The gap and ring resemble those in the HL Tau system, recently found by the ALMA long baseline campaign (ALMA Partnership et al. 2015). Our result shows that gaps and rings in the (sub)millimeter dust continuum can exist, not only in relatively young disks (0.1–1 Myr) but also in relatively old disks (3–10 Myr). One possible mechanism for opening a gap is the gravitational interaction between a planet and the gas (e.g., Lin & Papaloizou 1979; Goldreich & Tremaine 1980; Fung et al. 2014). Such an interaction may also produce the spiral density waves recently found in optical and near-infrared scattered light imaging of dust grains in protoplanetary disks (e.g., Muto et al. 2012). According to recent theoretical analyses of gap structure around a planet (Kanagawa et al. 2015a, 2015b, 2016), the depth and width of the gap are controlled by the planetary mass, the turbulent viscosity, and the gas temperature. The shape of the gap is strongly influenced by angular momentum transfer via turbulent viscosity and/or instability caused by a steep pressure gradient at the edges of a gap. The observed gap has an apparent width and depth of

 au and

, respectively. This is too shallow and too wide compared with that predicted by theory. However, the observations are limited to an angular resolution of ∼15 au, and the depth and width could be deeper and narrower in reality. For instance, if we assume that the gap depth times the gap width retains the value derived from the observations, it is possible for the gap to have a width and depth of

 6 au and

, which is similar to the GPI result (Rapson et al. 2015). Such a gap could be opened by a super-Neptune-mass planet, depending on the parameters of the disk, such as the turbulent viscosity (Kanagawa et al. 2015a, 2015b, 2016). If the gap in the larger dust grains is deeper than that in the gas, the planet could be lighter than super-Neptune mass. We note that a planet of even a few Earth masses, although it cannot open a gap in the gas, can open a gap in the dust distribution if a certain amount of pebble-sized particles, whose motions are not perfectly coupled to that of gas, are scattered by the planet and/or the spiral density waves excited by the planet (Paardekooper & Mellema 2006; Muto & Inutsuka 2009).


Morphology, colors, and star formation rate (SFR) primarily depend on the small-scale (1 Mpc) environment (Hogg et al. 2004; Kauffmann et al. 2004; Wetzel et al. 2012). This result has been extended by studying galaxy group samples; they show that colors and SF history most directly depend on the properties of the host dark matter (DM) halo (Blanton & Berlind 2007; Wei et al. 2010; Tinker et al. 2012), in agreement with the results of smooth particle hydrodynamic (SPH) simulations by Mazzei & Curir (2003, hereafter MC03). In this context, the investigation of the evolution of group members in the nearby universe acquires a great cosmological interest because more than half of galaxies reside in such environments. Furthermore, since the velocity dispersion of galaxies is significantly lower in groups than in clusters, the merger probability and the effects of interaction on galaxy evolution are much higher. Consequently, groups provide a zoom-in on phenomena driving the galaxy evolution before the galaxies fall into denser environments (e.g., Wilman et al. 2009; Just et al. 2010). Starting from the pioneering works of Toomre & Toomre (1972, and references therein), several papers contributed to shed light on the important role of mergers/interactions in galaxy evolution—Toomre (1977), Combes et al. (1990), Mihos & Hernquist (1994, 1996), Barnes & Hernquist (1996), Naab & Burkert (2003), Bournaud et al. (2005), and Di Matteo et al. (2007), to name a few—up to the most recent papers of Eliche-Moral et al. (2018) and Martin et al. (2018, and references therein). The dissipative merger simulations of Eliche-Moral et al. (2018) start from systems just formed, composed of a spherical nonrotating DM halo, and by a disk of gas particles with or without the presence of a stellar bulge. These simulations explore about 3–3.5 Gyr of evolution. Martin et al. (2018) focused on processes triggering galaxy transformations of massive galaxies (M > 1010 M⊙) exploiting cosmological hydrodynamic simulations by Kaviraj et al. (2017). These simulations, based on an adaptive mesh refinement code (RAMSES) and including the baryon treatment with stellar and active galactic nucleus (AGN) feedback, are able to resolve baryonic physics on kiloparsec scales, larger than we use in this paper (Section 3). They derived important statistical assessments about the processes that drive morphological transformation across cosmic time.


Over the past decade, many attempts to address this issue have been carried out, such as from the perspective of conventional nuclear physics and even exotic physics beyond the standard BBN framework (Angulo et al. 2005; Cyburt et al. 2008, 2016; Boyd et al. 2010; Pospelov & Pradler 2010; Fields 2011; Kirsebom & Davids 2011; Wang et al. 2011; Broggini et al. 2012; Coc et al. 2012, 2013, 2014; Cyburt & Pospelov 2012; Kang et al. 2012; Voronchev et al. 2012; Bertulani et al. 2013; Hammache et al. 2013; He et al. 2013; Kusakabe et al. 2014; Pizzone et al. 2014; Yamazaki et al. 2014; Hou et al. 2015, 2017; Famiano et al. 2016; Damone et al. 2018; Hartos et al. 2018; Luo et al. 2019; Rijal et al. 2019; Clara & Martins 2020). However, despite the fact some solutions using exotic physics have succeeded in resolving this issue, it appears there is still no universally accepted solution in the academic community since validations of these mysterious exotic physics are beyond the capabilities of current science. Conversely, it seems more worthwhile to exclude any potential possibility of resolving the 7Li discrepancy from the perspective of nuclear physics. It is known that the majority of the primordial 7Li production arises from the decay of 7Be by electron capture during the 2 months after BBN stops. Thus, for the solution of the Li problem, reactions involving 7Be could be more significant than those involving 7Li. Therefore, many reactions that potentially destroy 7Be were investigated to solve this discrepancy over past 10 yr (Kirsebom & Davids 2011; Broggini et al. 2012; Hammache et al. 2013; Hou et al. 2015; Hartos et al. 2018). Meanwhile, enormous efforts have been made to refine the reaction rates of key BBN reactions in the past 20 yr (Smith et al. 1993; Descouvemont et al. 2004; Serpico et al. 2004; Cyburt & Davids 2008; Neff 2011; Pizzone et al. 2014; Tumino et al. 2014; Hou et al. 2015; Barbagallo et al. 2016; Iliadis et al. 2016; Kawabata et al. 2017; Lamia et al. 2017, 2019; Damone et al. 2018; Rijal et al. 2019; Mossa et al. 2020), but the probability of solving or alleviating the 7Li problem by improving our knowledge of relevant nuclear reaction rates still cannot be eliminated. Recent experiments for key nuclear reactions like 7Be(n,p)7Li and 7Be(d,p)24He allow for a reduction of the 7Li production by about 12% (Damone et al. 2018; Rijal et al. 2019) compared to previous calculations. At present, nuclear uncertainties cannot rule out that some of the reactions destroying 7Li are indeed more efficient than those currently used (Boyd et al. 2010; Chakraborty et al. 2011).


Regardless of the cluster orientation, the spectral index observed for the halo at all available frequencies suggests that it is a USSRH. Despite the number of detected USSRH is still low, radio halos with steep indices are being discovered more and more frequently in the last years thanks to the improved observational capabilities of low-frequency instruments such as GMRT, MWA (Murchison Widefield Array) and LOFAR (Shimwell et al. 2016; Wilber et al. 2018; Bruno et al. 2021; Di Gennaro et al. 2021; Duchesne et al. 2022). An in-depth analysis of all radio halos hosted in Planck clusters and observed in LoTSS, including A1550, has recently been presented in Botteon et al. (2022). USSRH are a prediction of turbulent re-acceleration models (Cassano et al. 2006; Brunetti et al. 2008), in which particles are re-accelerated by turbulence (Brunetti et al. 2001, 2017; Petrosian 2001; Brunetti & Lazarian 2011). On the other hand, the detection of such steep indices is not expected from hadronic (or secondary) models, in which the emission of halos comes from the production of secondary electrons from hadronic collisions between thermal and CR protons (Blasi & Colafrancesco 1999; Dolag & Enßlin 2000; Pfrommer et al. 2008). Given that the integrated spectral index observed for the USSRH with LOFAR is

α

54

MHz

144

MHz

∼
−
1.6

$ \alpha_{54\,\rm MHz}^{144\,\rm MHz} \sim -1.6 $

, we expect an index for the spectral energy distribution8δ = 2α − 1 = −4.2. If there is no break in the spectrum, the energy budget for these particles would be untenable (Brunetti et al. 2008). Therefore, a break at low energies (∼GeV) should exist, suggesting a possible interplay between radiative losses and turbulent re-acceleration during the lifetime of emitting electrons (Brunetti & Jones 2014). Moreover, re-acceleration models predict that a large fraction of halos associated with clusters of masses between 4 and 7 × 1014 M⊙ should exhibit steep spectra (Cassano et al. 2010, 2012; Brunetti & Jones 2014; Cuciti et al. 2021). The mass of A1550 of ∼6 × 1014 M⊙ estimated from Planck Collaboration XXVII (2016) falls in this range9.


Shown in Figures 2(a) and (b) are the z-component of the electric field, Ez, and y-component of the magnetic field, By, respectively. Panels 2(c)–(e) show the transversally averaged (in the yz-plane) electric and magnetic field components,

. The energy distribution (total of jet+ambient) and average energy along the x-direction for electron and ion species are shown in Figures 2(f)–(i). All panels are at t∗ = 500. Where high-speed jet particles interact with the ambient medium (behind the RS at x∗ ≤ 340) or scattered ambient particles blend with the upstream ambient (in front of the FS at 430 ≤ x∗), particles distribution becomes strongly anisotropic. Anisotropies result in the Weibel instability which generates current filaments in these regions with currents in the x-direction. According to Ampere’s law, these current filaments are encircled by transverse magnetic fields, and we see that

 in Figure 2(c). The transverse electric fields are related to the magnetic fields via

 where

 is the velocity of the electron (ion) carrier. The carriers move roughly at the speed of light in the x-direction,

. Therefore, the transverse electric field components are Ey = Bz, and Ez = −By, as are observed in the simulation results for

 (Figures 2(a)–(e)). Additionally, there is a longitudinal ambipolar electric field within the RS transition region, 140 ≤ x∗ ≤ 340 for t∗ = 500 (Figure 2(c)). This electric field is generated by the density gradient and different mobilities of electrons and ions (Forslund & Shonk 1970; Forslund & Freidberg 1971; Hoshino 2001; Choi et al. 2014). The magnetic fields act to isotropize the momentum distribution, while the electric fields function to thermalize, and accelerate the particles afterwards. In Figures 2(f)–(i), the shocked region lies between x∗ = 340 and x∗ = 430. Within the RS transition region (140 ≤ x∗ ≤ 340) jet electrons are trapped by the ambipolar electric field and effectively accelerated up to γe = 200 by the transverse electric fields (Figures 2(f) and (g)). A tenuous population of these electrons convect upstream due to reflection by the magnetic fields in the shocked region (ellipse in Figure 2(f)). On the other hand, jet ions are slowed in the RS transition region (140 ≤ x∗ ≤ 340) by 40% from the initial Lorentz factor γi = 10, due to the effect of the ambipolar electric field. In the shocked region, jet electrons have been fully thermalized and are well merged with the thermalized ambient electrons. Thus, only a single electron population is present in the hot shocked region (Figure 2(f)). On the other hand, the kinetic energy of jet ions is transferred to the heating of ambient particles by means of the electromagnetic fields generated by the ion Weibel instability (Figure 2(h)). Full thermalization of the two ion populations (jet and ambient) has not yet occured (demands a longer simulation time), i.e., the two populations are distinguishable in Figure 2(h). Electrons located in the FS transition region (430 ≤ x∗ ≤ 500), predominantly ambient electrons, also undergo the Weibel instability and are thermalized by the jet upstream kinetic energy. In this region, penetrating jet ions interact with ambient particles and are slowed down gradually by 50% from the jet front Lorentz factor γi = 10 to a minimum value of γi = 5 (Figures 2(h) and (i)).


Where

 is the straggle or standard deviation of the implantation distribution, D is the diffusion coefficient, and t is the irradiation time. In the case of gas implantation via ion irradiation, it is often assumed that irradiation damage produces enough defects to serve as trapping sites that the diffusion coefficient is negligible (Martynenko 1977; Scherzer 1983). In Section 2.5, an estimate of the effective diffusion in the presence of damage produced by relativistic light ion impacts is made. Lattice defects caused by ion irradiation also provide nucleation for fixed helium bubbles (Kornelsen 1972). We can assume additional trapping due to the interaction of hydrogen with helium in the material, since helium bubbles serve as trapping sites for hydrogen atoms through synergistic effects (Hayward & Deo 2012), leading to a lower critical dose in the case of mixed hydrogen and helium exposure (Guseva & Martynenko 1981). At temperatures well below the melting point, such as that expected for a relativistic interstellar spacecraft, high-flux hydrogen irradiation damage effects have been reproduced in low-flux experiments (Gao et al. 2019). However, for an interstellar probe, the flux is sufficiently low that diffusion effects may play a non-negligible role. Classical diffusion of individual gas atoms could lead to lower local gas atom concentrations as the implantation distributions widen, or increased gas atom concentrations around trapping sites such as grain boundaries; the effect of diffusion is discussed in Section 2.5. Helium bubble nucleation and migration could lead to increased gas atom concentrations and damage near surfaces, grain boundaries, and other defects (Nakamura et al. 1977; Goodhew 1983; Lane & Goodhew 1983). No single theoretical framework exists to summarily treat these effects. However, simple models such as that presented here immediately offer compelling mitigation strategies. To perform this analysis, we find implantation profiles of ISM gas atoms at relativistic speeds using an ion–material interactions code, calculate critical concentrations for blistering onset for hydrogen and helium individually assuming a worst-case scenario of negligible diffusion, and show the effect of non-negligible diffusion on local gas concentrations.


It is important to note that the above procedure differs from the observational one, whereby Voigt profiles are fit to the absorption features in transmission spectra (e−τ) in order to extract column densities and Doppler broadening parameters. This approach enables deblending of multiple-component absorption, and takes into account the broadening of the lines due to the instrumental profile, enabling accurate recovery of their column densities. The technique was originally devised under the premise that the intervening absorption lines in QSO spectra arose from discrete absorbing clouds, but this picture was challenged early by the smoothly distributed IGM captured in cosmological hydrodynamical simulations (e.g. Cen et al. 1994; Hernquist et al. 1996; Miralda-Escudé et al. 1996; Theuns et al. 1998; Davé et al. 1999). Absorbers that have a large spatial extent take part in the Hubble expansion, which leads to a line profile that deviates from a Voigt profile. Voigt profile fitting the transmission spectra will therefore glean slightly different results to simply taking peaks in the τ distribution. In particular, we might expect a larger number of absorbers, and some differences in the derived Doppler broadening parameters and column densities. As a check, we derived Voigt profiles for lines recovered in some of the sight-lines extracted from the simulation using our τ peak method, and plotted these on top of the transmission spectrum. An example is shown in Fig. 3. The transmission spectrum is shown as the black dashed line, and the predicted Voigt profiles are shown in green. More structure is apparent in the real spectrum, which would yield a larger number of Voigt components in Voigt profile fitting, and column densities and Doppler broadening parameters that differ slightly from those we have recovered via the τ peak method. This is an obvious caveat to our approach, which we bear in mind when interpreting our results later on.


Although wind observations are very common in AGN (see Elvis 2000; Veilleux et al. 2005; and Fabian 2012, for reviews), most studies concern ionised gas and uncertain spatial scales. In the past few years the situation changed drastically. Several fast (vOF of the order of 1000 km s-1), massive outflows of ionised, neutral and molecular gas, extended on kpc scales, have been discovered thanks to three techniques: 1) deep optical/NIR spectroscopy, mainly from integral field observations (IFU, e.g. Nesvadba et al. 2006, 2008; Alexander et al. 2010; Rupke & Veilleux 2011; Riffel & Storchi-Bergmann 2011; Cano-Diaz et al. 2012; Greene et al. 2012; Harrison et al. 2012, 2014; Liu et al. 2013a,b; Cimatti et al. 2013; Tadhunter et al. 2014; Genzel et al. 2014; Brusa et al. 2015a; Cresci et al. 2015; Carniani et al. 2015; Perna et al. 2015a,b; Zakamska et al. 2016); 2) interferometric observations in the (sub)millmetre domain (e.g. Feruglio et al. 2010, 2013a,b, 2015; Alatalo et al. 2011; Aalto et al. 2012; Cicone et al. 2012, 2014, 2015; Maiolino et al. 2012, Krips et al. 2011; Morganti et al. 2013a,b; Combes et al. 2013; Garcia-Burillo et al. 2014); and 3) far-infrared spectroscopy from Herschel (e.g. Fischer et al. 2010; Sturm et al. 2011; Veilleux et al. 2013; Spoon et al. 2013; Stone et al. 2016; Gonzalez-Alfonso et al. 2017). In addition, AGN-driven winds from the accretion disk scale up to the dusty torus are now detected routinely both in the local and in the distant Universe, as blue-shifted absorption lines in the X-ray spectra of a substantial fraction of AGN (e.g. Piconcelli et al. 2005; Kaastra et al. 2014). The most powerful of these winds, observed in 20–40% of local AGN (e.g. Tombesi et al. 2010) and in a handful of higher redshift objects (e.g. Chartas et al. 2009; Lanzuisi et al. 2012), have extreme velocities (ultra-fast outflows, UFOs, v 0.1-0.3c) and are made by highly ionised gas which can be detected only at X-ray energies.


Observations of large-scale magnetic fields offer clear insights about regular ordered patterns of the field lines, suggesting that mean-field dynamo processes are responsible for their order and structure, as well as the existence of additional transport processes carrying magnetic energy into huge regions of space (Clarke, Kronberg & B”ohringer 2001; Bonafede et al. 2010; Arlen et al. 2012; Govoni, F. et al. 2017; Han 2017). Magnetic fields cannot be directly observed, so their impact on radiation processes need to be considered (see, for instance, Rybicki & Lightman 1979, and references therein). In addition, observations of magnetic fields in voids provide bounds on their strength, depending on the analytical model: From the simplest ones, it is possible to obtain lower bounds, although when improving such models a bounded range of magnetic field strength values can be provided, ranging from 10−25 to 10−15 nG (Neronov & Vovk 2010; Tavecchio et al. 2010; Essey, Ando & Kusenko 2011; Takahashi et al. 2013). Other authors argue magnetic field strengths between 10−16 and 10−15 G (Einstein 1915; Einstein 1916; Hubble 1929; Hubble & Humason 1931; Bull et al. 2016). Magnetic fields from astrophysical voids are relevant since they could evidence truly cosmological magnetic fields, which could have served as seeds for magnetic fields in lower scales (such as galactic fields). Observations in galaxy clusters yield values of the order of 10−6 G (Boulanger et al. 2018). Interest in primordial magnetic fields generated during inflation has recently increased. This scenario has driven the search for phenomenologically viable mechanisms to explain the observed magnetic fields in a broad set of scales (Grasso & Rubinstein 2001; Brandenburg & Subramanian 2005; Demozzi, Mukhanov & Rubinstein 2009; Ferreira, Jain & Sloth 2013; Green & Kobayashi 2016). The latest data on reionization and the observed UV luminosity function of high-redshift galaxies place limits on the magnetic field strength due to its impact on the reionization process.


GCs formed in dwarf galaxies may differ from those found in the Galactic halo, depending on their age and metallicity. Dwarf galaxies show a wide variety of star formation histories (Hidalgo et al. 2011, 2013; Weisz et al. 2014) that are predicted to lead to variations in their metallicity distribution functions and chemical abundances. It has also been suggested these variations could be attributed to differences in the IMFs of these galaxies (McWilliam, Wallerstein & Mottini 2013). If the IMFs are the root cause of these differences, then this would also lead to differences in the age–metallicity relationship, which is observed by both Forbes & Bridges (2010) and Leaman et al. (2013). From observations, field stars in dwarf galaxies do exhibit different abundance ratios from MW field stars, e.g. lower [α/Fe] ratios and variations in neutron-capture element ratios at intermediate metallicities. However, these typically do not show up until [Fe/H] ∼ −1.5 (Shetrone, Bolte & Stetson 1998; Shetrone, Côté & Sargent 2001; Shetrone et al. 2003; Venn et al. 2004; Okamoto et al. 2012; Tolstoy, Hill & Tosi 2009; Frebel 2010). At metallicities below [Fe/H] = −1.5, the abundance variations between field and GC stars become less pronounced in dwarfs and the MW (Hill et al. 2000; Pritzl, Venn & Irwin 2005; Carretta et al. 2010; Letarte et al. 2010); a good example of this is M54, located at the heart of the Sagittarius (Sgr) dwarf accretion remnant. M54 has a much lower metallicity than the Sgr field stars (e.g. Carretta et al. 2010) and the [α/Fe] ratios resemble the field stars in the MW halo and its detailed chemical abundance ratios resemble the patterns seen in other GC systems (e.g. the Na–O anticorrelation; Carretta et al. 2009). Therefore, other than its physical association with the Sgr remnant, M54 does not stand out from other GCs in terms of its chemical abundance patterns, similar to the metal-poor GCs Terzan 8 and Arp 2 (both also kinematically and spatial associated with the Sgr stream; Mottini, Wallerstein & McWilliam 2008). On the other hand, Hodge 11 in the large magellanic cloud at [Fe/H] = −2.0 does have lower [α/Fe] than MW field and GC stars(Mateluna et al. 2012), and Ruprecht 106 has an anomalously low [α/Fe] ratio for a MW GC (Villanova et al. 2013).


The coalescence of a binary compact object system (either a neutron star (NS) binary or a stellar-mass black hole (BH) and NS binary) has been widely suggested to account for short-duration gamma-ray burst (SGRB) events (Eichler et al. 1989; Narayan et al. 1992; Nakar 2007; Berger 2014) that last typically less than 2 s in the soft γ–ray band (Kouveliotou et al. 1993). Since 2006, it has been suspected that mergers of compact objects could also produce the so-called long–short GRBs (also known as the supernova-less long GRBs, which are apparently long-lasting but do not show any signal of supernovae down to very stringent limits), which share some properties of both long- and short-duration GRBs (Della Valle et al. 2006; Gal-Yam et al. 2006; Gehrels et al. 2006; Zhang et al. 2007). Compact binary coalescence (CBC) is generally expected to be a strong source of gravitational wave (GW) radiation and such events are prime targets for some GW detectors like advanced Laser Interferometer Gravitational-Wave Observatory (LIGO)/VIRGO (Abadie et al. 2015; Acernese et al. 2015; Belczynski et al. 2010, 2016; see also the latest LSC–Virgo white paper at https://dcc.ligo.org/LIGO-T1400054/public). On 2015 September 14, the two detectors of the LIGO simultaneously detected a transient gravitational wave signal from the merger of two BHs (GW150914, Abbott et al. 2016). GW150914 is the first direct detection of gravitational waves and the first identification of a binary BH merger (Abbott et al. 2016). Surprisingly, the observations from the Fermi Gamma-ray Burst Monitor (GBM) at the time of GW150914 claimed a detection of a weak gamma-ray transient (i.e., GBM transient 150914) 0.4 s after GW150914 with a false alarm probability of 0.0022 (Connaughton et al. 2016). If true, this is the first GW/SGRB association (see, however, Savchenko et al. 2016 for some arguments). Li et al. (2016) compared GBM transient 150914 with other SGRBs and found that such an event is remarkably different in its prompt emission properties. The binary BH merger origin as well as its property of “distinguished” prompt emission suggest that GW150914/GBM transient 150914 is not a typical GW/SGRB association.


After the detection of V391 Peg b, some other planet or brown dwarf (BD) candidates orbiting sdB stars were proposed using different detection methods. From eclipse timing, about one-third of the known detached sdB/sdO + dM (dM = M-dwarf) post-common-envelope binaries (PCEB) are suspected to host planets/BDs: HW Vir (Beuermann et al. 2012 and references therein), HS 0705+6700 (alias V470 Cam, Qian et al. 2013 and references therein), HS 2231+2441 (Qian et al. 2010 and references therein; but see also Lohr et al. 2014), NSVS 14256825 (Almeida et al. 2013; Hinse et al. 2014 and references therein), NY Vir (Lee et al. 2014 and references therein), and 2M 1938+4603 (Baran et al. 2015). Interesting explorations on the origin of PCEB (and specifically sdB+MS/BD) circumbinary planets can be found in Zorotovic & Schreiber (2013), Schleicher & Dreizler (2014), Bear & Soker (2014), and Völschow et al. (2016). Very different planets or planetary remnants with terrestrial radii have been proposed from tiny reflection effects detected by the Kepler spacecraft in KIC 05807616 (Charpinet et al. 2011) and KIC 10001893 (Silvotti et al. 2014). However, none of these sdB planet/BD candidates has been confirmed with at least two independent detection methods. More robust detections of a few brown dwarfs (BDs) in eclipsing sdB binaries (also called HW Vir systems from the sdB+dM protoptype) were obtained by combining stellar radial velocities (RVs) with photometric measurements: J08205+0008, J1622+4730 and V2008-1753 have companion masses of about 71, 67, and 69 MJup, respectively (Geier et al. 2011; Schaffenroth et al. 2014a, 2015). At least two more sdB+BD eclipsing systems were recently found from the OGLE survey (Schaffenroth, in prep., priv. comm.). Finally, two more BD candidates in sdB binaries were found by combining radial velocities (RVs) with photometric reflection effects: CPD-64°6481 and PHL 457, with minimum masses of 50 and 28 MJup, respectively (Schaffenroth et al. 2014b).


The global magnetospheric structures for the oblique rotator are very similar to the aligned one. We show the structure of magnetic field lines and the distribution of the accelerating electric field E0 in the x–z plane for a 60° rotator with the pair multiplicity κ = {0, 1, 3} in Figure 4. As the pair multiplicity κ increases, the field structure tends to the force-free solution with an equatorial current sheet outside the LC. We observer a dissipative region where E > B outside the LC. The spatial extension of the dissipative region decreases with increasing pair multiplicity and the E0 region is more confined to the equatorial current sheet outside the LC as the pair multiplicity κ increases. In fact, the E0 distribution for the high κ solution is qualitatively similar to the FIDO one (see, e.g., Kalapotharakos et al. 2014; Cao & Yang 2019). We also compare the field structures for κ = 0 with Figure 1 of Contopoulos (2016) for α = 0° and α = 60° respectively. We find that the field structures are qualitatively very similar to those of Contopoulos (2016). For comparison, we also show the magnetic field lines and the E0 distributions for a 60° rotator with the pair multiplicity κ = 0 by implementing the AE formulation everywhere in Figure 5. The magnetospheric structure is very similar to the aligned one with a force-free zone bounded by a radiation zone. We observe a strong E0 distribution inside the LC, which is very different from those in the SG and OG models. A strong E0 region with E > B also appears outside the LC. We show the distributions of magnetic field lines and the accelerating electric field E0 in the x–z plane for a 30° rotator with the pair multiplicity κ = 3 in Figure 6. We see that the field structure is very close to the force-free one and the E0 region is restricted to only near the current sheet outside the LC for this high κ value. We also show the normalized Poynting flux L/Laligned as a function of radius r for a 90° rotator with different pair multiplicities in Figure 7. We see that the Poynting flux increases with increasing κ values and approaches the force-free solution for the high κ value. Our simulation shows a more than 1% dissipation rate outside the LC for a 90° dissipative rotator. A similar dissipation rate is also found by the PIC simulation for the aligned and perpendicular rotator (Philippov et al. 2015). In fact, the spectral numerical methods present an unphysical dissipation beyond the LC due to discontinuity in the current sheet. A higher resolution is necessary to catch the discontinuity in the current sheet and reduce the unphysical dissipation.


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

which according to
is about of 0.5−1 Gyr after the merger


In both models, as we will see, the thick disc scale length is about a factor of two shorter than that of the thin disc, in agreement with the results by Bovy et al. (2012a). The choice of presenting two mass models for the mass distribution of our Galaxy is mainly dictated by two reasons. First, the need to add a central bulge to the global gravitational potential to reproduce the rotation curve in the inner kpcs of the Milky Way strongly depends on the observational data with which one compares the theoretical curve: to reproduce the rise observed in the inner kpcs (see the observational data adopted by Caldwell & Ostriker 1981), Allen & Santillan (1991) introduced a central mass concentration, whose mass is about 15% of the disc mass. However, the central rise observed in the rotation of the molecular gas in the inner Galaxy (for more recent estimates see, for example,  Sofue 2012) may be an effect of non circular motions generated by large scale asymmetries like the bar, as has been shown recently by Chemin et al. (2015). Moreover, this feature is not reported in all the observational studies (see, for example, Reid et al. 2014). Secondly, there is growing evidence that the mass of any classical bulge, if present in the Milky Way, must be small (Shen et al. 2010; Kunder et al. 2012, 2016; Di Matteo et al. 2014, 2015). For these reasons, we prefer to present a second model, our Model II, which does not include any spherical central component, and which is still compatible with the rotation curve of the Galaxy, as given by Reid et al. (2014). Because it has been widely used in the last decades, and due to the facility of its implementation, we explicitly aim at generating Galactic models similar to the Allen & Santillan (1991) model, so to make any implementation of these new models, and any comparison with Allen & Santillan (1991), straightforward. As for the model proposed by Allen & Santillan (1991), Models I and II are axisymmetric and time-independent, and do not include stellar asymmetries such as a bar or spiral arms. No truncation is assumed for the discs, while the halo is truncated at 100 kpc, in agreement with the choice of Allen & Santillan (1991). As we describe in the following section, the analytic forms for the discs, halo, and bulge potentials are the same as those adopted by Allen & Santillan (1991). To allow an easy comparison with the Allen & Santillan (1991) model, in the following we will make use of the same system of units adopted by these authors: the potential is given in units of 100 km2/ s2, lengths are in kpc, masses in units of 2.32 × 107M⊙, time in units of 0.1 Gyr, velocities in units of 10 km s-1 and the vertical force in units of 10-9 cm s2. In these units, the gravitational constant G is equal to 1 and the mass volume density is in units of 2.32 × 107M⊙/ kpc3.


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).


There are two possibilities resulting in the different magnetic field strengths inferred from the polarimetric and molecular-line observations: (1) the rotational-to-gravitational energy βrot is overestimated, and (2) there are additional contributions in the polarized intensity from other mechanisms, such as dust scattering. In our MHD simulations, βrot is adopted to be 0.4% based on the observational estimates of the core mass of ∼1 M and the angular speed of the core rotation of 4 × 10−14 s−1. The angular speed was estimated based on the global velocity gradient along the major axis of the dense core observed with single-dish telescopes (Saito et al. 1999; Yen et al. 2011; Kurono et al. 2013). Numerical simulations of dense cores including synthetic observations show that the specific angular momentum derived from the synthetic images of the dense cores can be a factor of 8–10 higher than their actual specific angular momentum computed by the sum of the angular momenta contributed by the individual gas parcels in the dense cores (Dib et al. 2010). In addition, if there are filamentary structures in the dense core in B335, which could not be resolved with the single-dish observations, infalling motions along the filamentary structures could also contribute to the observed velocity gradient, leading to an overestimated angular speed of the core rotation (Tobin et al. 2012). We have also performed our simulations with a lower βrot, and we find that the rotational velocity on a 100 au scale in the simulations decreases with decreasing βrot. Thus, the discrepancy in the magnetic field strengths inferred from the field structures and the gas kinematics can be reconciled, if the core rotation in B335 is overestimated by a factor of a few in the observations, and these results would suggest a weak magnetic field of initial λ of 9.6 in B335. Further observations combining single dishes and interferometers to have a high spatial dynamical range and to map the velocity structures of the entire dense core in B335 at a high angular resolution are needed to study coherent velocity features and provide a better estimate of the core rotation.


A number of fundamental results have been rigorously proved in the mathematical literature concerning the asymptotic behaviour in time of some of the solutions of the porous medium equation and related equations (e.g. Kamin & Vázquez 1991; Bernis et al. 1993; Hulshof et al. 2001). What is of interest for us here is, primarily, the results that can be applied to the cylindrically symmetric case with diffusion coefficient which is proportional to the square of the dependent variable (n = 2, m = 3 in the notation of Eq. (7)). The most basic result, already mentioned in Sect. 4.4.1, is that initial conditions which have a finite nonzero flux integral (called ‘the mass’ in the mathematical literature for the PME) converge toward the ZKBP solution with the same flux integral (‘mass’) asymptotically in time (Vázquez 2007, Theorem 18.2); here, allowance is made for either a positive or negative flux integral by globally changing the sign of the ZKBP solution; also, ‘convergence’ is meant in the sense that the Lp norm of the difference between the actual solution and the ZKBP function tends to zero as t → ∞ faster than a negative power of the time with an exponent which is a function of n, m, and p (e.g. −1/3 for n = 2 and m = 3 in the L2 norm; see details in the book by Vázquez 2007). A complementary result is the following: when the initial condition has positive net flux and its negative part has compact support, then the whole solution evolves into a positive function after a finite time (Vázquez 2007, Theorem 18.29). Since we are dealing with signed functions which have zero flux integral, these results are of interest mainly because they impose a strict condition on the possible flux imbalance caused by numerical errors (as discussed in Sect. 4.4.1, final paragraph): if it is not small, the numerical solutions will approach the ZKBP solution in a comparatively short time. However, the flux imbalance in all the Bifrost experiments discussed in the present paper is small enough that they have not shown this behaviour even though they have been run until a very long diffusive time.


Deep learning (DL) is the paradigm of machine learning which uses multilayer neural networks. Neural networks are ML models inspired from the network of brain cells or neurons. The differentiating factor between DL and conventional ML is in the process of feature selection. In conventional ML, performance strongly depends on the features used. More often than not, new features are created for a task to better capture correlations in the data. This is called feature engineering. ML models perform poorly if not given good features to learn from. Deep learning eliminates this dependence by defining its own features to learn from, which are relevant to the task at hand. This makes deep learning a versatile, high performing solution for a variety of supervised ML tasks. The Artificial Neural Networks (ANNs) have found application in not only the traditional tasks in extragalactic astronomy such as photometric redshift estimation (e.g. Firth, Lahav & Somerville 2003; Tagliaferri et al. 2003; Collister & Lahav 2004; Vanzella et al. 2004; Sadeh, Abdalla & Lahav 2016; Bilicki et al. 2018; Pasquet et al. 2019), but also in more specialized problems such as predicting infrared luminosity of galaxies (Ellison et al. 2016), estimation of star formation properties (Surana et al. 2020), and ranking the quenching parameters of galaxies (Teimoorinia, Bluck & Ellison 2016). In particular, Convolutional Neural Networks (CNN) are now becoming increasingly popular in studies using imaging data on galaxies. The CNNs are deep learning models designed to extract features from images. They have provided state-of-the-art performance in majority of computer vision tasks in recent times (Krizhevsky, Sutskever & Hinton 2012; He et al. 2017). CNNs have been utilized for galaxy morphological classification e.g. classifying the optical morphologies broadly into spheroidal, disc, and irregular types (Huertas-Company et al. 2015) and even for classification of various radio galaxy morphologies (Wu et al. 2019). Recently, Ribli, Dobos & Csabai (2019) have shown that CNNs can be used to predict galaxy shapes, needed for weak lensing studies, from a wide field but shallow sky survey images using the ‘ground truth’ images from a deeper but narrower field survey. CNNs have also found application in the automated detection of features in sky-survey images such as galactic bars (Abraham et al. 2018) and strong gravitational lenses (Canameras et al. 2020; Li et al. 2020). Further, ML and CNNs in particular have also been used to detect outliers in large area Sloan Digital Sky Survey (SDSS) data (e.g. Baron & Poznanski 2017; Sharma et al. 2019). In this study, we want to determine the bulge to total luminosity ratio (B/T) of a galaxy using its optical multiband images as input. Due to the use of a data set with galaxy images labelled by $B/T\, \in [0,1]$ in a continuous space $\rm I\!R$, we have performed CNN based regression in this work.


Cepheids and RR Lyrae stars are primary standard candles of the cosmological distance ladder because they follow canonical relations linking their intrinsic luminosity to the pulsation period and/or the metal abundance. Specifically, for Cepheids the intrinsic luminosity (L) at any passband depends on the period (P) of light variation. This is traditionally referred to as the Cepheid period–luminosity relation or Leavitt law, after its discoverer Mrs Henrietta Swan Leavitt (Leavitt & Pickering 1912). Modern realisations of the Cepheid PL relations from optical to infrared passbands include, among others, the ground-based studies of Madore & Freedman (1991), Ripepi et al. (2012), and Gieren et al. (2013), works based on Hubble Space Telescope (HST) data such as those by Freedman et al. (2001), Saha et al. (2006), and Riess et al. (2011), and theoretical investigations such as those by Marconi et al. (2005). Among the most recent studies of the Cepheid period–luminosity relations are those based on Gaia trigonometric parallaxes of Galactic Cepheids (e.g. Clementini 2017; hereafter Paper I and references therein; Riess et al. 2016, 2018). For RR Lyrae stars the intrinsic luminosity (L) in the infrared passbands depends on P and possibly stellar metallicity (Z; PL – metallicity relation – PL(Z)), as first shown by Longmore et al. (1986) and later confirmed by (i) empirical studies of field and cluster RR Lyrae stars (e.g. Sollima et al. 2006, 2008; Borissova et al. 2009), (ii) theoretical models by Marconi et al. (2015) and Neeley et al. (2017), and (iii) the Gaia parallax-calibrated relations of Sesar et al. (2017), Paper I and references therein, and Muraveva et al. (2018a,b). In the visual passband, the luminosity L depends on Z in the form of the so-called RR Lyrae luminosity–metallicity relation (see e.g. Cacciari & Clementini 2003; Clementini et al. 2003; the pulsation models by Bono et al. 2003; the theoretical calibration by Catelan et al. 2004; or the Gaia-based relations in Paper I; Muraveva et al. 2018a and references therein). The predicted precision of the Gaia end-of-mission parallaxes for local Cepheids and RR Lyrae stars1 will allow us to determine the slope and zero-point of these fundamental relations with unprecedented accuracy, thus setting the basis for a global reassessment of the whole cosmic distance ladder. As a first anticipation of the Gaia potential in this field of the cosmic distance ladder and a first assessment of improved precision with respect to previous astrometric missions such as HIPPARCOS, and the dramatic increase in statistics compared to what is achievable, for instance, through measuring parallaxes with the HST, Gaia DR1 published parallaxes for more than 700 Galactic Cepheids and RR Lyrae stars, computed as part of the Tycho-Gaia Astrometric Solution (TGAS; Lindegren et al. 2016). A number of papers after Gaia intermediate data releases in 2016 and 2018 (Gaia Data Release 1 – DR1 and DR2, respectively) have discussed Gaia Cepheids and RR Lyrae stars, specifically presenting the released samples (Clementini et al. 2016, 2019), their parallaxes (e.g. Lindegren et al. 2016) and possible offsets affecting them (Arenou et al. 2017, 2018); and addressing in particular their use as standard candles (Casertano et al. 2017 and Paper I for Gaia DR1 and Riess et al. 2018; Muraveva et al. 2018a for Gaia DR2). In Paper I we have used TGAS parallaxes, along with literature photometry and spectroscopy, to calibrate the zero-point of the PL relations of classical and type II Cepheids, and the near-infrared PL and PL(Z) relations of RR Lyrae stars by fitting these relations through adopting different techniques that operate either in parallax or absolute magnitude space. In that paper, different sources of biases affecting the TGAS samples of Cepheids and RR Lyrae stars were discussed at some length, and the possible systematic errors caused in the inferred luminosity calibrations were analysed in detail.


To explain the observed evolution, the physical processes invoked have to result in a large growth in size but not in stellar mass, nor drastic increase in the star formation rate. Most plausible candidates are mass-loss driven adiabatic expansion (‘puffing-up’) (e.g. Fan et al. 2008, 2010; Ragone-Figueroa & Granato 2011) and dry mergers scenarios (e.g. Bezanson et al. 2009; Naab, Johansson & Ostriker 2009; Trujillo, Ferreras & de La Rosa 2011). In the former scenario, galaxies experience a mass-loss from wind driven by active galactic nuclei (AGNs) or supernovae feedback, which lead to an expansion in size due to a change in the gravitational potential. In the latter, mergers either major involving merging with another galaxy of comparable mass, or minor that involves accretion of low mass companions, have to be dry to keep the low star formation rate (Trujillo et al. 2011). Nevertheless, major mergers are not compatible with the observed growth in mass function in clusters as well as the observed major merger rates since z ∼ 1 (e.g. Nipoti, Londrillo & Ciotti 2003; Bundy et al. 2009). On the other hand, minor mergers are able to produce an efficient size growth (see e.g. Trujillo et al. 2011; Shankar et al. 2013). The rates of minor mergers are roughly enough to account for the size evolution only up to z ≲ 1 Newman et al. (2012), at z ∼ 2 additional mechanisms are required (e.g. AGN feedback-driven star formation Ishibashi, Fabian & Canning 2013). In addition, the effect of continual quenched galaxies on to the red sequence as well as morphological mixing (known as the ‘progenitor bias’) further complicates the situation (e.g van Dokkum & Franx 2001). Processes that are specific in clusters such as harassment, strangulation and ram-pressure stripping (e.g. Treu et al. 2003; Moran et al. 2007) might play an important role in quenching and morphologically transforming galaxies. Several studies have already shown that the progenitor bias has a non-negligible effect on the size evolution (e.g. Saglia et al. 2010; Valentinuzzi et al. 2010b; Carollo et al. 2013; Poggianti et al. 2013; Beifiori et al. 2014; Delaye et al. 2014; Belli, Newman & Ellis 2015; Shankar et al. 2015).


We note that during these encounters, energy is also lost due to tidal oscillations in the neutron star excited by the black hole (e.g., Press & Teukolsky 1977), and contributes to σ. To check whether we should include this effect in our calculations or whether it can be safely neglected, we approximate the tidal energy dissipated during a parabolic encounter according to the formalism presented in Press & Teukolsky (1977):
6

where RNS is the radius of the neutron star, Rmin is the periastron of the approach, and Tl are dimensionless values associated with each spherical harmonic l (see Press & Teukolsky 1977 for calculation of Tl). We only consider the quadrupole mode (l = 2), which dominates over the other modes (Press & Teukolsky 1977). We approximate the NS as a polytropic star of index n = 0.5 (e.g., Finn 1987), and use values from Table 1 of Kokkotas & Schafer (1995) to aid in the calculation of Tl. Note that since there is a minimum impact parameter, there is a minimum possible value of Rmin. For a parabolic encounter the relationship between the impact parameter and the periastron distance is:
7

Thus, combining Equations (4) and (7), we find the minimum possible Rmin to be:
8

Encounters with Rmin  Rmin(bmin) will result in a direct collision between the BH and NS instead of a bound binary. In Figure 1 we plot

—the ratio of energy lost to tidal oscillations to the energy lost to gravitational waves—as a function of Rmin, for a 5

 BH and a 1.4

 NS. We have also marked the region where

. We see that in the region of interest where

, i.e., where bound binaries can form,

, an extremely small value. We have verified (not shown to avoid clutter) that for larger BH masses,

 is even smaller. This is consistent with previous studies about NS–NS captures (e.g., Gold et al. 2012; Chirenti et al. 2017). Thus, we can safely neglect tides in our calculation of the capture cross section.


Another way of obtaining information about the nucleosynthesis processes involved in producing carbon is to compare it with other elements that are characterised by a well-known source of production, as in the case of oxygen. In Fig. 5, we show the variation of [C/O] as a function of [Fe/H], which serves as a first-order approximation to the evolution with time. To calculate the [C/O] ratios, two oxygen abundance indicators are used independently. At subsolar metallicities, the abundance ratios with both oxygen indicators are mostly negative and show an increasing trend towards higher metallicity. This is explained by the fact that oxygen is entirely produced by SNe Type II from massive progenitors, which started to release theiryields at earlier ages in the Galaxy and, hence, at lower metallicities (e.g. Woosley & Weaver 1995). The massive stars producing carbon at low metallicities might be less massive than those producing oxygen (i.e. having a longer life), explaining a delayed contribution of carbon, hence, the negative [C/O] ratios. Alternatively, this could be explained by increasing O/C yields for more massive progenitors of SNeII. Once metallicity starts to increase, low- and intermediate-mass stars release carbon and massive stars start to eject more carbon than oxygen (Carigi et al. 2005). The [C/O] ratio seems to have a constant rise towards higher metallicities when using the forbidden oxygen line. However, in the case when the O I 6158 Å line is employed, we do observe that the maximum in [C/O] takes places close to solar metallicity to then become flat or decrease. This suggests that low-mass stars mostly contribute to carbon around solar metallicity, whereas at super-solar metallicities, massive stars produce carbon together with oxygen, thereby flattening or even decreasing the [C/O] ratio. This trend is in agreement with the metallicity dependent yields from Carigi et al. (2005), which provide higher carbon as [Fe/H] increases from massive stars (i.e. also increasing the O production) but lower carbon from low and intermediate mass stars as [Fe/H] increases (i.e. less production of C). The turning point of increased relative production of carbon from massive stars takes place at A(O) ~ 8.7 dex (see Fig. 2 of Carigi et al. 2005) which equals to [O/H] ~ 0.0 dex. This observed behaviour of [C/O] is in contrast to the steady increase of [C/O] up to [Fe/H] ~ 0.3 dex found, for example, by Franchini et al. (2021). Nevertheless, the general trend we find when using the [O I ] 6300 Å line is similar to the reported by Franchini et al. (2021), who use also that oxygen indicator. All thick-disk stars present negative [C/O] ratios and when using the oxygen line at 6158 Å thin-disk stars with [Fe/H] ≲ –0.2 have [C/O]  0 as well. Thick-disk stars and low-metallicity thin-disk stars at the same metallicity have similar [C/O] ratios, meaning that the balance between different production sites for oxygen and carbon is the same among both populations, despite [C/Fe] and [O/Fe] being systematically higher for thick-disk stars at a given metallicity.


It is generally thought that the emission from radio to medium energy gamma rays is generated by the injected electrons through the synchrotron radiative mechanism. The high-energy photon emission mainly comes from inverse Compton (IC) scattering of the high-energy electrons on the background seed photons, which include the synchrotron background, the cosmic microwave background, and infrared photons in the PWNe (see, e.g., Zhang et al. 2008; Fang & Zhang 2010; Torres et al. 2013; Lu et al. 2020). On the other hand, it is also suggested that the gamma rays could be emitted by the hadronic processes. The relativistic protons accelerated in the Crab pulsar outer gap interact with the matter inside the nebula. and this process may contribute in the high-energy gamma-ray range (see, e.g., Cheng et al. 1990; Bednarek & Protheroe 1997; Khangulyan et al. 2020; Cao et al. 2021). Therefore, it has been long debated whether the high-energy emission from the PWNe is the leptonic or hadronic origin. The details of the high-energy radiation produced by leptonic process have been discussed for the Crab Nebula (see, e.g., Venter & de Jager 2007; Zhang et al. 2008; Martín et al. 2012), and that of the gamma-ray emission about the hadronic process have been also investigated (see, e.g., Cheng et al. 1990; Bednarek & Protheroe 1997; Bednarek 2003, 2007). However, with the establishment of more and more high-energy observatories, some telescopes have possessed the performance of observing photons of exceeding to the PeV from the astronomic objects. An increasing number of observational data has been reported by the different experiments. For example, the Amenomori et al. (2019) reported that the Tibet air shower array with the underground water-Cerenkov-type muon detector array observed the highest energy photons of exceeding 100 TeV with a 5.6σ statistical significance and pointed the measured spectrum with energy extended to the sub-PeV from the Crab Nebula have an absence of high-energy cutoff. Recently, more than 530 photons at energies above 100 TeV and up to 1.4 PeV from the 12 ultra-high-energy gamma-ray sources with a statistical significance greater than seven standard deviations were reported again by LHAASO (Cao et al. 2021). Together with the earlier investigations about the leptonic scenario, the radiative spectrum from the leptons has a cutoff around the sub-PeV region (see, e.g., Zhang et al. 2008; Martín et al. 2012; Zhang et al. 2020). It seems that the other components of gamma rays have a significant contribution.


As a comparison, we also use the publicly available mock catalogue JAdes extraGalactic Ultradeep Artificial Realizations (JAGUAR; Williams et al. 2018) to derive the relation between UV and [O iii] + H β luminosity of simulated z ∼ 8 galaxies. The JAGUAR mock catalogue has been produced by matching luminosity and stellar mass functions as well as the relation between the stellar mass and UV luminosity, mostly at z ≤ 4. The galaxy properties are then extrapolated up to z ∼ 15. The JAGUAR catalogue provides emission line fluxes and EWs for the main lines based on modelling with the beagle code (Chevallard & Charlot 2016; Chevallard et al. 2018a). We identify all galaxies from the fiducial JAGUAR mock in the redshift range 7.11  z  9.05 and we randomly select 1000 of them to match the absolute UV magnitude distribution of our sample, and then fit the UV-[O iii] + H β luminosity data. The result is shown in red in Fig. 6. Similarly to our sample, the z ∼ 8 galaxies from the JAGUAR catalogue exhibit a tight relation between UV and [O iii] + H β luminosity (Spearman rank correlation coefficient ρ = 0.73, standard deviation from null hypothesis σ > 40). However, the mock galaxies exhibit a significantly lower [O iii] + H β luminosity (∼0.5 dex) at a given LUV compared to the relation of our galaxies. The detailed reason for this discrepancy relative to the JAGUAR mock is unclear, but one possible reason is differences in the median physical properties. For instance, while the mock galaxies exhibit (3.6–4.5)$\mu$m colour similar to the ones from our sample at a given UV luminosity, the average F125W-3.6$\mu$m colour in JAGUAR is smaller by ∼0.3 mag compared to the observed F125W-3.6$\mu$m colour in our sample. This means that while (3.6–4.5)$\mu$m colour and EW([O iii] + H β) are on average similar between JAGUAR and our sample, the absolute [O iii] + H β line luminosity scales with the 3.6$\mu$m flux which is larger in our sample compared to the JAGUAR mock catalogue. Furthermore, JAGUAR models a small field comparatively to our data, therefore the overlap in UV luminosity is small.


Unlike the CMB, the 21 cm signal is highly non-Gaussian, because patchy, bubble-like structures of ionized hydrogen (H ii) regions are produced surrounding the ionizing sources. Thus, there is potentially a wealth of information in the 21 cm signal that is not contained in the 21 cm power spectrum, a two-point statistics of 21 cm brightness temperature fluctuations that is traditionally well studied in the literature. It is therefore essential to develop new methods that maximally exploit the full information in the 3D 21 cm images obtained by the SKA. Toward this goal, conventionally, new summary statistics that can only be measured with imaging have been proposed. These include the three-point correlation function (Hoffmann et al. 2019; Jennings et al. 2020), bispectrum (Yoshiura et al. 2015; Shimabukuro et al. 2016, 2017; Majumdar et al. 2018, 2020; Hutter et al. 2020; Saxena et al. 2020; Kamran et al. 2021), one-point statistics (Harker et al. 2009; Shimabukuro et al. 2015; Gorce et al. 2021), topological quantities such as the Minkowski functionals (Gleser et al. 2006; Chen et al. 2019; Kapahtia et al. 2021) and Betti numbers (Giri & Mellema 2021), the cross correlation between the 21 cm line and other probes, such as the CO line (Gong et al. 2011; Lidz et al. 2011), the C ii line (Gong et al. 2012; Beane & Lidz 2018), the kinetic Sunyaev–Zel’dovich (kSZ) effect (Ma et al. 2018; La Plante et al. 2020), and novel techniques such as the antisymmetric cross correlation between the 21 cm line and CO line (Zhou et al. 2021). Since those summary statistics are fully determined by the parameters in the reionization models (hereafter “reionization parameters”), in principle, Monte Carlo Markov Chain (MCMC) methods can be employed to constrain the reionization parameters from measurements of those statistics with futuristic 21 cm experiments (see, e.g., Watkinson et al. 2022), just as the MCMC analysis with the 21 cm power spectrum (Greig & Mesinger 2015, 2017, 2018).


The progenitors of SNe Ib/c have been thought to be Wolf-Rayet (W-R) stars with high initial masses (MZAMS ≳ 25 M⊙; Crowther 2007). Before core collapse, these stars usually have experienced severe mass loss through strong stellar winds or due to interaction with companion stars (van der Hucht 2006; Paxton et al. 2015). As the evolution of massive stars is usually dominated by binary evolution (Heger et al. 2003) and also depends largely on metallicity, rotation, and so on (Heger et al. 2003; Georgy et al. 2013, 2012), this makes the direct identification of their progenitors complicated (Smartt 2015). However, there are increasing studies suggesting that a lower-mass binary scenario is more favorable for most SNe Ib/c, considering the measured low ejecta masses (Eldridge et al. 2013; Lyman et al. 2016). In addition, the H/He envelopes of the progenitor stars are stripped by binary interaction. There are many detections of progenitor stars for SNe II. For example, most SNe IIP are found to originate from red supergiants (Smartt et al. 2009), while SNe IIL are typically from progenitors with somewhat warmer colors (see Smartt 2015, for a review), and SNe IIb are from those with higher effective temperatures such as yellow supergiants that have had their H/He envelopes partially stripped through binary interaction (e.g., SN 1993J; Podsiadlowski et al. 1993; Maund et al. 2004; Fox et al. 2014). Until recently, there has been only one report of the possible identification of a progenitor star for SNe Ib, namely, iPTF 13bvn, which was proposed to spatially coincide with a single W-R-like star identified on the pre-explosion Hubble Space Telescope (HST) images (Cao et al. 2013; Groh et al. 2013). But such an identification is still controversial (e.g., Bersten et al. 2014; Fremling et al. 2014; Eldridge et al. 2015; Eldridge & Maund 2016). Direct detection of progenitor stars is still elusive for SNe Ic, which prevents us from further testing the theoretical evolution of massive stars (Eldridge et al. 2013).


Figure 1 shows the main features of the butterfly diagram: (1) the onset of the cycle at mid-latitudes; (2) the sunspot drift toward the equator and its slowdown represented by a change in the slope of the butterfly wing (Maunder 1904; Li et al. 2001); (3) the tail-like attachment over the minimum phase that is more prominent when the activity is stronger, which might lead to the overlap of successive cycles; (4) the length of the overlap varies within 1–2 years. It characterizes only the minimum phase and it is confined at latitudes ≤15° (Cliver 2014). This feature is also seen in torsional oscillations shown in the bottom panels of Figure 1 (Wilson et al. 1988; Howe et al. 2009). The rate of drift of sunspots toward the equator slows as the sunspot band approaches the equator, and halts at about 8° latitude (Hathaway et al. 2003). The end of the migration does not correspond to the end of the activity because it produces the tail-like attachment. When the new cycle at mid-latitudes starts before the end of the old cycle at low latitudes, it causes successive cycles to overlap. FTD models driven only by the Babcock–Leighton mechanism (Chatterjee et al. 2004), or along with the α-turbulent effect operating in the bulk of the convection zone, currently have the best agreement with observations (Passos et al. 2014), because the length of the simulated overlap is short and it occurs only during the minimum at low latitudes. Conversely, thin-shell dynamo wave models (Moss & Brooke 2000; Schüssler & Schmitt 2004; Bushby 2006) or the thin-shell flux transport dynamo (Dikpati & Gilman 2001) tend to produce dynamo waves with too short a wavelength, leading to excessive overlap between adjacent cycles because this involves a wider range of latitudes. Furthermore they also fail to reproduce the tail-like attachment over the minimum phase. Moreover the direction of the migration of activity could also provide information on the nature of the α mechanism. Both formalisms make strong assumptions to initiate the sunspot cycle at mid-latitudes. The Babcock–Leighton FTD models assume that the deep equatorward meridional flow penetrates slightly below the convection zone to a greater depth than usually believed (Nandy & Choudhuri 2002), in order to prevent the onset and occurrence of a sunspot cycle above 45° as well as any other kind of cyclic activity. The same result is achieved with the αΩ dynamo wave by inhibiting the α-turbulent effect at higher latitudes (Schüssler & Schmitt 2004). Based on these assumptions, the magnetic activity in any type of FTD model starts at higher latitudes and then propagates only equatorward, while in the thin-shell αΩ dynamo wave the magnetic activity can propagate poleward as well as equatorward (Bushby 2006). These two branches are also clearly seen in the torsional oscillation pattern (e.g., Howe et al. 2009). This results from the solar-like differential profile, which is characterized by a sign change in

 in the tachocline at high and low latitudes (Ruediger & Brandenburg 1995). This sign change, however, has not yet been confirmed by helioseismic observations.


We have presented MoCA, a Monte Carlo code for Comptonisation in Astrophysics which includes polarisation. To our knowledge MoCA is the first code operating with single photons and including all special relativity and quantum effects. The main disadvantage of this approach is the long computing time, which implies the need to parallelise the code on clusters of computers. The advantage with respect to pure analytical models such as those available in XSPEC is that we can explore the totality of the parameters space for the Comptonising medium (i.e. thermal energy and optical thickness of the corona) without any restriction and this approach will also allow a better understanding of the whole process. We also included all corrections such as Klein–Nishina cross-section and scattering angle distribution. These effects, small below 100 keV, must nonetheless be taken into account when inferring the thermal energy of the corona from observations. In some sources it has been found that coronae can have extremely high energy cut-off (e.g. NGC 5506, Matt et al. 2015) and therefore thermal energy, which is inferred by measuring the curvature of NuSTAR spectra at high energy and in this context K–N effects cannot be neglected. From the polarimetric point of view we did not seen any deviation due to K–N effects, but this was to expected as we focussed our attention below 100 keV where these effects are small. However, one can imagine a scenario in which the thermal energy of the corona is few tens of keV and in that case we expect to see a difference both on the spectrum and the polarisation but we defer such investigation to future papers focussed on the exploration of the coronal parameters space. In its actual form the code is fast enough to explore different geometries of the corona with different parameters. Spectra can then be compared with those obtained by NuSTAR to derive coronal parameters, especially in the high optical depth regime where analytical models are not reliable. As already mentioned, much observational evidence points in the direction of compact coronae above or around the compact object. In order to properly treat such coronae, gravitational effects must be taken into account. We have recently included a ray-tracing routine to take into account GR effects: this new version of MoCA, and applications of the code to different astrophysical scenarios, will be discussed in future papers. Nonetheless we have shown the geometrical effect of more compact coronae on the spectra and the polarisation signal: the spectra become softer as the corona shrinks and the polarisation changes dramatically as we approach a more symmetrical shape of the corona. The study we performed will also be useful to quantify the impact of GR effects on compact coronae with respect to a purely geometrical effect.


The C/O ratio varies across exoplanets’ host star populations (Delgado Mena et al. 2010; Brewer & Fischer 2016; Brewer et al. 2017), and this variation is likely to be reflected in the composition of exoplanet atmospheres, assuming that they are formed with the same materials as their stars. Moreover, various processes in the protoplanetary disks and the planet formation process can affect the exoplanet compositions and have a significant impact on the final C/O ratio (Öberg et al. 2011; Mordasini et al. 2016; Espinoza et al. 2017; Madhusudhan et al. 2017). For these reasons, it is necessary to consider the effects of the C/O ratio on the atmospheric chemistry and the formation of aerosols. Numerous studies have been performed using chemical models (Madhusudhan 2012; Moses et al. 2013; Venot et al. 2015; Tsai et al. 2017; Heng & Lyons 2016; Goyal et al. 2018; Drummond et al. 2019), but corresponding laboratory experiments are still largely nonexistent. Laboratory investigations can provide essential insight into the effects of the C/O ratio on the atmospheric photochemistry and the formation of aerosols. In a previous work, we performed the first laboratory experiments dedicated to the study of the chemistry in hot Jupiter atmospheres (Fleury et al. 2019). This work focused on the chemistry in atmospheres with T > 1000 K and a C/O ratio of 1 (representing C enhancement compared to the solar value of 0.54), because chemical models predict that the abundances of hydrocarbon and nitrile species increase by several orders of magnitude in these atmospheres compared to atmospheres with a low C/O ratio (Venot et al. 2015). Therefore, they can be considered as better candidates for the formation of complex organic molecules with longer carbon chains. This first study revealed that photochemical aerosols could be produced at temperatures as high as 1500 K and that water could be efficiently formed through photochemical channels. In the present work, we performed new experiments to study the chemistry in hot Jupiter atmospheres at similar temperatures (1173–1473 K) but with lower C/O ratios. We used a gas mixture of H2, H2O, and CO that represents the simplest plausible atmosphere for a hot Jupiter with a C/O ratio  1. This new study, compared with our previous work, allows us to assess the evolution of the chemistry in hot Jupiter atmospheres as a function of the C/O ratio and atmospheric composition.


While these early studies demonstrated the viability of neutron star mergers as a major r-process site, they identified only one ejection channel: ‘dynamical ejecta’ that are tidally flung out by gravitational torques. Since they are never substantially heated, these ejecta carry their original β −equilibrium electron fraction from the original neutron star, Ye ≈ 0.05, and this enormous neutron-richness allows them to undergo a ‘fission cycling’ process (Goriely, Bauswein & Janka 2011; Korobkin et al. 2012), which produces a very robust r-process abundance distribution close to the solar pattern for A ≥ 130, but hardly any lighter r-process elements. Oechslin, Janka & Marek (2007) pointed out that there is a second channel of mass ejection that also happens on a dynamical time-scale: shock-heated matter from the interface where the stars come into contact. As of today, many more mass ejection channels have been discussed: matter that becomes unbound on secular time-scales (∼1 s) from the post-merger accretion torus (Beloborodov 2008; Metzger, Piro & Quataert 2008; Fernandez & Metzger 2013; Fernandez et al. 2015; Just et al. 2015; Siegel & Metzger 2017, 2018; Fernandez et al. 2019; Miller et al. 2019a), as MHD-driven winds (Siegel & Ciolfi 2015) and by viscous effects (Shibata, Kiuchi & Sekiguchi 2017; Radice et al. 2018a; Shibata & Hotokezaka 2019) from a long-lived neutron star merger remnant. Similar to the case of proto-neutron stars, the enormous neutrino luminosities (>1053 erg s−1) after a neutron star merger can also drive substantial matter outflows (Ruffert et al. 1997; Rosswog & Ramirez-Ruiz 2002; Dessart et al. 2009; Perego et al. 2014; Martin et al. 2015; Radice et al. 2018b). The secular torus ejecta contain approximately 40 per cent of the initial torus mass and, although the latter may vary substantially from case to case, they likely contribute the lion’s share to the total ejecta mass. Due to their different thermal histories and exposure times to neutrinos, the ejecta channels can have different electron fractions Ye and therefore different nucleosynthesis yields.1 For electron fractions below a critical value, $Y_{\rm e}^{\rm crit}\approx 0.25$ (Korobkin et al. 2012; Lippuner & Roberts 2015), lanthanides and actinides are efficiently produced, which, due to their open f-shells, have particularly high bound–bound opacities (Barnes & Kasen 2013; Kasen, Badnell & Barnes 2013; Tanaka & Hotokezaka 2013; Tanaka et al. 2020) and therefore lead to red transients that peak days after the merger. Ejecta with electron fractions above $Y_{\rm e}^{\rm crit}$, in contrast, only produce ‘lighter’ elements with lower opacities and thus result in bluer transients that peak after about 1 d. Opaque, low-Ye ejecta blocking the view on high-Ye ejecta can lead to a ‘lanthanide curtaining’ effect (Kasen, Fernández & Metzger 2015; Wollaeger et al. 2018), which will efficiently block blue light. Therefore, it is important to understand the layering, dynamics, interaction and potential mixing of different ejecta channels.


The most widely adopted parametrization of the observed Universe is based on the so-called Λ cold dark matter (ΛCDM) model (Peebles 1984), relying on the existence of cold dark matter and dark energy (Λ) associated with a cosmological constant (Carroll 2001) in a spatially flat geometry. Predictions from this model have been found to agree with most of the observational probes such as the cosmic microwave background (CMB; e.g. Planck Collaboration 2020), the baryon acoustic oscillations (BAO; e.g. Alam et al. 2021), and the present accelerated expansion of the Hubble flow, based on the distance modulus–redshift relation (the so-called Hubble–Lemaître, or simply Hubble diagram) of type Ia supernovae (SNe Ia; e.g. Riess et al. 1998; Perlmutter et al. 1999), where a dominant dynamical contribution, dubbed dark energy (DE) and related to the cosmological constant, should drive such an acceleration. However, the fundamental physical origin and the properties of DE are still unknown, as the interpretation of Λ is plagued by a severe fine-tuning issue to obtain the right amount of DE observed today. Moreover, the data sets listed above do not fully fit the evolution of DE ranging from early to late epochs (Benetti et al. 2019; Yang et al. 2021) and do not fully rule out a spatially non-flat Universe (Park & Ratra 2019; Di Valentino, Melchiorri & Silk 2020, 2021; Handley 2021; Yang et al. 2021). The latter possibility has raised a remarkable debate about the importance of properly combining CMB data to infer significant statistical interpretations from the analysis (Efstathiou & Gratton 2020; Planck Collaboration 2020) and, by extension, the importance of combining data sets that do not reveal manifest tension (Gonzalez et al. 2021; Vagnozzi, Loeb & Moresco 2021). Deviations from the spatially flat ΛCDM model would imply important theoretical and observational consequences and a change in our current understanding of cosmic evolution (e.g. Capozziello, Benetti & Spallicci 2020). Statistically significant deviations in this directions have already been found in cosmological analyses with high-redshift probes such as Gamma-Ray Bursts (see Dainotti, Cardone & Capozziello 2008; Dainotti et al. 2011b, 2013a,b, 2015, 2017, 2020a; Dainotti, Ostrowaki & Willingale 2011a; Dainotti et al. 2020b for the standardization of these sources as cosmological candles) and quasars (QSOs) combined with SNe Ia (Risaliti & Lusso 2019; Lusso et al. 2019, 2020; Bargiacchi et al. 2021). Such a joint analysis (SNe + QSO) makes use of the observed non-linear relation between the ultraviolet and the X-ray luminosity in QSOs (e.g. Steffen et al. 2006; Just et al. 2007; Lusso et al. 2010; Lusso & Risaliti 2016; Bisogni et al. 2021; Dainotti et al. 2022) to provide an independent measurement of their distance (see e.g. Risaliti & Lusso 2015, 2019; Lusso et al. 2020, for details). The methodology is complementary to the traditional resort to type Ia SNe to estimate the cosmological parameters, yet it extends the Hubble–Lemaître diagram to a redshift range currently inaccessible to SNe ($\mathit{ z}$ = 2.4–7.5). Within a model where an evolution of the DE equation of state (EoS) in form w($\mathit{ z}$) = w0 + wa × $\mathit{ z}$/(1 + $\mathit{ z}$) is assumed, the data suggest that the DE parameter is increasing with time (Risaliti & Lusso 2019; Lusso et al. 2020). Therefore, it is compelling to further study extensions of the ΛCDM model that could produce such behaviour of DE.


Over the past decade, many attempts to address this issue have been carried out, such as from the perspective of conventional nuclear physics and even exotic physics beyond the standard BBN framework (Angulo et al. 2005; Cyburt et al. 2008, 2016; Boyd et al. 2010; Pospelov & Pradler 2010; Fields 2011; Kirsebom & Davids 2011; Wang et al. 2011; Broggini et al. 2012; Coc et al. 2012, 2013, 2014; Cyburt & Pospelov 2012; Kang et al. 2012; Voronchev et al. 2012; Bertulani et al. 2013; Hammache et al. 2013; He et al. 2013; Kusakabe et al. 2014; Pizzone et al. 2014; Yamazaki et al. 2014; Hou et al. 2015, 2017; Famiano et al. 2016; Damone et al. 2018; Hartos et al. 2018; Luo et al. 2019; Rijal et al. 2019; Clara & Martins 2020). However, despite the fact some solutions using exotic physics have succeeded in resolving this issue, it appears there is still no universally accepted solution in the academic community since validations of these mysterious exotic physics are beyond the capabilities of current science. Conversely, it seems more worthwhile to exclude any potential possibility of resolving the 7Li discrepancy from the perspective of nuclear physics. It is known that the majority of the primordial 7Li production arises from the decay of 7Be by electron capture during the 2 months after BBN stops. Thus, for the solution of the Li problem, reactions involving 7Be could be more significant than those involving 7Li. Therefore, many reactions that potentially destroy 7Be were investigated to solve this discrepancy over past 10 yr (Kirsebom & Davids 2011; Broggini et al. 2012; Hammache et al. 2013; Hou et al. 2015; Hartos et al. 2018). Meanwhile, enormous efforts have been made to refine the reaction rates of key BBN reactions in the past 20 yr (Smith et al. 1993; Descouvemont et al. 2004; Serpico et al. 2004; Cyburt & Davids 2008; Neff 2011; Pizzone et al. 2014; Tumino et al. 2014; Hou et al. 2015; Barbagallo et al. 2016; Iliadis et al. 2016; Kawabata et al. 2017; Lamia et al. 2017, 2019; Damone et al. 2018; Rijal et al. 2019; Mossa et al. 2020), but the probability of solving or alleviating the 7Li problem by improving our knowledge of relevant nuclear reaction rates still cannot be eliminated. Recent experiments for key nuclear reactions like 7Be(n,p)7Li and 7Be(d,p)24He allow for a reduction of the 7Li production by about 12% (Damone et al. 2018; Rijal et al. 2019) compared to previous calculations. At present, nuclear uncertainties cannot rule out that some of the reactions destroying 7Li are indeed more efficient than those currently used (Boyd et al. 2010; Chakraborty et al. 2011).


ESA's Rosetta spacecraft orbited closely around comet 67P/Churyumov–Gerasimenko (hereafter 67P) during its 2 yr mission. From 2014 August to 2016 September, the camera system Optical, Spectroscopic, and Infrared Remote Imaging System (OSIRIS; Keller et al. 2007) onboard Rosetta observed the comet nucleus down to few centimetres per pixel providing detailed images of the comet's surface and activity. The study of 67P morphology has been performed with unprecedented detail thanks to the OSIRIS image spatial coverage and resolution. A first OSIRIS analysis on the nucleus structure, morphology and jets activity is reported in Sierks et al. (2015) and Thomas et al. (2015). The comet's surface revealed a wide diversity in morphology such as layering (Massironi et al. 2015), pits (Vincent et al. 2015), boulders and fractures (Pajola et al. 2015, 2016c; El-Maarry et al. 2015a), and high reflectivity boulders (Pommerol et al. 2015). During the entire mission, surface changes have been detected thanks to images acquired before and after perihelion passage (Groussin et al. 2015; Pajola 2017; El-Maarry 2017). The global spectrophotometric properties of the comet were investigated in detail by Fornasier et al. (2015) identifying three different groups of terrains depending on spectral slope values. In addition to a global characterization of the comet, OSIRIS allowed a detailed analysis of specific regions identified on the comet by means of high resolution and multifilter images (La Forgia et al. 2015; Deshapriya et al. 2016; Lucchetti et al. 2016; Oklay et al. 2016a; Pajola et al. 2016d; Oklay et al. 2017). We decided to follow the approach of characterizing at the highest resolution possible specific regions of interest (ROIs) on 67P. We therefore performed a detailed analysis on an area belonging to the Seth region (El-Maarry et al. 2015b) that is characterized by flat-floored and steep-walled circular depressions (El-Maarry et al. 2015b; Giacomini et al. 2016). We focused our attention on the circular niches of the Seth area (Fig. 1) and performed a multidisciplinary study of these structures investigating their geomorphological and spectrophotometric properties. In addition, thanks to images acquired pre- and post-perihelion, we conducted a comparative analysis to find if this area has been subjected to surface changes. This can be useful to provide constraints about Seth's niches properties as well as their possible origin being one of the interesting features located on the surface of 67P. Specifically, the circular niches deposits can be considered the result of landslide events that occurred on the comet surface, as the recent Aswan cliff collapse (Pajola 2017), where it has been reported the occurring of falling material from the adjacent cliff after the Rosetta perihelion passage. Hence, with this work we plan to understand if the formation of these Seth's circular niches occurred recently or if it is correlated to older events that have shaped the comet's surface.


The CO emission, tracking the bulk of the gas, reveals an almost face-on one-armed spiral, of which almost two full windings can be traced. What could be the origin of this spiral structure? As the majority of AGB stars are in binary systems, and perhaps all host planets, interaction between the outflow and a sufficiently massive and nearby companion may be the explanation of the observed CO morphology. The intricate emission features in the inner 2″ of the central CO emission maps is strongly reminiscent of hydrodynamical simulations of wind–binary interaction by Mastrodemos & Morris (1998) and Mohamed & Podsiadlowski (2012), where the latter authors performed tailored simulations for the Mira AB system in which the outflow of the AGB star Mira A is perturbed by the presence of its close companion Mira B. The wind–binary interaction that ensues leads to what is known as wind Roche-lobe overflow (WRLOF), where the slow AGB wind is confined to the star’s Roche lobe, while overflowing through the L1 Lagrange point. Gravitational interaction of the overflowing material with the companion produces an intricate feedback system where the stellar outflow material is ejected into the surrounding CSE through two distinct streams (through L2 and the stagnation point3) which combine to form an annular stream. As this stream travels outwards, it creates the larger scale spiral observed in the wind. The morphology resulting from this particular type of wind–binary interaction is shown in Fig. 3 in Mohamed & Podsiadlowski (2012). In Fig. 11 we show the emission pattern seen in the central regions of the CO channel at υ*. We compare this image with the bottom left panel of Fig. 3 in Mohamed & Podsiadlowski (2012), an opacity map of the interaction zone. Though the two properties that are compared differ in nature, they likely still trace the same global morphological structure. Indeed, several of the predicted morphological features can be identified in the data of EP Aqr. The bright central region with a north and southward hook-like extension are strikingly similar, as are the eastern and western crescent-shaped “voids”, the overall shape, and the morphological properties of the small-scale instabilities.


Detecting the halo response to the LMC-induced DM wake would be an exciting advancement in testing our assumptions about the properties of DM, as well as providing key constraints on the potential of the MW and the mass and orbital history of the LMC. However, the GC19 simulations give predictions for the response in the context of smooth MW DM and stellar halos. In reality, the MW stellar halo contains a wealth of substructure that is not yet phase mixed, in the form of stellar streams (e.g., Odenkirchen et al. 2001; Newberg et al. 2002; Belokurov et al. 2006; Grillmair 2006; Shipp et al. 2018; also see Newberg & Carlin 2016 for a recent review) and stellar clouds (e.g., Newberg et al. 2002; Rocha-Pinto et al. 2004; Jurić et al. 2008; Li et al. 2016). In addition, using a sample of MW halo main-sequence turnoff stars from the HALO7D survey (Cunningham et al. 2019a), Cunningham et al. (2019b) observed that the estimated parameters of the velocity ellipsoid (i.e.,

) were different in the different survey fields; these differing estimates could be interpreted as evidence that the halo is not phase mixed over the survey range (

 kpc). They also showed maps of the halo velocity anisotropy β in two halos from the Latte suite of FIRE-2 simulations (introduced in Wetzel et al. 2016), finding that the anisotropy can vary over the range

 across the sky. Some of the variation in the β estimates appeared to correlate with stellar overdensities in the halos, indicating that galactic substructure is at least in part responsible for the different velocity distributions. While some substructure in the halo can be clearly identified as overdensities in phase space and removed from analysis, the presence of velocity substructure in the halo could complicate attempts to detect signatures of the LMC-induced DM wake. For example, Belokurov et al. (2019) recently argued that the Pisces Overdensity (Sesar et al. 2007; Watkins et al. 2009; Nie et al. 2015) might be stars in the wake trailing the LMC in its orbit, because of their net negative radial velocities. However, it remains difficult to conclusively argue this scenario given that these stars could also be in Galactic substructure (or, perhaps, stars that are in substructure and have been perturbed by the DM wake).


LISA Pathfinder (LPF; Antonucci et al. 2011), a European Space Agency (ESA) mission that operated near the first Sun–Earth Lagrange point (L1) from 2016 January through 2017 July, is in an ideal orbit to make such measurements. However, LPF flew no instrumentation dedicated to micrometeoroid or dust detection. LPF’s primary objective was to demonstrate technologies for a future space-based observatory of millihertz-band gravitational waves. The key achievement of LPF was placing two gold-platinum cubes known as “test masses” into a freefall so pure that it was characterized by accelerations at the femto-g level (e.g., Armano et al. 2016, 2018b), the level required to detect the minute disturbances caused by passing gravitational waves. In order to reach this level of performance, the test masses were released into cavities inside the spacecraft and a control system was employed to keep the spacecraft centered on the test masses. This control system was designed to counteract disturbances on the spacecraft, including those caused by impacts from micrometeoroids. Shortly before LPF’s launch, it was realized that data from the control system, if properly calibrated, could be used to detect and characterize these impacts and infer information about the impacting particles (e.g., Thorpe et al. 2016). While such impact events have been reported by other spacecraft, LPF’s unique instrumentation makes it sensitive to much smaller and much more numerous impacts and allows the impact geometry to be more fully constrained. Early results from the first few months of LPF operations suggested that such events could indeed be identified and were roughly consistent with the pre-launch predictions of their effect on the control system (e.g., Thorpe et al. 2017). In this paper we present results from the first systematic search for micrometeoroid impacts in the LPF data set. Our data set consists of 4348 hr of data in both the nominal LPF configuration and the “Disturbance Reduction System” (DRS) configuration, in which a NASA-supplied controller and thruster system took over control of the spacecraft (Anderson et al. 2018). Our data set corresponds to the times when LPF was operating in a “quiet” mode, without any intentional signal injections or other disturbances. During this period, we have identified 54 impact candidates using our detection pipeline and manual vetoing. We have characterized the properties of this data set and compared it to several theoretical models for the underlying dust population.


On the other hand, observational data are used to test these models. Among the most frequently used are the cosmic microwave background radiation (CMB; Planck Collaboration XIII 2016; Aghanim et al. 2018), baryonic acoustic oscillations (BAO; Eisenstein et al. 2005; Blake et al. 2012; Alam et al. 2017; Bautista et al. 2017), Type Ia supernovae (SNe Ia; Scolnic et al. 2018), and observational Hubble data (OHD; Jimenez & Loeb 2002; Moresco et al. 2016; Magaña et al. 2018). Consistency in the cosmological parameters among different techniques, rather than more accurate measurements, is desirable to better understand the nature of DE. In the last years, several efforts have been made by the community to include gravitational lens systems in the study of the Universe’s evolution. Some of the pioneers are Futamase & Yoshida (2001) and Biesiada (2006), who used only one strong-lens system to study some of the most popular cosmological models. Grillo, Lombardi & Bertin (2008) introduced a methodology to estimate cosmological parameters using strong-lensing systems (SLS; see also Jullo et al. 2010; Magaña et al. 2015, 2018). They apply the relation between the Einstein radius and the central stellar velocity dispersion, assuming an isothermal profile for the total density distribution of the lens (elliptical) galaxy. Their simulations found that the method is accurate enough to obtain information about the underlying cosmology. They concluded that the stellar velocity dispersion and velocity dispersion of the isothermal lens model are very similar in the w cold dark matter (wCDM) model. Biesiada, Piórkowska & Malec (2010) used the same procedure comparing a distance ratio, Dobs, constructed from SLS observations such as the Einstein radius and spectroscopic velocity dispersion of the lens galaxy, with a theoretical counterpart, Dth. By using a sample containing 20 SLS, they demonstrated that this technique is useful to provide insights into DE. Cao et al. (2012) updated the sample to 80 systems and proposed a modification that takes into account deviations from sphericity, i.e. from the singular isothermal sphere (SIS). Later on, Cao et al. (2015) considered lens profile deviations due to the redshift evolution of elliptical galaxies by using spherically symmetric power-law mass distributions for the lenses and also increased the compilation up to 118 points. They also explore the consequences of using aperture-corrected velocity dispersions on the parameter estimations. Some authors have pointed out the need for a sufficiently large sample to test DE models with higher precision (Yennapureddy & Melia 2018). For instance, Melia, Wei & Wu (2015) have emphasized that a sample of ∼200 SLS can discern the Rh = ct model from the standard one. Qi et al. (2018) simulated strong lensing data to constrain the curvature of the Universe and found that, by increasing the sample (16000 lenses) and combining with compact radio quasars, it could be constrained with an accuracy of ∼10−3. Recently, Leaf & Melia (2018) have revisited this cosmological tool with the largest sample of SLS (158) until now, including 40 new systems presented by Shu et al. (2017). The authors proposed a new approach to improve this technique by introducing in the observational distance ratio error (δDobs), a parameter σx to take into account the SIE scatter and any other source of errors in the measurements. In their analysis, they excluded 29 SLS that are outside the region 0  Dobs  1, and the system SL2SJ085019−034710 (Sonnenfeld et al. 2013b), which seems to be an extreme outlier for their models. Their results show that a $\sigma _x = 12.2{{\ \rm per\ cent}}$ provides more statistically significative cosmological constraints. Finally, Chen, Li & Shu (2018) used 157 SLS to analyse the Lambda cold dark matter (ΛCDM) model. They considered a lens mass distribution ρ(r) = ρ0r−γ and three possibilities for the γ parameter: a constant value, a dependence with the lens redshift (zl), and a dependence with both the surface mass density and the lens redshift. They concluded that although Ω0m, used as the only free parameter in ΛCDM scenario, is very sensitive to the lens mass model, it provides weak constraints that are also in tension with Planck measurements.


Close to the end of their formation, stars are surrounded by a gas and dust disk, from which planets form. Magnetic fields are known to play a key role in the star-disk system (Johns-Krull 2014). It is believed that the inner regions of the disk are significantly ionized by the stellar radiation and that accreting material flows along magnetic channels that connect the disk to the star (Koenigl 1991). Very long and intense X-ray flares in star-forming regions might occur in such long channels (Favata et al. 2005), but this is highly debated (Getman et al. 2008). These flux tubes might resemble those observed in the solar corona and diagnosed in the stellar coronae, but on a much larger scale. On the Sun we see the so-called coronal loops on the scale of several thousand kilometers in active regions, but some faint large-scale structures can extend up to ∼1 R⊙ (Reale 2014). Most solar flares occur in active region loops, but the long-lasting ones can involve more and more loops aligned in arcades. The other stars are so distant that we cannot resolve the flaring regions, but it is supposed that they occur in loops and even in arcades. Whereas the duration of solar flares typically ranges from a few minutes to several hours, stellar flares can be very intense, more than the solar bolometric luminosity, and long-lasting, including longer than one day, in very active stars. Several such gigantic coronal flares have been surveyed in star-forming regions (Favata et al. 2005) and where they occur is a big question. Magnetic instabilities in flux tubes were proposed to be the origin of the flaring activity also in T Tauri stars (Birk 1998; Birk et al. 2000), and long-lasting stellar flares might be expected to involve loop arcades (Getman et al. 2008), like those on the Sun. In long-lasting solar flares, the duration is mainly due to the progressive involvement of more and similar loops, therefore duration is not directly linked to the size of the flaring structures. This might also be the case for giant stellar flares. However, if a single stellar loop were flaring, the cooling time of the confined plasma would be proportional to the loop length (Serio et al. 1991; Reale 2014), and day-long flares would correspond to giant loops, as long as they possibly connected the star with the disk (Hartmann et al. 2016). There are ways to distinguish between a pure cooling in a single loop and a decay only due to progressive reduction of the energy release in a loop arcade (Reale et al. 1997), but the explanations are contested and the uncertainties are large (Getman et al. 2008). Several studies (Favata et al. 2005; Giardino et al. 2007) find results compatible with long magnetic channels in pre-main sequence (PMS) stars, but the derivation of the loop length is based on the assumption of a flare occurring in a single loop (Reale 2007).