore, the purity demands are lower since the signal of a point-like source is restricted to a small portion of the sky, hence reducing the background significantly. Consequently, the minimum required total charge deposited in the PMTs of the IceCube detector by an event is lowered to 1500 p.e. compared to 6000 p.e. (Aartsen et al. 2013a), resulting in a higher signal efficiency at lower energies. In addition, only down-going tracks are used, and cuts are imposed that select well-reconstructed track-like events (Aartsen et al. 2016a). For

 events at energies smaller than 200 TeV, the effective area of the analysis is bigger than for vertically through-going tracks (δ  −30°; Figure 1). For energies up to 1 PeV, the effective area is smaller, but a higher purity is achieved. The angular resolution for starting tracks is shown in Figure 2 (dashed) and is ∼1° in the interesting energy region; the reconstruction is worse than for through-going events (solid), due to a smaller lever arm for tracks starting within the fiducial volume of the detector.


The leading FRB source model invokes magnetars as the power source to produce repeating bursts. There are two versions of this model. One version invokes rapidly spinning young magnetars that are produced in extreme stellar transients such as GRBs and SLSNe. The main motivation is that the host galaxy of FRB 121102 resembles those of LGRBs and SLSNe (Metzger et al. 2017; Nicholl et al. 2017; Wadiasingh & Timokhin 2019). The fact that the hosts of all other FRBs do not resemble that of FRB 121102 disfavors the simplest version of this proposal. A possible fix of this proposal is to introduce rapidly spinning magnetars born from binary neutron star (BNS) mergers (Margalit et al. 2019; Wang et al. 2020). In order to make this scenario work, one needs to require that rapidly spinning magnetars made from BNS mergers should be much more abundant than those made from LGRBs and SLSNe. Comparing the event rate densities of BNS mergers, LGRBs, and SLSNe (e.g., Sun et al. 2015; Abbott et al. 2017; Nicholl et al. 2017), this may be possible if a significant fraction of BNS mergers leave behind stable neutron stars (e.g., Gao et al. 2016). However, if this fraction is very low, as required if GW170817 leaves behind a black hole (Margalit et al. 2019), the fast magnetar model may fail to explain the small fraction of LGRB/SLSN-like hosts in FRB samples. The second version of the magnetar model invokes emission (e.g., giant flares) from slowly rotating magnetars like the ones observed in the Galaxy (e.g., Popov & Postnov 2010; Katz 2014; Kulkarni et al. 2014). The births of these magnetars do not require extreme explosions such as GRBs and SLSNe (e.g., Beniamini et al. 2019). If this is the case, the host galaxy distribution may be more analogous to that of SNe II. All FRBs but FRB 121102 are consistent with this scenario (Figure 4). In order to interpret FRB 121102, the more extreme channel of forming rapid magnetars is still needed. So we conclude that the magnetar model would work only if both fast magnetars produced in extreme explosions and slow magnetars produced in regular channels (Beniamini et al. 2019) can produce FRBs. In any case, since the birth rate of these magnetars is very high (Beniamini et al. 2019), an additional factor is needed to select a small fraction of magnetars to produce FRBs (e.g., Ioka & Zhang 2020).


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.


A number of semi-analytic models (SAMs) have attempted to reproduce submm number counts (e.g. Granato et al. 2000; Fontanot et al. 2007; Somerville et al. 2012). One such model is the galform (SAM), which has been tuned to successfully reproduce the number counts of 850 μm and $\mathrm{1.1 \, mm}$ selected galaxies.1 However, in order to achieve this good agreement galform invokes a top-heavy initial mass function (IMF). Early versions of the model used a flat IMF above $1 \, \mathrm{M_{\odot }}$, in sub-L* mergers (Baugh et al. 2005; Swinbank et al. 2008). This is required to produce sufficiently bright submm emission during frequent low-mass merger events. Later versions of the model used a more moderately top-heavy IMF in starbursts, triggered by disc instabilities rather than mergers, and found similarly good agreement with the number counts (Cowley et al. 2015, 2019; Lacey et al. 2016; Park et al. 2016). However, such IMF variability is still controversial, particularly extreme forms and any dependence on merger state (Bastian, Covey & Meyer 2010; Hopkins 2013; Krumholz 2014), and is inconsistent with the constraints on the IMF in massive star-forming galaxies that are significantly less extreme (e.g. Tacconi et al. 2008), though there is tentative evidence of a bottom-light/top-heavy IMF in both local star-forming region analogues (Motte et al. 2018; Schneider et al. 2018) and some gravitationally lensed high-redshift starbursts (Zhang et al. 2018). Safarzadeh, Lu & Hayward (2017) showed that a variable IMF is degenerate with a number of other modelling processes in SAMs, such as the form of stellar feedback. They highlight that taking in to account dust mass allows for a good fit to the number counts without resorting to a variable IMF. Most recently, the shark SAM (Lagos et al. 2018) is able to broadly reproduce the 850 μm counts (whilst slightly overestimating the bright-end counts compared to S2CLS; Geach et al. 2017) using a fixed Chabrier (2003) IMF (Lagos et al. 2019). They attribute the good agreement to their use of physically motivated attenuation curves obtained from a self-consistent galaxy evolution model (eagle; Trayford et al. 2020).


The pulsational analysis presented in this work makes use of full stellar evolution models of pre-WDs generated with the LPCODE stellar evolution code. LPCODE computes in detail the complete evolutionary stages leading to WD formation, allowing the WD and pre-WD evolution to be studied in a consistent way based on the evolutionary history of progenitors. Details of LPCODE can be found in Althaus et al. (2005, 2009, 2013) and references therein. Here, we mention only those ingredients employed which are relevant for our analysis of low-mass, He-core WD and pre-WD stars (see Althaus et al. 2013, for details). The standard Mixing Length Theory (MLT) for convection in the version ML2 is used (Tassoul et al. 1990). In this prescription, due to Bohm & Cassinelli (1971), the parameter α (the mixing length in units of the local pressure scale height) is set equal to 1, while the coefficients a,b,c that appear in the equations for the average speed of the convective cell, the average convective flux, and the convective efficiency (see Cox 1968), have values a = 1,b = 2,c = 16. We emphasize that the results presented in this work are insensitive to the prescription of the MLT employed. In particular, we have also used the ML1 (α = 1,a = 1 / 8,b = 1 / 2,c = 24, Böhm-Vitense 1958) and ML3 (α = 2,a = 1,b = 2,c = 16, Tassoul et al. 1990) recipes, and we obtain the same results than for ML2. The metallicity of the progenitor stars has been assumed to be Z = 0.01. It is worth mentioning that the pulsation stability results presented in this paper do not depend on the value of Z2. Radiative opacities for arbitrary metallicity in the range from 0 to 0.1 are from the OPAL project (Iglesias & Rogers 1996). Conductive opacities are those of Cassisi et al. (2007). The equation of state during the main sequence evolution is that of OPAL for H- and He-rich compositions. Neutrino emission rates for pair, photo, and bremsstrahlung processes have been taken from Itoh et al. (1996), and for plasma processes we included the treatment of Haft et al. (1994). For the WD regime we have employed an updated version of the Magni & Mazzitelli (1979) equation of state. The nuclear network takes into account 16 elements and 34 thermonuclear reaction rates for pp-chains, CNO bi-cycle, He burning, and C ignition. Time-dependent diffusion due to gravitational settling and chemical and thermal diffusion of nuclear species has been taken into account following the multicomponent gas treatment of Burgers (1969). Abundance changes have been computed according to element diffusion, nuclear reactions, and convective mixing. This detailed treatment of abundance changes by different processes during the WD regime constitutes a key aspect in the evaluation of the importance of residual nuclear burning for the cooling of low-mass WDs.


We used the stellar population synthesis code STARLIGHT (Cid Fernandes et al. 2004) to describe the age distributions and metallicities of the stellar populations that fit the integrated light spectrum of object I. For object H it is not possible to perform the stellar population synthesis because of the low signal-to-noise ratio of the spectrum. This code is extensively discussed in Cid Fernandes et al. (2004; 2005), and it is built upon computational techniques that originally were developed for empirical population synthesis with additional ingredients from evolutionary synthesis models. This method was used by Krabbe et al. (2011) and Faúndez-Abans et al. (2012, 2013) and has been successful in describing the stellar population in interacting galaxies. Briefly, the code fits an observed spectrum with a combination of N single stellar populations (SSPs) from the models of Bruzual & Charlot (2003). These models are based on a high-resolution library of observed stellar spectra, which allows for detailed spectral evolution of the SSPs across the wavelength range of 3200−9500 Å  with a wide range of metallicities. We used the Padova 1994 tracks, as recommended by Bruzual & Charlot (2003), with the Salpeter initial mass function (Salpeter 1955). Extinction is modeled by STARLIGHT as due to foreground dust, using the Large Magellanic Cloud average reddening law of Gordon et al. (2003) with RV = 3.1, and parametrized by the V-band extinction AV. The SSPs used in this work cover 15 ages, t = 0.001, 0.003, 0.005, 0.01, 0.025, 0.04, 0.1, 0.3, 0.6, 0.9, 1.4, 2.5, 5, 11, and 13 Gyr, and three metallicities, Z = 0.2Z⊙, 1 Z⊙, and 2.5 Z⊙, adding to 45 SSP components. The fitting is carried out using a simulated annealing plus Metropolis scheme, with regions around emission lines and bad pixels excluded from the analysis. The upper panel of Fig. 4 shows the observed spectrum corrected for reddening and the model stellar population integrated light spectrum of object I. The synthesized spectrum fits the observed one closely. The results indicate that this galaxy is an old object, with 100% of the flux contribution at λ5870 Å being provided by a stellar population of 13 Gyr, with Z = 2.5Z⊙ and AV = 0.36. This spectrum is dominated principally by the bulge signal and does not represent the whole galaxy; the age estimate then is valid only for its bulge.


Molecular line ratio diagnostics are often used to investigate the physics and chemistry of the ISM in all of these environments. For example, as the gas chemistry located in the central, nuclear regions of galaxies is believed to be dominated by X-rays produced by the AGN, in so-called X-ray dominated regions (XDRs), the molecular content of the ISM surrounding such nuclei will greatly differ from that in starburst regions (Usero et al. 2004; García-Burillo et al. 2010). Hence, line ratios of specific molecules have been proposed as indicators of certain energetic or physical processes, for example HCN/HCO+ as a tracer of AGNs (Loenen et al. 2007), HCN/HNC as a mechanical heating tracer (Hacar et al. 2020), and HCN/CO as a density tracer (Leroy et al. 2017). In particular, the “submillimeter-HCN diagram”, first proposed in Izumi et al. (2013) and later expanded upon in Izumi et al. (2016), is a very notable example of the use of molecular line ratios as a probe of AGN-galaxies; this diagram makes use of two line ratios, HCN(4−3)/HCO+(4−3) and HCN(4−3)/CS(7−6), where all of the molecules involved are considered tracers of dense gas. Izumi et al. (2016) observed a clear trend that AGNs, including the Seyfert composite galaxy NGC 1068, tend to show higher HCN/HCO+ and/or HCN/CS than in SB galaxies as long as the observations were at high enough spatial resolutions to separate energetically discrete regions. Izumi et al. (2016) propose a scenario where it is the high temperature that is responsible for the HCN enhancement, whereby neutral-neutral reactions with high reaction barriers are enhanced (Harada et al. 2010), thus leading to the possible enhancement of HCN and the depletion of HCO+ via newly available formation and destruction paths, respectively. However, while of course higher gas temperatures are expected in AGN-dominated regions, these are not unique to these environments, as starburst regions and/or regions where outflows dominate can also harbour high enough temperatures for such enhancement to occur. Additionally, the higher temperatures could increase HCN excitation, relative to HCO+ and CS, without necessarily changing their relative abundances (Imanishi et al. 2018a). Finally, infrared radiative pumping is also a possible explanation of the HCN intensity enhancement relative to HCO+ and CS. Infrared pumping is a result of the emission of 14 μm infrared photons due to the presence of hot dust around AGN. These photons vibrationally excite HCN to the ν2 = 1 state. Upon de-exciting from this state back to the vibrational ground state, ν = 0, the HCN line intensities are thus pumped to higher fluxes (Imanishi et al. 2018a). However, we note that it is also not unlikely that the 12 μm infrared photons can similarly vibrationally excite HCO+, thus nullifying the extent of this effect (Imanishi et al. 2016).

Background Background Motivation Motivation Motivation


The (sub)mm wavelength range is of particular interest for various reasons: at sufficiently long wavelengths, disks become optically thin, and an estimate of their dust mass can be directly obtained (via some assumptions) by simply measuring their flux (e.g., Beckwith et al. 1990). Although the bulk of the disk mass in the system is in gaseous phase, fiducial (or measured, when available) gas-to-dust ratios provide an indirect estimate of the total mass in the disk. This is a crucial parameter for planet formation theories because it determines the available reservoir for this process. Using this method, surveys of star-forming regions with (sub)mm facilities such as SMA and ALMA have determined that protoplanetary disks have typical masses of 0.1%–0.5% of that of their host star (e.g., Andrews & Williams 2005; Andrews et al. 2013; Pascucci et al. 2016). On the other hand, dust growth represents the initial stage of planet formation; the observed spectral index at these wavelengths can be linked to the dust opacity in the disk, Informative of its properties and grain sizes (e.g., Miyake & Nakagawa 1993; Draine 2006). In fact, the comparison of the millimeter spectral index of the interstellar medium (ISM) with that of protoplanetary disks has already revealed significant dust growth in these disks, implying the presence of mm/cm-sized grains in many of them (e.g., D’Alessio et al. 2001; Lommen et al. 2010; Ricci et al. 2010a, 2010b; Ubach et al. 2012). The combination of the mm spectral index with additional information at other wavelengths, such as the spectral index at near/mid infrared (IR) wavelengths or silicate features may also point to links between the evolution of the inner and outer regions of the disks. As an example, Lommen et al. (2010) identified a tentative correlation between the strength of the 10 μm silicate feature and the 1–3 mm spectral index for a sample of T Tauri and Herbig Ae/Be stars, suggesting a connection between the evolution of the inner and outer regions of disks, although a later study by Ricci et al. (2010b) found no signs of such a correlation for disks in the Taurus and Ophiuchus star-forming regions. Despite the obvious interest of this wavelength regime, disks have relatively weak emission at millimeter wavelengths and many of them currently lack this type of data (or, at least, sufficient observations to provide robust estimates of their spectral indices).


The period candidates of other three ULXs may range from ∼100 to ∼600 d. Apart from noise and artefacts, all the candidate periods are only significant in a specific epoch. This suggests that they are not associated with any stable mechanism such as orbital motion. Instead, such long-term (> 100 d) X-ray quasi-periodic variations are likely related to superorbital periods that are thought to be due to radiation-driven warping of accretion discs (Ogilvie & Dubus 2001) or tidal interaction-induced disc precession (Whitehurst & King 1991). Alternatively, mass transfer rate-related events such as X-ray state changes and disc instability can also cause long-term modulations (Kotze & Charles 2012). In particular, there are two intermittent quasi-periodicities for both NGC 5408 X-1 and M81 X-6, suggesting that the quasi-periods are changing or evolving. They resemble some Galactic X-ray binaries that show similar behaviour (e.g. Cyg X–2 and SMC X–1; Kotze & Charles 2012) and it has been suggested that a warped disc could lead to an unstable steadily precessing disc, causing quasi-periodic behaviour (Ogilvie & Dubus 2001). We note that there are many uncertainties on the physical parameters of ULXs. To determine the origin of superorbital periods of ULXs, one has to know at least the mass ratio between the companion and the compact star (q = MC/MX) and the binary separation. Unfortunately, it is very difficult to get these parameters for ULXs. For the three ULXs discussed here (i.e. excluding ESO 243-49 HLX-1), only M81 X-6 has better constraints on the black hole mass and the nature of the companion. The masses of the black hole and companion star are estimated (MX = 18 M⊙, MC = 23 M⊙) such that q can be derived. In this case, we can rule out a tidal interaction-induced disc precession scenario that requires q  0.25–0.33 (Whitehurst & King 1991). For a warped disc, the binary separation and the mass ratio suggest that M81 X-6 lies in the intermediate instability zone for radiation-driven warping in X-ray binaries (see fig. 1 of Kotze & Charles 2012). The quasi-periodic variability may represent the switching time-scale between a warped disc and a flat disc.


To minimize the undesired effects caused by observing different lunar locations, we conducted observations according to the following procedure. On each observing night, we first pointed the Nayuta telescope toward the crater Grimaldi (selenographic coordinate: 68.6°W, 5.2°S) in the waxing phase and the crater Neper (84.5°E, 8.8°N, east of Mare Crisium) in the waning phase, after correcting the pointing error measured using a nearby star. Both craters are near the lunar edge (distances ≲2′). Then, we scanned the Moon along the RA axis until the edge of the Moon was placed near the center of the field of view (FOV). An example of the observed (and reduced) images is shown in Fig. 2. Our target locations are not on a major maria and near sites repeatedly observed in previous Earthshine photometry because they were expected to have roughly comparable albedos (Qiu et al. 2003; Pallé et al. 2004; Montañés-Rodríguez et al. 2007). Half of the FOV was reserved for the sky, which allows the sky background intensities and their positional gradients to be measured. The position angle of the instrument (ϕinspa) was maintained at 90° from the equatorial north, as measured counter-clockwise, so that the long side of the FOV was aligned with the RA axis. Telescope tracking was conducted in accordance with the sky motion of the Moon, which was calculated at the Jet Propulsion Laboratory (JPL) Horizons system1. Because the tracking was not perfect, we shifted the telescope east or west with a typical interval of ~ 30 min so that the lunar edge remained near the center of the FOV. Features on the Moon were hardly recognizable in the raw images because of the dim Earthshine and strong scattered light from the day side of the Moon, though we were able to visually identify the lunar edge in most cases2. Despite our efforts, the actually observed location may have varied night by night even within one phase (waxing phase or waning phase), or on an hourly timescale during a single night. Possible impacts induced by different lunar locations (namely different degrees of depolarization) are discussed in Sect. 4.1.


Currently, about 20 Be stars are known to have stripped companions, most of which were confirmed with far-UV (FUV) spectroscopy, as it is in the FUV where the flux ratios are most favorable (Wang et al. 2021; Klement et al. 2022, and references therein). All of the spectroscopic FUV detections were found to be compatible with the sdO nature of the companions, while no firm case of a cooler sdB companion has been presented as of yet. There are also more than 160 confirmed and candidate Be X-ray binaries (BeXRBs; Raguzova & Popov 2005),
12

12
List updated at http://xray.sai.msu.ru/~raguzova/BeXcat/. which are mostly Be+NS systems that probably evolved in a similar fashion as the Be+sdO binaries but from more massive progenitor systems (Reig 2011). The Be primaries occupy a narrow range between O9 and B2 in spectral type (Reig et al. 2017). These systems are conspicuous due to the X-ray emission resulting from the (episodic) accretion of Be disk material onto the compact object, so that the sample is drawn from a much larger (partly extragalactic) volume. The WD companions to Be stars proved unexpectedly elusive (Meurs et al. 1992; Cohen et al. 1997), but several supersoft X-ray emission sources consistent with (early-)Be+WD systems undergoing a type II BeXRB outburst were recently detected in the Magellanic Clouds (Coe et al. 2020; Kennea et al. 2021). One Be star (MWC 656) was reported to have a BH companion, but this was recently shown to be questionable on the basis of new higher-quality spectra, which rather point toward another Be+sdO system (Rivinius et al. 2022). Object HD 93521 is the first candidate postmerger Be star (Gies et al. 2022), and members of a Be star runaway population were found using Hipparcos and Gaia astrometric catalogs (Berger & Gies 2001; Boubert & Evans 2018; Wang et al. 2022). On the other hand, Be stars have been found missing among B stars at high Galactic latitude (Martin2004, 2006). Meanwhile, not a single early-type Be star has been confirmed to have a close MS companion (Gies 2000; Bodensteiner et al. 2020). However, the recently studied case of the B6Ve star α Eri, which is a highly eccentric binary with an early A-type dwarf companion on a 7 yr orbit, appears to be the first confirmed case of a Be star that does not owe its nature to mass transfer in a close binary, as the presence of a close stripped companion was ruled out (Kervella et al. 2022b). This implies that two evolutionary channels—single and binary—indeed exist for the formation of Be stars.


During a GF, a huge amount of the magnetic energy E > 1046 erg is subsequently released at the surface of the magnetar within the first 1 s, leading to the formation of a hot fireball similar to the case of classic GRBs (Mészáros & Rees 2000), which might be rich in electron–positron pairs (Thompson & Duncan 1995; Fermi-LAT Collaboration 2021). The afterglow emission of a GRB is generally well explained by the synchrotron emission from electrons accelerated by the shock produced during a relativistic ejecta colliding with an external medium. The ejecta might be formed by the baryon-rich crust of the NSs undergoing a large-scale disruption, or a dense outflowing gas of electron–positron pairs and radiations (Thompson 2021). Zhang et al. (2020) prefer a high fraction of electron energy for GRB 200415A, which is consistent with the situation of an electron–positron pair-dominating ejecta. Chand et al. (2021) suggest that a baryon-poor outflow can explain the high-energy afterglow emission of GRB 200415A, while a baryon-rich outflow is also viable if the dissipation happens below the photosphere via internal shocks. Therefore, in this paper, we calculate two situations that the ejecta is dominated by electron–positron pairs or protons, respectively. Around the magnetar, a cavity might be created by the pulsar wind and the earlier SGR activity. Holcomb et al. (2014) suggest that the Poynting flux from a pulsar can evacuate a cavity with a size of sub-parsec. Such a cavity is also needed to explain the radio afterglow of the GF from SGR 1806-20 (Gaensler et al. 2005; Gelfand et al. 2005; Taylor et al. 2005; Granot et al. 2006). Fermi-LAT Collaboration (2021) suggests a cavity environment to explain the expansion of the ejecta before its colliding with the bow-shock shell for GRB 200415A. Calculations in Zhang et al. (2020) on the afterglow of GRB 200415A also suggest a low density of ambient medium, which is consistent with the cavity environment. In this paper, we discuss two situations for the CSM density, including a density of $n = 10^{-2}\ \rm cm^{-3}$, and a low density of $n = 10^{-5}\ \rm cm^{-3}$ for the situation of a cavity environment surrounding the magnetar.


For a star that moves with respect to its surrounding medium, the stellar motion adds an asymmetry to the wind velocity profile, since different parts of the wind face the ISM with different relative velocities. If the motion is supersonic, a bow shock arises at the interface where the ram pressure of the ISM and the stellar wind balance. The stand-off distance, i.e. the distance of the star to the apex of the shock front, is given by (1)\begin{equation} \label{R0_eq} R_0=\sqrt{\frac{\dot{M} v_w}{4 \pi \rho_0 v_*^2}}, \end{equation}R0=Ṁvw4πρ0v∗2,where vw is the terminal wind velocity, v∗ the stellar velocity with respect to the ISM, Ṁ the mass-loss rate, and ρ0 the density of the surrounding medium (Baranov et al. 1971). The density can be expressed in number density of hydrogen atoms (mH = 1.6727 × 10-27 kg), which follows roughly (2)\begin{equation} \label{density} n_{\rm H} = 2.0 \; {\rm e}^{-\frac{|z|}{100\,{\rm pc}}}, \end{equation}nH=2.0e−|z|100 pc,where z is the galactic height (Mihalas & Binney 1981) and nH is given in atoms per cm3. Wilkin (1996) demonstrated that the shape of the bow shock only depends on the stand-off distance, while Cox et al. (2012) showed that this assumption remains valid for viewing angles up to 70°. Above this value, the bow shock cone becomes broader. Therefore, we were able to use Eq. (1) to estimate the mass-loss rate from the binary system by measuring the stand-off distance. Generally, the ISM density and stellar velocity can be determined following Eq. (2) and Johnson & Soderblom (1987), respectively. While the error of the space motion is negligible, the ISM density value is only an estimate since the star could move through a dense cloud, which is not considered by Eq. (2). The respective values of these quantities for the three objects are given in Table 1. To obtain the space motion (v∗ ,LSR) with respect to the local standard of rest, we corrected the heliocentric motions from the solar motion vector (U,V,W)⊙ = (8.50 ± 0.29,13.38 ± 0.43,6.49 ± 0.26) km s-1 (Coşkunoǧlu et al. 2011). However, since the proper motions are quite small (a few mas yr-1), the correction for the solar motion has a large impact, especially on the PA of the motion. Interestingly, the PA of the corrected LSR motion is a worse match to the bow-shock orientation than the PA of the uncorrected motion (see Figs. 1, 3, and 5). Using (U,V,W)⊙ = (10.00 ± 0.36,5.25 ± 0.62,7.17 ± 0.38) km s-1 determined from the Hipparcos data by Dehnen & Binney (1998) leads to vLSR velocities which are slightly closer to the better matching heliocentric values. A similar discrepancy between bow-shock orientation and LSR motion (and less so with heliocentric motion) is found by Peri et al. (2012, 2015) for a large number of O- and B-type stars with bow shocks, as collected in the WISE E-BOSS survey. For this reason, we overplotted both the heliocentric and LSR motions on the WISE images in Figs. 1, 3, and 5.


A few possibilities are explored for XCl: (1) The most straightforward explanation is that the Cl radical itself is present in the coma. Although it seems unlikely due to its reactivity, the CN radical was discovered in the coma of 67P (Hänni et al. 2020); therefore, the Cl radical cannot be excluded as a possible candidate. (2) Next to CH3Cl (Fayolle et al. 2017), which is too low in abundance to be a suitable neutral candidate to explain the observations, no other volatile Cl-bearing species (e.g. like Cl2) have been identified in the DFMS mass spectra. However, as all results show that EII product ion fractions decrease as the complexity of the molecules increases (Schuhmann et al. 2019) and since DFMS spectra for masses >100 have not been thoroughly analysed for the whole mission, there is still a possibility for other Cl-bearing species to be present within the DFMS mass range, albeit with low abundances. (3) One cannot exclude a priori that there would be heavier (semi)volatile Cl-bearing parents at higher m/z outside the DFMS mass range that sublimate and are present in the coma. However, the fragmentation of such parents, e.g. of the form CxHyCl, in the DFMS ion source would create Cl-bearing fragments at lower m/z. Alternatively, heavier semivolatile chlorine-bearing neutrals might undergo photolysis rather than sublimation, releasing intermediate chlorine-bearing neutrals at lower m/z into the coma, which in turn fragment into smaller ions in the DFMS ion source. Both processes are improbable since apart from CH3Cl no Cl-bearing fragments at lower m/z have been found and the effect of photodissociation should be much more pronounced during perihelion when Rosetta was farther away from the comet and this was not observed. (4) A source of Cl may lie in compounds that decompose upon sublimation and/or ionization. As an example, it is known that ammonium perchlorate (NH4ClO4) releases HCl upon warming (Boldyrev 2006). Other perchlorates or compounds with oxidized states of Cl may provide relatively more Cl+ than HCl+ upon dissociation. According to Schauble, Rossman & Taylor (2003), molecules with oxidized Cl are enriched with 37Cl relative to non-oxidized species, which could also explain the somewhat higher coma isotopic values observed.


When dipolar species, such as CO, are allowed to deposit in the laboratory on a cold surface under vacuum to form a layer of approximately 4–5 monolayers (ML), the molecules may spontaneously orient throughout the layer such that their positive or negative ends protrude from the surface, creating a polarization potential (Balog et al. 2009; Cassidy et al. 2012; Field et al. 2013; Lasne et al. 2015; Rosu-Finsen et al. 2016). We have developed a semi-empirical description of this so-called spontelectric effect, which describes major features of the phenomenon. This description is based on a physical model that invokes spontaneous orientation of the dipolar molecules of which the film is composed. The analysis, described in detail in Field et al. (2013), is remarkably successful in describing important characteristics of spontelectrics, such as the deposition temperature dependence of the surface polarization charge, including the counterintuitive behavior of solid methyl formate (Plekan et al. 2011). In the case of CO, the positive oxygen end protrudes on average (Collings et al. 2014). The spontelectric effect is exemplified by a powerful spontaneous electric field in the film. It was first discovered for nitrous oxide using a direct measurement of surface potential, employing high-resolution, low-energy electron beams (Balog et al. 2009). Studies have now been extended to CO using RAIR (reflection-absorption infrared) spectroscopy (Rosu-Finsen et al. 2016). RAIR spectra were recorded of CO laid down to a thickness of 20 ML, dosed at 21, 22, and 24 K, on 50 ML of compact amorphous solid water (cASW) formed by dosing at 110 K, in order to reproduce as far as possible the nature of ices on grains in a prestellar core. The presence of spontaneous electric fields in the film of CO was demonstrated by measurements of Stark frequency shifts in RAIR spectra, the values of which are film deposition temperature dependent. These observations were used to derive the electric field in the film, yielding a temperature-averaged surface potential, ϕ, of 6.7 ± 0.5 mV per ML of CO (Rosu-Finsen et al. 2016). We note that experimental observations show that the surface potential for any spontelectric material depends on the temperature of the substrate at the moment of deposition, which is ≤26 K in the ISM. The subsequent cooling of the substrate to 10 K in the ISM has no effect on the surface potential.


Dramatic changes in the optical broad emission lines (BELs) of CLQs can also probe the physical origin of the lines themselves. Optical spectra of AGN are characterized primarily by a power-law continuum, BELs, and narrow emission lines (e.g., Vanden Berk et al. 2001). BELs are emitted from the broad-line region (BLR), which is believed to be the high velocity gas gravitationally bound to the central supermassive black hole (Peterson & Wandel 2000; Peterson et al. 2004). Results from reverberation mapping (Denney et al. 2009; De Rosa et al. 2018), and study of the quasar orientation and observed width of BEL profiles (Shen & Ho 2014; Storchi-Bergmann et al. 2017) imply that the geometry of the BLR is likely to be disk-like. Furthermore, study of BEL profiles suggests the BLR gas is structured as a smooth continuous flow (Laor et al. 2006). However, the exact origin of the BLR gas is still unclear (Elvis 2017). The disk-wind model suggests that the BLR gas comes from winds produced by the accretion disk (Emmering et al. 1992; Murray et al. 1995; Murray & Chiang 1997; Elitzur & Ho 2009; Elitzur et al. 2014; Elitzur & Netzer 2016). When the mass accretion rate drops below a certain limit (corresponding to a luminosity of ∼

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M

⊙

), winds can no longer be sustained, such that the observed BELs “disappear” (Elitzur & Ho 2009). Below the critical luminosity, AGN are expected to be “true” Type 2 AGN (i.e., Type 2 AGN intrinsically absent of BLRs). Although this critical luminosity is dependent on the black hole mass, the disk-wind model predicts that more detections of BEL disappearance are expected below the 1% Eddington ratio (Elitzur & Netzer 2016), and a double-peaked BEL profile, the signature of a rotating disk, should emerge in the quasar spectrum when the accretion rate drops (Elitzur et al. 2014). Single-epoch observations show that Eddington ratios of bright quasars with prominent BELs, are always higher than 1% (Kollmeier et al. 2006; Steinhardt & Elvis 2010). Yet it is still unknown whether Eddington ratios of individual quasars varying around 1% will display the appearance/disappearance of broad emission lines. By studying the Eddington ratios of the single changing-look AGN UGC 3223, Wang et al. (2020) find this object crosses the 1% Eddington ratio when its AGN type changes. However, it is still necessary to investigate whether this 1% Eddington ratio is associated with appearance/disappearance of BELs with multiepoch observations of CLQs, and with a larger data set. In this work, we also study the changes in the Eddington ratio distributions from the brightest to the faintest optical spectra of CLQs, investigate the possible connection with the 1% Eddington ratio, and study the BEL profiles of those CLQs.


Low redshift (z ≲ 3) LyC emitters have proven difficult to find (e.g., Leitherer et al. 1995; Bergvall et al. 2006; Siana et al. 2007, 2015; Vanzella et al. 2010; Leitet et al. 2011, 2013) due to the attenuation of LyC photons by the intergalactic medium (IGM). For decades, the only confirmed, directly detected LyC escape from a low-redshift galaxy was that in Haro 11 (Bergvall et al. 2006; Leitet et al. 2011), with Östlin et al. (2015) providing the first spatially resolved kinematic study of a confirmed LyC emitter. Only recently has significant progress been made in detecting local LyC emitters. Green peas (GPs) are the most prolific class of local LyC-leaking galaxies, showing a very high LyC detection rate. So far, 11 out of 11 GPs have been confirmed as leaking LyC by direct detection with HST/COS (Izotov et al. 2016a,b, 2018a,b). GPs are compact galaxies which appear green in SDSS g, r, i composite images due to their extremely high equivalent width of [O III] λλ4959, 5007 and their redshifts of only z ∼ 0.2 − 0.4. They have subsolar metallicities, low masses, and high specific star-formation rate, and represent good analogs to high redshift SFGs (e.g., Cardamone et al. 2009; Izotov et al. 2011; Bian et al. 2016). While much closer than their high-z cousins, GPs are nevertheless at distances that still present severe difficulties for a spatially resolved analysis of their properties. Pioneering work with integral field unit (IFU) observations of four GPs (Lofthouse et al. 2017) reveals that two are rotationally supported, with seemingly undisturbed morphology, and two are dispersion-dominated, leading the authors to conclude that mergers may not be a necessary driver of their star formation properties, and by extension, for their LyC emission. At the same time, Amorín et al. (2012) find complex, multicomponent Hα line profiles, with implied high-velocity gas in all five of their observed GPs. While these GPs are not confirmed LyC emitters, and therefore their kinematical properties cannot be taken as a necessary condition for LyC escape, outflows nevertheless seem to play a role in enabling LyC leakage. For example, Chisholm et al. (2017) find that outflows can be detected in all LyC emitters they examine, albeit with velocities not statistically different from a control sample of non-leakers. Further, Martin et al. (2015) find that ULIRG galaxies with the strongest outflows manifest the lowest column densities of neutral gas. Simulations also support the scenario of outflows enabling LyC escape. Trebitsch et al. (2017) find that mechanical feedback from supernovae explosions has a vital role in disrupting dense star-forming regions and in clearing low-density escape paths for the LyC photons. While IFU observations of a large sample of GPs would help settle some of these issues, the large distances to these galaxies would likely prohibit the detailed identification and characterization of the exact escape mechanisms of LyC. Much closer galaxies, showing the same characteristics as GPs, are required for such an endeavor.


The density structure and infall kinematics discussed here are probably dominated by the most massive object clearly identified in Fig. 2b. The envelope profile defined above (see Sect. 3.5) and their corresponding values given in Table 3 have been used in Eq. (A.2) to estimate a mass of ~3500 M⊙ in a 2.5 pc radius for the central UCH ii region. This region exhibits a steeper density gradient in its outer envelope, with qout ≃ −2.5, also illustrated by the cloud structure studied with probability density function (Schneider et al. 2015; Rayner et al., in prep.). It suggests compression by external forces, as shown in numerical simulations of Hennebelle et al. (2003) and outlined in Appendix A. This compressive process could be associated with stellar winds and ionisation shocks of nearby stars such as those observed in the M 16 massive YSO (young stellar object) by Tremblin et al. (2014b) or with globally infalling gas driven by the dynamic formation of Mon R2 cloud through force fall. Such infalling motions have in fact been observed in CO by Loren (1977), and clumps with similar masses generally exhibit active global infall. This is the case of SDC335, which has a mass of ~5500 M⊙ in 2.5 pc (Peretto et al. 2013) and contains a protostellar object of similar type B1 (Avison et al. 2015), as well as of DR21, Clump-14 (DR21-south, ~4900 M⊙), and Clump-16 (DR21-north, ~3350 M⊙) (Schneider et al. 2010a), all of which show supersonic infall (V = 0.5−0.7 km s-1). Moreover, the hourglass morphology of the magnetic field is thought to be a signature of global infall (Carpenter & Hodapp 2008), and Koch et al. (2014) suggest that Mon R2 is a super critical, quickly collapsing cloud. The expected infall velocity gradient, as observed in SDC13 (Peretto et al. 2014), seems to be a crucial ingredient for generating the filament crossing (Dobashi et al. 2014) characteristic of all the regions mentioned here and commonly observed in many places.


Typical corotating high-speed streams (HSS), i.e. those streams coming from the equatorial extension of polar coronal holes, are characterized by different regions. Within these regions, extending on daily scales, field and plasma parameters assume different average values and fluctuations have a different character. The first region is the stream interface (SI), located between fast and slow stream, right where the dynamical interaction between the two flows is stronger (Schwenn & Marsch 1990; Bruno & Carbone 2016). The SI is characterized by high level of pressure, both kinetic and magnetic, high temperature and low Alfvénicity. This region is followed by the trailing edge (TE), which is characterized by the highest wind speed and temperature, the lowest compressibility and the highest level of Alfvénicity. The TE is then followed by a wind that progressively slows down and becomes cooler. This region was first named rarefaction region by Hundhausen (1972). Interesting enough, these two portions of the stream are separated by a very narrow region across which Alfvénicity changes abruptly, from high to low (Bruno & Telloni 2015). Most of the time, the rarefaction region is followed by the heliospheric current sheet crossing characterized by low Alfvénicity, low temperature but higher field and plasma compressibility. Thus, moving from fast to slow wind, one observes a different turbulence, not only relatively to its power level, but also to its Alfvénic character, although the spectral slope at fluid scales clearly shows a persistent Kolmogorov-like scaling. In addition, besides the above-cited differences within the fluid regime, there are also clear differences at proton kinetic scales. Bruno, Trenchi & Telloni (2014), investigating the behaviour of the spectral slope at proton scales, up to frequencies of a few Hz, beyond the high-frequency break separating fluid from kinetic scales, recorded a remarkable variability of the spectral index (Smith et al. 2006; Sahraoui et al. 2010) within the following frequency decade or so. The steepest spectra corresponded to the TEs of fast streams, while the flattest ones were found within the subsequent slow wind regions. The same authors found an empirical relationship between the power associated with the inertial range and the spectral slope observed within this narrow region, which allowed us to estimate that this slope approaches the Kolmogorov scaling within the slowest wind and reaches a limiting value of roughly 4.4 within the fast wind. In addition, Bruno et al. (2014) suggested the possible role played also by Alfvénicity in this spectral dependence. Recent theoretical results seem to move in this direction attributing a relevant role to Alfvénic imbalance within the inertial range (Voitenko & De Keyser 2016).


The AR NOAA 12673 was highly flare productive.5

5

https://www.swpc.noaa.gov/products/solar-region-summary

 It appeared in the eastern limb of the Sun on August 28 as a simple α-type AR and gradually evolved into complex βγ-type on September 4. It became an even more complex βγδ-type on September 5 and remained so until its disappearance over the western limb of the Sun on September 10. In total, it produced 27 M-class and 4 X-class flares between September 4–10. Following the occurrence of the two X-class flares on September 6, reported in this article, it went on to produce an X1.3-class flare on September 7 and an X8.2-class flare on September 10 besides several M-class major eruptive events. During the occurrence of X-class flares, the complex AR had shown δ-sunspots, which are identified with a complex distribution of sunspot groups in which the umbrae of positive and negative polarities share a common penumbra (Künzel 1960). Such complex ARs are known to produce powerful flares (see e.g., Zirin & Liggett 1987; Sammis et al. 2000; Takizawa & Kitai 2015). It is noteworthy that AR 12673 was a rather compact region (Figure 9) that displayed more spatial extension in the north–south direction than the usual east–west direction. The δ-sunspots were concentrated at the central part of the AR where magnetic fields were very strong, and the magnetic field gradient across the PIL was extremely high (∼2.4 × 103 G Mm−1; see Figure 9(c)). Earlier studies have shown a close relationship between major flare activities and strong magnetic field, especially those with a high gradient and those that are highly sheared across the PIL (Hagyard et al. 1984; Zirin & Wang 1993; Schrivjer 2007; Barnes et al. 2016). The reported eruptive activities in AR 12673, thus, represent the capability of the AR in the rapid generation and storage of huge amount of excess magnetic energy in the corona. In this context, the evolution of normalized free magnetic energy during the X-class flares is noteworthy (Figure 12). We found that the free magnetic energy stored in the AR before the flaring activity was ∼82% of the potential magnetic energy. After the two X-class flares, it reduced to ∼70%. Our analysis, therefore, implies that a large amount of free magnetic energy was already stored in the AR before the flaring activities and that the large X-class flares essentially released only a small fraction of it.


Nearly half of the known exoplanets orbit within 0.1 au from their star. At such close distances, the nature and evolution of these planets is shaped by interactions with their host star (irradiation, tidal effects, and magnetic fields). In particular, the deposition of stellar X-ray and extreme ultraviolet radiation (XUV) into an exoplanet upper atmosphere can lead to its hydrodynamic expansion and substantial escape (e.g. Lammer et al. 2003; Vidal-Madjar et al. 2003; Lecavelier des Etangs et al. 2004; Yelle 2004; García Muñoz 2007; Koskinen et al. 2010; Johnstone et al. 2015). Atmospheric loss is considered as one of the main processes behind the deficit of Neptune-mass planets at close orbital distances (the so-called hot Neptune desert, e.g. Lecavelier des Etangs 2007; Davis & Wheatley 2009; Szabó & Kiss 2011; Lopez, Fortney & Miller 2012; Beaugé & Nesvorný 2013; Lopez & Fortney 2013; Owen & Wu 2013; Jin et al. 2014; Kurokawa & Nakamoto 2014; Lundkvist et al. 2016). These planets are large enough to capture much of the stellar energy, but in contrast to hot Jupiters are not massive enough to retain their escaping atmospheres (e.g. Hubbard et al. 2007; Lecavelier des Etangs 2007; Ehrenreich, Lecavelier des Etangs & Delfosse 2011). The missing hot Neptunes could have lost their entire atmosphere via evaporation, evolving into bare rocky cores at the lower radius side of the desert (e.g. Lecavelier des Etangs et al. 2004; Owen & Jackson 2012). This scenario is strengthened by the recent observations of warm Neptunes at the border of the desert, on the verge of (Kulow et al. 2014; Bourrier, Ehrenreich & Lecavelier des Etangs 2015; Ehrenreich et al. 2015; Bourrier et al. 2016; Lavie et al. 2017) or undergoing (Bourrier et al. 2018b) considerable mass-loss. Because they survive more extreme conditions than lower mass gaseous exoplanets, hot Jupiters are particularly interesting targets to study star–planet interactions. Their upper atmosphere can be substantially ionized because of stellar photoionization (e.g. Schneiter et al. 2016), which could help the formation of reconnections between the stellar and planetary magnetospheres that would enhance stellar activity (e.g. Cuntz, Saar & Musielak 2000; Shkolnik, Walker & Bohlender 2003; Ip, Kopp & Hu 2004; Shkolnik et al. 2008; although see Poppenhaeger & Schmitt 2011; Scandariato et al. 2013; Llama & Shkolnik 2015 for the difficulties to detect such signatures). Atmospheric escape of neutral hydrogen and metal species has been detected via transmission spectroscopy for several Jupiter-mass planets, bringing information about their upper atmosphere and the stellar environment (HD 209458b, Vidal-Madjar et al. 2003, 2004, 2008; Ehrenreich et al. 2008; Ben-Jaffel & Sona Hosseini 2010; Linsky et al. 2010; Schlawin et al. 2010; Vidal-Madjar et al. 2013; Ballester & Ben-Jaffel 2015; HD 189733b, Lecavelier des Etangs et al. 2010, 2012; Bourrier et al. 2013; 55 Cnc b, Ehrenreich et al. 2012; WASP-12b, Fossati et al. 2010; Haswell et al. 2012). Shocks could for example form ahead of hot Jupiters because of the interaction between the stellar wind and the planetary outflow or magnetosphere (Vidotto, Jardine & Helling 2010; Cohen et al. 2011; Llama et al. 2013; Tremblin & Chiang 2013; Matsakos, Uribe & Königl 2015)


A likely scenario for the geometry of the entire system is presented in Fig. 15. The locations of systems A–C are fairly well constrained by their metal absorption. System D is harder to place, as there is little to be learned from its negligibly low metal content. It is highly improbable that it is an outflow similar to C but with a higher velocity, or indeed that it has originated inside the galaxy at all. Instead we believe it to trace the metal-poor gas that acts as a replenishment mechanism of galaxies throughout the universe. It has indeed been suggested that GRB hosts in general are fuelled by recent metal-poor gas inflow (Michałowski et al. 2015, 2016). The −320 km s-1 blueshift implies that it is not accreting onto galaxy B – if B hosted the GRB, then D must be in front of it, moving away at that velocity. The degeneracy between local velocities and cosmological redshift allows for a range of locations of system D. One possibility is that it is located between the two galaxies and is currently falling towards galaxy A, or has fallen from behind galaxy B and is being slowed by its gravity. Another possibility is that it has fallen from a large distance into the potential well of the complex and is now in the foreground, having initially passed both galaxies without being accreted, having joined in a disk-like structure (e.g. Bouché et al. 2013). For any of these scenarios, system D represents a rare detection of such a metal-poor inflow in a GRB galaxy, or indeed in absorption at such a small projected distance to any galaxy. Systems with positive velocities relative to the gas local to GRBs, and thus inferred to be falling in towards the star-forming regions of their hosts, have been detected previously (e.g. Prochaska et al. 2008), but these have tended to be richer in metals than the gas observed here and may represent recycled gas that has previously been expelled from the galaxy. Ly-α emitters are thought to trace accretion in large halos around quasars (Cantalupo et al. 2014; Hennawi et al. 2015), and have been detected at impact parameters of ≥70 kpc from an LLS (Fumagalli et al. 2016). If system D is indeed bound to the host system, then it could be much closer to the galaxies than that. Finally, as discussed when defining the host complex in Sect. 3.1, it could in fact be a foreground system located ~1.5 Mpc from the interacting complex. While this would exclude it from being a direct accretion flow onto the complex, it represents at the very least a reservoir of cool metal-poor gas residing in the IGM near the dark matter halo, available for star formation at some future epoch.


In this work, we study the merger and post-merger evolution of near equal-mass BH–NS binaries. Before turning to the properties of the accretion discs formed in such mergers, we first provide a very brief overview of their formation. We do this by considering the fiducial system TNT.chit.0.35. In order to illustrate the disc formation process, we begin by summarizing the dynamical formation of the disc in Fig. 1, which reports the comoving magnetic energy density b2, the rest-mass density ρ, the electron fraction Ye, and the local fluid temperature T. The different rows correspond to meridional (top panels) and equatorial (bottom panels) views of the accretion disc around the BH, while the different columns correspond to different times after the merger. The general dynamics of this process have been studied extensively in purely hydrodynamical simulations (Etienne et al. 2009; Kyutoku et al. 2011; Foucart et al. 2012). In order for a massive disc to form during and after merger, tidal disruption has to occur outside of the innermost stable circular orbit (ISCO) of the BH (Shibata & Taniguchi 2011; Pannarale et al. 2011b). Starting from the left-hand panel, we can see that shortly after tidal disruption, an initial accretion disc begins to form around the BH. Originating from the cold NS matter, the initial disc is very neutron rich (Ye  0.05), but already reaches temperatures $T \lesssim 10\, \rm MeV$. The disc quickly grows in mass and size due to fall-back accretion from the tidal arm (middle column), begins to circularize and a steady accretion flow develops over time. As expected, this happens on the dynamical time-scales of the discs, which are proportional to the disc mass $M_{\rm disk}^b$, so that the lightest discs circularize first. Initially, the pure neutron matter is far out of beta-equilibrium under these conditions and will rapidly re-equilibrate via beta decay of neutrons, leading to an increasing protonization especially of the low-density parts of the disc. At the same time, the magnetic-field strength is increasing throughout the disc, exceeding $10^{14}\, \rm G$ locally. More details on the magnetic-field evolution will be given in Section 3.3. Finally, after more than $50\, \rm ms$ past merger, the disc has settled into an initial quasi-equilibrium, consisting of a very neutron-rich disc, probing rest-mass densities ${\lesssim}10^{11}\, \rm g\, cm^{-3}$. A disc formed by this process will then set the initial conditions for the long-term evolution in terms of the accretion flow and mass ejection (Fernández et al. 2015, 2017).


For a warped disc, the binary separation and the mass ratio suggest that M81 X-6 lies in the intermediate instability zone for radiation-driven warping in X-ray binaries (see fig. 1 of


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


Cluster radio relics are usually found in the outskirts of merging galaxy clusters. They exhibit elongated morphologies and high degrees of polarisation above 1 GHz (up to 70%, Ensslin et al. 1998; Bonafede et al. 2014; Loi et al. 2019; de Gasperin et al. 2022). The resolved spectral index in radio relics shows a gradient: it steepens towards the cluster centre and flattens towards the outskirts. Their size can reach up to ∼2 Mpc, and high-resolution observations have revealed filamentary structures within relics themselves (Di Gennaro et al. 2018; Rajpurohit et al. 2020, 2022a,b; de Gasperin et al. 2022). The Largest Linear Sizes (LLS) and radio powers of relics are correlated, as well as the integrated spectral index and the radio power (van Weeren et al. 2009b; Bonafede et al. 2012; de Gasperin et al. 2014). Relics trace ICM shock waves with relatively low (M   3) Mach numbers (Finoguenov et al. 2010; Akamatsu et al. 2013; Shimwell et al. 2015; Botteon et al. 2016). The acceleration of electrons is believed to proceed via diffusive shock acceleration (DSA) in the ICM (Ensslin et al. 1998; Roettiger et al. 1999), in which particles scatter back and forth across the shock front gaining energy at every crossing. Nevertheless, this mechanism has been shown to be rather inefficient in accelerating electrons from the thermal pool (Vazza & Brüggen 2014; Vazza et al. 2016; Botteon et al. 2020a; Brüggen & Vazza 2020; see Brunetti & Jones 2014 for a review). Recently, it has been suggested that seed electrons could originate from the tails and lobes (driven by AGN outflows) of cluster radio galaxies (Bonafede et al. 2014; van Weeren et al. 2017; Stuardi et al. 2019), which alleviates the requirements of high acceleration efficiencies at cluster shocks (e.g., Markevitch et al. 2005; Kang et al. 2012, 2017; Botteon et al. 2016; Eckert et al. 2016). In some cases, double relics have been detected on opposite sides of the cluster centre (e.g., Rottgering et al. 1997; van Weeren et al. 2010, 2012b; Bonafede et al. 2012; de Gasperin et al. 2015a). In these clusters it is possible to constrain the merger history, providing important information about the formation processes of relics.


While spatially extended optical jets and bipolar CO molecular outflows have been observed in numerous Class 0/I protostars (e.g. Reipurth & Bally 2001; Bally 2016, and references therein), near-infrared high-resolution spectroscopy and spectroimaging observations in the past two decades have made it possible to study the kinematics of the outflowing gas and physical properties at the base of the jet within a few hundred au of the driving source in Class 0/I protostars (e.g. Davis et al. 2001, 2003, 2011; Nisini et al. 2005, 2016; Caratti o Garatti et al. 2006; Takami et al. 2006; Antoniucci et al. 2008, 2011, 2017; Garcia Lopez et al. 2008, 2013). These microjets are bright in [Fe ii] forbidden and H2 rovibrational emission lines, hence showing the presence of forbidden emission-line (FEL) regions and molecular hydrogen emission-line (MHEL) regions in low-mass Class 0/I protostars. While multiple low- and high-velocity components are observed in both MHELs and FELs, the higher velocity gas is slightly further offset from the driving source than the slower gas, and the kinematics of the H2 emission differs from [Fe ii] emission, revealing complicated kinematic structures. Evidence of H2 emission from cavity walls is also seen in some protostars, suggesting the presence of a wide-angled wind. Strong emission in the well-known accretion diagnostics of Paschen and Brackett hydrogen recombination lines is observed in protostars, with the ratio of the accretion luminosity to bolometric luminosity spanning from ∼0.1 to ∼1. The mass accretion and loss rates for Class 0/I low-mass protostars span the range of 10−6–10−8 M⊙ yr−1, and the derived jet efficiencies (ratio between mass ejection and accretion rates) range between ∼1 per cent and 10 per cent (e.g. Davis et al. 2001, 2003, 2011; Nisini et al. 2005, 2016; Caratti o Garatti et al. 2006; Takami et al. 2006; Antoniucci et al. 2008, 2011, 2017; Garcia Lopez et al. 2008, 2013). These measurements are within the range predicted by the magnetohydrodynamic jet launching models (e.g. Frank et al. 2014).


The band-limited nature of repeaters makes it essential to have high-resolution simultaneous observations with a large bandwidth that can help detect variations in the spectrotemporal and even polarimetric structures of the bursts at different frequencies. The study of these contemporaneous frequency-dependent variations, as well as long-term changes with time, are paramount in investigating the origins of FRBs. FRB 20121102A was one of the first repeaters to be coherently corrected for dispersion, allowing unprecedented high-time-resolution studies, revealing drifting bursts and an unusually high rotation measure (RM; Gajjar et al. 2018; Michilli et al. 2018; Hessels et al. 2019). Subsequent observations have shown a considerable decrease in the RM value with time (Hilmarsson et al. 2021). This variation is similar to what has been observed for Galactic center magnetar PSR J1745-2900 (Desvignes et al. 2018). On the other hand, FRB 20180916B has a nominal RM and has shown a gradual variation in RM value over the two year timescale (Pleunis et al. 2021). These variations may be due to changes in the viewing geometry or the result of drastic variations in the circumburst environment facilitating magnetoionic disparities. Interestingly, both of these FRBs show a flat polarization position angle (PPA) that can be indicative of emissions at higher altitudes due to relativistic shock (Beloborodov 2017; Metzger et al. 2019). Contrarily, observation of microstructures challenges this idea, as in the case of FRB 20180916B, where such structures have been seen down to 2–3 μs (Nimmo et al. 2020) at 1700 MHz, hinting at emission much closer to the magnetosphere. Studying the evolution of such structures with frequency will aid in constraining the emission region and help to understand the local environment. Another aspect of repeater burst morphology that remains a mystery is the “sad trombone” effect, which describes a linear decrease in burst features with frequency and has been observed for FRB 20121102A and other repeaters; it will be interesting to undertake a comparative study of this behavior between frequencies.


These signatures are inconsistent with ultra-low frequency waves, which have circular polarization and a period similar to the upstream ion gyroperiod. The waves are also inconsistent with ion Weibel instability, which generates linearly polarized waves. Interaction of reflected ions with incoming solar wind electrons or ions can cause foot instabilities that excite waves in the whistler mode branch. Modified Two Stream Instability (MTSI) due to relative drift between reflected ions and incoming solar wind electrons (fast drift), and incoming solar wind ions and electrons (slow drift) has been frequently considered (Matsukiyo & Scholer 2003; Comişel et al. 2011; Umeda et al. 2012; Marcowith et al. 2016; Wilson et al. 2016; Muschietti & Lembège 2017; Hull et al. 2020). This instability, however, if excited, creates significant ion heating throughout the foot and suppresses the reformation process (Shimada & Hoshino 2005; Matsukiyo & Scholer 2006), rather than creating episodic enhancements that we show in the foot. Furthermore, Gary et al. (1987) indicated that (fast drift) MTSI becomes dominant at low electron beta (βe  0.5), while at higher βe more resonant electrons stabilize this instability through increased electron Landau damping. Electron data for the time period we discussed in this paper show βe ≥ 1.2, and therefore fast drift mode MTSI is most likely not significant. The slow drift mode of MTSI could be a more viable candidate at high β plasmas. Wave properties around 1.6 Hz in the middle interval (purple segment) of Figure 6, indicate that the wave is in propagation toward the ramp (

 = −0.66, −0.71, 0.22) with Vph‐sc = 34 km s−1 and λwave = 21.4 km  ∼ 12λe, where λe is the upstream electron inertial length. The plasma rest-frame frequency of the wave is about 8 Hz  ∼ 2.5flh. Since these characteristics are somewhat consistent with model predictions for drift mode of MTSI (Muschietti & Lembège 2017), we do not rule out the possibility of some waves at certain frequencies and during some intervals being generated by the slow drift mode of MTSI.


We repeated the plasma pressure calculation presented by Schwadron et al. (2011) and Fuselier et al. (2012) for the new ENA energy spectrum. The results for the downwind hemisphere and for the Voyager 1 region are summarized in Table 3. The measured intensity

j

ENA

 of neutralized hydrogen at a given energy translates into a pressure of the parent ion population in the heliosheath times the integration length along the line of sight, ΔP × l, in the following way:
3

Δ
P
×
l
=

4
π

3

n

H

m

H

v

j

ENA

(
E
)

σ
(
E
)

Δ
E

c

f

4

c

f

=

(
v
+

u

R

)

2

v

4

(

v

2

+
4

u

R

2

+
2

u

R

v
)
.

In Equation (3), ΔE denotes the width of the respective energy bin; for the typical radial velocity of solar wind in the flanks and the downwind hemisphere of the inner heliosheath, we assumed uR = 140 km s−1 as measured by Voyager 2, whereas uR = 40 km s−1 for the heliosheath in the Voyager 1 direction (Schwadron et al. 2011; Gloeckler & Fisk 2015). For the density of neutral hydrogen in the inner heliosheath a constant nH = 0.1 cm−3 is assumed (Schwadron et al. 2011; Gloeckler & Fisk 2015). The charge-exchange cross section between protons and neutral hydrogen decreases from (4 to 2) × 10−15 cm−2 for 0.015 to 1.821 keV (Lindsay & Stebbings 2005). The integration length l for ENA production in the plasma is approximately the thickness of the inner heliosheath. The part of Equation (3) without the velocity factor cf can be interpreted as stationary pressure. The total pressure or dynamic pressure is the stationary pressure times this factor. Integrating over all energy bins in Table 3, we obtain the total plasma pressure times heliosheath thickness as P × l = 304 pdyn cm−2 au for the downwind hemisphere and 66 pdyn cm−2 au for the Voyager 1 region (1 pdyn cm−2 au = 0.015 N m−1). If we want to put these numbers into the context of other studies, we face two problems. First, the uncertainty of the total pressure is large given the upper limits in the two lowest energy bins. Second, heliosheath plasma more energetic than 2 keV will produce ENAs that cannot be detected with IBEX-Lo. We therefore used the observed median j = 0 cm−2 sr−1 s−1 keV−1 for heliospheric ENAs in the two lowest energy bins of IBEX-Lo and relied on the study by Livadiotis et al. (2013). They compared the expected plasma pressure from a kappa distribution of protons with the plasma pressure derived from IBEX-Hi energy spectra: the energy range between 0.03 and 2 keV, roughly corresponding to the IBEX-Lo range, covered more than half of the total plasma pressure predicted from a kappa distribution. The authors found a total plasma pressure of P = 2.1 pdyn cm−2 for all sky directions except for the ENA Ribbon. Gloeckler & Fisk (2015) presented a multi-component plasma model for the heliosheath to explain Voyager and IBEX observations. At low energies they assumed the ENA energy spectra provided by Fuselier et al. (2012). They derived a total pressure of 2.5 pdyn cm−2 in all three plasma regions in the nose of the heliotail (Gloeckler & Fisk 2015). Pressure contributions from the slowed solar wind, magnetic pressure, and the pressure exerted from pickup ions and anomalous cosmic rays all had to be taken into account to obtain this total pressure.


The above qualitative theoretical reasoning raises the question about why would Swift J1834.9−0846 be the only magnetar so far powering a wind nebula, given that previous searches around individual magnetars have returned no sign of extended emission attributable to wind nebulae (e.g., Viganò et al. 2014). With only one observed so far, it is difficult to draw any firm conclusions. Nevertheless, Swift J1834.9−0846 has some interesting characteristics that are not shared with the entire magnetar population. First, the environment of Swift J1834.9−0846 is extremely crowded, with a Fermi GeV source, an H.E.S.S. TeV source, an SNR, a GMC, and an OH maser in its vicinity (Frail et al. 2013; H.E.S.S. Collaboration et al. 2015). The relationship between all these sources is unclear. However, it is tempting to speculate that environmental effects from such a rich field could be playing a role in the production of this wind nebula (e.g., triggering of pair cascade by external gamma rays from a nearby source; Shukre & Radhakrishnan 1982; Istomin and Sob’yanin 2011). Second, the Swift J1834.9−0846 X-ray luminosity in quiescence is

 erg s−1. Only five other magnetars (SGR 0418+5729, SGR 1745−2900, XTE J1810−197, Swift J1822.3−1606, 3XMM J185246.6+003317)20

20

http://www.physics.mcgill.ca/ pulsar/magnetar/main.html

 have luminosities ≲1032 erg s−1. Among these five, three have the smallest surface B fields measured (SGR 0418+5729, Swift J1822.3−1606, 3XMM J185246.6+003317;

 G), and only one source, SGR 1745−2900, has a rotational energy loss rate

 similar to Swift J1834.9−0846, while the rest have

 at least an order of magnitude lower. Hence, from an observational point of view, it seems that the combination of very weak X-ray luminosity, a magnetar-like B-field strength, and a somewhat large

 (properties that are only shared by the Galactic center magnetar SGR 1745−2900) may favor wind nebula production. Another possibility is that the Swift J1834.9−0846 magnetar/nebula system is an older analog to the Kes 75 system, where the central pulsar evolves into a magnetar while preserving its originial PWN.