adsabs/FOCAL huggingface dataset:
Our results appear to stand apart from similar studies in full spherical shells (e.g., Nelson et al. 2013; Hotta et al. 2016) in that the differential rotation is strongly quenched as a function of the magnetic Reynolds number. However, in Nelson et al. (2013) the values of Rm′ (=2πReM) correspond to a range of 8...32 in ReM in our units where the radial and latitudinal differential rotation decrease by about 30 per cent. This is roughly consistent with our results. On the other hand, Hotta et al. (2016) reached higher values of ReM than in the present study, but no strong quenching was reported. The reason might be that their models are rotating substantially slower than ours, leading to weaker magnetic fields and a weaker back-reaction to the flow. Furthermore, in these models, the differential rotation is strongly influenced by their SGS heat flux, which transports one third of the luminosity. Another obvious candidate for explaining the difference is the wedge geometry used in the current simulations. However, we note that earlier simulations with a similar setup did not show a marked trend in the energy of the differential rotation as the azimuthal extent of the domain was varied (see Table 1 of Käpylä et al. 2013). However, results of Boussinesq simulations of convective dynamos have shown a similar change as a function of the magnetic Prandtl number (Schrinner et al. 2012). The drop in the amplitude of the differential rotation was associated with a change in the dynamo mode from an oscillatory multipolar solution to a quasi-steady dipolar configuration (cf. Fig. 15 of Schrinner et al. 2012) that prevents strong differential rotation from developing. We do not find a strong dipole component in our simulations (see Sect. 4.2.3). However, the strong suppression of the differential rotation often coincides with the appearance of a small-scale dynamo (see Table 1 and the discussion in the Sect. 4.1.2) or a change in the large-scale dynamo mode as discussed above.
Citation: Hotta et al. (2016)
Citation function: Differences; Differences
Already soon after the identification of the optical counterpart of Cyg X-1, X-ray light curves were found to show a strong orbital modulation of the X-ray absorption column, NH, due to absorption of X-rays from the black hole in the stellar wind (Li & Clark 1974; Remillard & Canizares 1984; Bałucińska-Church et al. 2000; Poutanen et al. 2008; Miškovičová et al. 2016; Grinberg et al. 2015). Together with observations of the orbital modulation of optical lines from HDE 226868 (Gies & Bolton 1986a,b; Gies et al. 2003), these phenomena led to the picture of the stellar wind of HDE 226868 as a line driven wind or CAK wind after Castor et al. (1975, see also Friend & Castor 1982 and Morton 1967) with an asymptotic velocity of v∞ ~ 2000 km s−1 (Muijres et al. 2012). This wind is disturbed by the gravitational potential of the black hole, which leads to a focusing of the wind toward the black hole; see Friend & Castor (1982) and Gies & Bolton (1986b) for early models and El Mellah et al. (2019) for a modern treatment ofthis process. In addition, the wind is also affected by the X-rays from the compact object: strong orbital modulation of NH is seen during the canonical hard state of the black hole, where the X-ray spectrum is dominated by a Comptonized power law (Parker et al. 2015; Nowak et al. 2011; Wilms et al. 2006, and references therein). The typical bolometric luminosity of Cyg X-1 in this state is around 2 × 1037 erg s−1 (e.g., Wilms et al. 2006; Nowak et al. 1999), albeit with a large uncertainty due to our lack of knowledge of the UV spectral shape. Only very little NH modulation is seen during the thermally dominated X-ray soft state (Wen et al. 1999; Boroson & Vrtilek 2010), in which the typical bolometric luminosity is at most a factor two higher than in the hard state (e.g., Tomsick et al. 2014; Zhang et al. 1997). Optical spectra show the stellar wind to be strongly photoionized during the latter state (Gies et al. 2003, 2008).
Citation: Gies et al. 2003; Gies et al. 2003
Citation function: Background; Background
Low- and intermediate-mass stars are formed by the gravitational collapse of the parental giant molecular cloud (GMC), followed by the accretion process (Palla 1996). During the accretion phase, material is ejected as well via collimated bipolar jets. However, when a YSO reaches 8 M⊙, the radiative flux becomes so intense (using ϕ = L/4πd2, the ratio between the radiative fluxes of an O5 and a B3 star – masses of ∼40 and ∼8 M⊙, respectively – is ≈250) that it may interrupt the accretion flow. A process that constrains the outcoming radiation field to narrower angles may leave some room for the accretion process to continue in some directions. This seems to be the case for the outflows driven by young stars from a very broad mass range, as previous reported by several authors (Bachiller 1996; Bontemps et al. 1996; Shepherd & Churchwell 1996; Beuther et al. 2002; Wu et al. 2004). Outflows associated with high-mass objects are expected to be more energetic than the outflows observed in lower mass YSOs (Beuther et al. 2005; Zhang et al. 2005; López-Sepulcre et al. 2009), with velocities greater than ∼100 km s−1 (Martí, Rodríguez & Reipurth 1998). Some authors have found evidences that outflows associated with massive stars are scaled up versions of their low-mass counterparts (Vaidya et al. 2011; Codella et al. 2013) while other works have reported that no well-collimated outflows have been found towards MYSOs (Shepherd, Testi & Stark 2003; Sollins et al. 2004). Massive YSO outflows mapped in high-velocity CO lines have collimation factors R = length/width ∼2.05 ± 0.96 as compared to R ∼ 2.81 ± 2.16 for low-mass stars (Wu et al. 2004), indicating a weak tendency that outflows associated with massive stars are less collimated than those from low-mass stars as previously thought (Richer et al. 2000). Besides the degree of collimation, these massive outflows would work removing mass from the plane of the accretion disc, lowering the density on the plane and, therefore, facilitating the accretion flow to reach the stellar core as shown in the recent 3D simulations presented by Krumholz et al. (2009). Although these authors have not included the outflow activity on their simulations, they argue that the presence of outflows would decrease the star formation efficiency from 70 per cent (considering purely radiation effects) to 50 per cent.
Citation: Wu et al. 2004
Citation function: Background
As we all know, sunspots that can represent solar activity have a long history of research. There are many researches on sunspot variation and solar cycle prediction. The prediction of the sunspot cycle can be classified as following three types: The first is the time series method. In this method, the sunspot number is regarded as a function of time for statistical prediction, in which the observed value or smoothed value of the sunspot number is regarded as a non-stationary random time series (Solanki 2003). The traditional approach is to transform the non-stationary into stationary time series, such as the famous McNish–Lincoln method (McNish & Lincoln 1949) and Box–Jenkins method (Box & Jenkins 1990). These kinds of the method have not yet found an ideal way to accomplish the transformation. It can only get the general trend of solar activity with the help of the periodicity of the series itself but cannot get satisfying prediction accuracy. The neural network method developed in the last decade is also used in forecasting. For example, Chattopadhyay et al. (2011) found that the effect of the auto-regressive neural network (ARNN) model in forecasting the solar cycle is better than the traditional time series method. With the sunspot number from 1749 to 2019, Pala & Atici (2019) used a long-term memory algorithm and neural network auto-regressive (NNAR) algorithm to predict the solar activity in the 25th cycle, showing that sunspots will reach the peak between 2022 and 2023. Although the neural network method sometimes performs better than many classical time series methods, however usually it is just like a black box without physical meaning and the performance also needs further validated. Similar to the above method is the time curve fitting method, Wilson (1988,1992) proposed the concept of long and short sunspot cycles, by using three sinusoidal curves to fit the sunspot number curve and extrapolate the predicted value. Some people also used other curves or multiple line regression methods to predict the sunspot number (Nordemann 1992; Hathaway, Wilson & Reichmann 1994). In particular, Hathaway et al. (1994) proposed a simple function (HWR function) to approximate the solar cycle profile with a certain predictive function. This kind of method aims at the profile prediction of a solar cycle, which is suitable for the prediction of the solar activity in the declining phase of the cycle, but the prediction for the peak value of the cycle is poor and not robust.
Citation: McNish & Lincoln 1949
Citation function: Background
It is expected that photons and/or cosmic rays coming from various sources could trigger the ejection of methanol from the icy dust grains into the gas phase and participate in the overall gas-to-ice balance of these cold regions. For the specific case of photons, this mechanism is known as photodesorption. Several experimental studies have been conducted to quantitatively constrain these processes in order to explain astrophysical observations, especially for methanol-containing ices. In one of these experiments, heavy ion 136Xe23+ irradiations were performed on methanol in pure ice and embedded in a water-ice matrix (Dartois et al. 2019). A sputtering yield of methanol close to that of the main water-ice matrix (Dartois et al. 2018), which is ~ 104 sputtered molecule per incident ion, was estimated for each studied ice. When it was embedded in a CO2 ice, Dartois et al. (2020) found that this sputtering yield is about six times higher. Experimental studies of UV photodesorption in the 7–10.5 eV range were first conducted for pure methanol ice at 20 K by Öberg et al. (2009). More recent experiments by Bertin et al. (2016) in the 7–14 eV range suggested an efficiency for methanol UV photodesorption from pure methanol ice of ~ 10−5 molecule/photon. This desorption was found to be below the detection threshold (10−6 molecule/photon) when methanol was mixed with CO ice for a wide range of dilution factors (from 1 in 4 to 1 in 50; Bertin et al. 2016). Accordingly, Cruz-Diaz et al. (2016) derived an upper limit of 3 × 10−5 molecule/photon for the UV photodesorption yield of methanol from pure methanol ice from 8 to 130 K (in the 6.88–10.9 eV range). In addition to these previous mechanisms, chemical desorption, which is the desorption induced by exothermic reactions, is a possible route for explaining the gas-phase enrichment in the ISM (Cazaux et al. 2016; Minissale et al. 2016b; Ligterink et al. 2018). However, chemical desorption of methanol by H addition onto CO, H2 CO, and CH3 OH ices was not detected (upper limit 5%; Minissale et al. 2016b), and more generally, chemical desorption of several other molecules appears to be very low when experiments are made on water substrates (Minissale et al. 2016a).
Citation: Dartois et al. 2018
Citation function: Uses; Uses
We model the NS structures using the EOSs summarized in Table 1. These have been chosen to span a representative range of maximum NS masses, but all with maximum gravitational masses 





 given current observational limits (Demorest et al. 2010; Antoniadis et al. 2013). For each of these EOSs, we solve for the hydrostatic structure including general relativity and compute a grid of NS models spanning a wide range of masses. We calculate the structures for both a non-rotating NS and a maximally rotating NS for use in treating the merger remnant (further discussed in Section 3). For the latter case, we assume solid-body rotation because magnetic braking eliminates differential rotation on an Alfvén timescale of ∼10–100 ms (Baumgarte et al. 2000; Shapiro 2000). All models have been computed using the LORENE library and in particular its publicly available codes rotseq and nrotstar. The former can easily compute sequences of non-rotating NSs, while the latter has been used to compute equilibrium configurations at the mass-shedding limit (Stergioulas 2003). The different EOSs were imported in LORENE in a table format either produced by using a piecewise polytropic approximation (APR4, Endrizzi et al. 2016; H4, Kawamura et al. 2016; MS1, Ciolfi et al. 2017), provided by the EOS authors (SHT; Kastaun et al. 2016), or given by the Compose project (GM16

6

http://compose.obspm.fr/

; Glendenning & Moszkowski 1991; Douchin & Haensel 2001). This large grid of models gives us, for each mass and EOS, a relation between the gravitational mass
1





where e is the energy density and R is the NS radius, and the baryonic mass (Shapiro & Teukolsky 1983)
2





where m(r) is the enclosed mass at a given radius and ρ is the rest-mass density (both equations refer to non-rotating models, but similar ones can be derived for rotating models). The important point is that in general 





 because of the gravitational redshift factor. In the fast rotating limit, we find that Mg is greater by ≈1.18–1.20, consistent with the calculations by Breu & Rezzolla (2016).
Citation: Shapiro 2000
Citation function: Uses
In the present era of “high precision” cosmology (see Riess et al. 2016), it is important to exploit the full potential of Cepheids as precise extragalactic distance indicators for determining the expansion rate of the universe and setting constraints on cosmology models. To achieve these goals, a deeper understanding and characterization of Cepheids is needed. Recent discoveries such as circumstellar environments (Nardetto et al. 2016), infrared excesses (Mérand et al. 2015), and ultraviolet emission line variability and possible more recent X-ray emissions (Engle et al. 2014) show that some important aspects of Cepheids may not be well understood. Cepheids have also been found to show additional complications that include cycle-to-cycle variations in their light and radial velocity curves (see Evans et al. 2015b; Anderson 2016; Anderson et al. 2016; Smolec & Śniegowska 2016; Derekas et al. 2017). These newly discovered properties and time-dependent phenomena of Cepheids, unless better understood and accounted for, could place impediments on achieving the challenging goal of determining the local Hubble constant (H0) with a precision of ∼1%, as suggested by Suyu et al. (2012). Great efforts are being undertaken to achieve this level of precision, and hopefully resolve the developing “Hubble Discrepancy” (see Riess et al. 2016), where theoretical values of the Hubble constant (H0) derived via the Lambda-cold dark matter (Λ-CDM) cosmology model (Λ = the cosmological constant), including cosmic microwave background data (e.g., Planck, WMAP (Wilkinson Microwave Anisotropy Probe)), show a small (albeit statistically significant) disagreement with the value of H0  will be derived via standard candles (e.g., Cepheids, SNe Ia (type Ia Supernovae)). As discussed by Suyu et al. (2012) and more recently by Riess et al. (2016), improved measurements of H0 provide critical independent constraints on dark energy and the validity of the present Λ-CDM model.
Citation: Engle et al. 2014
Citation function: Motivation; Motivation; Motivation
Several scenarios have been suggested to explain the presence of magnesium-enriched gas in between the galaxies of a group. Magnesium atoms are α elements released in the ISM and CGM by core collapse supernovae. One can therefore naturally imagine that the magnesium gas residing outside galaxies has been ejected through strong supernovae-driven outflows. The P Cygni shape of the Mg II doublet in the integrated spectrum of galaxy A (see Fig. 2 top) is in good agreement with this scenario. Indeed, according to radiative transfer (RT) models, this line profile is a signature of resonant scattering in an optically thick outflowing medium where the redshifted emission corresponds to back-scattered photons reflected by the receding medium (Haiman et al. 2000; Dijkstra et al. 2006; Verhamme et al. 2006; Kollmeier et al. 2010). The Mg II lines of galaxies D and E seem to be redshifted compared to their systemic redshifts (Fig. 2), however, the continuum is too faint to detect significant absorption lines and, thus, too faint to detect the P Cygni profiles, namely, the presence of outflows for those galaxies. We detected > 2σ Mg II emission aligned with the minor axis of galaxy E and near galaxy D (Fig. 3). These regions could have been enriched via outflows after supernovae feedback. However, the S/N of our data is too poor in these areas to carry out a resolved kinematics study. The resolved Mg II map (Fig. 6) shows that the two regions along the minor axis of galaxy A have different velocities with opposite signs (≈ + 120 and −80 km s−1 with respect to the systemic redshift of galaxy A), providing additional evidence of the presence of an outflow. Even if we do not eventually find clear evidence of an outflow in the four other galaxies, we cannot exclude the deposition of metals in the CGM by past feedback processes explaining the observed diffuse Mg II emission detected around each of the galaxy members. In that scenario, we could imagine that the low-surface-brightness gaseous bridge in-between the galaxy subgroups [A,B,C] and [D,E] could be due to galactic wind transfer between the group members as observed in the FIRE simulation (Anglés-Alcázar et al. 2017).
Citation: Haiman et al. 2000
Citation function: Similarities
Small-amplitude oscillations are observed without any clearly associated driver and can persist for several hours (Wang et al. 2012; Nisticò et al. 2013; Anfinogentov et al. 2015). Their amplitudes are of the order of 100 km and they show no observable damping. Even though it is clear they must be driven by some small-scale abundant process, the exact excitation mechanism of this oscillation regime is unclear. Several excitation mechanisms have been proposed to explain the persistent nature of these oscillations. These include mechanisms acting at coronal heights such as the onset of thermal instability in the corona (Kohutova & Verwichte 2017; Verwichte & Kohutova 2017), the presence of siphon flows in the coronal loop (Kohutova & Verwichte 2018), self-oscillation of coronal loops due to interaction with quasi-steady flows (Nakariakov et al. 2016; Karampelas & Van Doorsselaere 2020), and footpoint-concentrated drivers. A commonly assumed excitation mechanism is associated with turbulent vortex flows which are ubiquitous in the lower solar atmosphere (Carlsson et al. 2010; Shelyag et al. 2011; Liu et al. 2019) and result in random footpoint buffeting. This is typically modelled as a stochastic driver at the footpoints (e.g., Pagano & De Moortel 2019). Such a driver however does not reproduce the observational characteristics of this type of oscillation, in particular the stability of the oscillation phase and lack of any observable damping (Nakariakov et al. 2016). It has been proposed that a footpoint driver of the form of a broadband noise with a power-law fall-off can in principle lead to excitation of eigenmodes along coronal loops (Afanasyev et al. 2020), but this has not been tested through MHD simulations. Finally, excitation of transverse oscillations through solar p-mode coupling has also been proposed in several studies (Tomczyk & McIntosh 2009; Morton et al. 2016, 2019; Riedl et al. 2019). As p-modes correspond to acoustic waves generated by turbulent flows in the solar convection zone, they are subject to acoustic cut-off at the β = 1 equipartition layer. High magnetic field inclinations can however lead to p-mode leakage into chromosphere even for frequencies below the cut-off frequency (De Pontieu et al. 2004), and to subsequent coupling to different magnetoacoustic wave modes (Santamaria et al. 2015). However, three-dimensional (3D) MHD simulations of coronal fluxtubes with a gravity-acoustic wave driver at the footpoints did not reveal any evidence of the acoustic driving leading to significant transverse displacement of the fluxtubes at coronal heights (Riedl et al. 2019).
Citation: Verwichte & Kohutova 2017
Citation function: Background
The expected pressure profile of the interstellar medium inside the Bondi sphere is p ∝ z−2 and, according to analytical and numerical models (Tchekhovskoy et al. 2008; Komissarov et al. 2009; Beskin et al. 2017), a jet propagating in a medium with such a pressure gradient develops the characteristic parabolic shape. Based on Eq. (3), estimating the Bondi radius requires the knowledge of the black hole mass and of the temperature of the accretion flow in the nuclear region. A detailed analysis of the X-ray emission in NGC 315 was carried out by Worrall et al. (2003, 2007) using sensitive Chandra data. The jet, which presents several X-ray-bright knots of synchrotron emission, is embedded in a hot gaseous atmosphere, as inferred from the presence of an X-ray thermal emission component in the spectrum. Such a feature is also detected within a circle of 1 arcsec radius centered around the nucleus, in addition to a dominant, mildly absorbed power-law component possibly associated to the jet. The nuclear hot atmosphere has a temperature 



k
T
=
0
.

44

−
0.04


+
0.08




$ kT=0.44^{+0.08}_{-0.04} $


 keV. Assuming a black hole mass of 1.3 × 109 M⊙, we then estimate rB = 92 pc (dark blue vertical line in Fig. 3). Even if the black hole mass was smaller by a factor of several, this radius would still be much larger (two orders of magnitude) than the distance at which we observe the jet shape transition from parabolic to conical, zt = 0.58 ± 0.28 pc. This result would also not be affected by varying the gas temperature within the given small uncertainty. By assuming slightly different parameters, a similar value for the Bondi radius in NGC 315 was recently obtained by Inayoshi et al. (2020), who estimated an uncertainty of 50%3. Then, if the transition from parabolic to conical shape in NGC 315 is induced by a change in the external pressure gradient, this must occur not in the proximity of the black hole sphere of influence, but in the vicinity of the black hole, on sub-parsec scales. Based on the analysis of the VLA data in Fig. 3, right panel, we also note that no discontinuity is observed in the expansion profile after the jet crosses the Bondi radius: the jet shape is close to conical both inside and outside the Bondi sphere.
Citation: Worrall et al.; 2007
Citation function: Uses; Uses