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SAO/NASA Astrophysics Data System 13

2020AandA...635A.121M__Casassus_et_al._2018_Instance_1

As scattered light imaging is sensitive to the stellar irradiation, it allows one to search for misalignments between various disk regions. While studying the morphology of the innermost disk region is challenging due to its very small radial extent, often marginally resolvable by optical interferometry (Lazareff et al. 2017), scattered light imaging of the outer disk can indirectly reveal the presence of a misaligned inner disk. In this scenario, depending on the misalignment angle, the outer disk image will show narrow shadow lanes (e.g., Pinilla et al. 2015; Stolker et al. 2016; Benisty et al. 2017; Casassus et al. 2018), broad extended shadows (Benisty et al. 2018) or low-amplitude azimuthal variations (Debes et al. 2017; Poteet et al. 2018). In some cases, studies of the CO line kinematics support a misalignment between inner and outer disk regions (Loomis et al. 2017; Pérez et al. 2018). The exact origin of such a misalignment is still unclear. In the case of T Tauri stars, if the stellar magnetic field is inclined, it can warp the innermost edge of the disk, which would then rotate at the stellar period (AA Tau; Bouvier et al. 2007). Alternatively, inner and outer disk regions can have different orientations if the primordial envelope had a different angular momentum vector orientation at the time of the inner/outer disk formation (Bate 2018). Other scenarios involve the presence of a massive companion/planet that is inclined with respect to the disk. If the companion is massive enough, the disk can break into two separate inner and outer disk regions, that can then precess differently and result in a significant misalignment between each other (e.g., Nixon et al. 2012; Facchini et al. 2013; Nealon et al. 2018; Zhu 2019). A clear example of such a scenario is the disk around HD 142527, in which an M-star companion was detected (Biller et al. 2012), likely on an inclined and eccentric orbit (Lacour et al. 2016; Claudi et al. 2019). Dedicated hydrodynamical simulations successfully reproduce most of the observed features in this disk (eccentric cavity, spiral arms, misaligned inner disk and shadows; Price et al. 2018).

2021MNRAS.507.4367C__Mazzali_et_al._2005_Instance_1

In the previous section, we discussed how SNe Ia with blueshifted Na i D have broader light curves but this may be driven by their preference for late-type host galaxies and their differing environments. One of the main aims of this study was to investigate, for the first time, if there is a connection between the presence of broad high-velocity Ca ii features in SNe Ia and the presence of narrow blueshifted Na i D features. As previously discussed, high-velocity Ca ii features are ubiquitous in SNe Ia and we confirm this here finding that the vast majority of the sample require a high-velocity Ca ii component. We also identify high-velocity Si ii features in 12 SNe Ia in our sample, although these features are not clearly ‘detached’ from the photospheric component. The reason that high-velocity features are interesting to investigate in connection with blueshifted Na i D features is that there are suggestions that the high-velocity Ca ii features may be at least partially due to ejecta–CSM interaction (e.g. Mazzali et al. 2005; Tanaka et al. 2006). Therefore, if a link between these quantities was identified, it would provide evidence that blueshifted Na i D features are due to CSM also rather than contamination. We remind the reader that the two probes do explore different distances from the SN, with Ca ii originating at significantly shorter distances than Na i D. However, when we examine the strength of the high-velocity Ca ii components (parametrized through the ${\rm Ca\, {\small{\rm II}}}~R_{\textrm {HVF}}$ of Childress et al. 2013a) compared to the blueshifted${\rm Na\, {\small{I}}}~\text{D}_{2}$ pEQW (15), we do not identify any clear trend between them. In particular, we also find that a number of SNe Ia with no Na i D absorption at all have very high ${\rm Ca\, {\small{\rm II}}}~R_{\textrm {HVF}}$ values, which is difficult to interpret in the context of a common CSM origin for both strong blueshifted Na i D features and strong high-velocity Ca ii components, as with this interpretation of their origin a strong signature would be expected in both probes rather than singularly. We also find no correlation between the pEQW of the high-velocity Ca ii features and the pEQW of the blueshifted Na i D features.

2019MNRAS.486.4671M__Webb_&_Hundhausen_1987_Instance_1

CMEs are known for large-scale expulsion of magnetized plasma structures from closed magnetic field regions on the Sun. They were first detected in the coronagraphic images taken in 1971 by NASA’s OSO-7 spacecraft (Tousey 1973). However, some definite inferences for the solar wind (Eddington 1910; Birkeland 1916; Biermann 1951) as well as CMEs from the Sun (Chapman & Ferraro 1931; Eddy 1974) were made decades before their formal discovery. Following OSO-7, a series of spacecraft (Skylab, Helios, P78-1 Solwind, SOHO, Coriolis, and STEREO, etc.) have observed thousands of CMEs leading to a vast literature (Munro et al. 1979; Howard et al. 1985; Gosling 1993; Hundhausen 1999; Gopalswamy et al. 2000; Schwenn 2006; Vourlidas et al. 2010; Chen 2011; Wang et al. 2011; Webb & Howard 2012; Mishra & Srivastava 2013; Mishra et al. 2017; Harrison et al. 2018). CMEs have been observed to occur often having spatial and temporal relation with solar flares, eruptive prominences (Munro et al. 1979; Webb & Hundhausen 1987; Zhang et al. 2001; Gopalswamy et al. 2003) and with helmet streamer disruptions (Dryer 1996). Unlike CMEs from the Sun, to observe stellar CMEs are challenging because the close stellar environment cannot be spatially resolved. Although stellar CMEs have not yet been directly detected in Thomson-scattered optical light from other stars, it is believed that the extreme X-ray flares observed on stars may be in conjunction with extreme stellar CMEs (Houdebine, Foing & Rodono 1990; Wheatley 1998; Leitzinger et al. 2011; Aarnio, Matt & Stassun 2012; Osten & Wolk 2015; Vida et al. 2016). Indeed, the stellar X-ray flare, helmet streamers, and prominences observed on T Tauri Stars have shown similarities with those observed on the Sun (Haisch, Antunes & Schmitt 1995; Massi et al. 2008). The CMEs and flares themselves may not be causally related, they both seem to be involved with the reconfiguration of complex magnetic field lines within the corona caused by the same underlying physical processes, e.g. magnetic reconnection (Priest & Forbes 2002; Compagnino, Romano & Zuccarello 2017). But, even for the sun, it has been noted that not all flares are accompanied by CMEs and not all CMEs by flares (Munro et al. 1979; Harrison 1995; Yashiro et al. 2008b; Wang & Zhang 2008).

2021AandA...646A..67W__Widrow_et_al._2014_Instance_1

We consider time-varying dynamical effects, breaking the assumption of a steady state, to be the most probable reason for this unexpected result. We do see how local phase-space substructures can bias our result for individual stellar samples, most clearly for those of area cell B8. However, explaining the steep gravitational potential close to the Galactic mid-plane, inferred for almost all stellar samples, requires a phase-space structure that spans the whole spatial volume that is studied in this work. Indeed, spatially large time-varying phase-space structures are present in the Galaxy, for example in the form of phase-space spirals and ridges (Gaia Collaboration 2018c; Antoja et al. 2018), and Galactic plane mirror asymmetries (Bennett & Bovy 2019, especially prominent for heights |z|≳400 pc). In order to produce a steep gravitational potential at low |z|, there could be a breathing mode in the stellar disk (Widrow et al. 2014; Monari et al. 2016) which is currently in its most compressed state. Such a configuration would not be detectable by comparing the mean vertical velocities above and below the mid-plane, because the breathing oscillation would be at a turning point between contraction and expansion. Mass estimates close to the mid-plane and under the steady state assumption would be biased towards more massive results, as the stellar disk would have a smaller scale height and larger vertical velocities. In order to explain our results, such a breathing mode would have to be large enough for the stellar number density in the mid-plane to oscillate with a relative amplitude of ∼5%. Monari et al. (2016 see for example Fig. 4) have shown that a local breathing mode could be created by a spiral arm that passes close enough to the Sun, inducing a net motion away from (towards) the Galactic mid-plane for stars on the outside (inside) of the spiral arm. Furthermore, such a close passage of a spiral arm is indicated by some dynamical models of the horizontal motions within the Galactic disk (for example Siebert et al. 2012). The dynamics of the Galactic disk are very complex and probably contains a combination of breathing and bending modes created by the last impact of the Sagittarius dwarf galaxy (for example Laporte et al. 2019); it is very plausible to think that the famous phase-space spirals in the z–w plane are a local manifestation of these larger scale perturbations. The phase-space spiral structure could have an effect on our analysis, especially for the stellar samples with higher Zlim.. On a larger scale, the bending and breathing modes of the Galaxy have been shown to affect dynamical mass measurements of the Galactic disk, especially at greater heights (Banik et al. 2017; Haines et al. 2019).

2016AandA...587A.159G__Tian_et_al._2014_Instance_1

One has to be sure to rule out cases where inorganic chemistry can mimic the presence of life (“false positives”). Potential abiotic ozone production on Venus- and Mars-like planets has been discussed by Schindler & Kasting (2000, and references therein). While this is based on photolysis of e.g., CO2 and H2O and is thus limited in extent, a sustainable production of abiotic O3 which could build up to a detectable level has been suggested by Domagal-Goldman & Meadows (2010) for a planet within the habitable zone of AD Leonis with a specific atmospheric composition. Indeed, other studies confirm that abiotic buildup of ozone is possible (e.g., Hu et al. 2012; Tian et al. 2014); however, detectable levels are unlikely if liquid water is abundant, as e.g. rainout of oxidized species would keep atmospheric O2 and O3 low (Segura et al. 2007), unless the CO2 concentration is high and both H2 and CH4 emissions are low (Hu et al. 2012). False-positive detection of molecules such as CH4 and O3 is discussed by von Paris et al. (2011). Seager et al. (2013) present a biosignature gas classification. Since abiotic processes cannot be ruled out for individual molecules (e.g. for O3), searches for biosignature molecules should search for multiple biosignature species simultaneously. It has been suggested that the simultaneous presence of O2 and CH4 can be used as an indication for life (Sagan et al. 1993, and references therein). Similarly, Selsis et al. (2002) suggest a so-called “triple signature”, where the combined detection of O3, CO2 and H2O would indicate biological activity. Domagal-Goldman & Meadows (2010) suggest to simultaneously search for the signature of O2, CH4, and C2H6. Of course, care has to be taken to avoid combinations of biosignature molecules which can be generated abiotically together (see e.g. Tian et al. 2014). The detectability of biosignature molecules is discussed, e.g. by von Paris et al. (2011) and Hedelt et al. (2013). In particular, the simulation of the instrumental response to simulated spectra for currently planned or proposed exoplanet characterization missions has shown that the amount of information the retrieval process can provide on the atmospheric composition may not be sufficient (von Paris et al. 2013).

2017MNRAS.469S..39F__Youdin_&_Goodman_2005_Instance_1

At the disc age t ≈ 1.5 Myr, the decay of 26Al has sufficiently advanced to not melt forming planetesimals and dust from the inner disc has enriched the outer disc (Ciesla 2011). At rh = 30 au, we get τ ≈ 0.1 Myr and t0 ≈ 0.5 Myr because ρ(1 au)/ρ(30 au) = 170 and 1800 in the MMSN and Nice models, respectively (Zsom et al. 2010). The number of fractals is still larger than that of pebbles of similar sizes: (3) \begin{equation} N_{\rm F} / N_0 \gtrsim \exp [-(t - t_0)/\tau ] \gtrsim N_{\rm P} / N_0 \approx \phi _{\rm F} / \phi _{\rm P} \end{equation} with ϕF ≈ 10−5, the others being restructured into cm-sized pebbles (Weidling et al. 2009; Zsom et al. 2010). Any further pebble growth would require so longer times at so high relative speeds (Davidsson et al. 2016) to destroy or convert into pebbles all the remaining fractals. The input of dust from the inner disc triggers a streaming instability in the outer disc (Youdin & Goodman 2005; Johansen et al. 2007), lasting ≈0.1 Myr and concentrating pebbles and fractals sticking to pebbles into filaments where the comet nucleus collapses by gravity at speeds 1 m s−1 into a randomly packed (Onoda & Liniger 1990; Song, Wang & Makse 2008) aggregate of cm-sized pebbles, with the fractals stored in the voids among the pebbles (Fulle et al. 2016b). The voids fill (40 ± 5) per cent of the nucleus volume (Onoda & Liniger 1990; Song, Wang & Makse 2008). This scenario is confirmed by the pebbles observed in the pristine terrain at the final Philae landing site, which have a narrow differential size distribution (Fig. 2) peaked at diameters of $6^{+4}_{-2}$ mm (Poulet et al. 2016). For polydisperse spheres, the void volume occupies a lower percentage of the total, e.g. 32 per cent for a log-normal size distribution with a standard deviation of 0.3 (Baranau & Tallarek 2014). These facts allow us to model the 67P nucleus as composed of pebbles of the same size, with a void volume filling (37 ± 5) per cent of the total nucleus volume. A random packing of spheres does not result in the densest possible configuration, i.e. small pebbles will not always occupy available void spaces, because access to these voids is not necessarily available during the formation of the cometary nucleus.

2018MNRAS.475.1160H__Zschaechner_et_al._2016_Instance_1

Given their potentially profound impact on both the stellar and gas properties, galaxy mergers are often invoked to explain the presence of metals and low ionization species in the CGM at large impact parameters (e.g. Farina et al. 2013, 2014; Johnson, Chen & Mulchaey 2015b). Indeed, merger-induced tidal torques can give rise to observed stellar tidal features extending to many tens of kpc (e.g. Hernández-Toledo et al. 2006; Patton et al. 2011, 2013; Casteels et al. 2014). Although some gas asymmetries coexist with asymmetries in the stellar profiles of interacting galaxies, tidally induced asymmetries can persist even longer in H i gas (e.g. Lelli, Verheijen & Fraternali 2014a,b; Scott et al. 2014). Such tidal debris can potentially diffuse into the CGM contributing a significant gas and metal mass. In addition to morphological disturbances, galaxy mergers can trigger vigorous galactic outflows associated with feedback from the enhanced star formation (e.g. Martin 2005; Rupke, Veilleux & Sanders 2005a; Strickland & Heckman 2009; Hayward & Hopkins 2017) and AGN activity (e.g. Rupke, Veilleux & Sanders 2005b; Veilleux et al. 2013; Zschaechner et al. 2016; Woo, Son & Bae 2017) which can populate the CGM with metals while giving rise to multiphase absorbers (e.g. Borthakur et al. 2013; Bird et al. 2015; Heckman & Borthakur 2016; Bordoloi et al. 2017; Heckman et al. 2017). Besides, merger-induced shocks and feedback can increase the CGM's internal energy; numerical simulations of galaxy mergers show that CGM gas can be significantly heated (to X-ray emitting temperatures) through shocks and feedback processes (e.g. Cox et al. 2004, 2006b; Sinha & Holley-Bockelmann 2009). Furthermore, studies of galaxies in dense environments (higher merger probability) show different CGM properties when compared to a matched isolated galaxy sample: The CGM of galaxies in groups show distinct kinematics (Pointon et al. 2017) and ionic covering fractions (Johnson, Chen & Mulchaey 2015a; Burchett et al. 2016). Despite the possibly significant influence of galaxy mergers on the CGM, the details of the interplay between galaxy mergers and the CGM remain relatively unexplored and we are currently lacking clear and quantitative predictions of how the CGM will be affected during the merger process. Current observations of the impact of galaxy mergers on the CGM are limited to a few case studies (e.g. Keeney et al. 2011; Johnson et al. 2014). A systematic survey targeting the CGM of kinematic galaxy pairs (i.e. COS-Pairs: Bordoloi et al. in preparation) is needed to place observational constraints on the effect of mergers on the CGM.

2016AandA...585A..48G__Mostardi_et_al._(2013)_Instance_1

On the same SSA22 field, Nestor et al. (2011) and Nestor et al. (2013) show nine LBGs and 20 Lyman-α emitters (LAEs) with LyC detection out of a sample of 41 LBGs and 91 LAEs (all spectroscopically confirmed). They started from a different narrowband image centred at ~3640 Å, which is deeper than that used by Iwata et al. (2009), at ~3590 Å. A careful analysis of their LBG detections, however, shows that even in this case the LyC emission for many z ~ 3 sources is offset by 0.4–1.0 arcsec. The observed ratio between the 900 and 1500 Å rest-frame emission is difficult to reconcile with that expected by standard stellar populations (Vanzella et al. 2012a). The HST images in I814W filter available for a few of them show the presence of clearly separated galaxies, sometimes fainter at 1500 Å than in the ionizing continuum (i.e. their C16 object), with a resulting escape fraction well exceeding 1000% if estimated in the LyC position. For the majority of them, no HST imaging is available but even ground-based images often show the presence of slightly offset emission in LyC, w.r.t. the non-ionizing continuum. Similar conclusions can be reached for the Mostardi et al. (2013) sample where they adopt the same analysis as in Nestor et al. (2013). At z ~ 2.8, they found four LyC emitters out of 49 LBG galaxies and seven LyC emitters out of 91 LAEs. In this case the lack of high spatial resolution data from HST for the majority of the sample prevents any detailed analysis about possible contamination by interlopers/foregrounds. These conclusions have been strengthened by recent observations by Siana et al. (2015), who found no convincing detection in their deep HST imaging with WFC3-UVIS of five LyC emitters extracted from the sample of Nestor et al. (2011), or by Mostardi et al. (2015), who only found one robust LyC emitter after a reanalysis of a sample of 16 galaxies by Mostardi et al. (2013). More interestingly, the only robust candidate LyC emitter by Mostardi et al. (2015), the galaxy MD5b, has an observed ratio FUV/FLyC = 4.0 ± 2.0, equivalent to a relative escape fraction of 75%, when assuming complete transmission of the IGM. Instead, if a mean value of ⟨ exp(−τIGM) ⟩ = 0.4 at z ~ 3.1 is adopted, following Inoue et al. (2014), the relative escape fraction of MD5b turns out to be 188%. Imposing the constraint of a physical value for the relative escape fraction of \hbox{$f^{\rm rel}_{\rm esc}<100\%$}fescrel100%, the LoS of MD5b must be very transparent, exp(−τIGM) > 0.75, which corresponds to a probability 10-4, according to Inoue et al. (2014). This could be an indication that this galaxy is also a low-z contaminant, similar to the other cases studied by Mostardi et al. (2015).

2022MNRAS.513..232N__Haywood_et_al._2013_Instance_1

There are a plethora of data available in the form of spectra, astrometric, and photometric information, as well as multiwavelength maps with the advent of large-scale spectroscopic (Apache Point Observatory Galactic Evolution Experiment/APOGEE: Eisenstein et al. 2011, RAdial Velocity Experiment/RAVE: Steinmetz et al. 2006, Gaia-ESO: Gilmore et al. 2012, Large Sky Area Multi-Object Fiber Spectroscopic Telescope/LAMOST: Cui et al. 2012, Galactic Archaeology with HERMES/GALAH: De Silva et al. 2015, Abundances and Radial velocity Galactic Origins Survey/ARGOS: Ness et al. 2012), astrometric (Hipparcos: Perryman et al. 1997, Gaia: Gaia Collaboration 2016), and photometric surveys (Two-Micron All Sky Survey/2MASS: Skrutskie et al. 2006, Sloan Digital Sky Survey/SDSS: Stoughton et al. 2002, Vista Variables in the Vía Láctea/VVV: Minniti et al. 2010, the SkyMapper Southern Survey : Wolf et al. 2018). These surveys have enabled the chemo-dynamic characterization of stellar populations in the Milky way that constitute different Milky Way components like thin disc, thick disc, halo, bulge, etc. For example, star count observations in the solar neighbourhood (Yoshii 1982; Gilmore & Reid 1983) led to the discovery of the thick disc, followed by its characterization as the old α-enhanced population in the double sequence exhibited by the solar neighbourhood stars in the [α/Fe] versus [Fe/H] plane (Fuhrmann 1998; Bensby, Feltzing & Lundström 2003; Reddy, Lambert & Allende Prieto 2006; Adibekyan et al. 2012; Haywood et al. 2013). At present, data from large-scale spectroscopic surveys (Anders et al. 2014; Hayden et al. 2015; Weinberg et al. 2019) have led to the discovery of this trend at different galactocentric radius, R, and average height, |Z|, across the Galaxy shedding light on the disc formation and evolution scenarios. In addition, many age determination methods have been developed that uses these survey data to provide valuable information about the star formation histories and age metallicity relation of disc stellar populations (Casagrande et al. 2011; Bedell et al. 2018; Lin et al. 2020; Nissen et al. 2020). Secular processes such as radial migration (Sellwood & Binney 2002; Schönrich & Binney 2009; Minchev & Famaey 2010), which leads to the mixing of stars across the Galaxy, are also being explored using a combination of accurate phase space information from Gaia (Gaia Collaboration 2018) and chemistry and age information of stars from large-scale spectroscopic surveys (Buder et al. 2019). The discovery of streams and dynamically different stellar populations in the Milky Way halo, considered to be the result of past accretion/merger events (Belokurov et al. 2018; Helmi et al. 2018; Ibata, Malhan & Martin 2019; Myeong et al. 2019) using the Gaia data and their further exploration with chemistry from large-scale spectroscopic surveys (Buder et al. in preparation) is another example. Multiple components in the Bulge metallicity distribution function discovered by multiple individual and large-scale spectroscopic observations, are being studied in detail to understand the origin of the Bulge and its connection with the Milky Way bar and Galaxy evolution (Ness et al. 2013; Rojas-Arriagada et al. 2017, 2020). There are many upcoming surveys [4-metre Multi-Object Spectroscopic Telescope/4MOST: de Jong et al. (2019), Sloan Digital Sky Survey/SDSS-V: Kollmeier et al. (2017), WEAVE: Dalton et al. (2018)] that will further improve our understanding of the formation and evolution of the Milky Way and its components.

2017ApJ...840...98J__Nataf_et_al._2010_Instance_1

Figure 1 shows our synthetic color–magnitude diagrams (CMDs) for the two RCs at four different metallicity regimes, which are almost identical to those presented in Figure 1 of Paper I, but here ( ) CMDs are added in the right panels. We recall that in our models, the fRC and bRC are produced by G1 and G2, respectively, where G1 follows the standard helium-enrichment parameter (i.e., ), while G2 is substantially enhanced in helium abundance (Y = 0.406). Following Paper I, we have also adopted a 0.2 dex difference in metallicity between G2 and G1, and assumed 12 and 10 Gyrs for the ages of G1 and G2, respectively, with the same population ratio for the two RCs (Nataf et al. 2010; Paper I and references therein). It is evident from these models that in the metal-rich population like the bulge, highly helium-enhanced stars are not placed on the very blue HB as in the metal-poor GCs, but are instead placed on the bRC. We refer the reader to Section 2 of Paper I for a detailed description of the RC features of our models at four different metallicity regimes in Figure 1. Here, we note from the ( ) CMDs that the variation of the overall RC features on metallicity is similar to that in ( ) CMDs. Panel (f) in ( ) CMD can also naturally reproduce the observed double RC feature, i.e., ∼0.5 mag difference with almost negligible color difference between the two RCs, as is the case in panel (b) in ( ) CMD. Hereafter, we refer to these models in panels (b) and (f) constructed at as “ ,” to distinguish them from the models with different input parameters presented below. The CMDs for these reference models, down to the MS luminosity level, are further presented in Figure 2. This figure confirms that, in the position and width of the lower red-giant-branch (RGB) and main-sequence turn-off (MSTO), our models are not inconsistent with the observed CMDs of bulge population by Clarkson et al. (2008, 2011) and Brown et al. (2010).

2015ApJ...799...42D__Skemer_et_al._2014_Instance_1

The Large Binocular Telescope (LBT) consists of two 8.4 m aperture optical telescopes on a single ALT-AZ mount installed on Mount Graham in southeastern Arizona (at an elevation of 3192 m) and operated by an international collaboration among institutions in the United States, Italy, and Germany (Hill et al. 2014; Veillet et al. 2014). Both telescopes are equipped with deformable secondary mirrors which are driven with the LBT's adaptive optics (AO) system to correct atmospheric turbulence at 1 kHz (Esposito et al. 2010; Bailey et al. 2014). Each deformable mirror uses 672 actuators that routinely correct 400 modes and provide Strehl ratios exceeding 80%, 95%, and 99% at 1.6 μm, 3.8 μm, and 10 μm, respectively (Esposito et al. 2012; Skemer et al. 2014). The LBTI is an interferometric instrument designed to coherently combine the beams from the two 8.4 m primary mirrors of the LBT for high-angular resolution imaging at infrared wavelengths (1.5–13 μm; Hinz et al. 2012). It is developed and operated by the University of Arizona and based on the heritage of the Bracewell Infrared Nulling Cryostat on the MMT Hinz et al. (2000). The overall LBTI system architecture and performance will be presented in full detail in a forthcoming publication (P. M. Hinz et al., in preparation). In brief, the LBTI consists of a universal beam combiner (UBC) located at the bent center Gregorian focal station and a cryogenic Nulling Infrared Camera (NIC). The UBC provides a combined focal plane for the two LBT apertures while the precise overlapping of the beams is done in the NIC cryostat. Nulling interferometry, a technique proposed 36 yr ago to image extra-solar planets (Bracewell 1978), is used to suppress the stellar light and improve the dynamic range of the observations. The basic principle is to combine the beams in phase opposition in order to strongly reduce the on-axis stellar light while transmitting the flux of off-axis sources located at angular spacings which are odd multiples of 0.5λ/B (where B = 14.4 m is the distance between the telescope centers and λ is the wavelength of observation). Beam combination is done in the pupil plane on a 50/50 beamsplitter which can be translated to equalize the pathlengths between the two sides of the interferometer. One output of the interferometer is reflected on a short-pass dichroic and focused on the Nulling Optimized Mid-Infrared Camera (NOMIC) (Hoffmann et al. 2014). NOMIC uses a 1024 × 1024 Raytheon Aquarius detector split into two columns of eight contiguous channels. The optics provides a field of view (FOV) of 12 arcsec with a plate-scale of 0.018 arcsec. Tip/tilt and phase variations between the LBT apertures are measured using a fast-readout (1 Hz) K-band PICNIC detector (PHASECam) which receives the near-infrared light from both outputs of the interferometer. Closed-loop correction uses a fast pathlength corrector installed in the UBC (see more details in Defrère et al. 2014).

2021AandA...652A..30S__Santini_et_al._2019_Instance_1

This paper is the fourth in a series. In our first work (Merlin et al. 2018) we presented an accurate and conservative technique to single out passive galaxies at high redshift by means of spectral energy distribution (SED) fitting with a probabilistic approach. We selected 30 z > 3 candidates in the GOODS-S field. Passive galaxy candidates, while being relatively easy to select from photometric surveys once the technique is established, need to be confirmed by other means. This is usually achieved through spectroscopic observations. However, spectroscopy becomes particularly difficult and time consuming at z > 3, where only few candidates have been confirmed so far (Glazebrook et al. 2017; Schreiber et al. 2018a; Tanaka et al. 2019; Valentino et al. 2020; Forrest et al. 2020a,b; Saracco et al. 2020; D’Eugenio et al. 2020). In our second work (Santini et al. 2019, S19 hereafter), we used a complementary approach and looked for evidence of the lack of star formation as seen in the sub-millimetre regime to confirm the passive classification of the high-z candidates selected in GOODS-S. At that time, we could confirm 35% of the targets on an individual basis adopting conservative assumptions, and we validated the sample as a whole in a statistical sense. In our third work (Merlin et al. 2019, M19 hereafter), we extended the search for passive galaxies to the entire CANDELS sample and selected 102 z > 3 candidates over the five fields. In the present work, we first confirm the passive nature of these candidates, adopting the method presented in S19 and taking advantage of the richer ALMA archive, which includes observations that were still proprietary at the time of our previous work. We then analyse the emergence and mass growth of this peculiar class of galaxies by means of two powerful statistical tools: the stellar mass function (SMF) and the stellar mass density (SMD). Very few studies so far have pushed the analysis of the SMF of passive galaxies at z > 3 (Muzzin et al. 2013; Davidzon et al. 2017; Ichikawa & Matsuoka 2017; Girelli et al. 2019) because of the difficulty in assembling statistically meaningful samples of candidates at such high redshift.

2015AandA...584A.103S__Chamel_et_al._2007_Instance_1

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

2022AandA...660A..56A__Myers_2009_Instance_1

Recent theoretical models clearly demonstrate the role of expanding H I shells for the formation of molecular clouds (e.g., Hennebelle et al. 2008; Heitsch et al. 2009; Inoue & Inutsuka 2009). In particular, these models highlight the importance of multiple compressions for the formation of magnetized filamentary molecular clouds (e.g., Inutsuka et al. 2015; Iwasaki et al. 2019). The typical timescale of such compressions is estimated in the Galactic disk to be on average ~1 Myr (McKee & Ostriker 1977; Inutsuka et al. 2015). Thus, the formation of molecular clouds may last from a few million years to ~10 Myr or more (see, e.g., Kobayashi et al. 2017). These successive compressions may continuously alter the density, velocity, and magnetic field structures of the clouds, and also have a strong impact on the formation of new generations of filaments and consequently that of stars. While the first generation of stars form and impact their (local) surroundings (through outflow, jets, winds, and ionizing radiation), new cold matter is continuously assembled replenishing the sites of star formation (i.e., filaments and hubs). This matter replenishment may be channeled from within the cloud itself through molecular filaments toward dense ridge-like main filaments (Schneider et al. 2010; Palmeirim et al. 2013) or toward hubs (Myers 2009; Peretto et al. 2013, 2014; Treviño-Morales et al. 2019). Matter can also be brought into the system (the cloud) by a new event of external collision (e.g., Fukui et al. 2018b). Arzoumanian et al. (2018) identified, in position-velocity (PV) diagrams, extended structures with mean line-of-sight (LOS) velocities offset with respect to, and connected to the velocity of, a filament. They suggested a multi-interaction scenario where sheet-like extended structures interact, in space and time, with a star-forming filament, and are responsible for its compression or disruption. Arzoumanian et al. (2018) also identified a bent velocity structure in the PV space. They showed that such a V- or Λ-shaped velocity structure can result from the filament formation process by shock compression as proposed by the theoretical model of Inoue et al. (2018). In this model, a filament is formed due to convergence of a flow of matter generated by the bending of the ambient magnetic field structure induced by an interstellar shock compression (see also Inoue & Fukui 2013; Vaidya et al. 2013). This velocity structure has also been observed toward another filament, the Musca filament (Bonne et al. 2020). More recently, in a theoretical study, Abe et al. (2021) proposed a classification of filament formation mechanisms resulting from the variation in the relative importance between the shock velocity, the turbulence, and the magnetic field strength (see also the theoretical study by Chen et al. 2020a).

2022MNRAS.510.5676I__Faisst_et_al._2017_Instance_1

The above finds an explanation within the context of the two-phases formation scenario of ETGs (Oser et al. 2010). According to this scenario, in the early (z ≳ 2) dissipative phase driven by wet major mergers (e.g. Hopkins et al. 2008) and/or violent disc instability (e.g. Barro et al. 2013; Dekel & Burkert 2014; Zolotov et al. 2015; Tacchella et al. 2016), the proto-ellipticals are compact objects that form stars in situ in a very intense regime until this process is stopped abruptly by gas shock heated in massive haloes, and/or rapid gas exhaustion and Supernova/AGN feedback. As the compact quiescent galaxy ages and reddens (red nugget), during the second non-dissipative phase (z ≲ 1−2), more mass is assembled by the accretion of ex situ stars through merging with other galaxies, which are also mostly quiescent. Dry minor or intermediate mergers contribute little to the mass growth but promote substantial size growth (the accreted stars tend to be deposited mostly in the external regions), making spheroidal galaxies less compact and likely also less concentrated (Bezanson et al. 2009; Naab, Johansson & Ostriker 2009; Oser et al. 2010; Trujillo, Ferreras & de La Rosa 2011; Bluck et al. 2012; Johansson, Naab & Ostriker 2012; Hilz, Naab & Ostriker 2013; Shankar et al. 2013; van Dokkum et al. 2015; Wellons et al. 2016; Faisst et al. 2017; Furlong et al. 2017; Hill et al. 2017; Genel et al. 2018, among others). Dry mergers can also be major, though they are less frequent and relevant only to the most massive galaxies (e.g. Bundy et al. 2009; López-Sanjuan et al. 2012; Rodríguez-Puebla et al. 2017). These mergers increase substantially the mass, while the size increases approximately proportional to the mass increase (e.g. Nipoti, Londrillo & Ciotti 2003; Johansson et al. 2012; Nipoti et al. 2012; Hilz et al. 2013) in a such way that the shift in the mass–size relation is small, affecting in a lesser degree the compactness and concentration of the merged galaxies. Although the second phase of ETG formation, driven by dry mergers, is a reliable explanation for the growth in size and the puffing-up of massive ETGs, there is still an intense debate as to whether or not this mechanism is enough to describe observational inferences (e.g. López-Sanjuan et al. 2012; Newman et al. 2012; Nipoti et al. 2012; Sonnenfeld, Nipoti & Treu 2014; Man, Zirm & Toft 2016; Frigo & Balcells 2017, see for a discussion Zanisi et al. 2021 and more references therein). An alternative or complementary mechanism suggested for the apparent strong growth in size of massive ETGs is quasar feedback, which removes huge amounts of cold gas from the central regions, inducing an expansion of the stellar distribution (Fan et al. 2008, but see Trujillo et al. 2011).

2021MNRAS.500.4042S__Snellen_et_al._2010_Instance_1

To highlight the evidence for CO2 in the transmission spectrum, in Fig. 18 we show three atmo model atmospheres: our best-fitting model from the free-chemistry retrieval; a model with all of the same parameters as the best fit, except the CO2 abundance, which is set to zero; and a third model with both CO and CO2 abundances set to zero. The strong absorption feature centred on the 4.5 $\, \mu$m Spitzer channel disappears in the latter two models. With only Spitzer photometry, however, the contribution of CO to the 4.5 $\, \mu$m point complicates the interpretation of the C/O ratio. Theoretical models have found that CO should be the dominant carbon-bearing molecule for hydrogen-dominated atmospheres above 1000 K (e.g. Lodders & Fegley 2002; Heng & Lyons 2016), and CO has been detected at high resolution in hot Jupiter atmospheres (e.g. Snellen et al. 2010). In our equilibrium chemistry retrieval, CO is at least 100× more abundant than CO2 (see Fig. 12). However, at 4.5 $\, \mu$m the CO2 opacity is much stronger and dominates over the CO contribution even though CO has much higher VMR concentrations. In the free-chemistry retrieval, the CO VMR is not constrained by the data – only an upper limit to the CO is found, as very high values affect the mean molecular weight, and the data are consistent with no CO contribution. The lack of a CO feature in the WFC3 data further pushes the free-chemistry retrieval to prefer CO2 over CO. With CO constrained through chemistry in one retrieval and unconstrained in the free case, the C/O ratios obtained are vastly different. In the chemical-equilibrium case, a supersolar C/O is found (see Fig. 13) while in the free case a subsolar C/O ratio is found (Fig. 14). This finding highlights the extreme sensitivity and degeneracies of measuring the C/O ratio with a free-chemistry retrieval model, as all major molecular species have to be well constrained by the data. For a hot Jupiter such as WASP-127b, we consider a scenario with all of the carbon found in CO2 and little to none in CO to be thermochemically implausible, as no obvious non-equilibrium mechanism would deplete CO by many orders of magnitude while enhancing CO2. This situation is unlike CH4, where dynamical mixing and vertical quenching can dramatically enhance CH4 (e.g. Cooper & Showman 2006; Moses et al. 2011; Tsai et al. 2017; Drummond et al. 2018a, b). With only one photometric data point at 4.5 $\, \mu$m, it is currently impossible to fully disentangle the contribution of both CO and CO2 in a model-independent way. Further transmission spectroscopy observations of WASP-127b at high resolution with the JWST will clarify which is the dominant carbon-bearing molecule in WASP-127b’s atmosphere, and allow stronger constraints to be placed on its carbon-to-oxygen ratio.

2021ApJ...922..268J__Bonoli_et_al._2021_Instance_1

In the last decade, several projects have focused on collecting the deepest images of clusters and groups of galaxies. Among the most recent examples, we can mention the Frontier Fields initiative (Lotz et al. 2017), the Hyper Suprime-Cam Subaru Strategic Program (Aihara et al. 2018a, 2018b, 2019), the Fornax Deep Survey (Iodice et al. 2016), or the upcoming Vera C. Rubin Observatory (Ivezic et al. 2008). In addition, technical development offers a new and barely explored perspective: a panchromatic, more detailed view of the universe through the use of narrowband filter systems, such as those from the J-PLUS and J-PAS surveys (Benitez et al. 2014; Cenarro et al. 2019; Bonoli et al. 2021). In this work we intend to show how the combination of these two characteristics, high-quality data and information at different wavelengths, can provide strong clues on the clusters’ formation and evolution through the analysis of the ICL. It is known that the ICL presents different morphologies, substructures, and ICL fractions depending on the wavelength used (Tang et al. 2018). Whereas the ICL fraction measured in single bands can be useful for establishing qualitative relations (e. g., evolution with redshift or halo mass), a panchromatic view of the ICL can give insights into its stellar populations or the dynamical stage of the cluster (e.g., Jiménez-Teja et al. 2019). With this aim, we use the recently acquired data from the Reionization Lensing Cluster Survey (RELICS; Coe et al. 2019), a Hubble Space Telescope (HST) Treasure program that has observed 41 very massive, strong-lensing clusters in three optical (ACS) and four infrared (WFC3) bands. In this work we present the potential of the RELICS data for ICL purposes, focused on the cluster WHL J013719.8–08284 (WHL0137 hereafter). Future work will apply the CICLE algorithm, which is free of a priori assumptions, extensively to the whole RELICS sample with two main aims: (1) analyze the optical and infrared ICL fractions to extract information about the stellar populations of the diffuse light, and (2) infer the dynamical stage of the clusters.

2020ApJ...891...28T__Antoja_et_al._2018_Instance_1

With the release of Gaia DR2 (Gaia Collaboration et al. 2018; Katz et al. 2019), the proper motions of billions of stars are now available to the astronomical community. Combining with radial velocities from large spectroscopic surveys, like the Sloan Digital Sky Survey (Eisenstein et al. 2011; Blanton et al. 2017), Gaia-ESO survey (Gilmore et al. 2012; Randich et al. 2013), and LAMOST Galactic spectroscopic survey (Deng et al. 2012; Zhao et al. 2012), the wealth of 6D information of billions of stars have challenged and even overthrown many aspects of our understanding of the Milky Way (MW). The discovery of snail shells in the phase-space distribution of MW disk stars (Antoja et al. 2018) has inspired debates about their origin: whether they are generated by the passage of a dwarf galaxy (probably the Sagittarius dwarf galaxy) through the MW disk (Binney & Schönrich 2018) or are the echo of the MW bar buckling (Khoperskov et al. 2019). Meanwhile, major accretion events begin to unveil themselves when the stellar distribution in various energy–momentum spaces (e.g., Myeong et al. 2019) is investigated. These major accretion events injected most of the materials from the progenitor dwarf galaxies into our MW, including globular clusters (GCs). As GCs are one of the oldest objects in our Galaxy, identifying and studying accreted GCs help us trace back the accretion history of our Galaxy. Though details of major accretion events, e.g., the number of accretion events and the GCs associated with each event, are still under debate (e.g., Helmi et al. 2018; Massari et al. 2019; Myeong et al. 2019), it is widely accepted that a substantial number of halo stars and GCs were accreted (e.g., Ostdiek et al. 2019). Along the same line, more and more substructures, e.g., stellar streams, are identified inside the MW (e.g., Malhan et al. 2018; Ibata et al. 2019b). An increasing number of stellar streams are suggested to be related to the debris of (inner-halo) GCs (e.g., Ibata et al. 2019a). Aside from these GC destruction events under the influence of Galactic potential, the dynamical relaxation of GCs (e.g., Weinberg 1994; Vesperini & Heggie 1997) also ejects member stars into the field. It would be of great interest to find such GC-ejected stars in order to estimate the mass loss from GCs to better understand the formation and evolution of our MW. To help achieve this goal, another characteristic of GCs is very helpful.

2016MNRAS.461.3982B__Scheeres_2002_Instance_1

Many studies have been done to understand the dynamics and origin of such systems since the discovery of the first binary asteroid system, Dactyl orbiting around (243) Ida in 1993 (Chapman et al. 1995). Based on the structure of ‘rubble pile’ asteroids (a collection of gravitationally bound boulders with a distribution of size scales and very little tensile strength between them), a model for how they can disrupt due to close flybys of a planet was developed. However, close encounters with the planets proved not to be enough for creation of the current population of binary systems (Margot et al. 2002; Walsh & Richardson 2008). Another model for their formation is by increasing their spin rates due to incident and remitted solar photons, known as the Yarkovsky–O'Keefe–Radzievskii–Paddack (YORP) effect. The YORP effect on contact binary asteroids has been studied (Bottke et al. 2002; Merline et al. 2002; Scheeres 2002; Walsh & Richardson 2006). Using a model with an ellipsoid and a sphere in a planar case, Scheeres (2007) studied fission limits (spin limit to occur a fission) and the stability of that kind of system for different initial conditions. After that, the stability of a binary system was analysed using a two-ellipsoid model (Scheeres 2009). Pravec et al. (2010) made a complete study about formation of asteroid pairs through rotation fission. Jacobson & Scheeres (2011) studied the creation of binaries and other observed near-Earth asteroid (NEA) systems, including doubly synchronous binaries, high-e binaries, ternary systems and contact binaries. That study analysed the dynamics of a binary system just after rotational fission. Using a two-ellipsoid model taking into account mutual gravitational interactions and tidal dissipation, they analysed the dynamics for different mass ratios of the system under a planar assumption. The current work follows from these results, but looks at more likely, non-planar initial configurations. This extension is significant, as non-planar cases must take into account the complete rotational motion (rotation, precession and nutation) of each body. Our results are compared with the results obtained by Jacobson & Scheeres (2011).

2021MNRAS.505.2111L__Melia_et_al._2017_Instance_1

Recently, quasars observed with multiple measurements, another potential cosmological probe with a higher redshift range that reaches to z ∼ 5, is becoming popular to constrain cosmological models in the largely unexplored portion of redshift range from z ∼ 2 to z ∼ 5. A sample that contains 120 angular size measurements in intermediate-luminosity quasars from the very long baseline interferometry (VLBI) observations (Cao et al. 2017a,b) has become an effective standard ruler, which have been extensively applied to test cosmological models (Li et al. 2017; Melia et al. 2017; Qi et al. 2017; Zheng et al. 2017; Xu et al. 2018; Ryan, Chen & Ratra 2019), measuring the speed of light (Cao et al. 2017a, 2020a), exploring cosmic curvature at different redshifts (Cao et al. 2019; Qi et al. 2019), and the validity of cosmic distance duality relation (Zheng et al. 2020). Then, Risaliti & Lusso (2019) put forward a new compilation of quasars containing 1598 quasi-stellar object (QSO) X-ray and ultraviolet (UV) flux measurements in the redshift range of 0.036 ≤ z ≤ 5.1003, which have been used to constrain cosmological models (Khadka & Ratra 2020b) and cosmic curvature at high redshifts (Liu et al. 2020a,c), as well as test the cosmic opacity (Geng et al. 2020; Liu et al. 2020b). Making use of this data to explore cosmological researches mainly depends on the empirical relationship between the X-ray and UV luminosity of these high-redshift quasars proposed by Avni & Tananbaum (1986), which leads to the Hubble diagram constructed by quasars (Risaliti & Lusso 2015, 2017; Lusso & Risaliti 2016; Bisogni, Risaliti & Lusso 2017). In general, the advantage of these two QSO measurements over other traditional cosmological probes is that QSO has a larger redshift range, which may be rewarding in exploring the behaviour of the non-standard cosmological models at high redshifts, providing an important supplement to other astrophysical observations and also demonstrating the ability of QSO as an additional cosmological probe (Zheng et al. 2021).

2017AandA...599A...8V__Carbone_&_Veltri_1990_Instance_1

At scales comparable with dp, a variety of observations in the solar wind have suggested that fluctuations may consist primarily of kinetic Alfvén waves (KAWs; Bale et al. 2005; Sahraoui et al. 2009). In the linear fluctuation terminology, KAWs are waves belonging to the Alfvén branch, at wavevectors k almost perpendicular to the ambient magnetic field B0, with \hbox{$k \sim d_{\rm p}^{-1}$}k~dp-1. A detailed discussion of the properties of KAWs can be found in Hollweg (1999; see also references therein for a more complete view of the subject), for example. In the last decades, KAWs have received considerable attention due to their possible role in a normal mode description of turbulence. Indeed, theoretical studies (e.g., Shebalin et al. 1983; Carbone & Veltri 1990; Oughton et al. 1994) have shown that the turbulent cascade in magnetized plasma tends to develop mainly in directions perpendicular to B0. Anisotropic spectra have been commonly observed in space plasmas, showing the presence of a significant population of quasi-perpendicular wavevectors (Matthaeus et al. 1986, 1990). The above considerations suggest that fluctuations with characteristics similar to KAWs are naturally generated by a turbulent cascade at scales close to dp. Many solar wind observational studies (Bale et al. 2005; Sahraoui et al. 2009; Podesta & TenBarge 2012; Salem et al. 2012; Chen et al. 2013; Kiyani et al. 2013), theoretical works (Howes et al. 2008a; Schekochihin et al. 2009; Sahraoui et al. 2012) as well as numerical simulations (Gary & Nishimura 2004; Howes et al. 2008b; TenBarge & Howes 2012) have suggested that fluctuations near the end of the magnetohydrodynamics inertial cascade range may consist primarily of KAWs, and that such fluctuations can play an important role in the dissipation of turbulent energy. Due to a non-vanishing parallel component of the electric field associated with KAWs, these waves have also been considered in the problem of particle acceleration (Voitenko & Goossens 2004; Décamp & Malara 2006). Particle acceleration in Alfvén waves in a dispersive regime has been studied both in 2D (Tsiklauri et al. 2005; Tsiklauri 2011) and 3D (Tsiklauri 2012) configurations. Recently, Vasconez et al. (2014) have studied collisionless Landau damping and wave-particle resonant interactions in KAWs.

2021MNRAS.500.2209Z__Miyama,_Narita_&_Hayashi_1987_Instance_1

There is growing observational and numerical evidence that star-forming regions may be in a state of global gravitational contraction; see Vázquez-Semadeni et al. (2019). The supersonic collisions of flows of warm diffuse atomic gas simulated with both self-gravity and cooling exhibit the hierarchical collapse of the turbulent medium, as was shown by Vázquez-Semadeni et al. (2007) and Naranjo-Romero, Vázquez-Semadeni & Loughnane (2015) for example. This implies that gravitational instability (GI hereafter) manifests itself in a wide range of sufficiently large scales during the evolution of molecular clouds. Theoretical work has revealed that flattened dense structures form as a result of large collisions of diffuse matter. Later on, they give birth to filaments, which then fragment into multiple cores. This scenario is provided by the dynamical instability of self-gravitating layers, cylinders, and spheres, respectively. The linear stability analysis of these idealized configurations (e.g. Ledoux 1951; Chandrasekhar & Fermi 1953; Bonnor 1956; Elmegreen & Elmegreen 1978; Nagasawa 1987; Fiege & Pudritz 2000) as well as the corresponding non-linear solutions (e.g. Larson 1969; Penston 1969; Miyama, Narita & Hayashi 1987; Inutsuka & Miyama 1997; Masunaga & Inutsuka 2000, and many others) confirms this view. At the same time, as was noted by Larson (1985), the specific geometry of self-gravitating objects is not crucial for the instability condition, which does not differ much from the basic one derived for the unbounded uniform medium. In the latter case, the study of GI goes back to Jeans (1902), who established that plane-wave perturbations on such a background having finite pressure are heavy sound waves propagating at the subsonic velocity, which vanishes as the wavelength approaches the value now referred to as the Jeans length. Perturbations with scale larger than the Jeans length are the growing and damping static waves. Thus, the critical scales for GI of realistic configurations mentioned above are always similar to the Jeans scale, which includes typical speed of sound and density chosen appropriately for the corresponding configuration. However, the most unstable scale for realistic configurations has a finite value in contrast to the Jeans result, when the largest growth rate (corresponding to the inverse free-fall time) manifests at the infinitely large scale. The largest growth rates for GI of realistic configurations are commonly the fractions of the inverse free-fall time.

2019ApJ...872...97L__Skokos_et_al._2002_Instance_1

We find a similar dichotomy of bars in previous studies (Elmegreen & Elmegreen 1985, 1989; Baumgart & Peterson 1986; Elmegreen et al. 1996; Regan & Elmegreen 1997; Kim et al. 2015). They investigated two types of bars: flat and exponential profiles in surface brightness. Flat bars have nearly constant light distributions along the bar, whereas those of exponential bars decrease exponentially. They also differ in their structures and the intensity contrast between the bar and the disk: flat bars are longer, wider, and stronger than exponential bars and have a higher contrast than exponential bars. Besides, Athanassoula (1992) explained that flat bars could have roughly rectangular orbits around the end of the bar through the stellar orbit calculation. These properties can also be explained by the locations of resonances (Lynden-Bell 1979; Contopoulos & Papayannopoulos 1980; Sellwood 1981; Contopoulos et al. 1989; Athanassoula 1992; Skokos et al. 2002). A flat density profile develops in crowding stellar orbits between the inner 4:1 resonance and corotation radius (Combes & Elmegreen 1993; Elmegreen & Elmegreen 1985; Elmegreen et al. 1996). Exponential bars end near the inner Lindblad resonance and do not have such crowding orbits (Lynden-Bell 1979; Elmegreen & Elmegreen 1985; Elmegreen et al. 1996). Therefore, flat bars and exponential bars may be expected to have different pattern speeds based on the value of , where Rcr and Rbar are the radius of the corotation resonance and the bar, respectively (Debattista & Sellwood 2000; Valenzuela & Klypin 2003). Although some studies reported the observational lack of slow bars (Debattista & Sellwood 2000; Aguerri et al. 2015), others showed that the pattern speed of bars roughly depends on the Hubble type: fast bars in early-type spirals and slow bars in late-type spirals (Aguerri et al. 1998; Rautiainen et al. 2008). More observational data will help us understand the relation between the density profile and the pattern speed of bars.

2019AandA...624A..92K___2017_Instance_1

In Sects. 4.3 and 4.4 we showed that our image simulations match the actual data very well and that the re-weighting factors are close to unity for subsets of galaxies split by size and S/N, with fluctuations of the order of 10−4 or less. Consequently, the estimates of the shear biases presented in Table 2 should be accurate. Nonetheless it is worthwhile to explore the robustness of our results and attempt to quantify the potential systematic uncertainties that may still be present. For instance, we simplified the galaxy morphologies by representing them with Sérsic profiles. Our input catalogue is incomplete at the faintest magnitudes, but the missing galaxies may still affect the estimate of the multiplicative bias. Varying star densities can affect the results (e.g. Hoekstra et al. 2015, 2017). In this section we therefore explore the sensitivity of the shear measurement bias to various assumptions and simplifications made in the simulations. These results help to assess the robustness of the calibration presented in the previous section. In the language described in Sect. 2.2, these tests correspond to various evaluations of the Δ terms from different simulations. As these tests can become computationally expensive very quickly, we use only 5 of the 13 PSF sets. This results in only minor changes in the mean residual bias values ( 0.005; see lower right panel of Fig. 13), and hence they remain a good representation of the data. The smaller volume of simulations naturally results in a larger statistical uncertainty, but we note that we are interested in determining the change in the mean values of the multiplicative biases of the different tomographic samples when we vary the inputs. The errors are tightly correlated among the different simulations within a tomographic bin, as there is a significant overlap among the input samples of galaxies in the different simulations and because the noise realisations in the images are identical in many of these simulations, unless explicitly mentioned otherwise. Hence, the shifts are not driven by noise. Our main objective is to ensure that uncertainties in the input quantities do not change the bias by more than 0.02. The results in this section indicate that they are indeed controlled under 0.02, but it appears that we cannot impose a tighter limit on the overall uncertainty at this time.

2022ApJ...930...32C__Pilbratt_et_al._2010_Instance_1

Most DSFGs are discovered in wide-area observations using single-dish far-IR (FIR)/millimeter instruments such as the Submillimeter Common-User Bolometer Array (SCUBA/SCUBA-2) on the James Clerk Maxwell Telescope (JCMT; e.g., Smail et al. 1997; Hughes et al. 1998; Chapman et al. 2005; Koprowski et al. 2017; Simpson et al. 2019), the Herschel Space Observatory 10 10 Herschel is an ESA space observatory with science instruments provided by European-led Principal Investigator consortia and with important participation from NASA. (Pilbratt et al. 2010; Eales et al. 2010; Oliver et al. 2012), and the AzTEC instrument (Scott et al. 2008; Aretxaga et al. 2011). Multi-wavelength follow-up of hundreds of survey-identified DSFGs reveals that most sit between 1 z 3 (Casey et al. 2012a; Magnelli et al. 2011, 2013; Gruppioni et al. 2013; Le Floc’h et al. 2005). While there exists a handful of individually studied DSFGs at redshifts as high as z ∼ 5–7 (e.g., Cooray et al. 2014; Strandet et al. 2017; Marrone et al. 2018; Zavala et al. 2018b; Casey et al. 2019; Reuter et al. 2020), such high-z systems have proven difficult to both identify and spectroscopically confirm. This is because (i) DSFGs at z > 4 are outnumbered by the dominant DSFG population at z ≈ 1–3 and (ii) there are serious physical and evolutionary degeneracies that make DSFG photometric redshifts highly uncertain (with precision often σ Δz/1+z ≳ 1; Casey 2020). This latter point is often seen as a benefit: their strongly negative k-correction means that the flux density of DSFGs at z > 1 remains constant with increasing z for λ obs ≳ 850 μm, meaning a DSFG at z ∼ 10 can be observed as readily as a DSFG at z ∼ 1 (Blain et al. 2002). However, when searching for high-z DSFGs, this negative k-correction is also a hindrance as it becomes difficult to identify redshifts for galaxies with only long-wavelength emission. This effect is further exacerbated by the (sub)millimeter color degeneracy between dust temperature and redshift. Thus, barring clear identification at other wavelengths, it is easy to confuse z ∼ 2 and z ∼ 6 DSFGs with solely submillimeter observations.

2020AandA...644L...7G__Magdis_et_al._2020_Instance_1

As in G18, we compiled existing constraints on the molecular gas fraction fgas of quiescent and pSB galaxies from recent literature, namely: local QGs consisting of the ATLAS3D (Young et al. 2011; Cappellari et al. 2013; Davis et al. 2014) and HRS (Boselli et al. 2014; Lianou et al. 2016) ETG samples as well as the samples of pSB galaxies (hereafter, the “low-z pSB” sample) of French et al. (2015) and Alatalo et al. (2016); at low and intermediate redshift, the ETG sample of Spilker et al. (2018) and the pSB sample of Suess et al. (2017); at intermediate and high redshift, constraints from Hayashi et al. (2018) on gas in z ∼ 1.46 cluster ETGs, as well as on individual galaxies from Sargent et al. (2015), Bezanson et al. (2019), and Rudnick et al. (2017). Given its size, we divided the ATLAS3D sample into high- and low-mass subsamples, choosing 5 × 1010 M⊙ as the cut-off mass. In addition, we also included fgas estimates derived from the (median) stacked FIR spectral energy distributions of ETGs at z ∼ 1.8 (G18; 977 galaxies), z ∼ 1.2, z ∼ 0.8, and z ∼ 0.5 (1394, 1536, and 563 galaxies, respectively; Magdis et al. 2020, hereafter M20). Finally, at higher redshift (z ∼ 3), we converted star formation rates (SFR) estimated from spectroscopy (Schreiber et al. 2018a; D’Eugenio et al. 2020) into gas masses assuming the star formation efficiency found by G18. As a consequence of our zmax = 3.5, we did not include higher-redshift quiescent galaxies (Glazebrook et al. 2017; Schreiber et al. 2018b; Tanaka et al. 2019; Valentino et al. 2020) in the analysis and considered z ∼ 3 galaxies as pSB. The dust-based estimates of G18 and M20 (and, by extension, the z ∼ 3 semi-constraints) assume a gas-to-dust ratio (G/D). It is dependent on metallicity, which is presumed to be solar or higher owing to both the relatively high gas-phase metallicity of MS galaxies at z ≲ 1 (e.g., Mannucci et al. 2010) and the already high stellar metallicities of QGs at z >  1 (Onodera et al. 2015; Estrada-Carpenter et al. 2019). Here we adopted an intermediate value between the solar and supersolar G/Ds used in M20, and we increased the error bars of these points to include both the solar and supersolar confidence estimates. These various samples, which are summarized with their selection criteria in Table B.1, combine into a nonhomogeneous dataset: some were specifically selected as ETGs, and others were based on varying degrees of quiescence. In particular, pSB galaxies are not necessarily truly quiescent and could, in principle, resume normal star formation. However, as a possible precursor of QGs, they provide useful, though not constraining (see Sect. 4), comparison samples for the model. Here we refer to all equally as either QGs or pSB galaxies, and we make the assumption that, on average, these different samples are not otherwise significantly biased with regard to their gas content compared to the full population, given each mass limit and type.

2018AandA...614A..31B__Nakajima_et_al._2012_Instance_1

High-redshift z ≳ 2 galaxies with prominent Lyman-α emission (Partridge & Peebles 1967), referred to as Lyman-α emitters (LAEs), have become the subject of intense research over the last two decades. Since LAEs are believed to be powered at least partially by ongoing star formation, they are expected to belong to a relatively low-mass and young class of actively star-forming galaxies. This suggests that the number density of LAEs is large enough to map out the large-scale structure of the high-redshift universe. Increasing numbers of LAEs are being observed by surveys including narrow-band imaging (e.g., Nakajima et al. 2012) and integral field unit spectroscopy such as MUSE (e.g., Bacon et al. 2006), amounting to about 104 emitters known to date since the first detections (e.g., Cowie & Hu 1998; Hu & McMahon 1996). These observations enable us to study galaxy clustering at somewhat small scales (e.g., Diener et al. 2017; Ouchi et al. 2010, 2018) as well as evolution of the Lyman-α luminosity function (e.g., Konno et al. 2016, 2018; Ouchi et al. 2008). A remarkable example for such surveys is the Hobby-Eberly Telescope Dark Energy Experiment (HETDEX; Adams et al. 2011; Hill et al. 2008). HETDEX will map out the three-dimensional distribution of nearly one million LAEs (Leung et al. 2017) at 1.9 z 3.5 over 400 deg2. This allows us to precisely measure the large-scale (≳10 Mpc) clustering of LAEs (Agrawal et al. 2017; Chiang et al. 2013) and of the Lyman-α intensity map (Saito et al., in prep.) with the primary scientific goal being to determine the cosmic expansion history via baryon acoustic oscillations (BAOs) and the growth of structure via redshift-space distortions (RSD). We refer the reader to Alam et al. (2017) and references therein for recent efforts of BAO and RSD measurements. In addition, these surveys offer exciting opportunities to study the connection between LAEs and their environment, including the cross-correlation between LAEs and the Lyman-α forest or other galaxy populations.

2021MNRAS.505.2561C__Michtchenko_et_al._2018_Instance_1

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

2021AandA...650A.205V__Jones_et_al._2021_Instance_1

The question of the evolution of exoplanet systems after the main sequence of their host is generally addressed by studying exoplanets around subgiants, RGB stars, and normal HB (RC) stars (hereafter the ’classical’ evolved stars). These classical evolved stars are typically very large stars, with radii ranging from ~ 5− 10 R⊙ to more than 1000 R⊙. This is much larger than hot subdwarfs, which have radii in the range ~ 0.1−0.3 R⊙ (Heber 2016). Their mass is typically higher than ~ 1.5 M⊙, compared to~0.47 M⊙ for hot subdwarfs. The transit and radial velocity (RV) methods are both challenging for these classical evolved stars because the transit depth is diluted and there are additional noise sources (Van Eylen et al. 2016). Another difficulty forthe question of the fate of exoplanet systems after the RGB phase itself is the difficulty of distinguishing RGB and RC stars based on their spectroscopic parameters alone, which is sometimes hard even with help of asteroseismology (Campante et al. 2019). As a consequence, only large or massive planets are detected around the classical evolved stars (Jones et al. 2021, and references therein). A dearth of close-in giant planets is observed around these evolved stars compared to solar-type main-sequence stars (Sato et al. 2008; Döllinger et al. 2009). This may be caused by planet engulfment by the host star, but current technologies do not allow us to determine whether smaller planets and remnants (such as the dense cores of former giant planets) are present. The lack of close-in giant planets may also be explained by the intrinsically different planetary formation for these intermediate-mass stars (see the discussion in Jones et al. 2021). Ultimately, the very existence of planet remnants may be linked to the ejection of most of the envelope on the RGB that occurs for hot subdwarfs, while for classical evolved stars, nothing stops the in-spiraling planet inside the host star, and in all cases, the planet finally merges with the star, is fully tidally disrupted, or is totally ablated by heating or by the strong stellar wind. In other words, the ejection of the envelope not only enables the detection of small objects as remnants, but most importantly, it may even be the reason for the existence of these remnants by stopping the spiraling-in in the host star.

2022MNRAS.515.4520F__Rodriguez-Franco_&_Cuevas_2013_Instance_1

The atmosphere above the Canary Islands, and especially the western islands Tenerife and La Palma, has been extensively characterized over the past 50 yr (Kiepenheuer 1972; McInnes & Walker 1974; Brandt & Righini 1985; Murdin 1985; Stickland et al. 1987; Whittet, Bode & Murdin 1987; Menéndez-Valdés & Blanco 1992; Jímenez, González Jorge & Rabello-Soares 1998; Mahoney, Muñoz-Tuñón & Varela 1998; Maring et al. 2000; Torres et al. 2002; Rodríguez et al. 2004; Alonso-Pérez et al. 2007; Basart et al. 2009; Lombardi, Zitelli & Ortolani 2009; Delgado et al. 2010; Vernin et al. 2011; Cuevas et al. 2013; Rodriguez-Franco & Cuevas 2013; Varela et al. 2014; Laken et al. 2016; Vogiatzis et al. 2018; Hidalgo et al. 2021). Due to the combination of large-scale atmospheric circulation on the descending branch of the Hadley cell (see e.g. Palmén & Newton 1969; Rodríguez et al. 2004), and the ‘Trade’ or ‘Alisios’ winds coming from the Azores high area, a stable and strong temperature inversion layer (TIL) appears (Font 1956; Huetz-de-Lemps 1969), whose top is typically found at heights around 1200 m a.s.l. in summer and 1800 m a.s.l. in winter (Torres et al. 2002; Carrillo et al. 2016). Whenever the temperature inversion is able to separate two well-defined regimes, the moist marine boundary layer (MBL) and above it, the dry free troposphere (FT), the phenomenon is called an ‘Alisio’ inversion. This happens about 80 per cent of the time (Torres et al. 2002). Under these conditions, the FT can be characterized as ‘ultraclean’, i. e., the concentration of accumulation mode particulate matter is lower than 50 cm−3 (Pennypacker, Diamond & Wood 2020). The clean atmospheric conditions are one of the reasons why the Canarian observatories ‘Observatorio del Teide (OT)’, located at ∼2400 m a.s.l., and ‘Observatorio del Roque de los Muchachos (ORM)’, located between ∼2100 and ∼2400 m a.s.l., are known to belong to the best astronomical sites worldwide.

2021AandA...656A.126K__Lowry_et_al._2007_Instance_1

Very small asteroids (VSAs) are objects with diameters D 150 m. They often rotate with periods shorter than 2 h enabling us to study their internal structure by comparing the centrifugal force with the material forces holding them together (Holsapple 2007). Because of their small sizes, VSAs are sensitive to the YORP effect (Rubincam 2000), which is a torque induced on the rotating body by the thermal radiation emitted by its surface complemented by a torque produced by scattered sunlight. YORP can either spin up or slow down the asteroid rotation as well as change the obliquity of its spin axis ϵ, which is an angle between the normal to the asteroid orbital plane and its rotation axis. While the fast rotation has been observed for many VSAs, their spin axes were not determined except for one object1, (54 509) YORP (whose name is the same as the name of the effect itself). (54 509) was the first asteroid for which the effect of YORP has been observed (Lowry et al. 2007; Taylor et al. 2007). The obliquity of the (54 509) spin axis is ϵ = 173°, which means it is nearly perpendicular to the asteroid orbital plane. Such an orientation of the spin axis was found as an end state of the YORP evolution in the simulations performed by Čapek & Vokrouhlický (2004) for objects with finite surface thermal conductivity. If their prediction is true, then for the VSAs that experienced a strong YORP effect for a long time (and the fastest rotating VSAs are such objects), we should observe spin axis obliquities close to 0° or 180°. Recently Golubov et al. (2021) have shown that for very small objects of a highly irregular shape, the transverse heat conduction (TYORP) can add new asymptotic states for the obliquity. For this reason it would be interesting to verify those predictions with observationsof VSAs. To do that, we should observe their lightcurves at many different positions on the sky to be able to determine their spin axes. For near-Earth objects (NEAs) this condition is met either by the Earth co-orbital asteroids – and (54 509) is an example of such objects – or by objects for which their close encounter with the Earth can be observed along a long arc on the sky.

2020MNRAS.499..710G__Assef_et_al._2015_Instance_1

An alternative possibility to reduce the fraction of X-ray sources in the WISE R75 wedge is to allow a large fraction of heavily obscured AGN. This can be achieved for example, by increasing the fraction of Compton-thick AGN in the X-ray luminosity function above the current assumption of 34 per cent, or by increasing the scatter in the relation that links X-ray obscuration to optical extinction beyond the adopted value of 0.5 dex. Relaxing the above model assumptions would allow heavily obscured and hence, X-ray faint, sources to be selected by the WISE R75 criteria. There is indeed evidence for a potentially large population of heavily obscured, possibly Compton-thick AGN, among the WISE population (e.g. Assef et al. 2015; Mountrichas et al. 2017; Yan et al. 2019). The SDSS spectroscopic follow-up programme presented by LaMassa et al. (2019) also revealed a non-negligible number of WISE R75 sources that are optically faint ($r\gtrsim 22$ mag), lie at redshifts $z\lesssim 1$ and are spectroscopically identified by their prominent [O ii] 3727 emission lines. These sources are prime candidates for heavily obscured AGN. Our baseline model predicts that the most heavily obscured, Compton-thick, WISE AGN are at low redshift, $z\lesssim 0.6$ (see Fig. 19) and relatively optically bright. The top panel of Fig. 22 demonstrates the latter point by plotting the r-band distribution of the Compton-thick AGN population predicted by the model. Observationally, the identification of such AGN needs to account for the relatively high level of contamination of the WISE R75 AGN selection by star-forming galaxies at redshifts $z\lesssim 0.6$ (s1ee Fig. 14). One approach to achieve this is via diagnostic optical emission-line ratios (e.g. Kewley et al. 2001) to separate star-forming galaxies from Seyfert-2s. Observations at hard X-rays can also provide useful information on the nuclear activity of a galaxy and the level of line-of-sight obscuration. The bottom panel of Fig. 22 shows the expected 2–10 keV X-ray flux distribution of the Compton-thick AGN predicted by the model. The expected fluxes have already been reached in deep X-ray survey fields, for example, COSMOS-Legacy (Civano et al. 2016). Study of the X-ray spectral properties of WISE-selected AGN in such fields can test the baseline model predictions for the demographics of Compton-thick AGN in the WISE R75 wedge.

2018AandA...613L...1D__Salafia_et_al._2015_Instance_1

Similarly to long bursts, short GRBs are thought to be produced by a relativistic jet with a typical half-opening angle θjet ~ 5–15 deg (Fong et al. 2016). However, whether or not BNS mergers can always efficiently produce a relativistic jet is still debated (Paschalidis et al. 2015; Ruiz et al. 2016; Kawamura et al. 2016). Given the small probability that our line of sight was within the jet half-opening angle, 1 − cos(θjet), it is unlikelythat the first short GRB associated to a GW event had a jet pointing towards the Earth. The extremely low γ-ray luminosity of GRB 170817A has been interpreted as being due to (i) the debeamed radiation of a jet observed off-beam (i.e. viewing angle θview > θjet), provided that the jet bulk Lorentz factor is significantly smaller than usually assumed (see, e.g. Pian et al. 2017). Alternatively, the jet could be (ii) structured, with a fast and energetic inner core surrounded by a slower, less energetic layer/sheath/cocoon (first proposed for long GRBs – Lipunov et al. 2001; Rossi et al. 2002; Salafia et al. 2015 – and only recently extended to short GRBs – Kathirgamaraju et al. 2018; Lazzati et al. 2017a; Gottlieb et al. 2017; Lazzati et al. 2017b; Lyman et al. 2018; Margutti et al. 2018; Troja et al. 2018a). In this scenario the faint, off-beam emission is due to the slower component, which originates from the interaction of the jet with the merger dynamical ejecta or the post-merger winds. Recently, Mooley et al. (2018) suggested the possibility that (iii) the jet was not successful in excavating its way through the ambient medium and that GRB 170817A was due to its vestige, a quasi-isotropic cocoon with a velocity profile. Last but not least, (iv) a jet-less interpretation of GRB 17017A could still be viable: an isotropic fireball, expanding ahead of the kilonova ejecta, which could account for both the low luminosity of the γ-ray event and the properties of the EM component in the radio and X-ray bands (Salafia et al. 2017). In this case, all BNS mergers should have this kind of faint, hard X-ray counterpart. All of the above scenarios have relatively clear predictions for the temporal and spectral evolution of the electromagnetic emission from X-rays to the radioband. A comprehensive discussion of the possible physical scenarios leading to the observed broad-band emission of GW 170817/GRB 170817A can be found in Nakar & Piran (2018). Recent radio and X-ray observations(Mooley et al. 2018; Margutti et al. 2018; Ruan et al. 2018; Troja et al. 2018a), carried out until ~ 110−115 d after the event, indicate that the source flux is steadily rising and that the spectral energy distribution (SED) over these bands is consistent with a single power-law component. These results disfavour interpretation (i) reported above (an off-beam homogeneous jet).

2015MNRAS.454.1644L__Kotze_&_Charles_2012_Instance_2

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

2020AandA...644A..59K__Pastorello_et_al._2019_Instance_1

The year 2020 marks the 350 yr anniversary of the discovery of the eruption of Nova 1670 (or CK Vul) made by European astronomers (Shara et al. 1985). Their observations, predominantly performed with a naked eye, traced the object’s evolution on the sky in 1670−1672. From the archive records, we know that the eruption was rather unusual, in particular it was very much unlike classical novae. The light curve of CK Vul displayed three peaks and the star was described as reddish in the later stages of the eruption (Hevelius 1671; Shara et al. 1985). These characteristics resemble closest the behavior often observed in (luminous) red novae (Kato 2003; Tylenda et al. 2013), a modern category of eruptive stars known from our and other galaxies (e.g. Pastorello et al. 2019). Red novae are recognized as manifestations of on-going mergers of non-compact stars such as main-sequence dwarfs, sub-giants, or red giants (Soker & Tylenda 2003; Tylenda & Soker 2006; Tylenda et al. 2011; Pastorello et al. 2019). While the number of known red novae, mainly extragalactic ones, is quickly rising (e.g., Stritzinger et al. 2020), we know only a few red-nova remnants that are decades old (Kamiński et al. 2018a). The remnant of the 1670 eruption of CK Vul, as a candidate post-merger site, could be the oldest (counting from the onset of the eruption) known object of this type and as such offers the opportunity to investigate a merger aftermath centuries after the stellar coalescence. The nature of the progenitor system of CK Vul has been debated. Eyres et al. (2018) proposed that the seventeenth-century merger took place between a white dwarf and a brown dwarf, but there is little quantitative evidence to support this. Based on the analysis of the source’s chemical and isotopic composition, including the unique presence of the radioactive isotope of 26Al, Kamiński et al. (2018b) found that the progenitor system of CK Vul included at least one red-giant-branch (RGB) star with a fully developed helium core.

2015AandA...582A..22L__Todorov_et_al._2014_Instance_1

Dust distribution:We employed the standard flared disk model with well-mixed gas and dust, which has been successfully used to explain the observed SEDs of a large sample of young stellar objects and BDs (e.g., Wolf et al. 2003; Sauter et al. 2009; Harvey et al. 2012a; Joergens et al. 2013; Liu et al. 2015). The structure of the dust density is assumed with a Gaussian vertical profile (1)\begin{equation} \rho_{\rm{dust}}=\rho_{0}\left(\frac{R_{*}}{\varpi}\right)^{\alpha}\exp\left[-\frac{1}{2}\left(\frac{z}{h(\varpi)}\right)^2\right], \label{dust_density} \end{equation}ρdust=ρ0R∗ϖαexp−12zh(ϖ)2,and the surface density is described as a power-law function (2)\begin{equation} \Sigma(\varpi)=\Sigma_{0}\left(\frac{R_{*}}{\varpi}\right)^p, \end{equation}Σ(ϖ)=Σ0R∗ϖp,where ϖ is the radial distance from the central star measured in the disk midplane, and h(ϖ) is the scale height of the disk. The disk extends from an inner radius Rin to an outer radius Rout. To the best of our knowledge, among our sample, there are five objects that have been identified as binary systems so far. They are 2M1207 (a~55 AU, Chauvin et al. 2004), J04221332+1934392 (a~7 AU, Todorov et al. 2014), J04414489+2301513 (a~15 AU, Todorov et al. 2014), USD161833 (a~134 AU, Bouy et al. 2006), and USD161939 (a ~ 26 AU, Bouy et al. 2006), where a refers to the separation within the system. The disks around individual components in binary systems are expected to have truncation radii of the order of (0.3 − 0.5)a (Papaloizou & Pringle 1977). We adopted 0.5 a as the disk outer radii for 2M1207, USD161833, and USD161939. For the close pairs (a ≲ 15 AU, J04221332+1934392 and J04414489+2301513), dynamical simulations of star-disk interactions suggest that individual disks are unlikely to survive (e.g., Artymowicz & Lubow 1994). Disk modeling is complicated in those close multiple systems. For simplicity, we assume that the emission is associated with circumbinary disks of 100 AU in size. For other objects, we fix Rout = 100 AU in the modeling because the choice of this parameter value makes essentially no difference to the synthetic SEDs in the simulated wavelength range (Harvey et al. 2012a). The scale height follows the power-law distribution(3)\begin{equation} h(\varpi) = H_{100}\left(\frac{\varpi}{100\,\rm{AU}}\right)^\beta,\\ \end{equation}h(ϖ)=H100ϖ100 AUβ,with the exponent β characterizing the degree of flaring and H100 representing the scale height at a distance of 100 AU from the central star. The indices α, p, and β are codependent through p = α − β. We fix p = 1, which is the typical value found for T Tauri disks in the sub-millimeter (e.g., Isella et al. 2009; Guilloteau et al. 2011), since only spatially resolved data can place constraints on this parameter (e.g., Ricci et al. 2013, 2014). Dust properties:We assume the dust grains to be a homogeneous mixture of 75% amorphous silicate and 25% carbon with a mean density of ρgrain = 2.5 g cm-3 and the complex refractive indices given by Jäger et al. (1994, 1998), and Dorschner et al. (1995). Porous grains are not considered because the fluxes at wavelengths beyond ~ 2 μm are almost independent of the degree of grain porosity in low-mass disks, as shown by Kirchschlager & Wolf (2014). The grain size distribution is given by the standard power law dn(a) ∝ a-3.5da with minimum and maximum grain sizes amin = 0.1 μm and amax = 100 μm, respectively. The choice of the minimum value for the grain size, amin, ensures that its exact value has a negligible impact on the synthetic SEDs. Since there is no information about the maximum grain sizes of our target disks, as provided, e.g., by the (sub)millimeter spectral index, we adopt amax = 100 μm. The Herschel/PACS far-IR observations are sensitive to dust grains with this assumed sizes. Strong grain growth up to millimeter sizes as detected in some BD disks (e.g., Ricci et al. 2012, 2013, 2014; Broekhoven-Fiene et al. 2014) would remain undetected in our data and could affect the disk mass. Our prescription for the dust properties is identical to those used in Liu et al. (2015).

2021AandA...649A..58L__Bemporad_et_al._(2018)_Instance_3

The leading edges of the transients normally leave bright traces in the images of visible light, inspiring many methods that were developed to derive their locations and velocities, such as the icecream cone model (Fisher & Munro 1984), the graduated cylindrical shell (GCS) model (Thernisien 2011), geometric triangulation methods (Liu et al. 2010), mask-fitting methods (Feng et al. 2012), and trace-fitting methods including the point-p, fixed-Φ, harmonic mean, and self-similar expansion fitting methods (e.g., Sheeley et al. 1999; Howard et al. 2006; Davies et al. 2012; Möstl & Davies 2013). To derive the velocity distribution inside one transient rather than only at its leading edge, some other techniques have been proposed. Colaninno & Vourlidas (2006) applied an optical flow tool to extract the velocity vector of a coronal mass ejection (CME) in digital images. Feng et al. (2015) derived the radial velocity profiles of the whole CME from the spatial distribution of its density given by the mass continuum equation. A cross-correlation method was applied to derive continuous 2D speed maps of a CME from coronagraphic images by Bemporad et al. (2018). In their work, the radial shift pixel by pixel is determined by maximizing the cross correlation between the signal in a radial window at one frame and the signal in a radial shifted window at the previous frame, and the radial speed just equals the radial shift over the time interval between the two frames. Ying et al. (2019) improved this cross-correlation method by analyzing data in three steps: forward step (FS), backward step (BS), and average step (AS). In the FS (BS), the 2D velocity map between the current and the previous (next) frame is constructed with almost the same method as Bemporad et al. (2018). In the AS the average, velocity is obtained from the FS and BS. The velocities derived by all these methods are the component of the flow velocity vector projected onto the POS. This may underestimate the velocity especially for transients that do not propagate in the POS. Methods such as the polarizaition ratio technique (Moran & Davila 2004; DeForest et al. 2017) or the local correlation tracking (LCT) method (Mierla et al. 2009) can derive the 3D geometric information of the whole transients, but not the velocity distribution. Bemporad et al. (2018) chose the propagating direction averaged over the whole CME derived by the polarization ratio technique to correct the radial speed in the 2D maps, but the key information along the LOS is still lacking.

2020ApJ...898..143S__Shao_&_Li_2018_Instance_1

In Figure 6, we plot the histogram distributions of the calculated number of Galactic BH–He XRBs as a function of the companion mass, BH mass, orbital period, and orbital eccentricity for Models A–D. Since the progenitor systems containing a BH and an OB star possess similar distributions of the OB star masses and the orbital periods in all models, the BH–He XRBs also have similar distributions of helium masses and orbital periods under the assumptions of different models. The masses of the helium stars are mainly distributed in the range of , and the orbital periods are distributed in a wide range (∼0.1–1000 days) with a peak at ∼30–100 days. Some recent investigations showed that the mass-transfer process in the lobe-filling BH binaries is more stable than previously expected (e.g., Pavlovskii et al. 2017; Shao & Li 2018); the maximal mass ratio of the donor star to the BH for stable mass transfer can reach as high as ∼6. Hence only a few BH–He XRBs with orbital periods of ≲1 day are produced in our calculations. Remarkably, such wind-fed XRBs in close orbits may be bright enough to be detected. In nearby galaxies, there also exist a couple such close XRBs (Esposito et al. 2015, and references therein). During the evolution of the progenitor binaries, mass accretion of the BH can increase its mass. We see that the BH masses in the BH–He XRBs are distributed in a broad range of . Model A predicts that BHs are more likely to possess masses of since more binaries survive from BH formation via direct collapse without natal kicks, while Models B–D tend to produce light BHs with mass distributions that peak at due to the IMF. Differently from the BH systems with an OB star, almost all BH–He XRBs have relatively low eccentricities of ≲0.4, whose distribution has two distinct peaks at ≲0.1 and ∼0.3. Tides and mass transfer between binary components tend to circularize the orbits during the progenitor system evolution. Our obtained low eccentricities can coincide with the observations of Galactic W-R−O binaries (van der Hucht 2001) that are also post-mass-transfer systems; although, we do not include a detailed treatment for the orbital evolution of mass-transferring eccentric binaries (e.g., Sepinsky et al. 2009; Dosopoulou & Kalogera 2016).

2017MNRAS.470..626M__Robitaille_et_al._2012_Instance_1

The opacity file used for the PAH/VSG population (draine_opac_new.dat) was computed by Bruce Draine (Draine & Li 2007) and is available at https://github.com/hyperion-rt/hyperion-pah/tree/master/PAH-Legacy/input. It was computed taking two lognormal size distributions, one with grain radii below 20 Å and another for grains with radii from 20 to 200 Å. PAHs constitute a mass fraction of 4.5 per cent of this grain population and they have the size-dependent fractional ionization of the diffuse ISM (Wood et al. 2008). The presence of the 8.6-μm feature in the Spitzer IRS spectrum of MPI13 implies the existence of ionized PAHs. The PAH/VSG grain population is not in thermal equilibrium with the radiation field. However, their emissivities can be computed based on specific energy absorbed in each grid cell and assuming that the emissivity is the function of the mean intensity of the radiation field, which approximates the spectral shape to the first order (Robitaille et al. 2012; Whitney et al. 2013). The code uses the pre-computed emissivity tables of Draine & Li (2007) for different specific energy. As a Monte Carlo energy pocket is absorbed by a PAH/VSG, the pocket is reprocessed, sampling a new frequency from pre-computed emissivity tables. A detailed discussion on this can be found in Wood et al. (2008). The updated version of Hochunk3D is hence quite useful in modelling the circumstellar shells around evolved stars, which have thermally fluctuating dust. The code uses the method of Lucy (1999) to calculate the dust temperature very efficiently by summing up photon path-lengths. The temperature of the grid cell is updated when an iteration of simulation is completed, and the temperature converges in three to five iterations; this efficiency is due to the fact that flux is conserved exactly across all surfaces (Lucy 1999). By constructing Monte Carlo radiation fields that are rigorously divergence-free, rapid convergence is achieved with the temperature correction procedure in this method. We have performed the Lucy temperature correction of dust with five iterations.

2022ApJ...929..186L__Ferraro_et_al._2018b_Instance_2

Our group is addressing this problem by combining a variety of complementary perspectives: (i) by constructing a new generation of high-quality star density profiles derived from star counts instead of surface brightness (see Lanzoni et al. 2007a, 2010, 2019; Miocchi et al. 2013; Pallanca et al. 2021); (ii) by investigating the population of stellar exotica (Ferraro et al. 2001, 2003, 2015, 2016; Pallanca et al. 2010, 2013,2014, 2017; Cadelano et al. 2017, 2018, 2020) and their connection with the dynamical evolution of the parent cluster (see Ferraro et al. 2009, 2012, 2018a, 2019; Lanzoni et al. 2016); (iii) by characterizing the three-dimensional (3D) global velocity space through the analysis of the velocity dispersion profile and rotation curve from resolved star spectroscopy (Lanzoni et al. 2013, 2018a, 2018b; Ferraro et al. 2018b) and proper motions (PMs; see Raso et al. 2020). The determination of GGC internal kinematics from resolved star velocities is particularly relevant and challenging. In this context we promoted the ESO-VLT Multi-Instrument Kinematic Survey (hereafter the MIKiS survey; Ferraro et al. 2018b, 2018c), a project specifically designed to characterize the kinematical properties of a sample of GGCs in different dynamical evolutionary stages from the radial velocities (RVs) of hundreds of individual stars distributed over the entire radial range of each stellar system. To this end, the survey fully exploits the spectroscopic capabilities of different instruments currently available at the ESO Very Large Telescope (VLT): originally designed to use the adaptive optics (AO) assisted integral-field spectrograph SINFONI, the multiobject integral-field spectrograph KMOS, and the multiobject fiber-fed spectrograph FLAMES/GIRAFFE, it has been recently complemented with individual projects and an ongoing large program (PI: Ferraro) fully exploiting the remarkable performances of the AO-assisted integral-field spectrograph MUSE.

2022AandA...667A.131B__Izumi_et_al._(2016)_Instance_2

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

2020MNRAS.499.1788W__Malhotra_et_al._2001_Instance_2

Owing to their brightness at rest-frame FIR wavelengths, the ionized and neutral species of Carbon, Nitrogen, and Oxygen are powerful diagnostic lines for tracing the ISM of nearby and distant galaxies. When combined with photodissociation region (PDR) models (Tielens & Hollenbach 1985), measurement of the emission from different lines provide a means to constrain quantities such as the ionization rate and metallicity of the ISM of galaxies. Multilevel FIR transition lines have now been widely surveyed in local galaxies, originally by ISO (e.g. Luhman et al. 1998; Malhotra et al. 2001) and more recently using the PACS spectrometer (Poglitsch et al. 2010) on the ESA Herschel Space Observatory (Pilbratt et al. 2010). The [C ii]158 μm line is typically the brightest in star-forming galaxies, arising from ionized, and even neutral gas where it is the main coolant (Wolfire et al. 2003). Another commonly observed FIR line tracer of the ionized ISM is [N ii]122 μm. It has the advantage that it can be found associated with lower excitation gas, close to that observed in our own Galaxy (e.g. Goldsmith et al. 2015; Herrera-Camus et al. 2016). Another major coolant of the ISM is [O i]63 μm (Wolfire et al. 2003). Owing to its high excitation temperature and critical density, it can dominate the cooling in regions of starburst activity (Kaufman et al. 1999; Kaufman, Wolfire & Hollenbach 2006; Brauher, Dale & Helou 2008; Croxall et al. 2012; Narayanan & Krumholz 2017). When combined with measurements of the [C ii]158 μm line intensity and FIR luminosity, the [O i]63 μm line intensity can constrain the FUV field, G, and the gas density using PDR models. The luminosity in these FIR lines generally exhibits a deficit in the most FIR luminous galaxies compared to the trend expected from lower luminosity galaxies (e.g. Malhotra et al. 2001; Graciá-Carpio et al. 2011; Díaz-Santos et al. 2017). This has made the emission from lines like [O i]63 μm more challenging to detect at high-redshifts.

2022AandA...661A..10B__Santos_et_al._2008_Instance_1

It is also possible that these clusters have a smaller extent and can just be missed by our extent selection as our detection algorithm sets the extent to zero if it is smaller than 6 (Brunner et al. 2022). Following the method presented in Ghirardini et al. (2021a), we estimated several dynamical properties of the clusters in the point source sample and compared them with the extent-selected sample presented in Ghirardini et al. (2021a). In Fig. 5 we compare the distributions of the core radii (Rcore) constrained by the V06 model and the concentration parameter (cSB) between these two samples. The concentration parameter is defined as the ratio of the surface brightness within 0.1 R500 to the surface brightness within R500 (Ghirardini et al. 2021a; Santos et al. 2008; Maughan et al. 2012). Intuitively, the expectation is that the smaller the core radius, the more compact the cluster. The left panel of Fig. 5 clearly shows that the clusters in the point source sample have relatively smaller core radii, hence the emission is more concentrated in a smaller area. Consistently, the concentration of the point source sample shows a clear excess in higher values than the extent-selected sample, indicating that a significantly larger fraction of cool-core clusters and clusters host a central AGN. We performed the same experiment by applying cuts in flux 1.5 × 10−14 ergs s−1 cm−2 and in detection likelihood to test whether the clusters are missed by the extent selection because they are fainter and/or more compact than the extent-selected clusters. The distribution of number density, core radius, and concentration parameter remains the same, indicating that the population of clusters in the point source catalog is more compact than the extent-selected sample. The extent-selected sample does not show a clear bias toward cool-core clusters or clusters with a central AGN, but contains the fraction of cool-cores is similar to that of SZ surveys (Ghirardini et al. 2021a). In this sample, we observe the opposite trend.

2019MNRAS.482.2731Z__Perley_et_al._2008_Instance_1

We searched the literature carefully and selected all published GRB-DLAs and QSO-DLAs sightlines conforming to our requirements which are as follows. The object must have spectral energy distributions (SEDs) and optical spectroscopic data available with measurements of AV, column densities of Zn ii and Fe ii, or of S ii and Si ii. The GRBs are selected only if they had their optical extinction derived from simultaneous SED fitting to X-ray-to-optical/NIR data using either a single or broken power law (see Greiner et al. 2011; Zafar et al. 2011; Schady et al. 2012; Covino et al. 2013; Bolmer et al. 2018; Zafar et al. 2018a for discussion on AV determination). This is a reliable method to determine extinctions at higher redshifts where the intrinsic slopes are constrained by the X-ray data. Note that there is some degeneracy between broken power-law break frequency (see De Cia et al. 2011) and extinction, which could lead to inference of grey dust for some instances (Watson et al. 2006; Perley et al. 2008; Friis et al. 2015). However, overall a fixed spectral break change (Δβ = 0.5) between the optical and X-ray slopes is preferred for GRBs (Greiner et al. 2011; Zafar et al. 2011; Japelj et al. 2015). For QSO-DLAs, reddening must be determined either from QSO colours or extinctions through template fitting to the QSO SED. Those methods are less robust than the X-ray-supported GRB fits, but are widely adapted. We refer the reader to Zafar et al. (2015) and Krogager et al. (2015, 2016) for more discussions on AV determination for QSO. The requirement for the pairs of elements (Zn ii and Fe ii, or of S ii and Si ii) are in order to be able to derive depletions. Here, Zn and S are volatile elements and Fe and Si are refractory elements (e.g. Ledoux, Bergeron & Petitjean 2002; Draine 2003; Vladilo et al. 2011; De Cia et al. 2016). Defined this way our initial sample consists of 28 GRBs and 32 QSO-DLAs with the required measurements available. To this we add sources where only part of the required measurements are complete but where limits have been determined for the rest. The vast majority of the elemental abundances have been determined via detailed spectral line fitting, for eight GRBs elemental abundances (or limits) were derived from rest-frame equivalent widths of non-saturated lines provided by Fynbo et al. (2009) as described in Laskar, Berger & Chary (2011) and Zafar & Watson (2013, see references to Table 1).

2017AandA...607A.103G__Gradie_&_Tedesco_(1982)_Instance_1

In Fig. 4, 29 observations of 15B-type asteroids are presented. As shown by Gil-Hutton & Cañada-Assandri (2012) and Gil-Hutton et al. (2014), in any plot of the observations of B-type objects a dispersion always appears in the whole phase angle range and it seems that these objects do not follow a single phase-polarization curve. The reason is probably that the asteroids classified by Bus & Binzel (2002) as members of this taxonomic class include objects that belong to the old F-class, originally proposed by Gradie & Tedesco (1982) and included in the Tholen taxonomy, due to its flat spectrum and low albedo; but several F-class objects have phase-polarization curves that are characterized by a comparatively low value of the inversion angle (Belskaya et al. 2005). Therefore, we also include in Fig. 4 nine observations of five asteroids classified as F-class by Tholen (1989) but included in other taxonomic types by Lazzaro et al.: (426) Hippo (X-type), (530) Turandot (C-type), (762) Pulcova (Cb-type), (778) Theobalda (C-type), and (877) Walkure (without a Bus taxonomy classification). The group of objects formed by (142) Polana, (213) Lilaea, (877) Walkure, and (1021) Flammario, all of them belonging to the old F-class, have observations indicating a low value of inversion angle, and the single observation of (314) Rosalia also indicates a similar polarimetric behavior even though it was not classified using the taxonomy of Tholen. On the other hand, the asteroids (62) Erato, (372) Palma, (635) Vundtia, and (981) Martina have measurements well below the theoretical phase-polarization curve for this taxonomic class, so it would seem to have a polarimetric behavior that corresponds more to a C-type than to a B-type. The apparent discrepancies might be indicative of some heterogeneity in the polarimetric properties within this particular taxonomic type, or alternatively, there might be some misclassification of the objects as could be the case for (635) Vundtia and (981) Martina, which were classified as C- and CFU-type by Tholen (1984).

2018ApJ...864..158L__Drake_et_al._2006_Instance_1

However, if the magnetic curvature contraction term (last term in Equation (15)) is dominated by the compression term (second-to-last term in Equation (15)) so that , the contraction or merging is strongly compressible (∇ · VE 0), resulting in an increase in flux-rope magnetic energy instead. In this case the plasma environment is doing positive mechanical work on the flux-rope structure, pushing magnetic field lines together to enhance flux-rope magnetic energy. Although it appears that small-scale flux ropes tend to contract or merge predominantly incompressibly in discussions of particle simulations (e.g., Drake et al. 2006; Dahlin et al. 2016), and it is also thought of as intrinsically incompressible in its manifestation as the quasi-2D turbulence component in N i MHD theory of solar wind turbulence (Zank et al. 2017), there is observational evidence to the contrary. For example, when primary current sheets associated with interplanetary coronal mass ejections (ICMEs) interact with the heliospheric current sheet, these structures are disturbed and several small-scale flux-rope structures may be formed when turbulent magnetic reconnection occurs in these structures. The flux ropes, being trapped between the converging heliospheric current sheet and the primary current sheets of ICMEs, experience compression, which may lead to efficient particle acceleration (e.g., Khabarova et al. 2015). However, it is possible that the particles are bounded in space because they cannot escape easily the region filled with small-scale flux ropes, which implies more efficient acceleration. Furthermore, in N i MHD theory of quasi-2D magnetic island turbulence, incompressible flux ropes can be compressed by large-scale density and flow velocity gradients in the nonuniform solar wind (Zank et al. 2017; see also discussion of Equation (69) in Section 8.2). Closer to the Sun, Guidoni et al. (2016) discuss the possibility of strong plasma compression during magnetic island contraction for islands propagating sunward during a solar flare event.

2015ApJ...806..152S__Ransom_et_al._2005_Instance_1

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

2019ApJ...872..143B__Seckel_et_al._1991_Instance_2

The gamma-ray emission from the solar disk due to CR cascades in the solar atmosphere is denoted as a disk component. This secondary gamma-ray produced by the hadronic interaction of cosmic ray with the solar surface was first proposed by Dolan & Fazio (1965). While only upper limits were obtained by early measurements over the range 20 keV–10 MeV (Peterson et al. 1966). A detailed theoretical model for gamma-rays from the collision of cosmic ray with the solar atmosphere was presented by Seckel et al. (1991). The predicted gamma-ray flux at energies from 10 MeV to 10 GeV has a large uncertainty, being sensitive to the assumptions about the cosmic-ray transport in the magnetic field near the Sun. Gamma rays from the Sun were first detected by the Energetic Gamma-ray Experiment Telescope (Orlando & Strong 2008). The measured flux from 100 MeV to 2 GeV was within the range of the theoretical predictions. The Fermi collaboration (Abdo et al. 2011) reported the detection of high energy gamma-rays at 0.1–10 GeV from the quiescent Sun using the first 1.5 yr data. However, the measured solar disk emission flux was about a factor of seven higher than that predicted about this disk component by a “nominal” model (Seckel et al. 1991). This mismatch motivated Ng et al. (2016) to analyze 6 yr of public Fermi-LAT data. The obtained gamma-ray spectrum follows a simple power-law shape (α = −2.3) in 1–100 GeV without any evident high energy cutoff. For the flux in 1–10 GeV, a significant time variation of the solar disk gamma-ray flux that anticorrelates with solar activity was discovered, suggesting that the solar magnetic field would play an important role. An updated analysis with 9 yr of Fermi-LAT data, from 2008 August 7 to 2017 July 27, was performed, and Tang et al. (2018) confirmed these results and extended the gamma-ray spectrum up to >200 GeV. Notably, the bright gamma-ray flux above 100 GeV is dominant only during solar minimum at the end of Cycle 23 (Linden et al. 2018). The HAWC measurements in periods of high solar activity may support these findings (Albert et al. 2018a). Data collected from 2014 November to 2017 December, the second half of solar cycle 24, have been used to set strong upper limits on the flux of 1–100 TeV gamma-rays from the solar disk, about 10% of the maximum gamma-ray flux estimated by Linden et al. (2018). The HAWC 95% upper limit at 1 TeV is about 13% of the flux extrapolated from the solar minimum Fermi-LAT gamma-ray spectrum.

2019ApJ...887..118C__Joshi_et_al._2011_Instance_1

Generally, solar eruptions release their prestored magnetic energy via three phases, namely the precursor, impulsive, and gradual phases (Zhang et al. 2001; Zhou et al. 2016). Thereinto, the latter two, jointly termed as the main phase, correspond to the impulsive acceleration of the erupting CME flux rope, while the less-studied precursor phase includes key information on the eruptive structure and its trigger process. In the past, the limited observations showed that precursor activities of solar eruptions can come in various forms, in which preceding flux emergence (e.g., Palacios et al. 2015; Yan et al. 2017; Yang & Chen 2019) or cancellation (e.g., Green et al. 2011; Yang et al. 2016; Chen et al. 2018) in magnetograms, pre-eruption brightenings in Hα/EUV images (e.g., Bi et al. 2012; Dudík et al. 2016; Wang et al. 2017; Awasthi et al. 2018; Chen et al. 2018), as well as nonthermal processes in microwave or hard X-ray wavelength (Joshi et al. 2011; Altyntsev et al. 2012; Chen et al. 2017) are most common ones. In recent years, in the precursor phase of many CME/flare eruptive events, a type of new progenitor of CME flux ropes, namely hot channels (HCs, Cheng et al. 2011, 2013; Zhang et al. 2012), have been detected in the AIA high-temperature passbands (e.g., 131 and 94 Å). These newfound features commonly exist in the low corona with obvious helical/twisted fine fields or distinct writhed elbows (Cheng et al. 2013; Zhang et al. 2015), and they are found to be directly related to the occurrence of CMEs (Li & Zhang 2013; Patsourakos et al. 2013) and major eruptive flares (Nindos et al. 2015). To date, HCs have been evidenced as MFRs by some joint remote-sensing and in situ observations (e.g., Song et al. 2015), but for many HC eruption events, two problems are still elusive: (1) whether a corresponding MFR already exists before the HC eruption or is newly/partially built up during the eruption; and (2) whether its loss-of-equilibrium is initially facilitated by the preflare reconnection beneath/above the HC (Antiochos et al. 1999; Moore et al. 2001) or directly triggered the ideal magnetohydrodynamic (MHD) instabilities (Török et al. 2004; Kliem & Török 2006).

2015MNRAS.454.1468K__Winckel_2003_Instance_3

Owing to their dusty circumstellar environments, a large mid-infrared (mid-IR) excess is a characteristic feature of post-AGB stars and a detection of cold circumstellar material using mid-IR photometry can be used to identify these objects. The first extensive search for these objects was initiated in the mid-80's using results from the Infrared Astronomical Satellite (Neugebauer et al. 1984) which enabled the identification of post-AGB stars in our Galaxy (Kwok 1993). The Toru$\acute{\rm n}$ catalogue (Szczerba et al. 2007) for Galactic post-AGB stars lists around 391 very likely post-AGB objects. The Galactic sample of post-AGB stars have been found to be a very diverse group of objects (Van Winckel 2003). Studies showed that the majority of the optically visible Galactic post-AGB stars could be classified based on their spectral energy distributions (SEDs) into two groups: shell-sources and disc-sources (Van Winckel 2003). The shell-sources show a double-peaked SED with the hot central star peaking at shorter wavelengths while the cold, detached, expanding dust shell peaks at longer wavelengths. This type of SED is considered to be characteristic of objects that follow the single-star evolution scenario mentioned above. The disc-sources do not show two distinct flux peaks in the mid-IR but they do display a clear near-infrared (near-IR) excess indicating that circumstellar dust must be close to the central star, near sublimation temperature. It is now well established that this feature in the SED indicates the presence of a stable compact circumbinary disc, and therefore these sources are referred to as disc-sources (de Ruyter et al. 2006; Deroo et al. 2007; Gielen et al. 2011a; Hillen et al. 2013). The rotation of the disc was resolved with the ALMA array (Bujarrabal et al. 2013a) in one object and using single dish observations Bujarrabal et al. (2013b) confirmed that disc rotation is indeed widespread. Moreover, these disc-sources are confirmed to be binaries and show orbital periods between 100 and 2000 d (Van Winckel et al. 2009; Gorlova et al. 2014). In contrast, for the Galactic shell-sources long-term radial velocity monitoring efforts have not yet resulted in any clear detected binary orbit (Hrivnak et al. 2011), which either confirms the single-star nature of these objects or introduces a possibility that these systems can have companions on very wide orbits.

2019AandA...623A.140G__Dong_et_al._2018_Instance_1

Planet formation occurs in disks around young stellar objects. Interactions between planets and disks are very complex. Young planets are expected to cause rings, cavities, spirals, and disturbances in the velocity field and other features in the disk, which in turn may be used to infer the presence of these young planets. In the past few years, much evidence about this phase of planet formation has been accumulated because high-resolution images in the millimeter and sub-millimeter wavelength ranges have been provided by the Very Large Array (VLA) and the Atacama Large Millimeter Array (ALMA; see e.g. the case of HL Tau; ALMA Partnership 2015), and by high-contrast imagers such as the Gemini Planet Imager (GPI, Macintosh et al. 2014) and SPHERE (Spectro- Polarimetic High contrast imager for Exoplanets REsearch, Beuzit et al. 2008; see, e.g., Avenhaus et al. 2018). The literature on indirect evidence of the presence of planets is now becoming very rich, and nearby young stars surrounded by gas-rich disks are intensively studied for this purpose. In most cases, available data cannot fully eliminate alternative hypotheses, or the data have ambiguous interpretations (see, e.g., Bae et al. 2018 and Dong et al. 2018), although strong indirect evidence of the presence of planets from local disturbances of the velocity field have recently been considered for the case of HD 163296 (Pinte et al. 2018; Teague et al. 2018). In general, small grains are thought to be more strongly coupled with gas and are thus less sensitive to radial drift and concentration that can strongly affect large grains (see the discussion in Dipierro et al. 2018). For this reason, observations at short wavelengths provide an important complementary view of what can be seen with ALMA. On the other hand, a direct detection of still-forming planets embedded within primordial gas-rich disks, which is expected to be possible with high-contrast imaging in the near infrared (NIR), is still scarce; remarkable cases are LkCa-15 (Kraus & Ireland 2012; Sallum et al. 2015) and PDS-70 (Keppler et al. 2018; Müller et al. 2018; Wagner et al. 2018). In particular, in this second case, a clear detection of an accreting planet in the cavity between the inner and outer ring was obtained, making it an archetype for planet formation and planet-disk interactions. However, many cases remain ambiguous; a classical example is HD 100546 (see, e.g., Quanz et al. 2013a, 2015; Currie et al. 2014, 2015; Rameau et al. 2017; Sissa et al. 2018).

2019AandA...623A..16S__Olano_&_Poeppel_1987_Instance_1

Figure 10 (see also Figs. A.4 and A.5) compare the spatial distributions of the Hα and 857 GHz emission in the Taurus–California–Perseus region (e.g., Taurus, Auriga, California, and Perseus). The 857 GHz dust emission traces each molecular cloud and exhibits a hole-like structure. This hole-like structure can also be seen in HI emission, as shown in Figs. A.4 and A.5. The Hα emission fills the hole-like structure seen in the 857 GHz dust emission near the center of the field. The Taurus, California, and Perseus molecular complexes traced by the 857 GHz dust emission are distributed at the edge of the hole-likestructure. Lim et al. (2013) also found evidence of a shell-like structure using dust extinction and 12 CO (1–0) maps. The hole-like structure may result from the expansion of a large-scale supershell produced by a supernova in the Per OB2 association that compresses the Taurus cloud from the far side (Olano & Poeppel 1987; Bally et al. 2008). An Hα absorption feature is detected toward the Taurus cloud (see Figs. 10 and A.6), suggesting that the Taurus cloud lies at the front surface of the large-scale supershell produced by the Per OB2 association. The distance to the Per OB2 association is estimated to be 340 pc from the Sun (Cernis 1993), while the distance to the Taurus cloud is ~140 pc (Elias 1978). These distances are consistent with the Taurus cloud lying in front of the Per OB2 association. The B211/B213 filament also appears to be in front of the HI shell (see Fig. 10 in Chapman et al. 2011). This morphology suggests that the B211/B213 filament may have formed as a result of an expanding supershell. This may provide another reason for the different initial gas velocities for the northeastern and southwestern sheet components in addition to large-scale acceleration by the gravitational potential of the B211/B213 cloud (see Sect. 4.2.3). The Local Bubble surrounding the Sun might also compress the Taurus cloud from the opposite direction. The Local Bubble surrounding the Sun was produced by supernovae (Snowden et al. 1998; Sfeir et al. 1999), and the wall of the Local Bubble is located close to the Taurus cloud (Könyves et al. 2007; Lallement et al. 2014).

2019ApJ...883...73C__Ruffolo_et_al._2012_Instance_1

Assuming that the force-field approach to the solution of the Parker (1965) cosmic-ray transport equation is valid, the connection between historic cosmic-ray intensities and the solar properties they encountered lies in the effective diffusion coefficient that is assumed in this approximation. Establishing such a connection, however, is no simple task. Many theories have been proposed to describe the scattering of cosmic rays in the heliosphere. The most likely candidates for this task, given their reasonable agreement with observations and numerical simulations of cosmic-ray diffusion coefficients, are the quasilinear theory (QLT) of Jokipii (1966), the weakly nonlinear theory (WNLT) of Shalchi et al. (2004b), and the nonlinear guiding center theory (NLGC; or one of its variants; see, e.g., Matthaeus et al. 2003; Shalchi 2006, 2009, 2010; Ruffolo et al. 2012). Shalchi (2009) provides in depth theoretical treatments of most the abovementioned theories. These scattering theories all require as a key input an expression for the power spectrum of the turbulent fluctuations of the HMF. These spectra depend upon basic turbulence quantities, such as the magnetic variance, and various correlation scales. Turbulence power spectra are discussed in detail by, e.g., Batchelor (1970) and Matthaeus et al. (2007), whereas more background on the abovementioned turbulence quantities can be found in, e.g., Matthaeus & Goldstein (1982), Petrosyan et al. (2010), Matthaeus & Velli (2011), and Bruno & Carbone (2013). These basic turbulence quantities have been observed to show a marked dependence on the solar cycle at Earth (see, e.g., Smith et al. 2006b; Burger et al. 2014; Zhao et al. 2018). It follows then that mean free paths derived from these scattering theories would be expected to depend on the solar cycle as well, and several studies have reported such a dependence. Chen & Bieber (1993) find from an analysis of cosmic-ray anisotropies and gradients as observed by means of NMs, that larger mean free paths are associated with solar minima, and smaller mean free paths with solar maxima. The authors also report a mean free path dependence on solar magnetic polarity. Nel (2016) and Zhao et al. (2018) both extensively analyze spacecraft observations, using the turbulence quantities so calculated as inputs for expressions for diffusion coefficients derived from the QLT and NLGC theories. Both authors report that the resulting mean free paths display solar cycle dependences.

2018AandA...610A..44M__Krüger_&_Dreizler_(1992)_Instance_3

The first investigations of the rotational spectra of ethyl isocyanide were carried out in 1966 by Bolton et al. (1966). The spectra of the first vibrational and torsional excited states were measured in the centimeter wave domain (Anderson & Gwinn 1968). In this initial study, the dipole moment was determined to be μa = 3.79 D and μb = 1.31 D; this value is usually large for a molecule that includes a CN group. This causes dense and intense rotational spectra in the millimeter wave range and also in the submillimeter wave range up to 900 GHz (bQ lines). Anderson & Gwinn (1968) also observed some A–E splittings due to the internal rotation motion of the methyl group. The most recent spectroscopic study is from Krüger & Dreizler (1992) who reinvestigated the internal rotation measurements and also determined hyperfine coupling parameters due to the nitrogen quadrupole. As in our previous studies of ethyl cyanide isotopologs, it was not possible to observe internal rotation and hyperfine splittings due to our Doppler limited resolution. Our analysis was rather easy, starting from a prediction based on Krüger & Dreizler (1992) parameters. First, we analyzed and fit the most intense transitions, the aRh transitions, up to 330 GHz. These transitions were shifted only a few MHz from the initial predictions. Then bR and bQ lines were searched and included in the fit up to 330 GHz. Next, all the spectra were analyzed up to 990 GHz without difficulty. For the fitting, we employed ASFIT (Kisiel 2001) and predictions were made with SPCAT (Pickett 1991). The global fits included 6 transitions from Anderson & Gwinn (1968), 29 lines from Krüger & Dreizler (1992), and 2906 from this work. The maximal quantum numbers are J = 103 and Ka = 30. Both reductions A and S were tested. A reduction permits us to check theagreement of our new parameters set with those from Krüger & Dreizler (1992) (Table 1). Using S reduction slightly decreases root mean square from 30.3 to28.7 kHz. The condition numbers are nearly the same: 295 and 310 for the A and S reductions, respectively.The A reduction requires 23 parameters, but 5 additional parameters are required for the S reduction (Table 2). For this reason we used the A reduction even if this molecule is close to the prolate limit with kappa = −0.9521. Part of the new measurements are in Table 3. Owing to its large size, the complete version of the global fit Table S1 is supplied at the CDS. The fitting files .lin (S2), .par (S3), and the prediction .cat (S4) are also available at CDS.

2018ApJ...852L..20A__Prieto_et_al._2015_Instance_1

J0815+4729 is a main-sequence star (T eff = 6215 ± 82 K, log g = 4.7±0.5) with a metallicity of [Fe/H] ≤ −5.8 dex. Finding unevolved stars at this extremely low metallicity is very important since their stellar surface composition is not expected to be significantly modified by any internal mixing processes as in giant stars (Spite et al. 2005). J0815+4729 is similar to HE 1327–2326 in regard to its carbon enhancement, effective temperature, and metallicity. HE 1327–2326 is considered a turn-off/subgiant star, while J0815+4729 appears to be a dwarf. The ISIS spectrum of HE 1327–2326 indicates a metallicity of [Fe/H] ∼ −4.9 since the stellar Ca line is blended in that spectrum with the ISM features (Aguado et al. 2017b). However, the authors proposed a simple analysis taking into account the ISM absorption based on the UVES spectrum of HE 1327–2326. For J0815+4729, we require a high-resolution spectrum to clearly isolate the stellar Ca feature from possible additional ISM lines, and thus together with the detection of Fe lines, to establish the metallicity of this star. There are two other confirmed dwarf stars in this metallicity regime: one without any detectable carbon, J1029+1729 (Caffau et al. 2011), and another carbon-enhanced unevolved star, J1035+0641 (Bonifacio et al. 2015). The majority of extremely metal-poor stars shows overabundances of carbon, [C/Fe] > 0.7, and it appears that carbon-enhanced metal-poor (CEMP) stars split into two groups, with dramatically different carbon abundances (see, e.g., Beers & Christlieb 2005; Allende Prieto et al. 2015; Bonifacio et al. 2015 and references therein). The two carbon bands (high and low) studied have different origins. On the one hand, CEMP stars in the high-carbon band (A(C) ∼ 8.2) are probably produced by mass transfer from a binary companion, most likely an AGB star (Starkenburg et al. 2014). On the other hand, objects lying in the low-carbon band (A(C) ∼ 6.8) are thought to show the original carbon abundance inherited by the star from the ISM (Stancliffe 2009; Bonifacio et al. 2015; Abate et al. 2016). J0815+4729 has an abundance ratio of [ C / Fe ] ≥ + 5.0 dex corresponding to A(C) ∼ 7.7 dex (adopting [Fe/H] ≤ −5.8). In Figure 4 (bottom panel), we show the carbon abundance ratio [C/Fe] for all stars at [Fe/H] −4.5. All stars in this metallicity regime are considered to belong to the low-carbon band (Bonifacio et al. 2015), except for J0815+4729, which appears to be in between the low- and high-carbon bands. Both metallicity and carbon abundance are considered upper and lower limits, respectively. High-resolution spectra would be very useful to measure other elemental abundances and investigate the properties of the first supernovae. In particular, the barium abundance, or that of any other s-element, is not measurable from ISIS or OSIRIS spectra, and this is required to determine whether J0815+4729 is a CEMP-s, CEMP-r, or i-process star (Hampel et al. 2016). If we establish the abundance pattern, we will learn about the progenitor properties. Finally, the radial velocity accuracy from medium-resolution data is not enough to discard variations among different exposures, which would be indicative of binarity.

2022MNRAS.515.1086L__Naoz,_Farr_&_Rasio_2012_Instance_1

Regarding non-restricted hierarchical three-body systems, Krymolowski & Mazeh (1999) and Ford, Kozinsky & Rasio (2000) presented secular equations of motion (or Hamiltonian) up to the octupole order in semimajor axial ratio by using Hamiltonian perturbation techniques. Lee & Peale (2003) adopted both the octupole-level perturbation theory and direct numerical integrations to investigate the dynamical evolution for coplanar hierarchical planetary systems. In particular, the dynamics of apsidal resonance with critical argument of σ = ϖ1 − ϖ2 (ϖ1, 2 are the longitudes of pericentre) is studied and applied to some representative exoplanetary systems (Lee & Peale 2003). In a hierarchical planetary system with two comparable-mass planets orbiting a central star, Naoz et al. (2011) showed that orbits of the inner planet could flip from prograde to retrograde and back again due to the secular planet–planet interaction. Based on this behaviour, it becomes possible to form hot Jupiters on retrograde orbits by combining the eccentric ZLK effect and tidal friction (Naoz et al. 2011; Naoz, Farr & Rasio 2012; Teyssandier et al. 2013; Petrovich 2015; Petrovich & Tremaine 2016; Dawson & Johnson 2018). Naoz et al. (2013) re-derived the secular evolution equations for hierarchical three-body systems at the octupole-level approximation and found that orbital flips of inner planet are possible even at the quadrupole-level approximation. They pointed out that the relation h1 − h2 = π can be used to simplify the expression of Hamiltonian but the evolutions of H1 and H2 should be derived from the conservation of the total angular momentum rather than from the Hamiltonian canonical relations. Tan et al. (2020) explored the secular resonances with critical arguments arising in the Hamiltonian under the resonant Hamiltonian model, which is obtained by directly removing those terms involving short-period angles from the octupole-level Hamiltonian (i.e. only the secular and resonant terms are retained). It is of no problem when dealing with the quadrupole-order resonance (the ZLK resonance) because in this case the omitting terms are of octupole order. However, it may be inadequate to formulate the resonant model by directly removing those quadrupole-order periodic terms from the Hamiltonian when studying the octupole-order resonances. Hamers (2021) performed a semianalytic study about the properties of the ZLK oscillations at the quadrupole-level approximation, including the maximum eccentricities, time-scales of eccentricity/inclination oscillation and orbit flips. Naoz (2016) and Shevchenko (2016) reviewed various applications of the eccentric ZLK effect to a broad range of astrophysical systems, such as planetary and exoplanetary systems, stellar systems, and galaxies.

2018ApJ...854...26L___2015a_Instance_1

The hot emission line of Fe xxi 1354.09 Å and the cool emission line of Si iv 1402.77 Å have been used in many spectroscopic studies to investigate chromospheric evaporation (e.g., Tian et al. 2014, 2015; Li et al. 2015b, 2017a, 2017b; Brosius et al. 2016; Zhang et al. 2016a, 2016b). It is widely accepted that the forbidden line of Fe xxi 1354.09 Å is a hot (log T ∼ 7.05) and broad emission line during solar flares (Doschek et al. 1975; Cheng et al. 1979; Mason et al. 1986; Innes et al. 2003a, 2003b). Meanwhile, IRIS spectroscopic observations show that Fe xxi 1354.09 Å is always blended with a number of cool and narrow emission lines, which are from neutral or singly ionized species. Those blended emission lines can be easily detected at the position of the flare ribbon, including known and unknown emission lines, such as the C i line at 1354.29 Å, the Fe ii lines at 1353.02 Å, 1354.01 Å, and 1354.75 Å, the Si ii lines at 1352.64 Å and 1353.72 Å, and the unidentified lines at 1353.32 Å and 1353.39 Å (e.g., Li et al. 2015a, 2016a; Polito et al. 2015, 2016; Tian et al. 2015, 2016; Young et al. 2015; Tian 2017). In order to extract the hot line of Fe xxi 1354.09 Å and the cool line of C i 1354.29 Å (log T ∼ 4.0; Huang et al. 2014), we apply a multi-Gaussian function superimposed on a linear background to fit the IRIS spectrum at the “O i” window (e.g., Li et al. 2015a, 2016a), which has been pre-processed (i.e., IRIS spectral image deformation, bad pixel despiking and wavelength calibration) with the standard routines in Solar Soft Ware (SSW; Freeland et al. 2000). In short, the line positions and widths of these blended emission lines are fixed or constrained, and their peak intensities are tied to isolated emission lines from similar species. More details can be found in our previous papers (Li et al. 2015a, 2016a). On the other hand, the cool line of Si iv 1402.77 Å (log T ∼ 4.8) at the “Si iv” window is relatively isolated, and it can be well fitted with a single-Gaussian function superimposed on a linear background (Li et al. 2014, 2017a). Using the relatively strong neutral lines (i.e., “O i” 1355.60 Å and “S i” 1401.51 Å), we also perform an absolute wavelength calibration for the spectra at the “O i” and “Si iv” windows, respectively (Tian et al. 2015; Tian 2017). Finally, the Doppler velocities of Fe xxi 1354.09 Å, C i 1354.29 Å, and Si iv 1402.77 Å are determined by fitting line centers removed from their rest wavelengths (Cheng & Ding 2016b; Guo et al. 2017; Li et al. 2017a). As the hot Fe xxi line is absent in the non-flaring spectrum, the rest wavelength for the Fe xxi line (i.e., 1354.09 Å) is determined by averaging the line centers of the Fe XXI profiles which were used in the previous IRIS observations (Brosius & Daw 2015; Polito et al. 2015, 2016; Sadykov et al. 2015; Tian et al. 2015; Young et al. 2015; Brosius et al. 2016; Lee et al. 2017), while the rest wavelengths for the C i and Si iv lines, i.e., 1354.29 Å and 1402.77 Å, respectively, are determined from their quiet-Sun spectra (Li et al. 2014, 2015a).

2022ApJ...925...62K__Cohen_et_al._2020_Instance_1

We noticed that the fast wind stream had He++ beams moving ahead of the protons, low number densities, and low heavy-ion charge state ratios, which are characteristic of coronal-hole-originated solar wind. It should be noted that although the solar wind stream originates from the same coronal hole, the speed and the stream arrival time to 1 au varied from one Carrington rotation to another as was seen in Figure 2. The timing of the coronal-hole-originated solar wind observation aligned well with the PSP orbit. PSP started passing through the leading edge of the high-speed stream, continued as the high-speed stream corotates over the spacecraft, and ended passing the trailing edge of the stream as seen in Figure 3, standing as a perfect classic CIR/SIR event (Cohen et al. 2020). The decrease in crossing width and steepening of the increasing velocity profile at L1 versus at PSP is unlikely to be a product of the evolution of the coronal hole and is most likely a product of the faster wind stream overtaking the slower plasma at the interface during propagation (Burlaga 1974). Analyzing the SIR observed at ACE on 2018 November 4 to the corresponding SIR observed at PSP around 2018 November 15 indicated the enhancement of low-energy (keV) suprathermal ions (Figure 6) at the stream interface as the keV range ions are enhanced by local acceleration. The keV ions were enhanced again after the interface region, but for a short period of time compared to the PSP observation (Figure 8), while high-energy (MeV) suprathermal ions propagate along field lines to the inner heliosphere (Filwett et al. 2019) and are enhanced after the interface region in which articles accelerated at distant shocks dominated (see Desai et al. 2020; Joyce et al. 2021) and lasted longer in time in PSP. Minor enhancements in MeV particles were inconsistently observed over various Carrington rotations at ACE (not shown), but these were often at or just above the instrument noise floor. We cannot make any conclusive inferences as to the evolution of higher-energy ions in the coronal wind from these observations. In the future, as we approach solar maximum, it is likely that there will be more equatorial coronal holes observed during upcoming PSP orbits, which will allow us to compare more in depth the evolution of the solar wind structure at 1 au with PSP.

2015AandA...580A...5L__Voit_1991_Instance_1

The solutions in Sect. 3 show that it is possible for a high X-ray flux in galactic nuclei to alter the carbon ionization balance and reduce the C+ abundance and correspondingly the [C ii] luminosity. In the presence of a high flux of soft X-rays above 1 keV, a condition encountered in many AGNs (Stacey et al. 2010; Ebrero et al. 2009), the abundance of singly ionized carbon is reduced and converted into higher ionization states. This reduction occurs primarily in the hot highly ionized gas that fills most of the galactic central zone, to some degree in the dense ionized skins surrounding molecular clouds, less so in diffuse atomic hydrogen clouds, and very little, if at all, in the dense PDRs at the edge of the CO molecular cores. Furthermore, in galactic nuclei with a high X-ray flux we would expect to see an increase in the dust temperature and infrared luminosity (Voit 1991). Therefore in galactic nuclei, and in particular in AGNs, we expect a reduction in [C ii] emission relative to FIR/IR emission depending on the relative contribution of different ISM components to the [C ii] luminosity. Unfortunately the next most abundant ion, C2+, does not have fine structure FIR emission lines as its 2s2 ground state has spin zero, and C3+, although it has spin angular momentum due to its unpaired electron in the 2S level, does not have a nuclear spin or orbital angular momentum to break the degeneracy of the two electron spin states. Instead to test the effect of X-ray ionization on the carbon balance we would need studies of their UV emission. For example the [C iv] UV resonance lines are detected in extragalactic sources and have been used to trace the formation rate of massive stars (Leitherer & Lamers 1991; Robert et al. 1993). [C iv] UV absorption lines have also been used to study the properties of the Galactic Halo (Savage et al. 2000) where carbon is presumed to be collisionally ionized in hot coronal gas (Gnat & Sternberg 2007).

2022AandA...663A.172M___2012_Instance_1

We note that Pavesi et al. (2019) derived a lower κs parameter log(κs) ≈ −1 for HZ10, by observing the CO(2-1) line, which implies a low star formation efficiency for this source. The conflict between the two results can be explained by the fact that Pavesi et al. (2019) estimated the gas mass by adopting a large CO-to-Mgas conversion factor αCO = 4.5 M⊙ (K km s−1 pc2)−1, a value that is close to the Galactic conversion factor αCO = 4.36 M⊙ (K km s−1 pc2)−1 (Bolatto et al. 2013). Although the Galactic conversion factor is a derived value for Milky Way and normal, star-forming galaxies in the local Universe, it may not be applicable for more extreme environments of starburst galaxies at high-z (see Carilli & Walter 2013 for a review). The conversion factor depends on the physical conditions of the gas in the ISM (temperature, surface density, dynamics, and metallicity), as well as the star formation and associated feedback (Narayanan et al. 2011, 2012; Genzel et al. 2012; Feldmann et al. 2012; Renaud et al. 2019; see, e.g., Bolatto et al. 2013 for a review). It is typically in the range between 0.8 and 4.36 M⊙ (K km s−1 pc2)−1 (see, e.g., Bolatto et al. 2013; Carilli & Walter 2013, and Combes 2018 for reviews). Low metallicities (Z = 0.6 Z⊙ for HZ10) will drive αCO towards values higher than the Galactic value (Narayanan et al. 2012; Genzel et al. 2012; Popping et al. 2014), although αCO spans a broad range of values of αCO ∼ 0.4 − 11 M⊙ (K km s−1 pc2)−1, due to large uncertainties (see, e.g., Fig. 9 of Bolatto et al. 2013). On the other hand, high values of temperature, surface density, and velocity dispersion in a turbulent ISM of starbursts and merging systems will shift αCO towards lower values (Narayanan et al. 2011, 2012; Vallini et al. 2018). HZ10 has an extremely high value of the burstiness parameter log(κs) ∼ 1.4 and high total density of the [C II] emitting gas log(n) ∼ 3.35 cm−3, and it is also a multi-component system (Jones et al. 2017, Carniani et al. 2018a). Thus, for this source we assumed αCO = 0.8 M⊙ (K km s−1 pc2)−1, usually adopted for starburst galaxies (e.g., Downes & Solomon 1998; Bolatto et al. 2013). We obtained log(κs) = 0.53 ± 0.34, which is within the 2σ uncertainties of the log ( κ s ) = 1 . 43 − 0.53 + 0.38 $ \log{(\kappa_s)} = 1.43_{-0.53}^{+0.38} $ , estimated exploiting the C III] emission and the Vallini et al. (2020) model.

2019ApJ...875...61M__Jones_&_Boffin_2017_Instance_1

A substantial fraction of metal-poor stars that have recently evolved off the MS, e.g., giants and planetary nebulae (PNe), have been influenced by binary interactions. The IMF is significantly weighted toward low-mass stars (Bastian et al. 2010; Kroupa et al. 2013), and the MW star formation rate was ≈3 times larger ≈10 Gyr ago than it is now (Governato et al. 2007; De Lucia et al. 2014). Based on the measured IMF and modeled galactic star formation history, we estimate that ≈55% of MW giants and PNe have old, solar-type progenitors (τ* > 7 Gyr, M ≈ 0.8–1.2 ). Such old, low-mass giants tend to be metal-poor (Ratnatunga & Yoss 1991; Carollo et al. 2010; Mackereth et al. 2017). The metallicity trend therefore dramatically affects the properties of low-mass evolved stars. For example, the enhanced close binary fraction of metal-poor solar-type stars substantially strengthens the conclusion that the shaping of PN morphologies is the result of binary interactions (Moe & De Marco 2006; De Marco 2009; Jones & Boffin 2017). Providing further corroboration, Badenes et al. (2015) measured the delay-time distribution of bright PNe in the LMC and discovered two distinct populations of PN progenitors: an old channel (τ* = 5–8 Gyr) deriving from solar-type stars (M ≈ 1.0–1.2 ) and a young channel (35–800 Myr) evolving from late-B/early-A stars (≈2–8 ). According to the measured age–metallicity relation of the LMC (Olszewski et al. 1991; Pagel & Tautvaisiene 1998; Cole et al. 2005; Carrera et al. 2011; Piatti & Geisler 2013), the old, solar-type progenitors are metal-poor ([Fe/H] ≲ −1.0) and hence have a large close binary fraction of Fclose = 40%–50%. The young progenitors have a higher metallicity of [Fe/H] ≈ −0.4 but are sufficiently massive so that they also have a large close binary fraction of Fclose = 40%–60%. Meanwhile, evolved stars with intermediate masses (M ≈ 1.2–2.0 ) in the LMC have intermediate metallicities and therefore a smaller close binary fraction of Fclose ≈ 30%. If PNe derive from interactions in close binaries, then the variations in Fclose with respect to mass and metallicity can explain the observed bimodal mass/age distribution of PN progenitors in the LMC.

2019ApJ...875L..31H__Leary_et_al._2006_Instance_1

The recent detection of gravitational-wave (GW) emission from a merging neutron star binary (Abbott et al. 2017d) and merging black hole binaries (BHBs; Abbott et al. 2016a, 2016b, 2017a, 2017b, 2017c; The LIGO Scientific Collaboration & The Virgo Collaboration 2018) by the Laser Interferometer Gravitational-Wave Observatory (LIGO)/Virgo have ushered in an exciting new era of GW astrophysics. The astrophysical origin of the detected mergers is currently under debate, with numerous explanations proposed. These explanations can be very roughly divided into two main categories: mergers due to isolated binary evolution (e.g., Belczynski et al. 2016; de Mink & Mandel 2016; Mandel & de Mink 2016; Marchant et al. 2016), and mergers due to dynamical interactions (e.g., Portegies Zwart & McMillan 2000; Wen 2003; O’Leary et al. 2006, 2009, 2016; Antonini & Perets 2012; Kocsis & Levin 2012; Antonini et al. 2014; Antonini & Rasio 2016; Rodriguez et al. 2016; VanLandingham et al. 2016; Askar et al. 2017; Arca-Sedda & Gualandris 2018; Fragione & Kocsis 2018; Hoang et al. 2018; Randall & Xianyu 2018; Arca-Sedda & Capuzzo-Dolcetta 2019). Orbital eccentricity has been explored as a way to distinguish between these merger channels in both the LIGO/Virgo and Laser Interferometer Space Antenna (LISA) frequency bands. In contrast to mergers from isolated binary evolution, merging binaries from dynamical channels have been shown to have measurable eccentricities when they enter the LISA and/or LIGO/Virgo band, and can potentially be used as a way to distinguish between channels (e.g., O’Leary et al. 2009; Cholis et al. 2016; Gondán et al. 2018; Lower et al. 2018; Randall & Xianyu 2018; Rodriguez et al. 2018; Samsing 2018; Zevin et al. 2019). Unlike LIGO/Virgo, which can only detect merging BHBs in the final inspiral phase before merger, LISA will be able to detect eccentric stellar-mass BHBs for long timescales before they merge in the LIGO/Virgo band (e.g., O’Leary et al. 2006; Breivik et al. 2016; Nishizawa et al. 2016; Chen & Amaro-Seoane 2017; Nishizawa et al. 2017; D’Orazio & Samsing 2018; Kremer et al. 2019; Samsing & D’Orazio 2018). This provides us with invaluable insight into the dynamical evolution of eccentric binaries leading up to the merger, which has important implications about the astrophysical context in which merging binaries evolve.

2018MNRAS.479.4509R__Kingma_&_Ba_2014_Instance_2

After each step of calculations, the network should optimize the model based on its current and previous states to improve the subsequent mapping. Our model utilizes a computationally memory efficient optimization due to its dependence to only the first-order gradients, namely the ‘adaptive moment estimation’ (or Adam). For more details, we refer the readers to Kingma & Ba (2014). Adam optimization, compared to other gradient-based optimization, is very suitable for noisy and sparse gradients, and for simulated data that show very large scatter with respect to a given quantity of parameter (Kingma & Ba 2014). With this optimizer, we have to decide few parameters in advance. The learning step α and the parameters controlling the moving averages of the first- and second-order moments, namely β1 and β2 (both ∈[0,1)), respectively. For this purpose, we chose to minimize the MSE between the target and the prediction from the model: in what follows, we will alternatively call the MSE the ‘objective function’ f($\bf x$): with ${\bf x}$ the parameters of the model to be updated, such as weights and biases. At a given time t ≤ T, where T is the maximal learning time-step, we can update the parameters of the model as shown in the following: (11) \begin{eqnarray*} g_t &amp;=\nabla _\mathrm{ \text{$x$}} f(\mathrm{\text{$x$}}_{t-1}), \end{eqnarray*} (12) \begin{eqnarray*} \mu _{1,t} &amp;=\beta _1 \times \mu _{1,t-1} + (1-\beta _1)\times g_t, \end{eqnarray*} (13) \begin{eqnarray*} \bar{\mu }_{1,t} &amp;=\mu _{1,t}/(1-\beta _1^t), \end{eqnarray*} (14) \begin{eqnarray*} \mu _{2,t} &amp;=\beta _2 \times \mu _{2,t-1} + (1-\beta _2)\times g_t^2, \end{eqnarray*} (15) \begin{eqnarray*} \bar{\mu }_{2,t} &amp;=\mu _{2,t}/(1-\beta _2^t), \end{eqnarray*} (16) \begin{eqnarray*} \mathrm{\text{$x$}}_t &amp;=\mathrm{\text{$x$}}_{t-1} - \alpha _t \times \bar{\mu }_{1,t}/ (\sqrt{\bar{\mu }_{2,t}} + \epsilon), \end{eqnarray*} where $\alpha _t=\alpha \sqrt{1-\beta _2^t}/(1-\beta _1^t)$ is the time-step at t. Equation (11) shows the gradients of the objective function at t with respect to the model parameters. Equations (12) and (14) update the estimations of the first and second moments. Our moments are biased towards the initial values; thus, we require equations (13) and (15) to account for the corrections. Finally, we update the model parameters with equation (16).

2019ApJ...882..168P__Hildebrand_1983_Instance_1

In order to provide constraints on the gas masses in these galaxies independently from the CO measurements, we can use the Rayleigh–Jeans dust continuum emission. This will provide the first constraints to the αCO conversion factor in “normal” galaxies at z > 3 in the following. The Rayleigh–Jeans dust continuum emission has been used to estimate dust and gas masses, assuming an average emissivity and dust temperature for the dominant cold dust component and a constant dust-to-gas ratio (Hildebrand 1983; Eales et al. 2012; Bourne et al. 2013; Scoville 2013; Scoville et al. 2013, 2016, 2017; Groves et al. 2015). The dependence on cold dust temperature and dust-to-gas ratio may make the Rayleigh–Jeans method less reliable than at lower redshifts (e.g., Pavesi et al. 2018a). On the other hand, the opposing effects of increasing dust temperatures and decreasing dust-to-gas ratios that may occur in “normal” galaxies at high redshift may partially compensate for each other, as also found in recent simulations that are consistent with this approach to gas mass measurement (e.g., Liang et al. 2018; Privon et al. 2018). We here adopt Equations (10) and (13) of Scoville et al. (2016) to derive gas mass estimates based on our continuum flux measurements through the same assumptions that were used in those lower-redshift samples (Scoville et al. 2016, 2017). The 34 GHz upper limits imply 3σ gas mass limits of 2.8 × 1011 M⊙ for HZ10 and 1.6 × 1011 M⊙ for LBG-1, adopting the relation derived by Scoville et al. (2016, 2017). We also use the ∼230 GHz continuum fluxes to derive approximate estimates, although these measurements may not lie on the Rayleigh–Jeans tail and therefore may not accurately trace the cold dust component. These continuum measurements would imply gas masses of ∼1.3 × 1010 M⊙ for HZ4, ∼2.5 × 1010 M⊙ for LBG-1, ∼4.4 × 1010 M⊙ for HZ9, and ∼1.1 × 1011 M⊙ for HZ10, with dominant systematic uncertainties due to the extrapolation of the method to very high redshift.

2022ApJ...935..135B__Mathur_1990_Instance_2

All responses calculated in this paper only account for the direct response to a perturbing potential. In general, though, the response also has an indirect component that arises from the fact that neighboring regions in the disk interact with each other gravitationally. This self-gravity of the response, which we have ignored, triggers long-lived normal-mode oscillations of the slab that are not accounted for in our treatment. Several simulation-based studies have argued that including self-gravity is important for a realistic treatment of phase spirals (e.g., Darling & Widrow 2019a; Bennett & Bovy 2021). Using the Kalnajs matrix method (Kalnajs 1977; Binney & Tremaine 2008), we have made some initial attempts to include the self-gravity of the response in our perturbative analysis, along the lines of Weinberg (1991). Our preliminary analysis shows that the self-gravitating response is a linear superposition of two terms: (i) a continuum of modes given in Equation (12), dressed by self-gravity, that undergo phase mixing and give rise to the phase spiral; and (ii) a discrete set of modes called point modes or normal modes (see Mathur 1990; Weinberg 1991) that follow a dispersion relation. The continuum response can be amplified by self-gravity when the continuum frequencies, nΩ z + kv x , are close to the point-mode frequencies, ν. Depending on the value of k, the normal modes can be either stable or unstable. Araki (1985) finds that in an isothermal slab the bending normal mode undergoes fire hose instability below a certain critical wavelength if σ z /σ ≲ 0.3, while the breathing normal mode becomes unstable above the Jeans scale. In the regime of stability, the normal modes are undamped oscillation modes in absence of lateral streaming (Mathur 1990) but are Landau damped otherwise (Weinberg 1991). For an isothermal slab with typical MW-like parameter values, the point modes are strongly damped since their damping timescale (inverse of the imaginary part of ν) is of order their oscillation period (inverse of the real part of ν), which turns out to be of order the vertical dynamical time, h z /σ z . Moreover, the normal-mode oscillations are coherent oscillations of the entire system, independent of the vertical actions of the stars, and are decoupled from the phase spiral in linear theory since the full response is a linear superposition of the two. Based on the above arguments, we conclude that self-gravity has little impact on the evolution of phase spirals in the isothermal slab, at least in the linear regime. We emphasize that Darling & Widrow (2019a), who found their phase spirals to be significantly affected by the inclusion of self-gravity, assumed a perturber-induced velocity impulse with magnitude comparable to the local velocity dispersion in the solar neighborhood; hence, their results are likely to have been impacted by nonlinear effects. Moreover, the self-gravitating response of an inhomogeneous disk embedded in a DM halo, as in the simulations of Darling & Widrow (2019a), can be substantially different from that of the isothermal slab. We intend to include a formal treatment of self-gravity along the lines of Weinberg (1991) in future work.

2016ApJ...821...19F__Fulle_et_al._2010_Instance_1

GIADA characterizes individual dust particles by means of two independent sensors. At the instrument entrance the particle crosses a laser curtain, and is detected by photoelectric sensors (GDS, grain detection system) registering a signal (proportional to the particle cross section times the albedo) and the time at which the laser curtain is crossed. Then the particle hits the impact sensor (IS, with the same GDS cross section, A = 10−2 m2), which registers the individual particle impact momentum and its travel time from GDS to IS. The combination of GDS and IS measurements (GDS+IS particles) provides the particle mass and velocity, and constrains the particle bulk density by means of calibration curves (Della Corte et al. 2016) derived on the ground using cometary analogues (Ferrari et al. 2014). If the particle is too small to be detected by the GDS system, it may be detected by the IS sensor only (IS particles): in this case the particle momentum is converted to the mass assuming the mean value of the velocities of the GDS+IS particles in the same momentum bin, or assuming the velocities predicted by tail models (Fulle et al. 2010) if Ngds+is = 0 in that mass bin. The spacecraft velocities listed in Table 1 are always much lower than the dust velocities measured by GIADA. In this condition, in the Sun-facing coma (assumed to have uniform and R-dependent space density ρ), the dust flux from the nucleus surface corresponds to the dust flux at nadir-pointing GIADA scaled by the factor 2πR2/A. The dust number loss rate at the nucleus surface per GIADA detection is Qn = 2πR2(AΔt)−1, where Δt is the total dust collection time (Tables 5–8). In the same Tables, we show the mass loss rates Qm and the mean dust velocities already integrated in each mass bin, corresponding to the four GIADA collection periods considered in this paper: from 2015 February 19 to 28 (Table 5), from 2015 March 13 to 17 (Table 6), on 2015 March 28 (Table 7), and from 2015 August 23 to September 3 (Table 8). In Table 9 we show the data obtained during the first post-perihelion excursion at low phase angles (60 α 64°, 125 R 290 km). The R-values are too small to use the NAC DUST-MON sequences. The uncertainty affecting Afρ and the loss rates measured by GIADA and OSIRIS depends on the number of detections in each mass bin: an estimate of the relative error is given by N p − 1 / 2 and by ( N gds + is + N is ) − 1 / 2 . The dispersion of the dust velocities in Tables 2–4 provides the error affecting the dust velocities measured by OSIRIS, close to 30%. The relative error of the dust velocities provided by each GDS+IS detection is below 10%.

2022MNRAS.512..186K__Dutta_&_Bharadwaj_2013_Instance_1

A widely used statistical property of the sky brightness distribution is its power spectrum (Lazarian 1995; Bharadwaj & Sethi 2001, and others). As the redshifted 21-cm signal is expected to be faint and hard to detect with imaging, estimating its power spectrum or equivalently intensity mapping gives a possible probe of the evolution of the baryonic matter distribution over cosmic time. Bharadwaj & Sethi (2001) show that visibility correlation directly measures the power spectrum. This method and its variants (Datta, Choudhury & Bharadwaj 2007; Choudhuri et al. 2014; Choudhuri et al. 2016; Bharadwaj et al. 2019; Choudhuri et al.2019, and others) have been used to estimate the angular power spectrum of the diffused galactic foreground (Ghosh et al. 2012; Choudhuri et al. 2017b; Chakraborty et al. 2019a; Choudhuri et al. 2020) as well as the power spectrum of H i distribution in nearby galaxies (Dutta et al. 2009; Dutta & Bharadwaj 2013; Nandakumar & Dutta 2020). These works propagate the uncertainties in each visibility estimate and combine that with the sample variance error in measuring the power spectrum to quote uncertainties in the power spectrum estimates. In this work, we use the estimator discussed in Choudhuri et al. (2014), where visibilities are gridded before estimating the power spectrum. Given an angular field of view of θ0 to which the telescope is sensitive, it has been shown (Bharadwaj & Sethi 2001; Bharadwaj & Ali 2005; Choudhuri et al. 2014) that the visibilities in the nearby baselines remain correlated to a baseline separation of $\Delta U \lt \frac{1}{\pi \theta _0}$. The size of the uv-grids is chosen such that they are large enough to include a sufficient number of baselines in a given uv-grid and small enough to have all visibilities in the uv-grid correlated. In each uv-grid, they estimate the power spectrum by correlating visibilities only in nearby baselines, omitting the visibility autocorrelations. This drastically reduces the noise bias in estimates of the power spectrum in uv-grids. The contribution from each uv-grid within a given annulus in $U = \mid \vec{U} \mid$ is then combined, and the real part of it is used to quote the value of the isotropic power spectrum for the baseline separation U. We may schematically write it as (4)$$\begin{eqnarray} \mathcal {E} \lbrace P(U)\rbrace = \mathcal {R} [\langle \tilde{V}(\vec{U})^{*} \tilde{V}(\vec{U}+\Delta \vec{U}) \rangle]. \end{eqnarray}$$Here, the average is taken over the uv-grid first and then within the annulus, as explained above. Note that the power spectrum estimator here assumes that a perfect calibration is done and the gains are all unity. In such a case, the power spectrum estimate has no bias arising from instrumental noise, and its uncertainties can be written as (Ali et al. 2008; Dutta 2011) (5)$$\begin{eqnarray} \sigma _P^2 = \frac{P^2(U)}{N_\mathrm{ G}} + 2\frac{P(U)\sigma _N^2}{N_\mathrm{ B}} + 2\frac{\sigma _N^4}{N_\mathrm{ B}}, \end{eqnarray}$$where NG is the number of independent estimates of the power spectrum in a given annulus bin at U, NB is the total number of visibility pairs in the bin.

2022ApJ...929...19G__Yu_et_al._2019_Instance_1

From the theoretical point of view, we normalize the disk size to a fixed black hole mass and luminosity (thus to a fixed accretion rate; see Equations (14) and 17); therefore, we expect that the disk size τ 0 is independent of AGN properties. However, there are two possibilities that may cause the observed dependence in Figure 12. First, we adopt a constant virial factor f BLR when estimating the black hole mass using Equation (1). Dynamical modeling of broad-line regions (BLRs) generally showed that f BLR might depend on AGN properties and change from object to object (e.g., Pancoast et al. 2011; Li et al. 2013; Pancoast et al. 2014; Grier et al. 2017; Li et al. 2018; Williams et al. 2018). Observational calibrations of the virial factor also tend to support this conclusion (e.g., Ho & Kim 2014; Mejía-Restrepo et al. 2018; Yu et al. 2019). If there are systematic correlations between f BLR and luminosity L 5100, black hole mass M •, or accretion rate Ṁ , using a constant f BLR will lead to the apparent dependence of τ 0 on AGN properties. Such a bias can be eliminated once we have a solid understanding of the virial factor in the future. Second, it is possible that the disk’s temperature profile does change with the accretion rate. For example, the standard and slim-disk models predict different temperature profiles 12 12 However, we note that the temperature profiles of slim disks (T ∝ R −1/2) become different from those of standard disks (T ∝ R −3/4) only within the photon-trapping radius (e.g., Wang & Zhou 1999). (Shakura & Sunyaev 1973; Abramowicz et al. 1988). There are also other physical processes that may contribute to the correlation between β and Ṁ (see, e.g., Li et al. 2021; Kammoun et al. 2021). Figure 7 implies that τ 0 and β are highly anticorrelated. As a result, a shallower temperature profile (larger β) will lead to a smaller disk size τ 0 and vice versa. In our sample, we find that high-luminosity and high-mass black holes generally have low accretion rates, giving rise to the dependence of τ 0 on luminosity and black hole mass.

2022ApJ...939...43E__Protheroe_1999_Instance_1

The multimessenger observations of NGC 1068 dictate the need for multiple emission zones. In this work, we will capture its inner microparsecs referring to a spherically symmetric structure for the corona of the AGN as well as an outer starburst ring with a radius of ∼1 kpc (see Figure 1). Note, that in between these two emission sites NGC 1068 shows strong indications of a jet structure on scales of up to about 1 kpc (e.g., Wilson & Ulvestad 1982; Gallimore et al. 2004, 2006), which however is not included in this work. Due to mathematical convenience we treat both spatial regions as homogeneous. For particle acceleration processes that take place on considerably shorter timescales than the energy loss in these zones, we can disentangle these processes and only describe the steady-state transport of nonthermal, accelerated electrons and protons. Hereby, we suppose that in both zones some acceleration mechanism yields a differential source rate q(T) of relativistic protons and primary electrons that can be described by a power-law distribution in momentum space up to a certain maximal kinetic energy Tˆ , which depends on the competing energy-loss timescales in these zones. In the case of the starburst zone, we suppose that a certain fraction fSN of the total energy of the SN that releases about 1051 erg and occurs with an approximate rate 5 5 Supposing a SN rate νSN≃0.02[SFR/(1M⊙/yr)]yr−1 (note that Condon1992 suggested a value of 0.04 instead of 0.02 for normal galaxies) where the star formation rate SFR ≃ 17[L IR/(1011 L ⊙)] M ⊙yr−1. (Veilleux et al. 2005) νSN≃0.34[LIR/(1011L⊙)]yr−1 dependent on the IR luminosity L IR gets accelerated into CRs according to diffusive shock acceleration (DSA; e.g., Drury 1983; Protheroe 1999) by individual supernova remnants (SNR; e.g., Bell 2014, and references therein). In general, many starburst galaxies—NGC 253 is a prominent example—show a galactic superwind (e.g., Veilleux et al. 2005, and references therein) as a result of a large number of core-collapse SNe. These winds introduce another source of acceleration 6 6 Note that in these phenomena also stochastic diffuse acceleration may become relevant due to the presence of a turbulent plasma within the wind bubbles. (e.g., Anchordoqui et al. 1999; Romero et al. 2018), however, we are not aware of any observational indications of such a superwind in the starburst ring of NGC 1068. For the AGN corona, we suppose that a fraction f inj ≪ 1 of the mass accretion rate Ṁ=Lbol/(ηradc2) , with a radiation efficiency of η rad = 0.1 (Kato et al. 2008), goes into relativistic protons via stochastic diffuse acceleration (SDA; e.g., Lemoine & Malkov 2020, and references therein). For both zones, the nonthermal primary electrons are normalized by the nonthermal proton rates due to the requested quasi-neutral total charge number of the injection spectra of primary CRs above a characteristic kinetic energy of Tˇ≃10keV (Schlickeiser 2002; Eichmann & Becker Tjus2016; Merten et al. 2017). Note that this corresponds to a quasi-neutral acceleration site; however, CR transport can subsequently remove CR electrons and protons in different amounts from the nonthermal energy regime. Nevertheless their charge stays conserved. Transforming the source rates from momentum space into kinetic energy T, we obtain 1 qp(T)≡dNdVdTdt=qp,0Tˇ2+2TˇEp,0×T+Ep,0T2+2TEp,0T2+2TEp,0Tˇ2+2TˇEp,0−s/2exp[−T/Tˆ], for the injected nonthermal protons, and 2 qe±(T)=qe−,0Tˇ2+2TˇEe,0T+Ee,0T2+2TEe,0×T2+2TEe,0Tˇ2+2TˇEe,0−s/2exp[−T/Tˆ]+qe±2nd(T), for the nonthermal electrons (e −) and positrons (e +). Here, the latter term qe±2nd(T) introduces the source rate of secondary electrons and positrons that are generated by hadronic interaction processes, as discussed in the following. Thus, the steady-state behavior of the differential nonthermal electron and proton density n(T) in the AGN corona and the starburst zone, respectively, can be approximated by 3 −∂∂TTn(T)τcool(T)=q(T)−n(T)τesc(T). Here, τ cool refers to the total continuous energy-loss timescale, which in the case of the relativistic electrons is given by the inverse of the sum of the synchrotron (syn), inverse Compton (IC), nonthermal Bremsstrahlung (brems), and Coulomb (C) loss rates, according to 4 τcool(e)=[τsyn(e)−1+τic−1+τbrems−1+τC(e)−1]−1, and in the case of the relativistic protons we use 5 τcool(p)=[τsyn(p)−1+τC(p)−1+τpγπ−1+τBH−1+τpp−1]−1, including the photopion (π), Bethe–Heitler pairs (BH), and hadronic pion (pp) production loss rates. Proton synchrotron losses—as well as the associated radiation—are negligible for the considered environments. Note that these processes require additional information on the associated interaction medium, which is one of the following targets:(i)A magnetic field, which is assumed to be uniform on small scales (with respect to the particles’ gyro radius) and randomly oriented on significantly larger scales (due to isotropic Alfvénic turbulence).(ii)A photon target, which is in the case of the starburst zone dominated by the thermal IR emission due to the rescattered starlight by dust grains with a temperature θ dust and can be described by an isotropic, diluted modified blackbody radiation field 6 nIR(E)=Cdilπ2ℏc3E2exp(E/(kBθdust))−1EE0, where the dust clouds become optically thick above a critical energy E 0 = 8.2 meV (Yun & Carilli 2002). The constant dilution factor C dil is determined from the observed IR luminosity L IR according to the relation LIR/(πRstr2c)=∫dEEnIR(E) . 7 7 A more accurate approach to the IR photon spectrum has been proposed by Casey (2012), where a coupled modified greybody plus a mid-infrared power law has been used, but these modifications have no impact on our results. In case of the coronal region we used a parametrized model (Ho 2008) above 1 eV that accounts for the optical and UV emission by the disk as well as the Comptonized X-ray emission by hot thermal electrons in the corona. Hereby, the parameterization depends on the Eddington ratio (L bol/L Edd), i.e., the ratio of the bolometric over the Eddington luminosity, and we adopt the relation of Hopkins et al. (2007) to determine L bol based on the intrinsic X-ray luminosity L X between (2 and 10)keV.(iii)A thermal gas target with a given temperature θ, which is due to mathematical convenience assumed to be homogeneously distributed in both regions. For the starburst ring, Spinoglio et al. (2012) determine θ = 127 K, a gas density of n(H2) = 102.9 cm−3, and a molecular hydrogen mass of M(H2) ∼ 3.5 × 108 M ⊙.More details on the individual energy-loss timescales can be found in Appendix A.

2017MNRAS.471.3057M__Bovy_et_al._2016b_Instance_2

We have performed the first detailed dissection of the stellar populations of the Milky Way disc in age, [Fe/H] and $\mathrm{[ \alpha \mathrm{/Fe]}}$ space, bridging the gap between the detailed observational understanding of MAPs (e.g. Bovy et al. 2012b, 2016b) and the plethora of studies of co-eval stellar populations in simulated galaxies (e.g Bird et al. 2013; Stinson et al. 2013; Martig et al. 2014a). We have placed novel constraints on models for the formation of the Milky Way disc by combining detailed density models fit to the mono-age, mono-[Fe/H] populations of the low and high $\mathrm{[ \alpha \mathrm{/Fe]}}$ disc, with surface mass density contributions calculated on the basis of these density fits and stellar evolution models. We summarize our key results as follows: Radial and vertical profiles: The mono-age, mono-[Fe/H] populations of the $\mathrm{[ \alpha \mathrm{/Fe]}}$ poor disc are well fitted by a radially broken exponential, with a peak radius, Rpeak, that varies as a function of age and [Fe/H]. We find that the distance between Rpeak's of the low and high [Fe/H] populations increases with age, which we interpret as evidence for a decreasing [Fe/H] gradient with time (e.g. Anders et al. 2017). The radial variation of the stellar surface density of the high $\mathrm{[ \alpha \mathrm{/Fe]}}$ mono-age populations is found to have insignificant breaks, and they are better fit by a single exponential in this disc region. As these populations are the oldest, this may be a sign of the disc evolution washing out the density peak over time, or may point to a different formation scenario for high $\mathrm{[ \alpha \mathrm{/Fe]}}$ stars, where no density peak ever existed. These findings are in good agreement with earlier studies of MAPs (Bovy et al. 2016b). We measure an average high $\mathrm{[ \alpha \mathrm{/Fe]}}$ population scalelength of hR, in = 1.9 ± 0.1 kpc, and find scaleheights between 600 and 1000 pc, in good agreement with current measures of the $\mathrm{[ \alpha \mathrm{/Fe]}}$ rich disc scalelength and scaleheight (e.g. those outlined in Bland-Hawthorn & Gerhard 2016).Profile broadening: We show that the radial surface density profile of the low $\mathrm{[ \alpha \mathrm{/Fe]}}$ populations broadens with age in a given [Fe/H] bin, which we interpret as evidence of the gradual dispersal of mono-[Fe/H] populations, presumably due to radial migration and radial heating. The variation in shape of the broken exponential profile changes differently depending on the population [Fe/H], with low [Fe/H] populations inner profiles flattening faster, whereas the high [Fe/H] outer profiles flatten faster. We interpret this effect as tentative evidence for [Fe/H] dependent radial migration arising from pre-existing [Fe/H] gradients in the star-forming disc. We showed that our results qualitatively reproduce those of Hayden et al. (2015), finding a skewed MDF that varies as a function of R.Flaring: We find that flaring seems to be present in almost all mono-age populations, at differing levels. We have shown that the inverse flaring scalelength Rflare− 1 increases with age, meaning that the youngest populations flare most strongly. This finding appears inconsistent with that above, under the assumption that flaring is the result of radial migration. However, these results may be reconciled by invoking a more active accretion history in the early life of the disc, which could have suppressed flaring (e.g. Minchev et al. 2014b).The surface-mass density at R0: We have measured the surface mass density at the solar radius for each mono-age, mono-[Fe/H] population, finding a total surface mass density of $\Sigma _{R_0, {\rm tot}} = 20.0_{-2.9}^{+2.4}\mathrm{(stat.)}_{-2.4}^{+5.0}\mathrm{(syst.)}\ \mathrm{M_{{\odot }} \ pc^{-2}}$. Before allowing for systematics, this value is less than current estimates (e.g. Flynn et al. 2006; Bovy et al. 2012a; McKee et al. 2015), however, the systematic uncertainties are large, mainly due to a mismatch between the log g scales in APOGEE and the PARSEC models, and as such, we find our value to be consistent within the uncertainties. The relative contribution of high to low $\mathrm{[ \alpha \mathrm{/Fe]}}$ populations, $f_\Sigma$, is 18 per cent ± 5 per cent, which is consistent with existing measurements (e.g. Bland-Hawthorn & Gerhard 2016).The hZ distribution at R0: The shape of the mass-weighted hZ distribution found by this study is in good agreement with that of Bovy et al. (2012a), calling into question the existence of a vertical structural discontinuity in the Milky Way disc. The reconciliation of this finding with the discontinuity in chemical space (e.g. the bimodality in $\mathrm{[ \alpha \mathrm{/Fe]}}$ at fixed [Fe/H]: Nidever et al. 2014; Hayden et al. 2015) may shed new light on our understanding of the formation of the Galactic disc.The surface-mass density profile of the Milky Way: We have found the combined (from mono-age, mono-[Fe/H] populations at low and high $\mathrm{[ \alpha \mathrm{/Fe]}}$) surface-mass density-weighted profiles of the Milky Way disc as a function of $\mathrm{[ \alpha \mathrm{/Fe]}}$, age and [Fe/H], and found that the total surface density is also described by a broken exponential. We find that our results fail to determine the sign of the inner exponential to high significance out to ∼10 kpc, but detect a turnover to a declining exponential, at high significance, thereafter. We find evidence of a radial mean age and [Fe/H] gradient driven by the changing dominant population as a function of radius. A detailed comparison of these findings with numerical simulations is necessary for a proper interpretation. Our finding of a decline in stellar density may be consistent with that found in other studies (e.g. Reylé et al. 2009; Sale et al. 2010), albeit at shorter radii.

2019MNRAS.487..364L__Kruijssen_et_al._2018_Instance_1

It is well known that star formation is inefficient on galactic scales. The observed linear correlation between molecular gas surface density and star formation rate (SFR) surface density in normal star-forming galaxies suggests a nearly constant gas depletion time-scale around ∼2 Gyr, much longer than the dynamical time-scale of galactic discs (Kennicutt 1989; Bigiel et al. 2008; Saintonge et al. 2011; Leroy et al. 2013; Genzel et al. 2015; Tacconi et al. 2018). In contrast, the SFE on GMC scales shows a large variation ranging from less than a few per cent to nearly unity (Zuckerman & Evans 1974; Krumholz & Tan 2007; Wu et al. 2010; Evans, Heiderman & Vutisalchavakul 2014; Heyer et al. 2016; Lee, Miville-Deschênes & Murray 2016; Vutisalchavakul, Evans & Heyer 2016). The origin of this large scatter is usually explained as a combination of the time variability of the SFR during the course of cloud evolution and intrinsic scatter of SFEs due to the diversity of GMC properties (Feldmann & Gnedin 2011; Kruijssen & Longmore 2014; Lee et al. 2016; Grudić et al. 2018a; Kruijssen et al. 2018). For example, recent theoretical models and high-resolution magneto-hydrodynamics simulations suggest that the SFE depends on the local virial parameters of the cloud controlled by large-scale turbulence (e.g. Krumholz & McKee 2005; Padoan, Haugbølle & Nordlund 2012). However, it has recently been recognized that large-scale turbulence can only account for an ∼0.3 dex scatter, which is not sufficient to explain the observed SFE variations (Lee et al. 2016). Another source of variation comes from different stellar feedback channels that alter the dynamical states of the GMCs (Fall, Krumholz & Matzner 2010; Murray, Quataert & Thompson 2010; Dale et al. 2014; Myers et al. 2014; Raskutti, Ostriker & Skinner 2016; Kim et al. 2017; Grudić et al. 2018b). Previous studies found that GMC simulations adopting different stellar feedback mechanisms (stellar winds, ionizing radiation, or supernovae) lead to dramatically different final SFEs. The problem has recently been recognized to be more subtle than previously thought, since even small differences in numerical treatments, such as different radiative transfer schemes, massive star sampling, and momentum and energy deposition algorithms, can lead to drastic changes for the final SFE (Dale et al. 2005; Roškar et al. 2014; Raskutti et al. 2016; Grudić et al. 2018b; Kim, Kim & Ostriker 2018). Therefore, how the SFE depends on GMC properties and the strength of stellar feedback remains an open question.

2021ApJ...914L..38U__Dullemond_et_al._2001_Instance_1

Snow lines of abundant volatile species are one of the possible origins of the observed substructures in dust continuum emission (Zhang et al. 2015; Okuzumi et al. 2016; Pinilla et al. 2017). Recent disk surveys have shown that the locations of the observed gap/ring structures do not seem to be related to the radial locations of the snow lines (Huang et al. 2018; Long et al. 2018; van der Marel et al. 2019). To estimate the positions of the snow lines in observed disks, the temperature profile is often assumed to be a simple power law (Huang et al. 2018; Long et al. 2018), which is broken if the disk has shadows on the disk surface (Dullemond et al. 2001; Baillié & Charnoz 2014; Ueda et al. 2019). Figure 5 shows the radial positions of snow lines of H2O, NH3, CO2, H2S, C2H6, CH4, and CO at t = 0 and 11. For simplicity, we define the snow lines as the radial location where the midplane temperature reaches 150, 70, 55, 50, 40, 25, and 20 K for H2O, NH3, CO2, H2S, C2H6, CH4, and CO, respectively. We clearly see that the radial positions of snow lines move with time and multiple snow lines emerge even for a single species. This means that it is necessary to determine the disk temperature precisely when we evaluate the snow line locations. These snow lines would induce additional ring and gap structures and some of them would overlap with the TWI-induced substructures. It should be noted that the oscillation timescale of the TWI is much shorter than the dust radial drift timescale. Therefore, sintering-induced substructures would not coincide with the locations of the snow lines if the TWI is present. However, the dust-size variation can be induced by sublimation and recondensation, which potentially produces millimeter substructures. Since the radial locations of the snow lines are important not only for the substructure formation but also for the chemical composition of forming planets (e.g., Sato et al. 2016; Öberg & Wordsworth 2019), we should investigate how the TWI evolves in planet forming disks.

2018AandA...615A.148D__Weidner_et_al._(2010)_Instance_1

The last column in Table 1 reports the number of OB stars minus the “diffuse” population estimated from their density in the Reference field (22.5 stars per square degree): as is immediately seen, the M-star statistics is much larger than the OB star statistics. This can hardly be considered surprising, if an ordinary IMF (e.g., that from Weidner et al. 2010) is assumed for the star-formation region. In Fig. 19 we show the density ratio between M and OB stars, which provides a consistency test between our results and a plausible IMF: this ratio varies however by a large factor, close to 20, among our subregions. This might reflect differences in the respective IMFs, but also differences in completeness among the stellar samples considered for the various regions. We first note that the ratio between M and OB stars in NGC 6231 is dramatically lower than anywhere else in Sco OB1. We can indeed expect that M stars are detected less efficiently in the inner parts of NGC 6231, where the density of bright stars is very large, and their diffuse glare raises the limiting magnitude locally. As already discussed above in Sect. 4.1, this causes our sample of M stars in NGC 6231 to be highly incomplete. Moreover, we determined above that NGC 6231 is significantly more extincted, by almost half a magnitude in V, than Tr 24, and this implies a higher minimum detectable mass among NGC 6231 M stars compared to Tr 24 (see the MDA diagrams in Fig. 5); this effect reduces the completeness of the M-star sample in NGC 6231 more than in Tr 24. If Tr 24 is also slightly younger than NGC 6231, as we argue below, our M-stars in Tr 24 will reach down to lower masses than in NGC 6231, with a steep increase in the detected M-star population: adopting the IMF from Weidner et al. (2010), the predicted number of cluster M stars doubles considering the mass interval 0.25–0.5 M⊙ rather than 0.35–0.5 M⊙. If Tr 24 is younger than NGC 6231, moreover, its stars in the mass range 2.5–3 M⊙ might not have yet reached their ZAMS position as B stars, and therefore would not be counted among OB stars; this would further raisethe M/OB star ratio there by up to 30%. Therefore, the proportions of both M and OB stars that are detected in a young cluster will depend on their age and extinction, in accordance with the MDA diagrams, even for a fixed, spatially uniform photometric sensitivity. We estimated using the Weidner et al. (2010) IMF the expected range for the observed M/OB number ratio. Siess et al. (2000) predict that the latest-type B stars have a mass of ~ 3.5 M⊙ at 2 Myr, and ~ 2.2 M⊙ at 10 Myr, that is, in the range of ages expected for Sco OB1 clusters. The MDA diagrams of Fig. 5 predict that the lowest-mass stars we are able to detect using the available Sco-OB1 data have ~ 0.2 M⊙, even assuming the most favorable (and unlikely) circumstances of an age less than 2 Myr and negligible reddening. The extreme values found for the M/OB ratio are then ~ 3.8 for a minimum M-star mass as high as 0.35 M⊙ and an old age of 10 Myr, and ~ 20 for a minimum M-star massas low as 0.2 M⊙ and age of 2 Myr. These extremes are also shown as horizontal lines in Fig. 19. We note that the M/OB ratio in NGC 6231 falls well within this range; however, both Tr 24 regions are significantly richer of M stars than expected, by more than a factor of two and well above (statistical) errors. If true, then paradoxically this part of the OB association would form preferentially lower-mass stars. Of course, more detailed studies are needed to confirm this result. In the G345.45+1.50 region the M/OB ratio is highest, and far above predictions from the IMF: we may tentatively explain this since this region is very young, and some of its most massive members, like IRAS 16562-3959, are still in formation, thus decreasing the number of optically revealed OB stars. The lowest M/OB ratio in NGC 6231 is unlikely to be real, since as discussed above our M-star sample in this densest subregion is likely incomplete.

2022MNRAS.511.6218Z__Malkov,_Diamond_&_Sagdeev_2012_Instance_1

To further reduce the number of fitting parameters, we assume that no source is unique in terms of the acceleration spectral shape, i.e. all the source components have the same H spectral index α, He spectral index, maximum rigidity Rm, and He-to-H Galactic injection flux ratio χ. The spectrum of He seems harder than that of H as indicated by many spectral studies, implying, e.g. time-dependent particle acceleration (Zhang et al. 2017), or more efficient He injection in the shock acceleration process (Malkov, Diamond & Sagdeev 2012). It is thus assumed the He spectral index to be α − 0.077 according to the sub-TeV data fit (Aguilar et al. 2015b). To be consistent with the conventional scenario, we also consider that the nearby SNR can inject a total CR energy of 1050 erg into the ISM, i.e. (15)$$\begin{eqnarray} \int _{\text{GeV}}^{\infty }{\left[ \frac{{\rm d}R_{\text{H}}}{{\rm d}E}+\chi \left(E \right) \frac{{\rm d}R_{\text{He}}}{{\rm d}E} \right] N\left(E \right) E{\rm d}E}=10^{50}\text{ erg}, \end{eqnarray}$$where E is the particle kinetic energy, N refers to the injection spectrum of H with $N\left( \text{GV}\right) \approx 4.7\times 10^{52}\text{ GV}^{-1}$ (for the fitted parameters in Table 1). For brevity, we hereafter treat all of N and Q as quantities for H components. With these respects, in the naive approach, the spectral fitting parameters are $\chi \left( \text{GV}\right)$, α, Rm, ϕ, θ, MA, $\left| z/h \right|$, $\mu Q_{\text{d}}\left( \text{GV}\right)$, $Q_{\text{c}}\left( \text{GV}\right)$, tc, where $\theta =\arcsin \left(\rho /r_{\text{n}} \right)$ is the angle between the regular magnetic field line and the line-of-sight vector towards the nearby source. In the quasi-local approach equation (11), the alignment of the field line must uniquely be determined, leading to an additional angular parameter, which can be chosen as the angle φ around the line-of-sight vector (towards the nearby source with φ = 0° shown in Fig. 2d). The fitting results of the CR spectrum for models with the naive and quasi-local approach are shown with Figs 1(a) and 2(a), respectively.

2021AandA...655A..99D__Carigi_et_al._2005_Instance_3

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.

2021MNRAS.507.2766S__Sumiyoshi_et_al._2005_Instance_1

In order to make a linear analysis, first we have to prepare the PNS models as a background. The PNS properties depend on not only the density and pressure profiles but also the distributions of temperature (or entropy per baryon) and electron fraction inside the PNS, while such profiles can be determined only via the numerical simulation of the core-collapse supernova explosion. In this study, as in Sotani & Sumiyoshi (2019), we particularly adopt the profiles obtained via the numerical simulations performed by solving the general relativistic neutrino-radiation hydrodynamics under the spherical symmetry. In the simulations, hydrodynamics and neutrino transfer in general relativity are solved simultaneously (Yamada 1997; Yamada, Janka & Suzuki 1999; Sumiyoshi et al. 2005). To describe the neutrino transfer, the Boltzmann equation is directly solved with the multi-angle and multi-energy neutrino distributions for four species, νe, $\bar{\nu }_\mathrm{ e}$, νμ/τ, and $\bar{\nu }_{\mu /\tau }$, i.e. we implement six species of neutrinos by assuming μ-type and τ-type (anti-)neutrinos have identical distributions. For the collision term associated with neutrino emission, absorption, and scattering with leptons, nucleons, and nuclei handles, the basic neutrino reactions are adopted (Bruenn 1985; Sumiyoshi et al. 2005). The metric adopted in the numerical code is given by (1)$$\begin{eqnarray*} \mathrm{ d}s^2 = -\mathrm{ e}^{2\Phi (t,m_\mathrm{ b})}\mathrm{ d}t^2 + \mathrm{ e}^{2\Lambda (t,m_\mathrm{ b})}\mathrm{ d}m_\mathrm{ b}^2 + r^2(t,m_\mathrm{ b})(\mathrm{ d}\theta ^2 + \sin ^2\theta \mathrm{ d}\mathrm{ }\phi ^2), \nonumber\\ \end{eqnarray*}$$where t and mb denote the coordinate time and the baryon mass coordinate, respectively (Misner & Sharp 1964). In addition, mb is related to the circumference radius (r) via the baryon mass conservation, while the metric functions, Φ(t, mb) and Λ(t, mb), are evolved together with hydrodynamical variables in the numerical simulations (Yamada 1997). The numerical simulations for core-collapse supernovae have been done with 255 grid points in the radial mass coordinate, 6 grid points in the neutrino angle, and 14 grid points in the neutrino energy. The rezoning of radial mesh is made during the simulations to resolve the accreting matter. We remark that the radial grids of mass coordinate are non-uniformly arranged to cover not only the dense region inside the central object but also the region for accreting matter.

2019MNRAS.487.1210T__McNamara_&_Nulsen_2007_Instance_3

On larger scales, the clusters in which BCGs reside can generally be divided into two categories: cool core clusters, which exhibit very peaked surface brightness distributions at X-ray wavelengths, and non cool core clusters, with similar overall X-ray luminosities but with smoother, less peaked X-ray surface brightness distributions. Some authors (e.g. Hudson et al. 2010; Santos et al. 2010) define an intermediate category called moderate or weak cool core clusters. Since cool core clusters have short radiative cooling time-scales on the order of 108 yr in their centres (e.g. Voigt & Fabian 2004; McNamara & Nulsen 2007, 2012; Hlavacek-Larrondo et al. 2012), starbursts are expected to be common at the centre of such clusters. Indeed, the central cool gas in these clusters should condense onto the BCG, forming stars at rates of hundreds of solar masses per year (e.g. Fabian 1994). However, most BCGs are relatively quiescent and those that do show evidence of star formation generally tend to have star formation rates 1 order of magnitude smaller, on the order of $1-150 \, \mathrm{M_{\odot }\, {yr}^{-1}}$ (e.g. Donahue et al. 2007; Bildfell et al. 2008; O’Dea et al. 2008, 2010; Rawle et al. 2012). This mismatch between expected and observed star-forming rates, known as the cooling flow problem, is thought to be caused by active galactic nuclei (AGNs) feedback processes from the BCG. AGNs can release copious amounts of energy into the intracluster medium (ICM) through many ways, including: jetted outflows that inflate cavities, weak shocks, sound waves, or turbulence in the ICM (e.g. Markevitch & Vikhlinin 2007; McNamara & Nulsen 2007, 2012; Zhuravleva et al. 2014; Fabian et al. 2017). Alone, the energy released by jetted outflows appears to be on the same order as the energy needed to offset cooling (e.g. Rafferty et al. 2006; McNamara & Nulsen 2007; Hlavacek-Larrondo et al. 2012), therefore suggesting that AGN feedback is a good candidate for solving the cooling flow problem.

2020MNRAS.499.2575E__Pontzen_&_Governato_2014_Instance_1

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

2022ApJ...940L..18Z__Perna_et_al._2022_Instance_1

In addition to some fraction of BNS mergers masquerading as long GRBs, our sample used to constrain the DTD may suffer from other issues of incompleteness. As we rely on the modeling of host galaxies when constraining the DTD, we do not consider short GRBs that do not have a confident host association. Though the properties of short GRB hosts do not seem to deviate strongly as a function of the host-association confidence (see, e.g., Figure 4 of Nugent et al. 2022), neglecting these events may have a potential impact on both the low and high ends of our inferred DTD. For example, GRBs that are highly offset from their hosts may have afterglows with much lower luminosities, making precise localization (and therefore host identification) difficult (Perna et al. 2022). Such systems may have migrated over longer timescales to reach the highly offset locations of the burst and therefore may have longer delay times than the general population. Furthermore, the P cc method for host identification may incorrectly associate a GRB with a faint underlying host rather than a bright host at a larger offset, though Fong et al. (2022) predicted this to be an effect only at the ≲7% level. On the other hand, if such poorly associated GRBs are instead truly associated with faint galaxies that are below detection limits, we may be excluding additional systems with short delay times as these faint, low-mass galaxies are typically star-forming. Furthermore, though Swift can detect GRBs out to z ∼ 3, there is likely some fraction of short GRBs that occur beyond this horizon, when the universe was ≲2 Gyr old. Short GRBs that occur at these early stages in the history of the universe must have short delay times, and this selection effect may bias the general population in our analysis to longer delay times. This would lead to a larger inferred tmin , and, due to the correlation between tmin and α, more negative values of α. However, this population of high-redshift short GRBs is likely small; assuming the SFH from Madau & Fragos (2017), 10% of stars are born beyond z = 3, and the fraction of compact object binary mergers beyond this redshift will be even smaller due to the delay time between formation and merger.

2021AandA...653A.156F__Valentini_et_al._2007_Instance_1

In general, hybrid-kinetic models have proved to be capable of satisfactorily catching the main kinetic physics at play for a large number of problems, ranging from fluid and kinetic instabilities (Hellinger & Matsumoto 2000; Matteini et al. 2006; Califano et al. 2008; Henri et al. 2013; Kunz et al. 2014) to collisionless shocks (Lembège et al. 2009; Caprioli & Spitkovsky 2013; Weidl et al. 2016), dynamo effects (Rincon et al. 2016; St-Onge & Kunz 2018), MR (Le et al. 2016; Palmroth et al. 2017; Cerri & Califano 2017; Franci et al. 2017; Wang et al. 2019; Califano et al. 2020), and kinetic-scale turbulence (Servidio et al. 2015; Grošelj et al. 2017; Cerri et al. 2018, 2019; Hellinger et al. 2019; Wang et al. 2019). The main goal of our project is to investigate the possibility of improving the electron description in hybrid-kinetic models. The starting point is the “hybrid Vlasov–Maxwell” (HVM) code (Mangeney et al. 2002; Valentini et al. 2007), which is already equipped with a fluid model in which electrons are described as an isotropic, isothermal fluid with finite inertia. The HVM code has recently been upgraded in order to implement a more sophisticated model for the electron fluid. This includes evolution equations for the anisotropic (gyrotropic) electron pressures, p∥, e and p⊥, e (where ∥ and ⊥ refer to the local magnetic-field direction, b = B/|B|), and a Landau-fluid (LF) closure for the parallel transport of the gyrotropic electron thermal energy along field lines (i.e., parallel heat fluxes, q∥, e and q⊥, e). Hereafter, we refer to this new model as “hybrid Vlasov–Landau-fluid” (HVLF). The idea is to include within a hybrid description the relevant electron pressure-anisotropy effects and a fluid model for the electron-kinetic response that still holds in a nonlinear regime. Therefore, the LF model implemented in the HVLF code goes beyond the early attempts to include these effects in simplified settings (e.g., Hammett & Perkins 1990; Snyder et al. 1997; Passot & Sulem 2007) and is based on the approach presented by Sulem & Passot (2015).

2021AandA...655A..72S___2019_Instance_2

In this paper, we report on spectroscopic CH3CN, CH3OH (methanol), and dust continuum observations with the Atacama Large Millimeter/submillimeter Array (ALMA) at 349 GHz with an angular resolution of 0′′.1. We exploit the CH3CN (19K–18K) K-ladder, with excitation energies ranging from 168 K (for K = 0) to 881 K (for K = 10), to probe, at different radii, the physical conditions in the accretion disk of an early-type young star. We targeted the star-forming region G023.01−00.41, at a trigonometric distance of 4.59 $^{+0.38}_{-0.33}$ −0.33 +0.38 kpc from the Sun (Brunthaler et al. 2009), where we recently revealed the accretion disk around a young star of 104.6 L⊙, corresponding to a ZAMS star of 20 M⊙ (Sanna et al. 2019, their Fig. 1); the disk was imaged by means of spectroscopic ALMA observations of both CH3CN and CH3OH lines at 0′′.2 resolution inthe 230 GHz band. The disk extends up to radii of 3000 au from the central star where it warps above the midplane; here, we resolve the outer disk regions in two apparent spirals projected onto the plane of the sky. We showed that molecular gas is falling in and slowly rotating with sub-Keplerian velocities down to radii of 500 au from the central star, where we measured a mass infall rate of 6 × 10−4 M⊙ yr−1 (Sanna et al. 2019, their Fig. 5). The disk and star system drives a radio continuum jet and a molecular outflow aligned along a position angle of 57°, measured east of north (Sanna et al. 2016, their Fig. 2); their projected axis is oriented perpendicular to the disk midplane whose inclination with respect to the line-of-sight was estimated to be less than 30° (namely, the disk is seen approximately edge-on; Sanna et al. 2014, 2019). Previously, we also measured the average gas conditions over the same extent of the whole disk, by means of Submillimeter Array (SMA) observations of the CH3CN (12K–11K) emission, and we estimated a kinetic temperature of 195 K and CH3CN column density of 5.1 × 1016 cm−2 (Sanna et al. 2014, their Fig. 2 and Table 4).

2022ApJ...937...62L__Xiao_et_al._2022_Instance_1

However, the intrinsic effects caused by the unknown emission and acceleration mechanisms in the source could mitigate or enhance the LIV-induced time delay, which would impact the accuracy of the resulting constraints on LIV. A key challenge is then to distinguish an intrinsic time lag at the source from a delay induced by LIV. Long GRBs usually have significantly positive or negative intrinsic spectral lags and should not be used for LIV searches until reasonable progress is made on the modeling of the emission and acceleration mechanisms (Chen et al. 2005; Ukwatta et al. 2012; Bernardini et al. 2015), while short GRBs are consistent with null or negligible intrinsic spectral lag and are therefore an ideal tool to measure the LIV effect (Norris & Bonnell 2006; Bernardini et al. 2015, 2017; Xiao et al. 2022). Currently, in addition to short GRBs, active galactic nucleus (AGN) flares and gamma-ray pulsars are two other classes of astrophysical sources that have no significant intrinsic lag in general and are often used for LIV tests (Biller et al. 1999; Kaaret 1999; Aharonian et al. 2008; MAGIC Collaboration et al. 2017). It should be noted, however, that there is also evidence of intrinsic lags in some cases of AGN flares and pulsars. For example, MAGIC Collaboration during an observational campaign regarding Mkn 501 blazar found an indication of about 4 minutes time delay between the peaks at E 0.25 TeV and E > 1.2 TeV (MAGIC Collaboration et al. 2008), which may indicate a progressive acceleration of electrons in the emitting plasma blob. A robust method to study the correlations between arrival times and energy, based on a likelihood function built from the physical picture assumed for the emission, propagation, and detection of the photons was proposed by Martínez & Errando (2009). In the case of pulsars, there are some lags if the energy range is extended too much toward low energies (e.g., radio versus TeV). No real progress on the topic of intrinsic effects will be made without accurate models for production and acceleration mechanisms for each type of source. Perennes et al. (2020) first attempted to gain knowledge on source-intrinsic spectral lags of flaring AGNs at high and very high energies and on short timescales relevant for LIV searches, using leptonic AGN flare modeling. Concerning GRBs, some ways have been proposed to reduce the impact of intrinsic effects, e.g., fitting the observed spectral lags of statistical samples of GRBs at a range of different redshifts (Ellis et al. 2006; Bernardini et al. 2017; Xiao et al. 2022), or using only a limited observer-frame energy bands range corresponding to the fixed source-frame energy bands (Wei & Wu 2017). Anyway, there is no reason to think that the low and high-energy photons should be emitted simultaneously at the source, and while detecting distinct signals at different energy channels, we have no idea which one was sent first. Previous studies usually assumed that the intrinsic time delays are either an unknown constant for all GRBs considered or scale with the photon energy E according to some power-law function (Ellis et al. 2006; Biesiada & Piórkowska 2009a; Zhang & Ma 2015; Wei et al. 2017a; Acciari et al. 2020; Pan et al. 2020; Du et al. 2021).

2022MNRAS.510.4943S__Murray_&_Dermott_1999_Instance_1

The gravitational potential of an eccentric companion at the quadrupole order can be decomposed as a sum over circular orbits (e.g. Storch & Lai 2013; Vick, Lai & Fuller 2017): (5)$$\begin{eqnarray*} U\left(\boldsymbol{\mathbf {r}}, t\right) = \sum \limits _{m=-2}^2 U_{2m} \left(\boldsymbol{\mathbf {r}}, t\right) , \end{eqnarray*}$$(6)$$\begin{eqnarray*} U_{2m}\left(\boldsymbol{\mathbf {r}}, t\right) &amp;=&amp; -\frac{GM_2 W_{2m} r^2}{D(t)^3} Y_{2m}(\theta , \phi) e^{-imf\!\!\!\:(t)}, \\ &amp;=&amp; -\frac{GM_2W_{2m} r^2}{a^3}Y_{2m}\left(\theta , \phi \right) \sum \limits _{N = -\infty }^\infty \!\!F_{Nm}e^{-iN\Omega t} . \end{eqnarray*}$$Here, the coordinate system is centered on the MS star, (r, θ, ϕ) are the radial, polar, and azimuthal coordinates of $\boldsymbol{\mathbf {r}}$ respectively, $W_{2 \pm 2} = \sqrt{3\pi /10}$, W2 ± 1 = 0, $W_{20} = -\sqrt{\pi / 5}$, D(t) is the instantaneous distance to the companion, f is the true anomaly, and Ylm denote the spherical harmonics. FNm denote the Hansen coefficients for l = 2 (also denoted $X^N_{2m}$ in Murray & Dermott 1999), which are the Fourier coefficients of the perturbing function, i.e. (7)$$\begin{eqnarray*} \frac{a^3}{D(t)^3} e^{-imf\!\!\!\:(t)} = \sum \limits _{N = -\infty }^\infty \!\!F_{Nm} e^{-iN\Omega t}. \end{eqnarray*}$$The FNm can be written explicitly as an integral over the eccentric anomaly (Murray & Dermott 1999; Storch & Lai 2013): (8)$$\begin{eqnarray*} F_{Nm} = \frac{1}{\pi }\int \limits _{0}^{\pi } \frac{\cos \left[N\left(E - e\sin E\right) - mf(E)\right]}{\left(1 - e\cos E\right)^2}\,\,\mathrm{d}E. \end{eqnarray*}$$By considering the effect of each summand in equation (5), the total torque on the star, energy transfer in the inertial frame, and energy transfer in the star’s corotating frame (which is also the tidal heating rate) can be obtained (Storch & Lai 2013; Vick et al. 2017): (9)$$\begin{eqnarray*} T = \sum \limits _{N = -\infty }^\infty F_{N2}^2 T_{\rm circ}\left(N\Omega - 2\Omega _{\rm s}\right), \end{eqnarray*}$$(10)$$\begin{eqnarray*} \dot{E}_{\rm in} &amp;=&amp; \frac{1}{2}\sum \limits _{N = -\infty }^\infty \Bigg [ \left(\frac{W_{20}}{W_{22}}\right)^2 N\Omega F_{N0}^2 T_{\rm circ}\left(N\Omega \right) \\ &amp;&amp;+\, N\Omega F_{N2}^2 T_{\rm circ}\left(N\Omega - 2\Omega _{\rm s}\right) \Bigg ] , \end{eqnarray*}$$(11)$$\begin{eqnarray*} \dot{E}_{\rm rot} = \dot{E}_{\rm in} - \Omega _{\rm s} T . \end{eqnarray*}$$Here, dots indicate time derivatives.

2022MNRAS.514.2010M__Feng_&_Holder_2018_Instance_1

In the last few years, several experiments have reported upper limits on the power spectrum of 21-cm fluctuations during reionization (Parsons et al. 2014; Patil et al. 2017; Barry et al. 2019; Mertens et al. 2020; The HERA Collaboration 2021b) and the earlier cosmic-dawn era (Eastwood et al. 2019; Gehlot et al. 2019, 2020; Garsden et al. 2021; Yoshiura et al. 2021). Scenarios in which the bulk IGM is still colder than the cosmic microwave background (CMB) during reionization give rise to the strongest fluctuations and so will be the first models to be tested as upper limits continue to improve (e.g. Parsons et al. 2014; Pober et al. 2015; Greig, Mesinger & Pober 2016). Similarly, stronger-than-expected 21-cm signals can arise if the cosmic radio background has contributions other than the CMB (Feng & Holder 2018), e.g. synchrotron emission from accreting black holes (Ewall-Wice et al. 2018), star-forming galaxies (Mirocha & Furlanetto 2019), or from decaying particles (Fraser et al. 2018; Pospelov et al. 2018). Indeed, constraints from MWA, HERA, and LoFAR disfavour models with negligible X-ray heating at z ∼ 8–9 or very strong radio backgrounds (Ghara et al. 2020, 2021; Mondal et al. 2020; Greig et al. 2021a, b; The HERA Collaboration 2021a). Of course, the recent report of an absorption signal in the sky-averaged spectrum at z ∼ 17 from EDGES (Bowman et al. 2018) requires an even colder IGM (Barkana 2018; Boddy et al. 2018; Fialkov, Barkana & Cohen 2018; Kovetz et al. 2018; Muñoz & Loeb 2018) or a brighter background (Ewall-Wice et al. 2018; Feng & Holder 2018; Fialkov & Barkana 2019; Mirocha & Furlanetto 2019) than models in ΛCDM cosmologies generally predict. However, the most stringent power spectrum upper limits from The HERA Collaboration (2021b) are derived at sufficiently low redshifts relative to EDGES (z ≲ 10 versus z ≃ 18) that they cannot yet directly address the EDGES controversy (Hills et al. 2018; Bradley et al. 2019; Singh & Subrahmanyan 2019; Sims & Pober 2020; Tauscher, Rapetti & Burns 2020; Singh et al. 2021).

2021ApJ...908..220T__Milone_et_al._2020_Instance_1

Recently, our knowledge of the MW formation and evolution has been revolutionized by the massive amount of data products from the Gaia mission and large spectroscopic surveys. Several dwarf galaxies are suggested to have been accreted by the MW since its formation (e.g., Helmi et al. 2018; Belokurov et al. 2018; Myeong et al. 2019). It is expected that GCs formed in these dwarf galaxies mix up with GCs formed in situ, i.e., in the main MW progenitor, to form the current GC system. This formation history can be revealed by precise kinematics and dynamical modeling. Massari et al. (2019) placed Pal 13 into a group of GCs related to a dwarf galaxy, named Sequoia, which was likely accreted 9 Gyr ago (Myeong et al. 2019). The existence of MPs in Pal 13 supports the statement that this phenomenon is not unique in Galactic GCs formed in situ (e.g., Li & de Grijs 2019; Milone et al. 2020). Besides its proposed accreted origin, Pal 13 is suggested to be experiencing tidal stripping (Yepez et al. 2019; Piatti & Fernández-Trincado 2020). Hamren et al. (2013) found that Pal 13 has lost a considerable amount of mass, which is related to its low present-day cluster mass. Recently, the discovery of tidal tails in Pal 13 was reported by Shipp et al. (2020). Using the RR Lyrae stars, these authors estimated the initial luminosity to be L⊙, which is significantly larger than the current luminosity estimated by B11, L⊙. Due to its large distance, several detailed features of Pal 13 have been discovered just recently (e.g., Piatti & Fernández-Trincado 2020; Shipp et al. 2020) and more are still waiting for further investigation, the estimation of total mass loss or initial mass of Pal 13 should be illuminative for future studies. Besides Pal 13, several remote, low mass GCs, including Whiting 1 and Eridanus GC, also show evidence of tidal tails or extra-tidal structures(Carballo-Bello et al. 2014; Myeong et al. 2017). These features agree with the predictions of dynamical evolution of low mass GCs: it is difficult to keep stars in a shallow potential well. The relatively large stellar mass lost from low mass GCs also complicates the discussion of MPs and especially of the lower mass limit needed to maintain MPs. On the other hand, the N-rich stars located in these GCs are lost to the field, contributing to the rare N-rich field stars (e.g., Martell & Grebel 2010; Fernández-Trincado et al. 2017; Tang et al. 2019; Fernández-Trincado et al. 2019; Tang et al. 2020).

2022MNRAS.515.2256V__Drury_1983_Instance_1

SBs are complex environments. The hot rarefied plasma is spanned by primary shocks, which decay into turbulent motions and MHD waves. Collective effects such as re-acceleration processes contribute to the acceleration of particles. SBs are also delimited by a forward shock which expands into the ISM. While the size of this shock can be very large (R ≳ 100 pc, Weaver et al. 1977), it is too slow (u ∼ 30 km s−1) to accelerate PeV protons, even assuming large magnetic fields in the ISM. More generally, in the case of strong shocks, the velocity u is identified as the velocity of the shock, as it is precisely the velocity jump at the shock discontinuity which drives the acceleration of the particles via the DSA mechanism (e.g. Drury 1983). Inside SBs, there are two types of strong primary shocks: the time-dependent SNR shocks which expand after a supernova (SN) explosion, and the wind termination shocks (WTS) which surround the individual massive stars, or the entire stellar cluster if it is compact enough. SNRs and WTSs have a typical velocity of several 1000 km s−1. In the low-density SB interior, SNRs expand to a typical radius of a few tens of pc before reaching the Sedov–Taylor phase, which is generally larger than the radius of the WTS, even in the case of a WTS powered by a very massive compact cluster. The size of the latter shock depends on the mechanical power of the stellar cluster and is typically of the order of 10 pc (Weaver et al. 1977). It is believed that, due to in situ CR acceleration, efficient magnetic field amplification takes place upstream of SNR shocks, leading to a magnetic field of up to several tens of μG. On the other hand, it is less clear if the magnetic field can be as efficiently amplified in stellar winds. In both cases, the maximum energy is generally inferred to be only ∼1 PeV (e.g. Gupta et al. 2020; Morlino et al. 2021; Vieu et al. 2022b). While this is slightly larger than the maximum energy achieved at isolated SNR, atypical conditions would be required for protons to be accelerated well beyond PeV by primary shocks embedded in SBs. A promising situation combining the advantages of both WTS and SNRs might be that of an SNR expanding in a wind profile close to a compact cluster. A powerful cluster may indeed convert a substantial amount of its mechanical energy into turbulence, amplifying magnetic fields up to hundreds of μG in its vicinity, such that particles could be accelerated up to 10 PeV by a powerful SNR propagating in the wind, even in the absence of additional magnetic-field amplification. Considering a similar scenario of SNR shock – cluster wind interaction with efficient turbulence generation, Bykov et al. (2015) found that proton energies up to 40 PeV could be reached in the case of fast shocks (U = 104 km s−1).

2016AandA...588A..44Y__Lehtinen_&_Mattila_(1996)_Instance_1

All coreshine observations were obtained with the IRAC instrument on board the Spitzer observatory and are gathered in Paladini (2014) and Lefèvre et al. (2014), from which we selected 21 starless cores in the Taurus-Perseus, Chamaeleon, Cepheus, and L183/L134 regions (see their Fig. 9 and Table 1). Lefèvre et al. (2014) summarised their results in two figures that we reproduce here: the 4.5 to 3.6 μm ratio, that they name “coreshine ratio”, as a function of the 3.6 μm intensity and the 2.2 to 3.6 μm ratio, that they name the “near-IR to mid-IR ratio”, as a function of the coreshine ratio. For the model cloud, we use the parameters of the control cloud defined in Sect. 3. After convolving our models with a 10″ FWHM Gaussian kernel to simulate the data analysis presented in Lefèvre et al. (2014) and following the Lehtinen & Mattila (1996) prescription to take into account the part of the ISRF+CM light that can be transmitted through the cloud3, we compute the synthetic photometry for each pixel along a radial cut through our model clouds. The results are shown in Figs. 10a and b, which present the coreshine ratio as a function of the 3.6 μm intensity for the 21 starless cores of Lefèvre et al. (2014) and in Fig. 10c and d, which displays the near- to mid-IR ratio as a function of the coreshine ratio for the Taurus-Perseus and L183/L134 regions. The CMM model has to be ruled out to explain coreshine since it only marginally fits the coreshine ratio and fails to reproduce the near- to mid-IR ratio (green areas in Fig. 10). On the contrary, the CMM+AMM and CMM+AMMI models can explain the coreshine observations (magenta and blue areas in Fig. 10, respectively). Based on the results presented in Sect. 3 and Fig. 7, the cloud parameters are as important as the dust model for explaining the dispersion in the observations. Varying the cloud external radius from Rout = 0.3 to 0.1 pc (Figs. 10a and c) and the central density from ρC = 104 to 5 × 105 H/cm3 (Figs. 10b and d) seem enough to explain most of the observed scatter.

2022ApJ...933..243F__Hajela_et_al._2022_Instance_1

Recently, Fraija et al. (2021a) presented the afterglow light curves generated by the deceleration of sub-relativistic masses ejected from the merger of BCOs and the death of massive stars. The authors assumed that a PL velocity distribution describes the isotropic-equivalent kinetic energy of these masses and that the sub-relativistic ejected masses were decelerated, in turn, by a stratified-density environment. As a particular case, to explain the multiwavelength observations of the gravitational event GW170817/GRB 170817A at ∼900 days, they constrained the parameter space of the synchrotron light curves of a sub-relativistic mass ejected during the merger of two NSs and decelerated in a constant-density environment. The synchrotron radiation of the sub-relativistic material was consistent with the faster blue KN afterglow. Inspired by the new observations of this GW event at 3.3 yr after the initial merger (Hajela et al. 2022), in this paper, we extend the synchrotron model presented in Fraija et al. (2021a), including the continuous energy injection from the central engine (either a spinning magnetized NS or BH remnant) into the blast wave through a numerical approach and analytic arguments. In addition, we apply the current model to potential candidates of sGRB events with evidence of a KN. The paper is organized as follows: Section 2 presents the dynamical evolution of the afterglow when the central engine continuously injects energy into the blast wave. We show an analytical solution and numerical approach. In Section 3, we show a synchrotron model with energy injection from a spinning magnetized NS and BH remnants. Section 4 shows the analysis of the multiwavelength light curves using typical values of the GRB afterglow. In Section 5, we apply our model to several potential candidates, including GW170817/GRB 170817A, and finally, in Section 6, we summarize. We consider the convention Qx=Q10x in cgs units and assume for the cosmological constants a spatially flat universe Λ cold dark matter model with H 0 = 69.6 km s−1 Mpc−1, ΩM = 0.286, and ΩΛ = 0.714 (Planck Collaboration et al. 2016).

2019AandA...622A..60C__Drake_et_al._2013a_Instance_1

In order to establish the completeness and purity of the RR Lyrae stars confirmed by the SOS Cep&RRL pipeline and to estimate the number of new discoveries by Gaia, we performed a deep and careful comparison with the literature. As a first step, the catalogue of 140 784 confirmed sources was cross-matched against all major catalogues of known RR Lyrae stars that are available. We primarily used the OGLE catalogues for RR Lyrae stars (version IV of the survey, Soszyński et al. 2014, 2016), but we also used RR Lyrae stars by CTRS (Drake et al. 2013a,b, 2014, 2017; Torrealba et al. 2015), ASAS (Pojmanski 1997; Richards et al. 2012), ASAS-SN (Jayasinghe et al. 2018), ATLAS (Tonry et al. 2018), IOMC (Alfonso-Garzón et al. 2012), LINEAR (Palaversa et al. 2013), NSVS (Kinemuchi et al. 2006), Pann-Stars (PS1 Sesar et al. 2017), and from the works based on Kepler/K2 (Debosscher et al. 2011; Nemec et al. 2011; Molnár et al. 2015a,b, 2016) and on the Simbad database (Wenger et al. 2000). These cross-matches returned a list of 88 578 known RR Lyrae stars in our sample of 140 784 stars. The SOS Cep&RRL confirmed RR Lyrae stars were also cross-matched against catalogues of candidate RR Lyrae stars discovered by the VVV survey (Gran et al. 2016; Minniti et al. 2017; D. Minniti, priv. comm.) in the MW disc and bulge. This returned 319 VVV cross-identified sources in the MW disc and 222 in the MW bulge. We thus confirm these VVV candidates. For known RR Lyrae stars in GCs, the main reference was the catalogue of Clement et al. (2001), which was updated to the latest literature as described in Garofalo et al. (in prep.). For variables in dSphs, we used the following references: Kaluzny et al. (1995), Clementini et al. (2005), Kinemuchi et al. (2008), Dall’Ora et al. (2012) and Garofalo et al. (2013). These latter cross-matches returned a list of 1986 further known RR Lyrae stars. At the end of this cross-match procedure, of the 140 784 RR Lyrae stars that are confirmed by the SOS Cep&RRL pipeline, 90 564 were shown to be known previously, and 50 220 are new discoveries by Gaia.

2021ApJ...909..172Z__Read_&_Lebonnois_2018_Instance_2

Atmospheric superrotation is characterized by eastward wind at the equator, which means the atmosphere there has a higher angular momentum than the solid surface. Atmospheric superrotation is a common phenomenon across the universe. In the solar system, superrotation exists in the atmospheres of Venus, Titan, Saturn, and Jupiter, as well as the stratospheric atmosphere of Earth during the westerly phase of the quasi-biennial oscillation (e.g., Kraucunas & Hartmann 2005; Schneider & Liu 2009; Lutsko 2018; Read & Lebonnois 2018). In order to maintain atmospheric superrotation, there must be momentum transports from higher latitudes to the equator against friction or other processes, according to angular momentum conservation (Hide 1969; Held 1999; Showman et al. 2013). This up-gradient transport into the jet can result from Rossby waves, coupled Rossby–Kelvin waves, mixed Rossby–gravity waves, wave–jet resonance, barotropic instability, or baroclinic instability (Suarez & Duffy 1992; Del Genio & Zhou 1996; Joshi et al. 1997; Lee 1999; Williams 2003; Kraucunas & Hartmann 2005; Schneider & Liu 2009; Caballero & Huber 2010; Mitchell & Vallis 2010; Showman & Polvani 2010, 2011; Showman et al. 2010; Liu & Schneider 2011; Arnold et al. 2012; Pinto & Mitchell 2014; Tsai et al. 2014; Wang & Mitchell 2014; Laraia & Schneider 2015; Lutsko 2018; Read & Lebonnois 2018; Pierrehumbert & Hammond 2019). For example, Kraucunas & Hartmann (2005) suggested that in an Earth-like atmosphere, equatorial superrotation can be generated by equatorward stationary eddy momentum convergence, which is associated with zonal variations in the diabatic heating at low latitudes. Mitchell & Vallis (2010) studied the transition from current Earth-like atmospheric circulation to an equatorial superrotation state. They found that during the spin-up period, superrotation is generated by equatorward momentum convergence associated with both barotropic and baroclinic instabilities.

2021AandA...649A.168D__Oh_&_Escuti_2008_Instance_1

A critical property of the phase mask is that it needs to be able to image subapertures in off-axis interferograms. The off-axis interferograms are rather large, with size scaling with λ/Dsub, where λ is the wavelength and Dsub is the diameter of the subaperture. Therefore, imaging multiple interferograms onto separate locations (so as to avoid overlap) on the detector requires large phase tilts. This makes it difficult to manufacture classical phase implementations of a HAM phase mask for transmissive pupil planes. A solution is offered by liquid-crystal diffractive phase masks as they have an unbounded continuous phase (Escuti et al. 2016). This property enables the creation of steep phase ramps that efficiently diffract light into a single order without scattering (Oh & Escuti 2008). In Fig. 1, no noticeable second-order diffraction is seen for any off-axis interferogram. Another advantage of liquid-crystal masks is that it is possible to manufacture almost any phase pattern (Kim et al. 2015), meaning there is more design freedom. We exploited this by combining phase ramps into a single phase pattern that images a single subaperture onto multiple locations in the focal plane. This was done through multiplexing the phase ramps, and the mathematical description of multiplexing can be found in Doelman et al. (2018). An example of multiplexed subapertures is seen in the third column of Fig. 1, where multiple baselines connected to a single aperture are imaged onto different interferograms. In addition, liquid-crystal masks are diffractive because they apply a different kind of phase delay to incoming light that is independent of wavelength. These phase delays are called “geometric phase delays” and are discussed in greater detail in Sect. 3. Due to this diffraction, the location of an imaged subaperture changes with wavelength. The advantage of the diffractive nature is that, together with the right subaperture combination and fringe orientation, the wavelength smearing enables low-resolution spectroscopy. However, each interferogram can then only consist of 1D combinations of subapertures. This limits the design freedom significantly and also greatly increases the number of off-axis interferograms. As shown in Fig. 1, the fringe direction of all off-axis interferograms is orthogonal to the smearing direction for 1D combinations of subapertures. Lastly, a specific property of liquid-crystal diffractive phase masks is that they produce two off-axis interferograms for a single phase ramp with opposite location in the focal plane. This can be seen in Fig. 1, where all interferograms have an identical counterpart. The aforementioned properties have a large impact on the design of the HAM mask, which we discuss next.

2018MNRAS.473.2144B__Nakamura_&_Li_2007_Instance_1

The problem of the balance between driving and decay for magnetized turbulence is most acute in molecular clouds. Since these have linewidths indicating the presence of supersonic flow, the fast dissipation of turbulence found by these simulations necessitates a mechanism to reinject the energy equally quickly. A number of candidates have been proposed, including internal feedback from H ii regions (Matzner 2002; Krumholz, Matzner & McKee 2006; Goldbaum et al. 2011) or protostellar outflows (Li & Nakamura 2006; Nakamura & Li 2007; Wang et al. 2010; Federrath et al. 2014a), driving of turbulence by ongoing accretion (Klessen & Hennebelle 2010; Goldbaum et al. 2011; Lee & Hennebelle 2016) or gravitational contraction on small scales (Federrath et al. 2011b; Sur et al. 2012), thermal instability driving (Koyama & Inutsuka 2002; Hennebelle & Inutsuka 2006) and injection of energy from external supernova shocks (Mac Low & Klessen 2004; Padoan et al. 2016a,b; Pan et al. 2016). Alternately, it is possible that the linewidths do not reflect turbulent motion at all, and instead indicate global gravitational collapse (Ballesteros-Paredes et al. 2011; Zamora-Avilés & Vázquez-Semadeni 2014). Each of these proposals, however, faces challenges – internal feedback must maintain large linewidths without destroying the clouds in which they occur, driving by accretion faces the problem of what happens when the accretion eventually ends, thermal instability seems unlikely to be a viable mechanism in molecule-dominated galaxies that lack a significant warm phase, and external driving requires efficient coupling between the low-density external medium and the dense clouds. The view that clouds are in global collapse is hard to reconcile with the observed very low rates of star formation found even in gas at densities ≳ 105 cm−3 (Krumholz & Tan 2007; Federrath & Klessen 2012; Krumholz, Dekel & McKee 2012; Evans, Heiderman & Vutisalchavakul 2014; Padoan et al. 2014; Salim, Federrath & Kewley 2015; Usero et al. 2015; Heyer et al. 2016; Vutisalchavakul, Evans & Heyer 2016).

2021MNRAS.507.5053E__Johnston_et_al._2006_Instance_1

Multiwavelength observations of the GC indicate that the number of pulsars in the central few parsecs should be high (Wharton et al. 2012) and conditions are highly favourable for relativistic binaries (Faucher-Giguère & Loeb 2011). The dense nuclear star cluster surrounding Sgr A* (see e.g. Genzel, Eisenhauer & Gillessen 2010, for a review) contains a majority of older late-type stars, but contrary to expectations, massive young main-sequence stars (Ghez et al. 2003) and possible neutron star progenitors such as Wolf–Rayet stars (Paumard et al. 2001). The presence of neutron stars is further indicated by large numbers of X-ray binaries, possible pulsar wind nebulae, X-ray features such as the ‘cannonball’ and compact radio variables (Muno et al. 2005; Wang, Lu & Gotthelf 2006; Zhao, Morris & Goss 2013, 2020). Despite this only six radio pulsars have been discovered within half a degree of Sgr A* (Johnston et al. 2006; Deneva, Cordes & Lazio 2009; Eatough et al. 2013c; Shannon & Johnston 2013) even after many dedicated searches at multiple wavelengths (Kramer et al. 1996a, 2000; Klein et al. 2004; Klein 2005; Deneva 2010; Macquart et al. 2010; Eatough et al. 2013a; Siemion et al. 2013). Hyperstrong scattering of radio waves in the GC has been the principal explanation for the scarcity of detected pulsars (Cordes & Lazio 1997, 2002; Lazio & Cordes 1998a,b), however, scatter broadening measurements of PSR J1745−2900 in Spitler et al. (2014) and Bower et al. (2014) appear to contest this.1 Other authors have noted that the lack of GC pulsars is expected under a certain set of conditions and considering the sensitivity limits of existing pulsar surveys (Chennamangalam & Lorimer 2014; Liu & Eatough 2017; Rajwade, Lorimer & Anderson 2017). Alternatively, the scarcity of detected pulsars might be caused by a more complex scattering structure towards the GC (Cordes & Lazio 1997; Lazio & Cordes 1998a, b; Johnston et al. 2006; Schnitzeler et al. 2016; Dexter et al. 2017).

2021MNRAS.503..324M__Zhao_et_al._2019_Instance_3

We first determined the orbital parameters for RS Ser, V449 Per, and V1095 Her. Further, we updated the parameters for V593 Cen and MR Del. Using the formula f = (Ωin – Ω1)/(Ωin – Ωout), we calculated the contact factors f for RS Ser, V593 Cen, and V1095 Her as 6.5 per cent, 40 per cent, and 53 per cent, respectively. RS Ser is a contact binary with a small temperature difference of 131 K and a low contact factor. For V593 Cen, we updated the orbital parameters using more complete light curves. The orbital inclination of 83°.18 is similar to the result (82°.6) obtained by Zhao et al. (2019). The temperature of the secondary component (15 284 K) is higher than the previous result of 15 099 K. However, the mass ratio of 0.6 is lower than the previous result of 1.05 (Zhao et al. 2019). More spectroscopic observations are required to confirm the mass ratio. We confirmed that V593 Cen is an early-type contact binary with a deep contact factor as well as a black hole candidate. For MR Del, we revised the absolute parameters using its full light curve and the published radial velocities, which are similar to those published previously (Zhao et al. 2019; Pribulla et al. 2009; Djurašević et al. 2011). V1095 Her is also a contact binary with a deep contact factor of 40 per cent and a temperature difference of about 172 K. Looking over our four complete light curves, we found no evident starspot activity, and estimated variations exist over a long-term time-scale of years. V449 Per is an interesting target for detecting extra-solar and brown dwarfs using the minima timing variability of a low-mass eclipsing binary (Pribulla et al. 2012). Additional minima with higher precision are required to study its periodic variation further. Our physical parameters for RS Ser, V593 Cen, and V1095 Her are based on a light curve with a q-search-determined mass ratio. The nature of these parameters is speculative and preliminary. Radial velocities may eventually come to rescue them and provide a more definitive determination.

2022ApJ...935..137K__Tegler_et_al._1995_Instance_1

In Figure 2, the reduced AKARI IRC spectra of all protostars are presented; the absorption features of the H2O, CO2, and CO ices are clearly detected. All of our targets show deep and broad absorption features of H2O ice in the wavelength range 2.7–3.4 μm. In the case of AFGL 7009S, strong extinction toward the source saturates the absorption features throughout the wavelengths from 2.7 to 3.6 μm, including H2O ice. Other ice features, such as CH4 (Lacy et al. 1991; Boogert et al. 2004) and CH3OH (Grim et al. 1991; Brooke et al. 1996), were observed at 3.3−3.5 μm, but it is difficult to extract their absorption profiles from the blended features due to the low spectral resolution of AKARI IRC. An absorption feature of the CO2 ice around 4.27 μm was clearly detected toward all targets. At the wavelength around 4.6 μm for Perseus 1 and 3, RNO 91, and AFGL 7009S, there is a hint for another ice component overlapping with the CO absorption feature at 4.67 μm. Many near-infrared observations have revealed the same feature on the blue wing part of the CO absorption feature (Tegler et al. 1995; Chiar et al. 1998; Whittet et al. 2001; van Broekhuizen et al. 2005; Aikawa et al. 2012), which was suggested as an absorption feature of XCN ice. Lacy et al. (1984) and Pendleton et al. (1999) reported that the 4.62 μm absorption feature of XCN ice consists of a nitrile group and an unknown component “X”. Many laboratory studies have suggested that UV photolysis or cosmic ray irradiation of ice mantle could make the solid state OCN− on grain surfaces (Lacy et al. 1984; Grim & Greenberg 1987; Bernstein et al. 2000; Palumbo et al. 2000; Hudson et al. 2001; van Broekhuizen et al. 2004). In addition to these ice components, there are some absorption features around 4.8 and 4.9 μm. For Perseus 3 and the background star, the absorption features with a peak position around 4.78 μm are likely associated with 13CO ice (Boogert et al. 2002; Pontoppidan et al. 2003). We also detected another absorption feature at 4.83 μm toward the low-luminosity targets. However, we could not find any corresponding ice features from previous studies. The 4.9 μm absorption feature detected toward all targets was identified as solid carbonyl sulfide (OCS) ice. OCS ice can be produced when the interstellar ices containing CO and CO2 are exposed to UV photons or cosmic rays (Palumbo et al. 1997).

2020AandA...643A..35P__Irwin_et_al._2007_Instance_1

In order to achieve high photometric accuracy and be sensitive to low amplitude undulations, we adopted techniques from the exoplanet community, with the purpose of eliminating the systematic errors. When performing differential photometry (Sect. 3), accurate bias-subtraction and flat-fielding are of major importance. According to Irwin et al. (2007), the Poisson noise is 200 e− for a typical detector with a gain of a few e− ADU−1 and the flat illumination level is of 20 000 ADU pixel−1 = 40 000 e− pixel−1. Thus, a typical photometric aperture with a radius of 3 pixels contributes ∼1 mmag photon noise. For this reason, we obtained a considerable amount of biases (150−300 frames) and twilight flat-fields (25−100 frames) each night to reduce the Poisson noise to less than 0.2 mmag (Irwin et al. 2007). The bias frames were averaged together using the minmax in the reject option of the zerocombine task in IRAF with a view of keeping radiation events out of the master bias frame. The master flat frame was the result of combining all the frames using a median mode. The median value is an excellent way of removing the effects of hot pixels and cosmic rays, so these extreme values do not affect the calculation, as they would, if they would averaged. The reject option was set to avsigclip, in which case the “typical” sigma would have been determined from the data itself rather than an a priori knowledge of the noise characteristics of the CCD. Other related issues that can limit the photometric precision are: (i) the positioning of the telescope, (ii) fringing issues, and (iii) the differential variations on the quantum efficiency of the pixels. With the aim of minimizing the contribution of these effects, we repositioned each star almost on the same pixel of the detector using the autoguiding system of each telescope. The read-out-noise of the detectors are insignificant, as it can be as low as a few e− ( 10 e− for RISE2 and Andor Zyla cameras).

2021AandA...654A..88W__Bordoloi_et_al._2014_Instance_1

The CGM (see Tumlinson et al. 2017, for a detailed review) is now understood to be a key component in disentangling the feedback processes in active galaxies. It links the smaller-scale interstellar medium (ISM) of the galaxy to the larger-scale intergalactic medium (IGM), not only in a geometrical way but also by acting as the reservoir fueling star formation and the central black hole, where the feedback interacts with the galactic environment and where the gas recycling during galaxy evolution is controlled. This complex environment is multiphase and has been observed in numerous surveys (e.g., Tumlinson et al. 2013; Bordoloi et al. 2014; Peek et al. 2015; Borthakur et al. 2015) at low redshift. A prominent feature of the CGM around active galaxies is the Lyα (Lyman-α) emission line, which is also ubiquitously observed at high redshift (e.g., Haiman & Rees 2001; Reuland et al. 2003; van Breugel et al. 2006; Villar-Martín 2007; Humphrey et al. 2013; Cantalupo et al. 2014; Wisotzki et al. 2016, 2018; Arrigoni Battaia et al. 2018, 2019; Nielsen et al. 2020). Lyα is the transition of the hydrogen electron from the 2p orbit to its ground state. It can happen primarily through collisional excitation and recombination (see Dijkstra 2014, 2017, for a detailed review of Lyα emission mechanisms and radiative transfer). In extragalactic studies, the recombination production of Lyα emission can be generated by photoionization by young stars and/or AGN (fluorescence). This fluorescence emission on larger scales (CGM and IGM) can also be due to UV background radiation. Additionally, collisional excitation can play an important role in the emission seen in outflows and infalling gas (Ouchi et al. 2020). The bright Lyα emission line, along with other UV lines excited by the central or background sources, provides a useful tool for studying the galactic environments in the early Universe. Additionally, H I and metal absorption features observed in the CGM are powerful tracers of feedback signatures as well as tracers of infalling pristine gas (e.g., low metallicity absorption in a z ∼ 2.7 submillimeter galaxy; Fu et al. 2021). The sensitive integral field spectrographs on the largest ground-based telescopes, such as MUSE (Multi-Unit Spectroscopic Explorer; Bacon et al. 2010, 2014) and KCWI (Keck Cosmic Web Imager; Morrissey et al. 2012; see Cai et al. 2019 for observation of Lyα halos with KCWI), are perfectly suited for mapping these UV features as they move into the optical band for high-redshift sources.

2015AandA...584A..75V__Essen_et_al._(2014)_Instance_2

The data presented here comprise quasi-simultaneous observations during secondary eclipse of WASP-33 b around the V and Y bands. The predicted planet-star flux ratio in the V-band is 0.2 ppt, four times lower than the accuracy of our measurements. Therefore, we can neglect the planet imprint and use this band to measure the stellar pulsations, and most specifically to tune their current phases (see phase shifts in von Essen et al. 2014). Particularly, our model for the light contribution of the stellar pulsations consists of eight sinusoidal pulsation frequencies with corresponding amplitudes and phases. Hence, to reduce the number of 24 free parameters and the values they can take, we use prior knowledge about the pulsation spectrum of the star that was acquired during von Essen et al. (2014). As the frequency resolution is 1/ΔT (Kurtz 1983), 3.5 h of data are not sufficient to determine the pulsations frequencies. Therefore, during our fitting procedure we use the frequencies determined in von Essen et al. (2014) as starting values plus their derived errors as Gaussian priors. As pointed out in von Essen et al. (2014), we found clear evidences of pulsation phase variability with a maximum change of 2 × 10-3 c/d. In other words, as an example after one year time a phase-constant model would appear to have the correct shape with respect to the pulsation pattern of the star, but shifted several minutes in time. To account for this, the eight phases were considered as fitting parameters. The von Essen et al. (2014) photometric follow-up started in August, 2010, and ended in October, 2012, coinciding with these LBT data. We then used the phases determined in von Essen et al. (2014) during our last observing season as starting values, and we restricted them to the limiting cases derived in Sect. 3.5 of von Essen et al. (2014), rather than allowing them to take arbitrary values. The pulsation amplitudes in δ Scuti stars are expected to be wavelength-dependent (see e.g. Daszyńska-Daszkiewicz 2008). Our follow-up campaign and these data were acquired in the blue wavelength range. Therefore the amplitudes estimated in von Essen et al. (2014), listed in Table 1, are used in this work as fixed values. This approach would be incorrect if the pulsation amplitudes would be significantly variable (see e.g., Breger et al. 2005). Nonetheless, the short time span of LBT data, and the achieved photometric precision compared to the intrinsically low values of WASP-33’s amplitudes, make the detection of any amplitude variability impossible.

2019MNRAS.485.5453S__Sahal-Brechot_et_al._1996_Instance_1

Here we give a simple estimate for the population of 2s-state hydrogen atoms by considering a three-level system (1s, 2s, and 3p). We set the rate equation for the 2s population as (1) \begin{eqnarray*} n_{\rm H,1s}C_{\rm 1s,2s}+n_{\rm H,3p}A_{\rm 3p,2s}-n_{\rm H,2s}A_{\rm 2s,1s}=0, \end{eqnarray*} where nH, j is the number density of hydrogen atoms in the state j; Cj, k and Aj, k are the collisional excitation rate and spontaneous decay rate for the transition from j to k, respectively. For the bound states, we use the notation j = njlj, where nj is the principal quantum number of the state j. Similarly, lj = 0, 1, 2, 3, ..., nj − 1 (equivalently: s, p, d, f,...) is the orbital angular-momentum quantum number of the state j. Here we suppose that the depopulation term of 2s-state atoms is dominated by the spontaneous transition at the rate of A2s, 1s ≃ 8.2 s−1. In reality, the collisional transition from 2s to 2p can be a subdominant process for depopulation. For a collision at a velocity ∼108 cm s−1, which is a typical velocity scale for young SNR shocks, the cross-section is ∼10−13 cm2, giving a reaction rate ∼10−5 cm3 s−1 (e.g. Janev, Langer & Evans 1987; Sahal-Brechot et al. 1996). Thus, if the density is ∼106 cm−3, the collisional depopulation becomes important. Note that we assume no strong radiation field inducing the radiative transition from 2s to any other state.1 The occupation number of 3p, nH, 3p, depends on the absorption of Ly β. Here we assume an isotropic radiation field for Ly β. Then, we obtain the rate equation for 3p as (2) \begin{eqnarray*} n_{\rm H,1s} \left(C_{\rm 1s,3p}{+}\int _0^{\infty } \frac{4\pi \sigma _\nu ^{\rm 1s,3p}}{h\nu }I_\nu {\rm d}\nu \right) {-}n_{\rm H,3p} \left(A_{\rm 3p,1s}+A_{\rm 3p,2s}\right){=}0, \nonumber \\ \end{eqnarray*} where h, ν, $\sigma _\nu ^{\rm 1s,3p}$, and Iν are the Planck constant, frequency, absorption cross-section for the transition from 1s to 3p and the specific intensity, respectively. The intensity is set to be (3) \begin{eqnarray*} I_\nu =S_\nu (1-{\rm e}^{-\tau _\nu }) =\frac{ \frac{h\nu }{4\pi } A_{\rm 3p,1s} n_{\rm H,3p} }{\sigma ^{\prime } n_{\rm H,1s}} (1-{\rm e}^{-\tau _\nu }), \end{eqnarray*} where Sν and τν are the source function and optical depth, respectively. σ′ is a combination of physical constants relevant to the radiative absorption cross-section. Thus, we derive the occupation number of 2s as (4) \begin{eqnarray*} n_{\rm H,2s}=\frac{C_{\rm 1s,2s}}{A_{\rm 2s,1s}} \left[ 1+\frac{A_{\rm 3p,2s}}{ {\rm e}^{-\tau _0}A_{\rm 3p,1s}+A_{\rm 3p,2s} } \frac{ C_{\rm 1s,3p} }{ C_{\rm 1s,2s} } \right] n_{\rm H,1s}, \end{eqnarray*} where τ0 is the optical depth at the line centre. Here we assume a narrow line profile function ϕν for which we can approximate as $\int _0^{\infty }(1-{\rm e}^{-\tau _\nu }){\rm d}\nu \approx 1-{\rm e}^{-\tau _0}$. The terms in the brackets [...] indicate the contribution of the combination of the absorption and cascades. Note that roughly say, the ratios are A3p, 2s/(A3p, 1s + A3p, 2s) ≃ 0.118, C1s, 3p/C1s, 2s ∼ 2–10, and C1s, 2s/A2s, 1s ∼ 10−9np, where np is the proton number density. Thus, if Ly β is in the optically thick limit, nH, 2s is enhanced roughly at most 10 times compared with the optically thin case. The absorption coefficient of H α at the line centre becomes (5) \begin{eqnarray*} k_0({\rm H\alpha }) &amp;=&amp; \sigma _0({\rm H\alpha }) n_{\rm H,2s} \nonumber \\ &amp;&amp;{\sim} 10^{-23}\!-\!10^{-22}{\rm cm^{-1}} \!\left(\frac{T_0}{6000{\rm K}}\right)^{-\frac{1}{2}}\! \left(\frac{n_{\rm H,1s}}{ {\rm 1cm^{-3} } }\right)\! \left(\frac{n_{\rm p}}{ {\rm 1cm^{-3} } }\right)\!, \nonumber \\ \end{eqnarray*} where σ0(Hα) is the radiative cross-section of H α at the line centre for given temperature T0. Thus, if the SNR shock interacts with somewhat dense clump with a density of ∼30 cm−3 and a size of ∼1 pc, the H α emission can be scattered. Note that the H β emission can also be scattered but its absorption coefficient is about quarter of the H α coefficient. The interaction between the shock and a dense clump is implied by the ripple of an SNR shock with a length-scale of ∼10 per cent of SNR radius (e.g. Ishihara et al. 2010; Williams et al. 2013, 2016; Miceli et al. 2014; Sano et al. 2017; Tsubone et al. 2017, and see the discussion of Shimoda et al. 2015). Note that according to magnetohydrodynamic simulations performed by Inoue, Yamazaki & Inutsuka (2009) and Inoue et al. (2012), even if the shock propagates into a simulated ISM having density contrast ranging in ∼1–30 cm−3 as a consequence of thermal instability, the scale length of rippling is ∼10 per cent of the length of sides of simulation box.

2019AandA...622A..91G__Genzel_&_Stutzki_1989_Instance_1

Among the species studied in this work, the detection of broad (Δv > 30 km s−1) line-wing H2O (312-221) emission implies the presence of shocked gas activity (e.g., van Dishoeck et al. 2011, 2013). Indeed, we only detect H2 O and CH3OH rotationally excited lines toward Orion S and BN/KL (Figs. 2 and B.1). Both species are abundant in the ice mantles that coat grains in cold dark clouds (e.g., Gibb et al. 2004). After the onset of protostellar outflows, high-velocity shocks sputter these grain mantles and heat the gas to high temperatures. Both effects enhance the abundance of gas-phase H2 O and CH3OH (e.g., Draine 1995; Jiménez-Serra et al. 2008). In OMC-1, the low- and high-velocity outflows from BN/KL plunge into the ambient molecular cloud (Genzel & Stutzki 1989) producing hot (from Tk ≃ 200–2000 K) and dense (up to n(H2) ≃ 106–107 cm−3) post-shocked gas in H2 Peaks 1 and 2 (e.g., González-Alfonso et al. 2002; Goicoechea et al. 2015b). In our HIFI maps, these extreme conditions can be inferred from the moderately extended emission of the HCN J = 13–12 line (Fig. B.2), a rotational transition with a critical density close to 1010 cm−3 (ncr = Aul∕γul(Tk), where γul is the collisional de-excitation rate coefficient in cm3 s−1), around BN/KL outflows. Interestingly, the observed HCN to HCO+ J = 6–5 line intensity ratio is ≥2 toward BN/KL, and 1 almost elsewhere (see Fig. B.4, right). This may reflect a change in the chemistry between the extended PDR cloud component and the shocked gas in BN/KL outflows. It may also reflect the much stronger mid-IR (MIR) radiation from the BN/KL region that favors the radiative pumping of HCN through its vibrational levels and enhances the high-J rotational emission (e.g, Carroll & Goldsmith 1981; Ziurys & Turner 1986). Finally, the maps show that both HCO+ and HCN J = 6–5 lines display widespread emission outside the main star-forming sites (see Fig. B.1). This suggests that the gas density of the extended cloud layers traced by HCO+ and HCN J = 6–5 lines is moderately high.