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SAO/NASA Astrophysics Data System 13

2017ApJ...845..160P__Lyubarsky_2009_Instance_1

It is a pressing question as to how the radiation that is observed in relativistic jets in active galactic nuclei (AGNs) is generated (e.g., Blandford & Königl 1979; Marscher 1980; Zensus 1997; Laing & Bridle 2002; Honda 2010; Levinson & Rieger 2011; Mościbrodzka et al. 2011; Ito et al. 2013; Mason et al. 2013; Potter & Cotter 2013; Hovatta et al. 2014; Scott & Stewart 2014; Shih & Stockton 2014; Wang et al. 2014; Turner & Shabala 2015; Asada et al. 2016; Hirotani et al. 2016; Koay et al. 2016; Khabibullin et al. 2016; Prieto et al. 2016). Although there is a common consensus that the emitters are energetic particles, how these particles are accelerated to such high energies, how they dissipate their energy, and how they are transported with the jets themselves are still the subject of feverish investigation (e.g., Blandford & Eichler 1987). It is argued that relativistic outflows from black holes are associated with accretion flows (Blandford 1976; Fender et al. 2004; Meier 2005; Ferreira et al. 2006; Trump et al. 2011; Pu et al. 2012; Wu et al. 2013; Ishibashi et al. 2014; Sbarrato et al. 2014). In the case of collimated relativistic jets, magnetic fields must play an important role (Camenzind 1986a, 1986b, 1987; Fendt & Greiner 2001; Vlahakis & Königl 2004; Komissarov et al. 2007; Lyubarsky 2009; Nakamura & Asada 2013; Homan et al. 2015), and it is argued that jets are powered at the expense of the black hole, wherein energy is extracted from a reservoir of rotational energy from the black hole itself, either by electromagnetic means (Blandford & Znajek 1977; Komissarov 2004, 2005; Toma & Takahara 2014) or through magnetohydrodynamical processes (Phinney 1983; Takahashi et al. 1990; Koide et al. 2002; McKinney & Gammie 2004; Hawley & Krolik 2006). Models have been proposed for both of these cases, and they both in principle possess certain testable predictions. In particular, for the latter, numerical GRMHD simulations (e.g., McKinney 2006) and analytical GRMHD studies (e.g., Takahashi et al. 1990; Pu et al. 2016) consistently show the presence of a stagnation or separation surface (a separatrix). This surface separates the (inner) inflow region from the (outer) outflow region, both of which follow the same global, black-hole-threading magnetic field lines. The relatively slow radial velocities near the stagnation surface imply a high concentration of fluid particles. If energetic particles are injected in the vicinity of the stagnation surface or near the black hole event horizon, they must accumulate in high concentrations near the stagnation surface, provided that the cooling timescale is not significantly shorter than the dynamical timescale of the jet fluid flow. This surface is a unique feature of relativistic GRMHD jets and in contrast to an ideal force-free magnetic jet (e.g., McKinney & Narayan 2007; Tchekhovskoy et al. 2008; Broderick & Loeb 2009).

2019ApJ...875L...7D__Saikia_et_al._2018_Instance_1

The first transit of HD 21749b was identified by both the MIT Quick Look Pipeline (which searches for planet candidates in the 30 minutes Full Frame Images) and the Science Processing Operations Center (SPOC) pipeline based at the NASA Ames Research Center (Jenkins et al. 2016). No other matching transits were found in the publicly released data from sectors 1 and 2. After TOI 186.01 was alerted, we searched for archival spectroscopy of this very bright star and found 59 High-accuracy Radial-velocity Planet Searcher (HARPS) radial velocities (RVs) in the European Southern Observatory (ESO) archive (see Section 2.4). A periodogram of these RVs showed a clear signal at 35.57 days, but the TESS photometry and the RHK index (Boro Saikia et al. 2018) indicate a stellar rotation period of around 35 days, calling for caution. If the strongest period in the RVs did correspond to the planet, then we expected to see additional transits in sectors 3 and 4. Once the sector 3 data were released, we discovered that a momentum dump29 29 “Momentum dumps” consist of resetting the momentum wheel speed every 2.5–3 days and are used to mitigate the noisier-than-expected measurements of the spacecraft momentum wheel speeds at higher wheel speeds (see https://archive.stsci.edu/files/live/sites/mast/files/home/missions-and-data/active-missions/tess/_documents/TESS_Instrument_Handbook_v0.1.pdf for details). Momentum dumps require brief interruptions to Fine Pointing mode, during which an increase in the flux dispersion is noticeable in the science data, so data acquired during these intervals are excluded from our analysis. occurred approximately 35.6 days after the sector 1 transit (see Figure 1). We did not let this unexpected turn of events foil our search efforts, and upon close inspection of the light curve we succeeded in recovering a partial transit (including egress) immediately following the momentum dump. Finally, we observed a third transit in sector 4, thus allowing for a robust ephemeris determination (see Section 3.3). Serendipitously, when applied to the first three sectors of the HD 21749 light curve, the SPOC pipeline yielded an additional planet candidate with a period of 7.9 days (TOI 186.02).

2019MNRAS.487.4473M__Hosokawa_&_Omukai_2009_Instance_1

The simulation model Run1-hr of Meyer et al. (2018) begins with the gravitational collapse of $100\, \rm M_{\odot }$ of pre-stellar rotating molecular material on to the stellar embryo. At the end of the free-fall collapse phase, the infalling material ends on a centrifugally-balanced disc, from which the gas is subsequently transferred to the growing protostar. Fig. 1 plots the accretion rate history on to the central massive protostar (solid blue line, in $\rm M_{\odot }\, \rm yr^{-1}$) together with the evolution of the protostellar mass (dashed red line, in $\rm M_{\odot }$) that is calculated as the integrated disc-to-star mass transfer rate through the sink cell. The vertical thin black line indicates the onset of disc formation, when the free-fall collapse of the envelope material on to the protostar stops and the star begins to gain its mass exclusively via accretion from its surrounding disc. After the initial infall of material, the collapse of the parent pre-stellar core material generates an initial increase of the mass flux through the inner boundary at a time ${\approx } 2\, \rm kyr$. The accretion rate then reaches the standard value predicted for MYSOs of $10^{-3}\, \rm M_{\odot }\, \rm yr^{-1}$ (Hosokawa & Omukai 2009) up to the onset of the disc formation happening at ${\approx } 12\, \rm kyr$. Variabilities in the accretion flow begin right after the disc formation and the accretion rate history exhibit numerous peaks of growing intensity as the MYSO becomes heavier. This is not caused by different inner boundary conditions, but simply reflects the time-dependent azimuthal anisotropies induced in the accretion flow by the disc evolution. The efficient gravitational instabilities produce complex substructures in the disc, such as overdense spiral arms in which gaseous clumps of various morphologies form at radii of ${\sim } 100\, \rm au$ and inward-migrate down to the central protostar. This produces luminosity outbursts via the mechanism revealed in Meyer et al. (2017) for massive stars. These outbursts are responsible for step-like increases in the stellar mass evolution (see thick dotted red line) due to the fast accretion of dense circumstellar material inside the sink cell. A more detailed description of the evolution of circumstellar discs around young massive stars irradiating their self-gravitating discs can be found in our precedent study (Meyer et al. 2018). In Fig. 1, several magenta dots mark the time instances of the chosen simulation snapshots considered in this study. Note that the young star becomes, by definition, a massive object when $M_{\star }=8\, \rm M_{\odot }$. Hence, our selected disc models are all in the high-mass regime.

2019AandA...621A.124B__Leconte_et_al._2010_Instance_1

A large proportion of systems where one planet or more is orbiting closer to its host star than Mercury to the Sun have been observed. Tidal interactions play a key role in the orbital configuration of these very compact systems since they are likely to circularize orbits, align spins, and synchronize periods (Zahn 1977; Mathis & Remus 2013; Ogilvie 2014). These interactions consists in an exchange of angular momentum between the orbit and the spins of the celestial bodies. This exchange is the consequence of the dissipation of tidal flows. The kinetic energy of these tidal flows is converted into heat through tidal dissipation. Since the planet is synchronized within a timescale of a few thousand years, the stellar tide drives the secular orbital evolution (Guillot et al. 1996; Rasio et al. 1996; Leconte et al. 2010). In this work, we neglect the impact of the dissipation in the radiative zone. In stellar convection zones, there are two kinds of tides and both are dissipated by the turbulent friction applied by convective eddies. On the one hand, the equilibrium tide is the large-scale velocity field associated with tidal deformation, the so-called tidal bulge. This nonwave-like entity corresponds to the hydrostatic adjustment of the star to the gravitational perturbation (Zahn 1966; Remus et al. 2012). The friction applied by convective motions delays the response of the star to the perturbation (e.g., Zahn 1989; Ogilvie & Lesur 2012; Mathis et al. 2016). This results in a lag angle between the axes of the tidal bulge and the line of centers. This angle increases with dissipation magnitude. Hansen (2012) calibrated its value for several stellar masses by constraining the dissipation using observations of planetary systems. Since lower-mass stars have deeper convective envelopes, they dissipate more energy than higher-mass stars. On the other hand, in rotating bodies such as stars, at low tidal frequencies, the Coriolis acceleration acting on this equlibrium tide excites inertial modes (Ogilvie & Lin 2007). Their ensemble, the dynamical tide, constitutes the wavelike part of the tidal response. Its dissipation strongly depends on internal structure since it arises from the reflection of inertial modes on the radiative, stably stratified core (Ogilvie 2013; Mathis 2015). The dynamical tide may also vary over several orders of magnitude with rotation since inertial waves are restored by the Coriolis force. At low frequencies, dissipation of the dynamical tide is several orders of magnitude higher than the dissipation of the equilibrium tide (Ogilvie & Lin 2007).

2018ApJ...866L...1S__Pecharromán_et_al._1999_Instance_4

It was found that the complex dielectric function from Pecharromán et al. (1999) for the sample obtained by heating bayerite at 1273 K, assuming a spheroid with depolarization parameters of (0.35, 0.003), produced an opacity with 11, 20, 28, and 32 μm features, so this component was included in the models. However, with only this component, the observed 20 μm features in the residual spectra were found to be wider than those in the models. By adding the opacity of the sample obtained by heating boehmite at 1173 K, the width of the 20 μm feature could be matched. This was done using the complex dielectric function for the sample obtained by heating boehmite at 1173 K from Pecharromán et al. (1999), assuming a spheroid with depolarization parameters of (0.35, 0.035). The complex dielectric functions of the samples obtained by heating bayerite and boehmite to various temperatures (Pecharromán et al. 1999) were derived by modeling the reflectance spectra of pellets obtained by pressing powders of these materials under great pressure. This method required Pecharromán et al. (1999) to assume an effective medium theory, such that a pellet is a mixture of one of their samples with a matrix of air. Pecharromán et al. (1999) noted that heating bayerite at 500°C eliminates the XRD pattern of bayerite, and they note that at 700°C, the infrared reflectance spectrum of the boehmite sample no longer shows OH− stretching bands. This must mean that the samples obtained from heating bayerite at 1273 K and from heating boehmite at 1173 K are no longer bayerite or boehmite, respectively. XRD performed by Pecharromán et al. (1999) of the sample of bayerite prepared at 1273 K suggests only θ-alumina was present, and their infrared and NMR spectroscopy confirms this. XRD of their sample obtained from heating boehmite to 1173 K (Pecharromán et al. 1999) suggests δ-alumina to be present, though some amounts of θ-alumina and α-alumina are present, as they deduce from XRD and infrared and NMR spectroscopy.

2022MNRAS.517.5541T__Mooley_et_al._2018a_Instance_1

GRB 170817A, the short GRB directly associated with the first gravitational wave signal GW170817 from a binary neutron star merger, is a great example presenting afterglow spectra with a beautiful single power-law extending from radio to X-ray (Abbott et al. 2017a, b; Goldstein et al. 2017; Savchenko et al. 2017). A successful launch of a relativistic jet in GRB 170817A is supported by the detection of superluminal motion of a compact radio source (Mooley et al. 2018b; Ghirlanda et al. 2019). The jet was observed from off-axis (i.e. the jet axis and the line of sight was misaligned), leading to a faint gamma-ray emission (Abbott et al. 2017c; Ioka & Nakamura 2018, 2019; Salafia & Ghirlanda 2022, and references therein) and early rising afterglow light curves (Alexander et al. 2017; Hallinan et al. 2017; Margutti et al. 2017; Troja et al. 2017). The slow rising afterglow light curves also rejected a uniform top-hat jet (Mooley et al. 2018a; Troja et al. 2018) and revealed a launch of a structured jet, while various jet structures can explain the observations (see, however, Lamb, Levan & Tanvir (2020) for an alternative scenario with a refreshed shock in an off-axis uniform jet). Often assumed structured jets are a Gaussian jet or a power-law jet (D’Avanzo et al. 2018; Gill & Granot 2018; Lamb & Kobayashi 2018; Lyman et al. 2018; Resmi et al. 2018; Ghirlanda et al. 2019; Lamb et al. 2019; Troja et al. 2019, 2020; Ryan et al. 2020). Other non-trivial jet structures such as hollow-cone jets and spindle jets are also possible candidates as recently discovered by Takahashi & Ioka (2020, 2021), while the consistency of a hollow-cone jet was also confirmed by Nathanail et al. (2020, 2021). As predicted by Troja et al. (2018), the rapid decline of the afterglow after the peak (Mooley et al. 2018c; Lamb et al. 2019; Hajela et al. 2020; Makhathini et al. 2021; Troja et al. 2020; Balasubramanian et al. 2021) revealed a distinctive signature of a successfully launched structured jet, setting it apart from a chocked jet scenario. The observed electron power-law index derived from the afterglow spectrum falls within the range predicted in a theory of particle acceleration at trans-relativistic shock (Keshet & Waxman 2005) as mentioned by Margutti et al. (2018) and is consistent with being constant in time within observational errors (D’Avanzo et al. 2018; Dobie et al. 2018; Margutti et al. 2018; Fong et al. 2019; Kilpatrick et al. 2022).

2016MNRAS.456.1901H__Yoo_&_Miralda-Escudé_2004_Instance_1

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

2016ApJ...827...93N__Neuhauser_et_al._2007_Instance_1

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

2018ApJ...855...48Q__Nagy_et_al._2017_Instance_1

The Orion Bar is probably the best studied PDR in our Galaxy. It is located between the Orion Molecular Cloud 1 and the H ii region excited by the Trapezium cluster, and is exposed to an FUV field a few 104 times the mean interstellar radiation field. Owing to its proximity (417 pc, Menten et al. 2007) and nearly edge-on orientation, the Bar provides an ideal laboratory for testing PDR models (e.g., Jansen et al. 1995; Gorti & Hollenbach 2002; Andree-Labsch et al. 2017) and a primary target for observational studies of physical and chemical structures of PDRs (e.g., Tielens et al. 1993; Walmsley et al. 2000; van der Wiel et al. 2009; Arab et al. 2012; Peng et al. 2012; Goicoechea et al. 2016; Nagy et al. 2017). Observations of various molecular spectral lines have shown that the emissions could be better interpreted with an inhomogeneous density structure containing an extended and relatively low density ( cm−3) medium and a compact and high-density ( cm−3) component (e.g., Hogerheijde et al. 1995; Young Owl et al. 2000; Leurini et al. 2006, 2010; Goicoechea et al. 2016). However, due to the scarcity of high-resolution observations capable of spatially resolving the density structure, the nature of the high-density clumps or condensations is still not well understood. Lis & Schilke (2003) mapped the Bar in H13CN (1–0) with the Plateau de Bure Interferometer (PdBI) at an angular resolution of about 5″, and detected 10 dense clumps. They proposed that the H13CN clumps are in virial equilibrium and may be collapsing to form stars. Goicoechea et al. (2016) performed Atacama Large Millimeter/submillimeter Array HCO+ (4–3) observations of the Bar and detected over-dense substructures close to the cloud edge, and found that the substructures have masses much lower than the mass needed to make them gravitationally unstable. These two interferometric observations both target molecular spectral lines. A high-resolution map of the dust continuum emission of the Bar, which is highly desirable in constraining the mass and density of the dense condensations, is still lacking. Here we report our Submillimeter Array (SMA) observations of the dust continuum and molecular spectral line observations of the Bar.

2015AandA...584A.103S__Gögelein_&_Müther_2007_Instance_1

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

2020AandA...641A.155V__Ceverino_et_al._2010_Instance_1

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

2016AandA...596A.113B__Perryman_et_al._1997_Instance_1

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

2018AandA...614A..48B__Keselman_&_Nusser_2012_Instance_1

The driving mechanisms and chronology of the buildup of bulges in late-type galaxies (LTGs) is an issue of key relevance to our understanding of galaxy evolution. According to our current knowledge on bulge demographics in the local universe, a large fraction of LTGs host pseudo-bulges (PBs; e.g., Gadotti 2009; Fisher & Drory 2011; Fernández Lorenzo et al. 2014) that substantially differ from classical bulges (CBs) in their spectrophotometric and kinematical characteristics. The latter resemble in many respects “old and dead” elliptical galaxies, lacking ongoing star-formation (SF), exhibit a spheroidal shape with inwardly steeply increasing surface brightness profiles (SBPs) being well approximated by the Sérsic (1963) fitting law with a high (≳3) exponent η, show stellar kinematics dominated by velocity dispersion (σ⋆) and obey the Kormendy (1977) scaling relations for normal elliptical galaxies (Fisher & Drory 2010). It is observationally established that CBs contain a super-massive black hole (SMBH) with a mass M∙ tightly correlating with their stellar mass ${\cal M}_{\star,\textrm{B}}, \sigma_{*}$M⋆,B,σ* and optical luminosity (Ho 2008; Kormendy & Ho 2013; see also Ferrarese & Merritt 2000). Traditionally, bulges were thought to invariably form early-on via violent quasi-monolithic gas collapse (Larson 1974) or mergers (Bender et al. 1992; Aguerri et al. 2001; Keselman & Nusser 2012) associated with vigorous nuclear starbursts (Okamoto 2012), with the disk gradually building up around them. Whereas this inside-out galaxy formation scenario appears consistent with important integral characteristics of CBs (e.g., their red colors), it does not offer a plausible explanation for the presence of PBs in present-day LTGs. These generally show ongoing SF, a significant degree of rotational support (Kormendy & Kennicutt 2004, for a review) and flatter/ellipsoidal shapes with nearly exponential SBPs (η≲2; e.g., Drory & Fisher 2007; Fisher & Drory 2010). Even though there is observational evidence that PBs also contain a SMBH (Kormendy et al. 2011; Kormendy & Ho 2013), in some cases revealing itself as an active galactic nucleus (AGN; e.g., Kotilainen et al. 2016; see Kormendy & Ho 2013 for a review), these do not follow the M∙ –σ* correlation for CBs, which appears to be consistent with a different formation route. Indeed, the prevailing concept on PB formation is that these entities emerge gradually out of galactic disks through gentle gas inflow spawning quasi-continuous SF and the emergence of a central bulge-like luminosity excess at their centers (e.g., Courteau et al. 1996; Carollo et al. 2001; Kormendy & Kennicutt 2004). Besides bar-driven gas inflow (e.g., Springel & Hernquist 2005), various other mechanisms, such as inward stellar migration, minor mergers with low-mass satellites, or a purely dynamical re-arrangement of the disk (Scannapieco et al. 2010; Guedes et al. 2013; Bird et al. 2012; Roskar et al. 2012; Grand et al. 2014; Halle et al. 2015) have been proposed as further contributors to PB growth along the Gyr-long secular evolution of LTGs.

2019AandA...630A..26M__Wozniakiewicz_et_al._(2012)_Instance_2

The majority of the particles collected by Stardust are olivine and pyroxene silicates with solar isotopic compositions, which suggests an origin in our solar system rather than an interstellar provenance. These polymineralic particles dominate those made of a single mineral even down to sizes smaller than 100 nm, indicating that the dust composition is surprisingly consistent at different scales and that the smallest subunits of the dust may be as small as tens of nanometers (Hörz et al. 2006; Zolensky et al. 2006). The sizes of these smallest single mineral impactors are similar to those of the nanocrystals determined by Rietmeijer (1993). As discussed above, they might also be existing in MIDAS dust particles and might be fused into the 100 nm features. Price et al. (2010) and Wozniakiewicz et al. (2012) investigated the sizes of particles smaller than 10 μm that impacted the aluminum foils of the Stardust probe. The distribution peaks at about 175 nm, but if we assume that the particles areagglomerates of smaller subunits, as indicated by their common polymineralic nature, then the subunit size distribution would peak at sizes below 100 nm (Price et al. 2010). A study of over 450 particles that do not seem to be agglomerates, that is, those that show single mineral impactors of silicate or sulfide, found geometric mean sizes of $532^{741}_{-310}$ 532 −310 +741 nm for the silicate particles and $406^{491}_{-222}$ 406 −222 +491 nm for the sulfides (Wozniakiewicz et al. 2013). These sizes are notably larger than the 175 nm (or less) found for the whole dataset. This large spread of subunit sizes could indicate a size distribution with a large width. No fits of these sizedistributions are available, but the figures in Wozniakiewicz et al. (2012) and Price et al. (2010) indicate that the differential sizes may follow a log-normal distribution. When we assume that the smallest subunit sizes are possibly between tens and hundreds of nanometers, the subunit size range found for MIDAS smallest features would be encompassed. The determination of the size distributions for the small Stardust particles and a detailed comparison to the distributions obtained for comet 67P could be the work of an interesting future project.

2021MNRAS.507..524M__Steidel_et_al._2016_Instance_1

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

2018ApJ...863..162M__Liu_et_al._2013_Instance_1

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.

2022AandA...663L...4L___2020_Instance_1

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

2016MNRAS.462.3945A__Ebbets_&_Savage_1982_Instance_1

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

2021AandA...653A.111R__Jones_et_al._(2021)_Instance_1

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

2017AandA...605A..88L__Cordiner_et_al._2015_Instance_1

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

2018MNRAS.473.1512A__Granot_&_Sari_2002_Instance_1

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

2018ApJ...863..162M__Liu_et_al._2013_Instance_3

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.

2021AandA...645A..96P__Marcantonio_et_al._(2018)_Instance_2

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

2022MNRAS.513.2349C__Bradač_et_al._2002_Instance_1

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

2021ApJ...913..115A__Linden_et_al._2012_Instance_1

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

2015MNRAS.451.2174T__Yoshida_et_al._2006_Instance_1

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

2018ApJ...864..158L__Zank_&_Matthaeus_1992_Instance_1

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

2019ApJ...874..166C__Dyks_et_al._2004_Instance_1

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

2018AandA...620A.122S__Švanda_et_al._(2014)_Instance_1

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

2020AandA...639A..46B__Štverák_et_al._(2009)_Instance_4

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

2021ApJ...916...64T__Lundstedt_et_al._2002_Instance_1

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

2021AandA...655A..99D__Bensby_&_Feltzing_2006_Instance_1

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

2016MNRAS.461.1719C__Fu_et_al._2012_Instance_2

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

2022MNRAS.517.1313M__Meidt_et_al._2018_Instance_1

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

2015ApJ...799..149J___2014_Instance_1

With our joint analysis of stellar mass fraction and source size, we find a larger stellar mass fraction than earlier statistical studies. In Figure 2, we compare our determination of the stellar surface density fraction to a simple theoretical model and to the best fit of a sample of lens galaxies by Oguri et al. (2014). The simple theoretical model is the early-type galaxy equivalent of a maximal disk model for spirals. We follow the rotation curve of a de Vaucouleurs component for the stars outward in radius until it reaches its maximum and then simply extend it as a flat rotation curve to become a singular isothermal sphere (SIS) at large radius (see details in the Appendix). The ratio of the surface mass density of the de Vaucouleurs component to the total surface mass density is shown as a dashed curve in Figure 2. We also show as a gray band the best fit for the stellar fraction in the form of stars determined by Oguri et al (2014) in a study of a large sample of lens galaxies using strong lensing and photometry, as well as the best model using a Hernquist component for the stars and an NFW halo for the dark matter with and without adiabatic contraction, also from Oguri et al. (2014). We have used the average and dispersion estimates for the Einstein and effective radii available for 13 of the objects in our sample from Oguri et al. (2014), Sluse et al. (2012), Fadely et al. (2010), and Lehár et al. (2000; see Table 1) as an estimate of RE/Reff in Figure 2. The average value and dispersion of the sample is RE/Reff = 1.8 ± 0.8. This also averages over the different radii of the lensed images. The agreement of our estimates with the expectations of the simple theoretical model and with estimates from other studies (Oguri et al. 2014) is quite good. For comparison, the estimate of Pooley et al. (2012; using the Einstein and effective radii estimates for 10 out of 14 of their objects from Schechter et al. 2014) seems somewhat lower than expected at those radii. The range of stellar mass fractions from MED09 for source sizes in the range 0.3–15.6 light days is also shown in Figure 2. In this case, the discrepancy between our estimate and their reported value of α = 0.05 is completely due to the effect of the source size. Although accretion disk sizes are known to be smaller in X-rays, recent estimates are in the range of 0.1–1 light-days, depending on the mass of the black hole (see Mosquera et al. 2013), and these finite sizes will increase the stellar surface densities implied by the X-ray data. Another possible origin for this discrepancy is that Pooley et al. (2012) use the macro model as an unmicrolensed baseline for their analysis. It is well known that simple macro models are good at reproducing the positions of images, but have difficulty reproducing the flux ratios of images due to a range of effects beyond microlensing. Recently, Schechter et al. (2014) found that the fundamental plane stellar mass densities have to be scaled up by a factor 1.23 in order to be compatible with microlensing in X-rays in a sample of lenses with a large overlap with that analyzed by Pooley et al. (2012). It is unclear how this need for more mass in stars at the position of the images found by Schechter et al. (2014) can be reconciled with the apparently low estimate of mass in stars at those radii by Pooley et al. (2012). Our estimate of the stellar mass fraction agrees better with the results of microlensing studies of individual lenses (Keeton et al. 2006; Kochanek et al. 2006; Morgan et al. 2008, 2012; Chartas et al. 2009; Pooley et al. 2009; Dai et al. 2010) that reported values in the range 8%–25%, and with the estimates from strong lensing studies (see for example Jiang & Kochanek 2007; Gavazzi et al. 2007; Treu 2010; Auger et al. 2010; Treu et al. 2010; Leier et al. 2011; Oguri et al. 2014) which produced stellar mass fractions in the range 30%–70% integrated inside the Einstein radius of the lenses.

2020MNRAS.499.2575E__Madau,_Shen_&_Governato_2014_Instance_1

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

2019ApJ...887..185S__Pushkarev_et_al._2010_Instance_1

The radio versus optical DCF peaks at non-zero lags (τdelay) with the optical emissions leading the radio emissions, implying that the radio and optical emission regions are not cospatial with the optical/IR emission region being closer to the base of the jet. This is consistent with the observation that optical emission regions are smaller than radio emission regions (Table 8). These observations are in accordance with standard shock-in-jet models where higher frequencies are emitted closer to the shock front while lower frequencies are produced from larger volumes that extend farther away from the shock (e.g., Marscher & Gear 1985; Marscher et al. 2008). This could also be understood as a manifestation of the position offset of optically thick features that can be interpreted as a frequency-dependent shift of the self-absorbed core of the jet (e.g., Lobanov 1998; Pushkarev et al. 2012). The linear separation of the V and 37 GHz emission region can be estimated using the relation (Pushkarev et al. 2010; Lisakov et al. 2017): 10 where βapp is the apparent jet speed and θ is the viewing angle. Using a range of 15–100 days for τdelay from Table 6, the maximum linear separation between the emission regions ( ) is estimated to be in the range of 4.06 (for Segment 6) to 27.04 pc (for Segment 1) and the corresponding projected separation varies from 0.18 to 1.20 pc using a viewing angle of θ = 13. The resulting angular separation is 0.022–0.151 mas. The simplest jet geometry is that of a conical jet. It cannot, however, explain the change in the separation between emission regions. In a conical jet geometry, the distance of an emission region from the central engine can be calculated using (Abdo et al. 2011) assuming the emission region fills the cross-section of the jet and the opening angle . The obtained dce is given in Table 8 and the separation comes out to be 1.4 pc. Thus, the conical jet model also severely underestimates the separation between the emission regions, as it does not take into account the jet collimation. An alternative model is that of an inhomogeneous curved jet, where synchrotron radiation of decreasing frequency is produced in an outer and wider jet region that changes orientation with time. It is possible that the long-term variability behavior of 3C 454.3 during our extended observation is dominated by geometrical effects that also lead to temporal delays between the radio and optical bands.

2015MNRAS.448..666S__Steidel_et_al._2014_Instance_1

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

2022MNRAS.512.3243S__Mason_&_Gronke_2020_Instance_1

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

2015AandA...584A.103S__Potekhin_et_al._2013_Instance_1

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

2016AandA...594A..64P__Judge_(2015)_Instance_2

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

2018MNRAS.478.2576M__Marleau_et_al._2017_Instance_1

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

2021ApJ...923..126S__Samsing_&_Ilan_2018_Instance_1

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

2021ApJ...916...68B__Chakrabarti_1989_Instance_1

In this paper, we presented our analysis of the spectral and timing behavior of GRS 1915+105 using the TCAF paradigm. For this purpose, the θ class data of the source as obtained by the LAXPC instrument of AstroSat satellite were used. To the best of our knowledge, this is the first time that the θ class data of AstroSat are analyzed with the TCAF paradigm. In this paradigm, different spectral states as well as the QPOs resulting from resonance oscillation of the Compton cloud arise out of the interplay between two types of accretion rates, namely, the Keplerian disk rate ( m d ) and the sub-Keplerian halo rate ( m h ). In the sub-Keplerian flow, shocks are formed when the Rankine–Hugoniot conditions are satisfied (Chakrabarti 1989, 1996; Chakrabarti & Das 2004). If the Keplerian disk rate increases, it increases the soft-seed photons, and the post-shock region (CENBOL, which acts as a “Compton cloud”) is cooled down rapidly, causing the shock to proceed toward the black hole. This has indeed been obtained in our spectral analysis. In the process, both the compressional heating timescale and the cooling timescale of the CENBOL decrease and become comparable (Chakrabarti et al. 2015), triggering shock oscillations that manifest as QPOs in the light curve. We first obtained the state of the system by estimating the photon index using the phabs*(diskbb+power-law) model. Throughout our analysis, the photon index (Γ) was found to be above 2.4, indicating that the source was either in the soft-intermediate state (SIMS) or in the soft state (SS). In the case of both orbits, Γ monotonically increases with the increase in photon flux, implying the transition from SS to SIMS. In the first orbit, Γ made an excursion from 2.5 to 3.0 during the 400 s span of spectral analysis, indicating the transition of the source from SIMS to SS. The gradual enhancement of the disk accretion rate and the consequent decline of the shock location were also obtained from the TCAF model fitting. The shock location moves from 39rs to 17rs, while the disk accretion rate increases from 0.77 to 0.84 M Edd . A similar trend was followed in the second orbit as well (Table 2). In the two orbits (02345 and 02346) we have analyzed, within a span of 400–500 s, the total flux changes by a factor of 3–5. The corresponding shock velocity is ∼2200 and 900 m s−1, respectively. This is significantly higher than the shock velocity obtained earlier in the case of transient sources, which lie within 10–20 m s−1 (Chakrabarti et al. 2008; Nandi et al. 2012; Debnath et al. 2013). This suggests the possibility that the local modulation of the accretion rate is due to feedback from the outflow (Chakrabarti & Manickam 2000; Chakrabarti & Nandi 2000). All these results indicate that GRS 1915+105 went through repeated microflares when it was in θ class.

2015ApJ...806..152S__Ransom_et_al._2005_Instance_3

One of the most astonishing characteristics of Liller 1 is the extremely large value of the collision rate parameter. Verbunt & Hut (1987) showed that Liller 1 has the second-highest value of stellar encounter rate (after Terzan 5; see also Lanzoni et al. 2010) among all star clusters in the Galaxy, thus suggesting that it represents an ideal environment where exotic objects, generated by collisions, can form. In fact, it is commonly believed that dynamical interactions in GCs facilitate the formation of close binary systems and exotic objects such as cataclysmic variables (CVs), low-mass X-ray binaries (LMXBs), millisecond pulsars (MSPs), and blue straggler stars (BSSs) (e.g., Bailyn 1992; Paresce et al. 1992; Ferraro et al. 2001, 2009a, 2012; Ransom et al. 2005; Pooley & Hut 2006). Moreover, Hui et al. (2010) found that clusters with large collisional parameters and high metallicity (see also Bellazzini et al. 1995) usually host more MSPs. Indeed, Terzan 5 hosts the largest population of MSPs among all Galactic GCs (Ransom et al. 2005). 6 6 Note that Terzan 5 is suspected to not be a genuine GC, because it harbors at least three stellar populations with different iron abundances (Ferraro et al. 2009a; Origlia et al. 2011, 2013; Massari et al. 2014). A strong γ-ray emission has been recently detected in the direction of Liller 1 by the Large Area Telescope (LAT) on board Fermi (Tam et al. 2011). This is the most intense emission detected so far from a Galactic GC, again suggesting the presence of a large number of MSPs. However, no direct radio detection of these objects has been obtained so far in this system (Ransom et al. 2005). The only exotic object identified in the cluster is the rapid burster MXB 1730-335, an LMXB observed to emit radio waves and type I and type II X-ray bursts (Hoffman et al. 1978). It seems to be located in the central region of Liller 1, but no optical/IR counterpart of this object has been found so far (Homer et al. 2001).

2021AandA...647A..49S__Murphy_et_al._2015_Instance_1

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.

2018ApJ...864..154D__Liu_et_al._2016_Instance_1

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.

2021MNRAS.503.2108P__Andresen_et_al._2017_Instance_2

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.

2015MNRAS.452.1112E__Illarionov_&_Sunyaev_1975_Instance_1

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).

2022ApJ...928....3A__Rappazzo_et_al._2019_Instance_1

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).

2021MNRAS.503.5367B__Lu,_Kumar_&_Zhang_2020_Instance_2

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).

2019MNRAS.485.5652D__Yamanaka_et_al._2009_Instance_1

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.

2017ApJ...850...81P__Muñoz-Jaramillo_et_al._2012_Instance_1

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).

2020AandA...641A.123H__Mikal-Evans_et_al._(2019)_Instance_4

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.

2021MNRAS.504.4626K__Kraljic_et_al._2020b_Instance_3

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).

2017ApJ...835..151A__Gaisser_et_al._2014_Instance_1

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.

2020ApJ...899L...6L__Margalit_et_al._2019_Instance_2

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).

2018ApJ...863..162M__Liu_et_al._2013_Instance_1

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.

2021MNRAS.502..772L__Cowley_et_al._2015_Instance_1

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).

2016AandA...588A..74C__Böhm-Vitense_1958_Instance_1

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.

2015AandA...574A..70F__Padova_1994_Instance_1

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.

2022AandA...667A.131B__Izumi_et_al._(2016)_Instance_3

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).

2017ApJ...849...63R__Pascucci_et_al._2016_Instance_1

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).

2015MNRAS.454.1644L__Kotze_&_Charles_2012_Instance_1

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.

2021AandA...653A..99T__Pallé_et_al._2004_Instance_1

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.

2022ApJ...940...86K__Cohen_et_al._1997_Instance_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.

2022MNRAS.517.3881L__GRB_2004_Instance_1

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.

2016AandA...587A..30M__Mihalas_&_Binney_1981_Instance_1

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.

2021MNRAS.508.1020D__Schuhmann_et_al._2019_Instance_1

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.

2016ApJ...832....1R__Plekan_et_al._2011_Instance_1

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.

2021ApJ...912...20J__Murray_et_al._1995_Instance_1

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 ∼ 4.7 × 10 39 M 7 2 / 3 erg s − 1 , where M 7 2 / 3 is the black hole mass in 10 7 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.

2019AandA...623A.145M__Lofthouse_et_al._2017_Instance_1

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.

2015AandA...584A...4D__Schneider_et_al._2015_Instance_1

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.

2017MNRAS.472.1052B__Bruno_&_Telloni_2015_Instance_1

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).

2018ApJ...869...69M__Künzel_1960_Instance_1

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.

2020MNRAS.493..559B__Kulow_et_al._2014_Instance_1

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)

2017AandA...607A.107W__Michałowski_et_al._2015_Instance_1

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.

2021MNRAS.506.3511M__Foucart_et_al._2012_Instance_1

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).

2015MNRAS.454.1644L__Kotze_&_Charles_2012_Instance_3

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.

2022AandA...666A.112L__Cormier_et_al._2015_Instance_1

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).

2022AandA...663A.105P__Bonafede_et_al._2012_Instance_2

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.

2021MNRAS.501.3781R__Nisini_et_al._2005_Instance_2

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).

2022ApJ...932...98S__hand,_FRB_2018_Instance_1

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.

2021ApJ...908...40M__Muschietti_&_Lembège_2017_Instance_1

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.

2016ApJ...821..107G__Gloeckler_&_Fisk_2015_Instance_4

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.

2016ApJ...824..138Y__yanin_2011_Instance_1

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.

2016MNRAS.457.3084M__Storm_et_al._2005_Instance_1

The populous blue LMC cluster NGC 1866 is already known to host an exceptionally rich sample of more than 20 Cepheids (Musella et al. 2006; Welch & Stetson 1993). One of these was identified by Musella et al. (2006) in a preliminary analysis of the proprietary BVI Very Large Telescope (VLT) data. It is unquestionable that such a unique sample of Cepheids – likely all members of the cluster and at the same distance, chemical composition and age – would constitute a milestone in our understanding of the Cepheid pulsational scenario. Indeed, it offers an unprecedented opportunity to investigate both empirical and theoretical estimates of the luminosity and colour of the pulsating structures and their relation with the observed periods. For this reason, many authors have studied the NGC 1866 Cepheids (Welch 1991; Welch & Stetson 1993; Gieren, Richtler & Hilker 1994; Walker 1995; Gieren et al. 2000; Storm et al. 2005; Testa et al. 2007) in both the optical and near-infrared bands, and have tested different methods to calibrate the PL relations in different filters. In Brocato et al. (2004), we already discussed the sample of the 23 known Cepheids in NGC 1866, concluding that unfortunately only 4–6 Cepheids had light curves accurate enough to allow a meaningful determination of their luminosities and colours. On the basis of such a tantalizing situation, we took advantage of assigned observing time at the ESO VLT to perform an accurate photometric investigation of the cluster field, with the aim of securing suitable data constraining the light curves of all the member Cepheids. Moreover, to get accurate information about radial velocities and chemical abundances of the stars in NGC 1866, we have performed FLAMES@VLT (Fibre Large Array Multi Element Spectrograph mounted on VLT) spectroscopic observations for 30 stars (19 belonging to the cluster and 11 to the LMC field), including three Cepheids (Mucciarelli et al. 2011; Molinaro et al. 2012). Mucciarelli et al. (2011) found that, as far as the chemical composition is concerned, the cluster stars are reasonably homogeneous. Indeed, they appear to share the same abundances within the uncertainties, and this property is independent of the evolutionary status. The average iron abundance is [Fe/H] = −0.43 ± 0.01 dex, with a dispersion σ = 0.04 dex. For the three spectroscopically investigated Cepheids Molinaro et al. (2012), adopting the same procedure used in Mucciarelli et al. (2011), found values fully consistent with the average iron content. Moreover, Molinaro et al. (2012) applied the CORS Baade–Wesselink (BW) method (Ripepi et al. 1997) to a sample of 11 Cepheids, using radial velocities obtained both from our FLAMES investigation and from literature data, and light curves based on a part of the UBVI data used in this paper complemented with K data by Testa et al. (2007). In this way, they obtained a direct estimate of the distance modulus of NGC 1866, μ0 = 18.51 ± 0.03 mag (see Molinaro et al. 2012, for details).

2020MNRAS.494.2465B__Hornik_1991_Instance_1

Here we demonstrate that, over a fixed time interval, the planar three-body problem can be solved by means of a multilayered deep artificial neural network (ANN; e.g. see LeCun, Bengio & Hinto 2015). These networks are designed for high-quality pattern recognition by mirroring the function of our brains (McCulloch & Pitts 1943; Rosenblatt 1985) and have been successfully applied to a wide variety of pattern recognition problems in science and industry, even mastering the game of Go (Silver et al. 2016). The abundance of real-world applications of ANNs is largely a consequence of two properties: (i) an ANN is capable of closely approximating any continuous function that describes the relationship between an outcome and a set of covariates, known as the universal approximation theorem (Cybenko 1989; Hornik 1991); and (ii) once trained, an ANN has a predictable and a fixed computational burden. Together, these properties lead to the result that an ANN can be trained to provide accurate and practical solutions to Newton’s laws of motion, resulting in major improvements in computational economy (Lee, Sode-Yome & Park 1991) relative to modern technologies. Our proof-of-principle method shows that an ANN can accurately match the results of converged solutions found using the arbitrary precision numerical integrator that, for computationally challenging scenarios, e.g. during multiple close encounters, can offer numerical solutions at a fraction of the time cost and CO2 expense. We demonstrate the importance of training an ANN on converged solutions. This enables the trained ANN to accurately predict particle locations even when a conventional ‘double-precision’ numerical integrator fails dramatically. By training an ANN that can accurately compute particle trajectories during close encounters, our work extends previous work training neural networks on an n-body-type problem (e.g. Quito, Monterola & Saloma 2001; Battaglia et al. 2016). Our findings also add to the growing body of literature that supports machine learning technologies being developed to enrich the assessment of chaotic systems (Pathak et al. 2018; Stinis 2019) and providing alternative approaches to classical numerical solvers more broadly (Hennig, Osborne & Girolami 2015).

2016ApJ...817...12P__Sur_et_al._2007_Instance_1

Large-scale magnetic fields with strength of the order of 1–10 μG have been observed in disk galaxies (e.g., Beck et al. 1996; Fletcher 2010; Beck 2012; Beck & Wielebinski 2013; Van Eck et al. 2015). The origin of these fields can be explained through mean-field dynamo theory (Ruzmaikin et al. 1988; Beck et al. 1996; Brandenburg & Subramanian 2005a; Kulsrud & Zweibel 2008). The conservation of magnetic helicity is one of the key constraints in these models, and also leads to the suppression of the α-effect. The operation of the mean-field dynamo automatically leads to the growth of magnetic helicity of opposite signs between the large-scale and small-scale magnetic fields (Pouquet et al. 1976; Gruzinov & Diamond 1994; Blackman & Field 2002). To avoid catastrophic suppression of the dynamo action (α-quenching), the magnetic helicity due to the small-scale magnetic field should be removed from the system (Blackman & Field 2000, 2001; Kleeorin et al. 2000). Mechanisms suggested to produce these small-scale magnetic helicity fluxes are: advection of magnetic fields by an outflow from the disk through the galactic fountain or wind (Shukurov et al. 2006; Sur et al. 2007; Chamandy et al. 2014), magnetic helicity flux from anisotropy of the turbulence produced by differential rotation (Vishniac & Cho 2001; Subramanian & Brandenburg 2004, 2006; Sur et al. 2007; Vishniac & Shapovalov 2014), and through diffusive flux (Kleeorin et al. 2000, 2002; Brandenburg et al. 2009; Mitra et al. 2010; Chamandy et al. 2014). The outflow of magnetic helicity from the disk through dynamo operation leads to the formation of a corona (Blackman & Field 2000). According to Taylor's hypothesis, an infinitely conducting corona would resistively relax to force-free field configurations under the constraint of global magnetic helicity conservation (Woltjer 1960; Taylor 1974; Finn & Antonsen 1983; Berger & Field 1984; Mangalam & Krishan 2000). In this paper, we include advective and diffusive fluxes in a simple semi-analytic model of a galactic dynamo that transfers magnetic helicity outside the disk and consequently builds up a force-free corona in course of time. We first solve the time-dependent dynamo equations by expressing them as separable in variables r and z. The radial part of the dynamo equation is solved using an eigenvector expansion constructed using the steady-state solutions of the dynamo equation. The eigenvalues of the z part of the solution are obtained by solving a fourth-order algebraic equation, which primarily depends upon the turbulence parameters and the magnetic helicity fluxes. Once the dynamo solutions are written out as parametric functions of these parameters, the evolution of the mean magnetic field is computed numerically by simultaneously solving the dynamical equations for α-quenching and the growth of large-scale coronal magnetic helicity. Since the large-scale magnetic field lines cross the boundary between the galactic disk and the corona, the magnetic helicity of the large-scale magnetic field in the disk volume is not well defined. Hence we use the concept of gauge-invariant relative helicity (Finn & Antonsen 1983; Berger & Field 1984; Berger 1985) to estimate the large-scale magnetic helicity in the disk and the corona. Here the gauge-invariant relative helicity for the cylindrical geometry is calculated using the prescription given in Low (2006, 2011). We then investigate the dependence of the saturated mean magnetic field strength and its geometry on the magnetic helicity fluxes within the disk and the corresponding evolution of the force-free field in the corona.

2022AandA...667A...5R__Bakx_et_al._2020_Instance_1

The characterization of the dynamics of galaxies just after (4 ≲ z ≲ 6) and within the Epoch of Reionization (z ≳ 6) is still in its infancy as spatially resolved observations targeting emission lines are available only for about ten of sub-mm sources (e.g., Rizzo et al. 2020; Lelli et al. 2021; Rizzo et al. 2021) and a handful of Lyman break galaxies (e.g., Jones et al. 2017; Fujimoto et al. 2021). Instead, the number of marginally resolved observations aiming at inferring the integrated properties of early galaxies has increased in the last five years (Smit et al. 2018; Le Fèvre et al. 2020; Bouwens et al. 2022). In their pioneering work, Smit et al. (2018) used marginally resolved Atacama Large Millimeter/Submillimetre Array (ALMA, Wootten & Thompson 2009) observations of the [C II]-158 μm emission line to show that two star-forming galaxies at z ∼ 7 have smooth velocity gradients and interpreted them as being rotationally supported disks. Similar results are obtained in studies of individual galaxies at z ∼ 6–8 (Bakx et al. 2020; Harikane et al. 2020). However, given the low angular resolution of these data, they can not rule out the possibility that one or more merging [C II]-bright satellites are mimicking the smooth gradient observed in the velocity maps (see discussion in Smit et al. 2018; Simons et al. 2019). Instead, a gradient in the velocity fields of two Lyman break galaxies at z = 6.1 and 7.1, combined with the identification of two compact components, has been interpreted as evidence of mergers (Jones et al. 2017; Hashimoto et al. 2019). Recently, Le Fèvre et al. (2020) and Romano et al. (2021) performed the first systematic morpho-kinematic analysis of a statistically significant sample of z ∼ 4–6 main-sequence galaxies from the ALMA Large Program to INvestigate [C II] at Early times (ALPINE) survey. After combining the morphological and kinematic analysis of the [C II] observations with the rest-frame UV and optical data, Romano et al. (2021) conclude that 23 out of the 75 ALPINE galaxies are merging systems. This large fraction of mergers in the ALPINE sample could imply a significant contribution of major mergers to the mass assembly in the early Universe (Le Fèvre et al. 2020; Romano et al. 2021), confirming previous results based on the count of close-pair galaxies in photometric surveys (Mantha et al. 2018; Duncan et al. 2019). However, due to the sensitivity and angular resolutions of the ALPINE observations, the kinematic characterization and, consequently, the merger fraction may be uncertain.

2015MNRAS.450.4364N__Wu_et_al._2004_Instance_2

Low- and intermediate-mass stars are formed by the gravitational collapse of the parental giant molecular cloud (GMC), followed by the accretion process (Palla 1996). During the accretion phase, material is ejected as well via collimated bipolar jets. However, when a YSO reaches 8 M⊙, the radiative flux becomes so intense (using ϕ = L/4πd2, the ratio between the radiative fluxes of an O5 and a B3 star – masses of ∼40 and ∼8 M⊙, respectively – is ≈250) that it may interrupt the accretion flow. A process that constrains the outcoming radiation field to narrower angles may leave some room for the accretion process to continue in some directions. This seems to be the case for the outflows driven by young stars from a very broad mass range, as previous reported by several authors (Bachiller 1996; Bontemps et al. 1996; Shepherd & Churchwell 1996; Beuther et al. 2002; Wu et al. 2004). Outflows associated with high-mass objects are expected to be more energetic than the outflows observed in lower mass YSOs (Beuther et al. 2005; Zhang et al. 2005; López-Sepulcre et al. 2009), with velocities greater than ∼100 km s−1 (Martí, Rodríguez & Reipurth 1998). Some authors have found evidences that outflows associated with massive stars are scaled up versions of their low-mass counterparts (Vaidya et al. 2011; Codella et al. 2013) while other works have reported that no well-collimated outflows have been found towards MYSOs (Shepherd, Testi & Stark 2003; Sollins et al. 2004). Massive YSO outflows mapped in high-velocity CO lines have collimation factors R = length/width ∼2.05 ± 0.96 as compared to R ∼ 2.81 ± 2.16 for low-mass stars (Wu et al. 2004), indicating a weak tendency that outflows associated with massive stars are less collimated than those from low-mass stars as previously thought (Richer et al. 2000). Besides the degree of collimation, these massive outflows would work removing mass from the plane of the accretion disc, lowering the density on the plane and, therefore, facilitating the accretion flow to reach the stellar core as shown in the recent 3D simulations presented by Krumholz et al. (2009). Although these authors have not included the outflow activity on their simulations, they argue that the presence of outflows would decrease the star formation efficiency from 70 per cent (considering purely radiation effects) to 50 per cent.

2021AandA...645A..95H__Boese_2000_Instance_1

Next we chose an optimum cut-off radius for the detector FOV. The PSPC has a circular FOV with a radius 57′. The PSPC entrance window has a rib support structure with an inner ring at a radius corresponding to 20′ (Pfeffermann et al. 1987; Hasinger & Zamorani 2000). Both the ROSAT telescope angular resolution and its vignetting function are roughly constant within the inner 20′ ring, but degrade significantly towards larger off-axis angles. The combined detector and telescope PSFs are described in detail in Boese (2000). To the first order, the PSF at each off-axis angle can be approximated by a Gaussian function with a half power radius (HPR) of 13, 22, 52, 93, 130, and 180″, at off-axis angles of 0, 12, 24, 36, 48, and 57′, respectively (at 1 keV). The vignetting function at 1 keV drops almost linearly to about 50% at an off-axis angle of 50′. Taking into account all these effects, the HPR of the overall RASS PSF is 84″ (Boese 2000). This means that the classical confusion limit (40 beams per source) is reached at a source density of about 15 sources deg−2, which is exceeded in the high-exposure areas of our survey. In addition, we need to optimally discriminate between extended and point-like X-ray sources, calling for an angular resolution that is as high as possible. We therefore have to reduce the detector FOV. The sharpest imaging is achieved within the inner 20′ of the PSPC FOV, corresponding to the inner ring-like rib of the PSPC support structure (see Fig. 1). However, there is a trade-off between image sharpness and the number of photons required for detection and image characterization. In particular in the outer areas of our survey, where the RASS exposure times drop significantly, a 20′ FOV radius does not provide sufficient exposure time. Taking into account the various competing factors in this trade-off, we made a few tests varying the FOV cut-off radius, and finally decided on an optimum FOV radius of 30′. The PSPC detector coordinates have a pixel size of 0.934″. We thus removed all X-ray events from the dataset, which are further than 1925 pixels from the PSPC centre pixel coordinate [4119,3929]. A similar cut had to be applied to the modified PSPC instrument map (MOIMP), which is used later for the construction of the survey exposure map.

2018AandA...613A..76J__Kennedy_&_Kenyon_2008_Instance_2

One of the most intriguing results from RV surveys is the observed scarcity of relatively close-in (a ≲ 0.5 AU) planets around post-MS stars. This observational trend has been attributed to the strong tidal torque exerted by the star as its radius grows during the giant phase. As a result, planets are expected to lose orbital angular momentum, thus moving inward until they are evaporated in the stellar atmosphere (Livio & Soker 1983; Sato et al. 2008; Villaver & Livio 2009; Kunitomo et al. 2011). On the other hand, the majority of the giant stars targeted by RV surveys are intermediate-mass stars (M⋆ ~ 1.5–3.0 M⊙), thus they are the post-MS counterpart of A and early F main-sequence stars. Therefore, their companions should not be directly compared to those orbiting solar-type stars. Based on this analysis, known planets orbiting field giant stars are expected to be born in different conditions from those around low-mass stars. In particular, these planets are formed in more massive disks (since Md ∝ M⋆; Andrews et al. 2013), from which they can efficiently accrete a significant amount of gas, becoming gas giants (e.g., Kennedy & Kenyon 2008). In addition, due to the higher gas accretion rate (Muzerolle et al. 2005) and higher irradiation, these disks have shorter dissipation timescales (Currie 2009; Kennedy & Kenyon 2009) and the snow line is located at a greater distance from the central star (Kennedy & Kenyon 2008). As a consequence, these planets are most likely formed at greater orbital distances and, due to the shorter disk timescale, inward migration is halted; they thus reach their final position at a relatively large distance from the parent star. For comparison, Currie (2009) predicted that only ~1.5% of intermediate-mass stars host giant planets with a ≲ 0.5 AU, while ≳7.5% of them host at least one gas giant at a ≳ 0.5 AU. Fig. 9 shows the mass versus the orbital distance of planets detected around giant stars (log g ≲ 3.5), via RV measurements (black dots) and by the transit method (red open circles). We note that values of the RV detected systems correspond to the minimum planet mass (Mp sini). The dotted line represents a radial velocity semi-amplitude of K = 30 m s−1 for a 1.5 M⊙ star, (corresponding to a 3-σ detection; e.g., Hekker et al. 2006). As can be seen, there is only one companion detected via RVs interior to 0.1 AU, and the rest of them reside at an orbital distance a ≳ 0.4 AU. As discussed above, this observational result might be explained by the engulfment of the innermost planets as the parent star evolves off the MS and becomes a giant star. However, since a similar trend is observed in less evolved subgiants whose radii have not yet reached a value where tidal interactions are strong enough to affect the orbits of their companions, Johnson et al. (2007) argued that this is probably explained by a different formation scenario between planets around low-mass stars and those formed in more massive disks. From Fig. 9 it is also evident that planets residing interior to ~0.1 AU are significantly less massive (Mp ≲ 1 MJ) than those orbiting at a greater distance. In fact, two of these transiting planets are well below the 3-σ detection threshold, thus they are not detectable via radial velocities. A similar trend is also observed in MS stars (Zucker & Mazeh 2002), which might be caused by a decrease in the type II migration speed with increasing planetary mass, i.e., d a∕dt ∝ M $_P^{-1}$ P−1 (Mordasini et al. 2009). This theoretical prediction naturally explains why the most massive planets are found at a ≳ 0.4 AU. On the other hand, the mass distribution of the parent stars of these two populations of planets are different. While the mean stellar mass of the RV detected planets is 1.78 M⊙, this value is only 1.38 M⊙ for the transiting systems and thus two distinct planet mass distributions are expected to be found. Moreover, a similar result is observed between the mass of planets orbiting subgiant and giant stars (planets around giant stars being significantly more massive than those around subgiants; see Jones et al. 2014). In fact, the mean mass of the subgiant parent stars is 1.5 M⊙, significantly lower than giant host stars. These results provide further observational support of a different formation and migration scenario for planets at different host star mass. This result suggest that the observed lack of planet around giant stars is mainly due to the primordial distinct formation scenario proposed by Johnson et al. (Johnson et al. (2007)).

2017ApJ...850...97B__Haynes_et_al._2011_Instance_1

The H i mass fraction of every gas particle in the baryonic runs is calculated based on the particle’s temperature and density and the cosmic UV background radiation flux while including a prescription for self-shielding of H2 and dust shielding in both H i and H2 (Christensen et al. 2012). This allows for the straightforward calculation of the total H i mass of each simulated galaxy. We create mock H i data cubes only for the 42 halos that contain . Specifically, we create mock data cubes that mimic ALFALFA observations (Haynes et al. 2011). After specifying a viewing angle (see below), our code considers the line-of-sight velocity of each gas particle. The velocity of each particle is tracked in the simulation by solving Newton’s equations of motion, but any turbulent velocity of the gas is not taken into account. Velocity dispersions in dwarf galaxies can be on the order of the rotational velocity, ∼10–15 km s−1 (e.g., Stanimirović et al. 2004; Tamburro et al. 2009; Oh et al. 2015). Dispersions are thought to be driven at least partially by thermal velocities or supernovae (Tamburro et al. 2009; Stilp et al. 2013a, 2013b). In our simulations, supernovae inject thermal energy, and the thermal state of the H i gas needs to be considered in the mock H i linewidth for a realistic comparison to observations. To account for the thermal velocity, the H i mass of each gas particle is assumed to be distributed along the line-of-sight in a Gaussian distribution with a standard deviation given by the thermal velocity dispersion, , where T is the temperature of the gas particle. After this thermal broadening is calculated, a mock H i data cube can be generated by specifying the spatial and velocity resolution. For all of our galaxies, we adopt a spatial resolution of 54 pixels across 2Rvir. In practice, this corresponds to a range of ∼1 kpc resolution in our lowest-mass galaxies up to ∼9 kpc resolution in our most massive galaxies. However, the spatial resolution plays no role in our study, since measurements of the VF are based on spatially unresolved H i data. For the velocity resolution, we match the ALFALFA specification of 11.2 km s−1 (two-channel boxcar-smoothed).

2021MNRAS.505.2561C__Hunt_et_al._2018_Instance_1

Several mechanisms have been suggested to explain the formation of moving groups. A common explanation is that these velocity structures are the remnants of open clusters, or formed by interactions with a bar (Eggen 1965; Dehnen 2000). One problem with the cluster formation idea is that stars in moving groups can have a variety of different ages and compositions, so it is unlikely that they all came from the same cluster (Eggen 1965; Kushniruk et al. 2020). Analysis of GALAH data (Quillen et al. 2018) indicates that some moving groups, such as the Hercules moving group, may be due to a resonant bar. It has also been suggested that moving groups could have been formed from perturbations due to the Magellanic Clouds via gravitational interactions (Dehnen 1998). Recent work also finds that transient spiral structure (Hunt et al. 2018) may lead to the formation of moving groups, as well as perturbations due to spiral arms in the MW (Michtchenko et al. 2018). Moving groups in Gaia data have also been identified and analysed in action space. In the (JR, Jz) plane there are at least seven overdensities that follow lines of constant slope in this plane, which correspond to known moving groups in the solar neighbourhood (Trick, Coronado & Rix 2019). It is likely that there may be multiple causal mechanisms at play in the formation of moving groups in the Milky Way. The analysis of Gaia DR2 data has revealed many facets of a Galaxy that are clearly out of equilibrium, including the so-called phase-space spiral (Antoja et al. 2018), and the Enceladus merger (Helmi et al. 2018), that have been interpreted as arising from interactions with dwarf galaxies. Analysis of Gaia DR2 data also led to the discovery of a new dwarf galaxy (Torrealba et al. 2019) that likely interacted with the Milky Way (Chakrabarti et al. 2019). However, the formation of moving groups due to dwarf galaxy interactions has not yet been studied with full N-body simulations. Motivated by these earlier works that indicate that the MW may has been perturbed by dwarf galaxies, we focus our study here in trying to understand if some of the moving groups in the Galaxy may have arisen from dwarf galaxy interactions.

2019AandA...632A..37B__Desidera_et_al._2015_Instance_1

Of the almost 4000 exoplanets known today, more than 3700 were discovered via the radial-velocity (RV) or transit methods1. According to the database, only three of them (V830 Tau b, Donati et al. 2016; K2-33 b, David et al. 2016, and TAP 26 b, Yu et al. 2017) are younger than 100 Myr. The main reason for this is the strong stellar activity of young stars, which makes it hard to find the subtle planetary signal in the large stellar variations. This is unfortunate for two reasons: first, planet formation takes place in young systems and at least gas giants need to form before the disk has dissipated after less than a few tens of millions of years (e.g., Ercolano & Pascucci 2017). Second, planets at large orbital distances (≳50 AU) are almost exclusively detected via direct imaging (DI), which is best applicable to young systems where the planets are still hot from their formation. Thus, in order to discover all planets in a system, one either needs to image old stars – which seems currently impossible given the already low detection rate around young stars probed by large DI surveys (e.g., Desidera et al. 2015; Lannier et al. 2016; Tamura 2016; Stone et al. 2018) – or try to minimize the impact of the stellar activity of young stars. A lot has been done to understand and characterize stellar activity (e.g., Dumusque 2018). Lindegren & Dravins (2003) further estimate the effects of stellar activity such as oscillation, granulation, meridional flow, long-term magnetic cycle, surface magnetic activity and rotation, gravitational redshift and many more on the RV measurement. Meunier & Lagrange (2019) then try to model the effect of this kind of activity signal on RV data of mature stars. With our data probing activity timescales of days to years, we are mainly probing the combined effect of stellar rotation, reconfiguration of active regions, and long-term magnetic cycles. Still, large uncertainties remain in the prediction and interpretation of any RV signal, in particular for young pre-main sequence stars. But since this is what we measure, knowledge about the typical RV variability is important, for example for developing and testing RV activity models or planning RV surveys. In this paper we therefore derive a model-free analytic relation between stellar jitter, stellar age, and lag, where lag denotes the timescale on which the jitter is measured.

2020MNRAS.493.4868L__Stephens_et_al._2019_Instance_1

Recently, polarized (sub)millimetre emission has been detected in an increasing number of discs by Atacama Large Millimeter/submillimeter Array (ALMA) with its high sensitivity and angular resolution. However, the origin of disc polarization remains uncertain, since grains do not have to be aligned with just the magnetic field (Kataoka et al. 2017; Yang et al. 2019). They may also be aligned in the direction of the radiative anisotropy (Lazarian & Hoang 2007a; Tazaki, Lazarian & Nomura 2017) or the drift velocity of the grains relative to the ambient gas (Gold 1952; Lazarian 1995; Lazarian & Hoang 2007b). Furthermore, even spherical grains can produce polarized emission by self-scattering of large grains in an anisotropic radiation field (Kataoka et al. 2015; Yang et al. 2016, 2017; Stephens et al. 2019). The scattering interpretation of the disc polarization is favoured in several targets (e.g. Stephens et al. 2014, 2017; Kataoka et al. 2016; Bacciotti et al. 2018; Dent et al. 2019; Girart et al. 2018; Harris et al. 2018; Hull et al. 2018; Lee et al. 2018). One way to gauge the effects of scattering and identify polarization from aligned grains would be to observe at multiple wavelengths since the efficiency for scattering for grains of given sizes decreases rapidly with the wavelength in the optically thin and small-particle (or Rayleigh scattering) limit. Indeed, in the disc of Class I protostar BHB 07-11, Alves et al. (2018) detected polarization with ALMA at three wavebands (Bands 3, 6, and 7 or ∼ 3, 1.3, and 0.87 mm, respectively) with consistent polarization orientations across three bands and increasing polarization fraction with wavelength, which is generally not expected for scattering-induced polarization. The rather high mean polarization fractions (∼7.9, 5.3, and 3.5 ${{\ \rm per\ cent}}$ for Bands 3, 6, and 7 respectively) are also higher than those typically produced in models of scattering-induced disc polarization ($\sim \!1 {{\ \rm per\ cent}}$). At least for this well-studied source, scattering is unlikely the main mechanism for producing the observed multiwavelength disc polarization and aligned grains are favoured.

2022ApJ...939..117Z__Malkan_&_Moore_1986_Instance_1

Blazars are a subclass of active galactic nuclei (AGNs) with relativistic jets of high-energy particles pointing near our line of sight (e.g., Urry & Padovani 1995). Their nonthermal emission is generally detected across the entire electromagnetic spectrum from radio to γ-ray bands. Blazars are subclassified into flat-spectrum radio quasars (FSRQs) and BL Lac objects (BL Lacs), according to the equivalent width of the emission lines in their optical spectrum (Stickel et al. 1991; Stocke et al. 1991; Marcha et al. 1996). These two subclasses of blazars are thought to be intrinsically different, perhaps based on their accretion mode (Dermer & Giebels 2016). FSRQs have high luminosity and a thin and radiatively efficient black hole accretion disk (Malkan & Moore 1986), while BL Lacs are powered by an advection-dominated, low radiative efficiency accretion flow (Dermer & Giebels 2016; Blandford et al. 2019). The jet emission is relativistically beamed (Ghisellini 2019), with a Doppler boosting factor corresponding to a bulk Lorentz factor of several to greater than 10 (Pushkarev et al. 2009). In both cases, the broadband spectra consist of two broad humps, one peaking in the IR-to-X-ray regime and the other peaking in the γ-ray regime. The low-energy peak is believed to be due to synchrotron emission, while the high-energy peak is likely due to inverse Compton scattering of low-energy photons of either the same synchrotron photons (for BL Lacs) or external photons from the disk/BLR (for FSRQs) (e.g., Dutka et al. 2017). However, some blazars might not necessarily be detected in γ-rays (e.g., Paliya et al. 2017). Indeed, a recent study showed that blazars undetected in γ-rays are likely to have relatively smaller Doppler factors and more disk dominance (Paliya et al. 2017). In the case of strong Compton scattering, the beaming of γ-rays could be larger than, e.g., that seen in the radio (Dermer 1995), leading to the possible nondetection (or reduced detection efficiency) of γ-rays from sources not seen exactly pole-on.

2022AandA...663A...4S__Jönsson_et_al._2020_Instance_1

The pioneer of spectroscopic surveys, the Radial Velocity Experiment (RAVE), produced its first data release (DR) 15 years ago (Steinmetz et al. 2006) and its final one, DR6, last year (Steinmetz et al. 2020a,b). In the meantime, other surveys have been operated, and survey designers have developed new methodologies, learning progressively from their own and one another’s experience on how to reduce biases and uncertainties in the automated determination of APs and abundances from massive datasets of stellar spectra with various resolutions and spectral coverages. Besides RAVE, other surveys have published successive DRs that are available for public use. At the time of writing, there is open access to the DR16 of the Apache Point Observatory Galactic Evolution Experiment (APOGEE; Jönsson et al. 2020), the DR3 of the Galactic ArchaeoLogy with HERMES project (GALAH; Buder et al. 2021), the DR6 of RAVE (Steinmetz et al. 2020a,b), the DR5 of the Large sky Area Multi-Object Fibre Spectroscopic Telescope (LAMOST; Luo et al. 2015, 2019), Sloan Extension for Galactic Understanding and Exploration (SEGUE; Yanny et al. 2009), and the Gaia-ESO Survey (DR3; Gilmore et al. 2012; Randich & Gilmore 2013). Additional versions of APs, based on different methods, are also provided for RAVE DR6 (Guiglion et al. 2020) and for LAMOST DR5 (Xiang et al. 2019). The next generation of optical and near-infrared spectrographs, wide-field and massively multiplexed, is in preparation and will soon provide even larger catalogues of APs and abundances, such as the WHT Enhanced Area Velocity Explorer (WEAVE; Dalton et al. 2012), the Multi-Object Optical and Near-infrared Spectrograph (MOONS) on ESO’s Very Large Telescope (Taylor et al. 2018), the 4-metre Multi-Object Spectroscopic Telescope (4MOST; de Jong et al. 2019), the Prime Focus Spectrograph (PFS; Takada et al. 2014), and the Maunakea Spectroscopic Explorer (MSE; McConnachie et al. 2016). In terms of numbers of stars, the most revolutionary survey will certainly be that of the Gaia space mission (Gaia Collaboration 2016), which will deliver in its DR3 in 20221 estimates of the physical properties, including metallicities, for millions of stars obtained with various methods through an astrophysical parameter inference system (Bailer-Jones et al. 2013).

2020AandA...639A..20K___2015_Instance_1

where D ϵ Dt $ \frac{D \epsilon}{Dt} $ is the advective derivative of the energy density, κ = κ 0 T 5 / 2 b ̂ $ {\boldsymbol{\kappa}} = \kappa_0 T^{5/2} \hat{{\boldsymbol{b}}} $ is the Spitzer conductivity along magnetic field lines, Qheat and Qcool are heating and radiative cooling rates, and ρ, v, T, P and g are the plasma density, velocity, temperature, pressure, and gravitational acceleration respectively. Such user-defined heating term usually has a form of an exponentially decreasing function along the vertical coordinate y; Q heat = c 0 exp ( − y λ ) $ Q_{\mathrm{heat}} = c_0 \exp (-\frac{\mathit{y}}{\lambda}) $ , where c0 is the peak heating rate and λ is the heating scale height. The user defined heating is therefore highly stratified, spatially smooth and steady (e.g. Müller et al. 2003; Fang et al. 2013, 2015; Xia et al. 2017). The need for the user-defined heating in the previous coronal rain simulations arises from the fact that they typically do not include any self-consistent dissipation mechanisms. They also do not cover complete lower solar atmosphere including chromosphere, photosphere, and convection zone, therefore omitting key physical processes in the lower atmosphere, such as magneto-convection, associated magnetic buffeting, braiding, and flows. Another drawback of several coronal rain simulations is the commonly used approximation that all of the plasma cooling (i.e. the process essential for modelling the thermal instability and catastrophic cooling) occurs via optically thin radiative losses. This approximation is perfectly valid in the corona but ceases to apply for cool plasma (below temperatures of a few 100 000 K). Such assumption means that regardless whether the radiative loss function is calculated from CHIANTI (Dere et al. 2019) or using scaling law approximations (e.g. Rosner et al. 1978), there is a cut-off temperature for the radiative cooling of the plasma condensations. Once electron recombination starts during the cooling process, the energy gained (which depends on the ionisation potential and thus also on the ionisation degree of the plasma) is expected to slow down the cooling rate. The thermal evolution of the plasma condensations is therefore not modelled correctly at low temperatures in the previous coronal rain simulations.