adsabs/FOCAL huggingface dataset:
The EWs were measured with the methods described in Vásquez et al. (2015). As in Paper I, we used the sum of the EWs of the two strongest CaT lines (λ8542, λ8662) as a metallicity estimator, following the Ca II triplet method of Armandroff & Da Costa (1991). Different functions have been tested in the literature to measure the EWs of the CaT lines, depending on the metallicity regime. For metal-poor stars ([Fe/H] ≲ −0.7 dex) a Gaussian function was used with excellent results (Armandroff & Da Costa 1991), while a more general function (such as a Moffat function or the sum of a Gaussian and Lorentzian, G + L) is needed to fit the strong wings of the CaT lines observed in metal-rich stars (Rutledge et al. 1997; Cole et al. 2004). Following our previous work (Gullieuszik et al. 2009; Saviane et al. 2012) we have adopted here a G+L profile fit. To measure the EWs, each spectrum was normalised with a low-order polynomial using the IRAF continuum task, and set to the rest frame by correcting for the observed radial velocity. The two strongest CaT lines were fitted by a G+L profile using a non-linear least squares fitting routine part of the R programming language. Five clusters from the sample of Saviane et al. (2012) covering a wide metallicity range were re-reduced and analysed with our code to ensure that our EWs measurements are on the same scale as the template clusters used to define the metallicity calibration. Figure 5 shows the comparison between our EWs measurements and the line strengths measured by Saviane et al. (2012) (in both cases the sum of the two strongest lines) for the five calibration clusters. The observed scatter is consistent with the internal errors of the EW measurements, computed as in Vásquez et al. (2015). The measurements show a small deviation from the unity relation, which is more evident at low metallicity. A linear fit to this trend gives the relation: ΣEW(S12) = 0.97 ΣEW(this work) + 0.21, with an rms about the fit of 0.13 Å. This fit is shown in Fig. 5 as a dashed black line. For internal consistency, all EWs in this work have been adjusted to the measurement scale of Saviane et al. (2012) by using this relation. In Table 3 we provide the coordinates, radial velocities, and the sum of the equivalent widths for the cluster member stars, both measured (“m”) and corrected (“c”) to the system of Saviane et al. 2012.
Citation: Saviane et al. 2012
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Citation: Saviane et al. (2012)
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Citation: Saviane et al. (2012)
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The total energy released at an SN type II explosion, 





, exceeds, by two orders of magnitude, the energy transferred to the ejecta, 





. The bulk of energy escapes with the neutrino emission. Typically, the pulsar rotation energy is small compared to the ejecta energy. However, it cannot be excluded that at some circumstances the pulsar can receive a larger fraction of the explosion energy. When it becomes comparable to the kinetic energy of the ejecta, the sizes of both the SNR and the PWN can significantly deviate from Sedov’s solution. To illustrate the implication of this scenario, we considered a specific case of HESS J1825−137, a very extended and bright in VHE gamma-ray PWN (Aharonian et al. 2006; Pavlov et al. 2008; Mitchell et al. 2017a, 2017b; H.E.S.S. Collaboration et al. 2018a). The X-ray and gamma-ray observations of HESS J1825−137 show that the shape of the nebula deviates from the spherically symmetric geometry. Thus, the 1D model described in Section 2 has limited capability for accurate quantitative predictions. However, in the framework of the suggested scenario, the hydrodynamic processes evolve differently from the conventional situations, and the 1D model should provide a correct estimate for the energy required for the PWN/SNR inflation. In the framework of the scenario, the pressure in the PWN is the main driving force of the SNR expansion. Independently on symmetry of the system and dimension of the used model, the PWN is nearly isobaric, thus the radius of the pulsar wind TS provides an observational constraint for the pressure in the nebula. Therefore, we adopt the following requirement: the present-day pressure in the nebula should be consistent with the pulsar wind TS located at 





 from the pulsar (Van Etten & Romani 2011). The 1D model allows us to obtain the radii of the SNR, PWN, and pulsar wind TS as functions of three key model parameters: density of the ISM, pulsar braking index, and the initial angular velocity of the pulsar. In the framework of this model, it is possible to reproduce the three key properties of the system: 





, 





, and 





, provided that (i) the pulsar obtains, at its birth, a significant fraction of the explosion energy (





), and (ii) the braking index is small, 





. The conventional Sedov-like solution implies a low density of the ISM, 





. In contrast, the model proposed in this paper requires a rather dense environment, 





. This agrees with the presence of dense molecular clouds reported in the vicinity of the source (Voisin et al. 2016). What concerns the requirement for the pulsar braking index, 





, is that it is small compared to the value measured for the Crab pulsar (Lyne et al. 1993), but it matches both the old (e.g., Vela; n = 1.4 ± 0.2; Lyne et al. 1996) and young (e.g., J1833–1034; n = 1.8569 ± 0.001; Roy et al. 2012) pulsars.
Citation: Mitchell et al.; 2017b
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Three additional products were detected using the QMS during the TPD of the processed ice analogs in Experiments 1−3: C2O,3

3
No other carbon chain oxide molecules were monitored with the QMS.
 H2C2O, and H15NCO. The TPD curves of these product main mass fragments (m/z = 40, C2O; m/z = 14, H2C2O; and m/z = 44, H15NCO4

4
We assign the thermal desorption detected for the m/z = 44 mass fragment at T > 100 K to H15NCO because it is the most stable isomer among HOC15N, H15NCO, HC15NO, and HO15NC. In previous experimental simulations where HNCO ice molecules were produced, the metastable isomers were only tentatively detected, if at all (see, e.g., Jiménez-Escobar et al. 2014). In addition, gas-phase observations of different astrophysical environments reveal an abundance of the metastable isomers with respect to HNCO of less than 1% (Quan et al. 2010 and references therein).
 in Experiment 3 are also shown in Figure 3. The C2O IR feature at 1989 cm−1 (Palumbo et al. 2008) was only tentatively detected at the edge of the hatched region in the top panel of Figure 1. The H2C2O and H15NCO IR features at 2129 cm−1 and 2260 cm−1, respectively (van Broekhuizen et al. 2004; Hudson & Ferrante 2020), overlapped with different CO and carbon chain oxide IR features in the hatched region and could not be unambiguously detected. In order to confirm the assignments of the m/z = 40, m/z = 14, and m/z = 44 desorptions in Figure 3, to C2O, H2C2O, and H15NCO (the thermal desorption detected at T > 100 K), respectively, we energetically processed four additional ice analogs composed of different combinations of isotopically labeled H2, CO, and N2 molecules with 2 keV electrons (Experiments 4−7; see Table 1),  and checked that the same thermal desorption was observed for the expected mass fragments, considering the shift in the mass of the main fragments according to their isotopic composition in each experiment. The results are shown in Figure 4, confirming the formation of these species upon electron irradiation of H2:CO:N2 ice analogs.
Citation: Palumbo et al. 2008
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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.
Citation: Zhao et al. 2019
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Situated in the closest vicinity of their host stars (0.05 AU) and having no counterparts in our Solar system, ultra-hot Jupiters (Arcangeli et al. 2018; Bell & Cowan 2018; Parmentier et al. 2018) are ideal testbeds for studying the impact of 3D effects on high-resolution spectra. There are two important reasons for this. First, ultra-hot Jupiters are accessible objects to observe. Their short orbital periods (1–2 d) and hot, extended atmospheres make them perfect targets for transmission spectroscopy (Hoeijmakers et al. 2019; Von Essen et al. 2019; Ehrenreich et al. 2020; Borsa et al. 2021), emission spectroscopy (Evans et al. 2017; Arcangeli et al. 2018; Mikal-Evans et al. 2020) and phase-curve studies (Zhang et al. 2018; Bourrier et al. 2020b; Mansfield et al. 2020). Secondly, ultra-hot Jupiters display extreme variations across their atmospheres, because they are expected to become tidally locked soon after their formation (Rasio et al. 1996; Showman & Guillot 2002). As a result, their atmospheres virtually consist two different worlds: a permanently irradiated dayside and a permanently dark nightside. The scorching, cloud-free dayside (T ≳ 2500 K) nearly resembles a stellar photosphere, where most molecules are dissociated1 and metals become ionized (Parmentier et al. 2018; Hoeijmakers et al. 2019). On the other hand, the nightside is substantially cooler (T ≲ 1000 K) and may even serve as a stage for cloud formation (Helling et al. 2019; Ehrenreich et al. 2020). Ultra-hot Jupiters also exhibit large differences in their thermal structures: the dayside is expected to show strong thermal inversions (Haynes et al. 2015; Evans et al. 2017; Kreidberg et al. 2018; Pino et al. 2020; Yan et al. 2020), whereas nightside temperatures are expected to monotonically decrease with altitude. Furthermore, ultra-hot Jupiters feature strong winds in the order of 1–10 km s−1 (Tan & Komacek 2019), which arise as a result of the continuous day-night forcing. Many observational studies have measured Doppler shifts due to winds on ultra-hot Jupiters (Casasayas-Barris et al. 2019; Bourrier et al. 2020a; Cabot et al. 2020; Ehrenreich et al. 2020; Gibson et al. 2020; Hoeijmakers et al. 2020; Nugroho et al. 2020; Stangret et al. 2020; Borsa et al. 2021; Kesseli & Snellen 2021; Rainer et al. 2021; Tabernero et al. 2021), yet inferring the underlying 3D circulation pattern is a formidable challenge.
Citation: Borsa et al. 2021
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For the launching of a two-sided jet, the BZ power depends on the BH spin (assumed to be aligned with the disc) and on the magnetic flux threading the BH horizon at rh (e.g. Tchekhovskoy, Narayan & McKinney 2011):1(1)$$\begin{eqnarray*}
L \approx \frac{10^{-3}}{c} \Phi _h^2\Omega _h^2f(\Omega _h)\approx 10^{51} M_{\, {\rm BH},5}^2 B_{h,15}^2a_{-0.1}^2 \, {\rm erg}~\, {\rm s}^{-1},
\end{eqnarray*}$$where Qx denotes the value of the quantity Q in units of 10x times its c.g.s. units, except for Mx, which is given in units of $\, {\rm M_{\odot }}$. $\Phi _h = 4\pi r_h^2|B_h|$ is the integrated magnetic flux on one hemisphere of the BH horizon, |Bh| is the value of the radial contravariant magnetic field on the horizon, $\, {M_{\rm BH}}$ and a are the BH mass and spin, Ωh = ac/(2rh) is the angular velocity at the BH horizon, and f(Ωh) ≈ 1 + 1.38x2 − 9.2x4, where $x \equiv 0.5a\left(1+\sqrt{1-a^2}\right)^{-1}$. A second necessary condition for a successful launching of a BZ jet is that the BZ jet power will be sufficient to overcome the accretion power of the infalling material along the jet path (Burrows et al. 2007; Komissarov & Barkov 2009). Prior to the jet launching, matter is free-falling on the BH quasi-spherically with a power on one hemisphere of 
(2)$$\begin{eqnarray*}
\dot{M}c^2 = 4\pi r_g^2 \rho _h \beta ^r_h c^3\approx 2\times 10^{51} M_{\, {\rm BH},5}^2\beta ^r_h\rho _{h,7} \, {\rm erg}~\, {\rm s}^{-1},
\end{eqnarray*}$$where ρh is the mass density on the horizon at the time of the jet launching, and $\beta ^r_h$ is the dimensionless radial velocity of the infalling material at that time, which is expected to be close to unity. Comparing the jet luminosity from equation (1) with the accretion power from equation (2), we get a necessary condition for the strength of magnetic flux on the horizon that allows a successful jet launching 
(3)$$\begin{eqnarray*}
\Phi _h \gtrsim \, {\Phi _{h,\rm min}}\approx 7\times 10^{27}\frac{\sqrt{\beta ^r_h\rho _{h,7}}}{a_{-0.1}}~{\rm G~{\rm cm^2}},
\end{eqnarray*}$$corresponds to minimal magnetic field on the horizon $\, {B_{h,\rm min}}\approx 1.4\times 10^{15}$ G for the same normalization of parameters. Note that since a and $\beta ^r_h$ are both of order unity, the minimal magnetic field is primarily dictated by the central density of the star, which in turn depends on the stellar mass, radius, and density profile.
Citation: Komissarov & Barkov 2009
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In the left-hand panel of Fig. 2, we compare the morphological types assigned by the STAGES collaboration for the galaxies in the whole OMEGA sample and the jellyfish candidates sample. The sample of jellyfish galaxies (JC345) is composed mainly of late-type spirals and irregulars. In the middle panel of Fig. 2, we show the distribution of SED types for both samples. Based on the SED types of the galaxies, of the 70 jellyfish galaxy candidates analysed, 66 were found to be part of the blue cloud and 4 as being dusty reds (IDs: 11633, 17155, 19108, 30604). However, contrary to what could be expected, dusty red galaxies are only a small portion of our sample of jellyfish candidates. One reason why we may not detect many dusty reds as jellyfish galaxies might be due to the fact that these galaxies, despite having relatively high SFRs (only four times lower than that in blue spirals at fixed mass, Wolf et al. 2009), have significant levels of obscuration by dust that might hamper the identification of the jellyfish signatures. Another reason for that is that we selected jellyfish galaxy candidates within a parent sample of H α-emitting galaxies that already had a low fraction of dusty red galaxies (≈15 ${{\ \rm per\ cent}}$). As these galaxies have low star formation, it is harder to perceive the morphological features of RPS. Dusty red galaxies have been previously studied in this same system (Wolf et al. 2009), and RPS was suggested to be the main mechanism acting in these galaxies (Bösch et al. 2013). While in Bösch et al. (2013) one of the main pieces of evidence suggesting the action of enhanced RPS were the existence of disturbed kinematics without disturbed morphologies, in our study we strongly base our selection on such morphological distortions. Both our jellyfish galaxy candidates and the dusty red galaxies show different characteristics that can be correlated to the effect of RPS. Nevertheless, they might be tracing different stages of the same phenomenon, where dusty red galaxies have more regular morphologies, but disturbed kinematics. Our sample of morphologically disturbed jellyfish galaxy candidates may be showing the stage where the features of RPS are the most visible and the SFRs are enhanced.
Citation: Wolf et al. 2009
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Molecular outflows are a common and essential component in the formation process of low- and high-mass stars. In the past 40 yr, astronomers have mapped outflows in the whole mass range of young stellar objects (YSOs; e.g., Frank et al. 2014; Bally 2016; Anglada et al. 2018; Ray & Ferreira 2021). Magnetohydrodynamical (MHD) simulations have shown that the magnetic field plays a crucial role in the launching of molecular outflows (e.g., Pudritz & Ray 2019), more significantly so in the case of massive YSOs (e.g., Matsushita et al. 2018). Here, for instance, the presence of a magnetic field leads to the formation of early outflows. These reduce the radiation pressure, which allows the protostar mass to grow further (Banerjee & Pudritz 2007; Rosen & Krumholz 2020). In addition, the intensity of the magnetic field may influence the collimation of the outflows in massive YSOs. The outflows are well collimated for weak fields and poorly collimated for strong fields (Hennebelle et al. 2011; Seifried et al. 2012). In case of strong magnetic fields, the structure of the outflows is determined by the large-scale geometry of the magnetic field lines (Matsushita et al. 2017). Recently, Machida & Hosokawa (2020) have found a strong dependence of the evolution of outflows in massive YSOs on the initial magnetic field strength of the prestellar cloud for different accretion rates. In their 3D MHD simulations, they grouped the results into three categories: successful outflows, failed outflows, and delayed outflows. In the successful outflows, the outflows appear only when the prestellar cloud is strongly magnetized (μ1 = 2,3), and after an evolution time of ~104 yr, they reach a distance from the central protostar of about 104 au. When the magnetic field is weak (μ ≥ 5), we have failed and delayed outflows; even though small outflows of about 100–1000 au are observed in both cases, only in delayed outflows they can overcome the ram pressure and can ultimately grow. In a massive YSO, a large molecular outflow is therefore formed only if the initial magnetic field strength is B0 ≳B0,cr = 10−4(Mcl∕100M⊙) G, where Mcl is the cloud mass (Machida & Hosokawa 2020).
Citation: Seifried et al. 2012
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Other important result we reported in Section 3.3 is the reversing behaviour of red and blue galaxies with respect to velocity and groupcentric distances segregation, with redshift. Regarding velocity segregation, the preceding paragraph provides a qualitative scenario. Now, to explain the spatial segregation, we should notice that our analyses in Sections 3.2 and 3.3 take into account galaxies within 2R/R200. One can reasonably assume that such objects at lower redshifts correspond to a mixture of descendants of galaxies at higher redshifts in the same radii and of infalling objects from outer radii. Thus, both survival and replenishment of galaxies should be expected over the time, and two important factors come into play: (i) the accretion rate of galaxies; and (ii) the orbital dependence of galaxy properties (e.g. Biviano & Katgert 2004; Iannuzzi & Dolag 2012). Indeed, regarding velocity segregation, it has also been interpreted as red and blue galaxies having different kinds of orbits, with the orbits of blue galaxies being more anisotropic than the red ones (e.g. Biviano & Katgert 2004). Recently, Biviano et al. (2016) verified that the anisotropy profile of z ∼ 1 clusters is nearly isotropic near the cluster centre, and increasingly elongated with radius. This result is consistent with a halo evolution through an initial phase of fast collapse and a subsequent slow phase of inside-out growth by accrection of field material (e.g. Lapi & Cavaliere 2009). Since the accretion rate of galaxies from the field is higher at higher redshifts (e.g. McGee et al. 2009), our sample at z ∼ 0.8 is expected to be more affected by recent infalls, which had less time to go deeper into the group potential. This could explain the development of a more marked difference between the mean groupcentric distance of red and blue galaxies (see Fig. 12). After ∼3 Gyr, part of these infalling galaxies may reach the R  2R200 region, at z ∼ 0.4, mixing with virialized and backsplash objects, and thus presenting a less pronounced radial segregation between red and blue galaxies.
Citation: Biviano & Katgert 2004
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In order to determine the fundamental parameters of DM Ori, we take a similar approach to our V409 Tau study (Rodriguez et al. 2015), namely, to infer stellar parameters using stellar evolution isochrone models. However, unlike the V409 Tau analysis, we are very limited in the available data for DM Ori. We are only able to use the temperature of 





 = 6740 ± 250 K determined from the fits to the SED, as well as the star’s B- and V-band magnitudes (See Table 1). The longer wavelength photometric data for DM Ori are not useful due to contamination at these wavelengths from the disk. Assuming DM Ori is associated with the OB1c or the Orion Nebula Cluster portion of the Orion association (Bally 2008) allows us to place priors on some of the parameters. We assume uniform priors for distance (250–550 pc), metallicity (





 dex), and age (0–6 Myr). We also assume an extinction of 





 based on our best-fit SED. We use both the Dartmouth Stellar Evolution Models (Dotter et al. 2008) and the MIST stellar evolution models (Choi et al. 2016) combined with a Markov-Chain Monte Carlo (MCMC) to derive posterior distributions for the mass, age, radius, and distance of DM Ori (see Figure 4). From this analysis, we determine the most probable parameter values based on the posterior medians with errors from their 68% confidence interval uncertainties. For the Dartmouth models, we find a mass of 





 M⊙, an age of 





 Myr, a radius of 





 R⊙, and a distance of 





 pc. Similarly, for the MIST models, we find a mass of 





 M⊙, an age of 





 Myr, a radius of 





 R⊙, and a distance of 





 pc. Both sets of stellar evolution models produce consistent parameters for DM Ori, and the inferred age and distance are consistent with it being associated with the Orion star-forming region. To rule out the possibility that DM Ori is actually a foreground star in our line of sight to the Orion association, we reran our analysis with distance prior of 0–550 pc. Our results are very similar to the distribution shown in Figure 4, providing strong evidence that DM Ori is not a foreground star with a distance  250 pc.
Citation: Dotter et al. 2008
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