adsabs/FOCAL huggingface dataset:
We derived stellar parameters, [C/M], and [N/M] using SLAM. To avoid bad fits at the edges of the parameter space, we exclude stars with spectral S/N in the g band less than 50, and metallicity less than −1.4. The derived C and N abundances are shown in Figure 8. Clearly, in the top panel, the CH-strong, CH-normal, and metal-poor field stars are separated, and their relative distribution in the N–C parameter space is similar to the case of APOGEE abundances (left panel of Figure 7): (1) metal-poor field stars form a sequence in the lower left of the top panel. As evolved stars ascend the RGB, C and N abundances may be changed by first dredge-up (Iben 1964, 1967) and extra mixing (Gratton et al. 2000; Charbonnel & Lagarde 2010). Given that for a typical halo/thick-disk star of 1 M⊙, the first dredge-up occurs around Teff = 5200 K (Boothroyd & Sackmann 1999), and most of our sample stars have 





 K and log g  2.5, we infer that most stars have already undergone first dredge-up. On the other hand, the C and N abundances of these stars could be altered by extra mixing. Stars with brighter K-band absolute magnitudes tend to have higher [N/Fe] and lower [C/Fe] (middle and bottom panels of Figure 8), which is consistent with extra-mixing theory and observation (Gratton et al. 2000; Charbonnel & Lagarde 2010); (2) CH-normal stars show an enhanced median N abundance and slightly depleted median C abundance. Clearly, the median N abundance of CH-normal stars is enhanced compared to that of normal metal-poor field stars with similar C abundances. In other words, the enhanced N abundances in CH-normal stars cannot be explained by the extra-mixing effect alone. We notice that a few CH-normal stars may have low N abundances, probably due to large uncertainties when spectra of a particular type are scarce in the training set, i.e., high-N metal-poor stars. The statistical similarity between APOGEE C and N abundances and LAMOST-derived C and N abundances further strengthens our statement above.
Citation: Charbonnel & Lagarde 2010
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Citation: Charbonnel & Lagarde 2010
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W51e2 is the strongest and best-studied HC HII region in the W51 Main cluster, and it is believed to be powered by an O8-type young star (e.g., Shi et al. 2010a). A number of interferometric studies conducted with varying angular resolutions, at centimetre (cm) and (sub)millimetre (mm) bands, identified molecular and ionized gas undergoing infall and rotation toward W51e2. VLA observations of the NH3 inversion lines (1, 1) and (2, 2) seen in absorption (1\hbox{$\farcs$}.̋1 beamsize) revealed radial infall on scales larger than 5000 AU toward the W51e2 core (Zhang & Ho 1997). Higher angular resolution observations of the (3, 3) NH3 absorption line (0\hbox{$\farcs$}.̋3 beamsize) showed signatures of rotation within 2000 AU based on a position-velocity (pv) diagram (Zhang & Ho 1997). Zhang et al. (1998) identified a velocity gradient in a CH3CN transition at 2 mm, deriving a position angle (PA) of 20 ± 20°. Keto & Klaassen (2008) imaged the H53α radio recombination line (RL) with the VLA (0\hbox{$\farcs$}.̋45 beamsize) and they claimed rotation in the ionized gas along the axis of a molecular bipolar outflow (oriented NW-SE) imaged with the SMA in the CO (2−1) line (1′′ beamsize), suggesting a simple inflow/outflow picture in a single high-mass young stellar object (YSO). However, higher resolution observations, using the SMA at the wavelengths of 0.85 mm (0\hbox{$\farcs$}.̋3 beamsize) and 1.3 mm (0\hbox{$\farcs$}.̋7 beamsize), revealed a more complex picture, by resolving W51e2 into three subcores (Shi et al. 2010a): W51e2-W, corresponding to the HC HII region, W51e2-E, located about 1′′ east of the HC HII region and corresponding to the brightest dust continuum source, and W51e2-NW, the weakest continuum component, located about 1′′  NW of the HC HII region. Shi et al. (2010b) imaged the CO (3−2) line (with a 0\hbox{$\farcs$}.̋7 beamsize) and established that the driving source of the powerful molecular outflow in this region is the protostellar core W51e2-E, and not the HC HII region W51e2-W, challenging the scenario proposed by Keto & Klaassen (2008). Etoka et al. (2012) used MERLIN to image the Class II 6.7 GHz CH3OH masers (typical signpost of HMSF), and found that the bulk of maser emission is indeed concentrated toward W51e2-E, and not the HC HII region W51e2-W. This further supports the scenario proposed by Shi et al. (2010a), where the ongoing star formation activity in the region is not concentrated on the HC HII region but toward its companion 1′′ to the east.
Citation: Zhang & Ho 1997
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Citation: Zhang & Ho 1997
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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: Rosen & Krumholz 2020
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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.
Citation: M20
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Citation: Magdis et al. 2020
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Citation: M20
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For these four remaining FHSC candidates (L1451-mm, MC35-mm, SM1N, and Oph A-N6) that have been observed at intermediate scales (few 100 au to few 1000 au) a final confirmation of their true evolutionary state requires higher resolution observations. For L1451-mm, the compact outflow needs to be resolved to investigate its morphology and kinematics, as a higher velocity component (an indication of protostellar nature) could be revealed by observations with a beam smaller than 100 au, similar to the case of B1b-N (Hirano 2019). An additional goal of high-resolution observations for L1451-mm and the remaining youngest candidates should be to investigate the temperature and density profiles of the envelope at scales from few au to 100 au. This is because simulations show that the temperature remains lower than ∼30 K even at several tens of au up to 100 au from the centre (Bate et al. 2014; Tomida et al. 2015; Hincelin et al. 2016; Young et al. 2019) during the FHSC stage. On the other hand, Class 0 sources show temperatures of 20–30 K or higher at scales of several 100 au (sufficient for thermal evaporation of CO) that results in the inner envelope and disc being easily detected using C18O observations (Yen et al. 2015, 2017; Stephens et al. 2018). This holds even in very low luminosity objects, for which the unexpected large extent of C18O is interpreted as evidence of a previous burst of accretion (Frimann et al. 2017; Hsieh et al. 2018). As for the density profile, simulations of the FHSC stage show a flat inner region, corresponding to the FHSC structure and extending up to ∼10 au (Tomida et al. 2013; Bate et al. 2014). For a protostar, on the other hand, the density profile should increase towards the central ∼1 au region (Young et al. 2019). Observations of the continuum emission with a resolution better than a few tens of au are likely required to model the emission and provide a density and temperature profile that can probe the relevant scales. Additional line observations with a similar resolution can also help to further distinguish between the different models. We note that, as pointed out in Young et al. (2019), distinguishing a dense core with only an FHSC and one that has recently formed a protostar but in which the FHSC structure is still present is likely not possible, even with high-resolution observations. Given the optically thick nature of the FHSC core, it is difficult to probe the physical properties within the FHSC structure. Despite this, finding a source with density and temperature profiles as well as with outflow properties consistent with the theoretical predictions will provide convincing evidence in support of a bona fide FHSC.
Citation: Bate et al. 2014
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Citation: Bate et al. 2014
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Other important components of the X-ray sky are X-ray binaries, the hot gas present in our own Galaxy and the extragalactic hot gas. Amongst these X-ray sources, the large hot gas reservoir (Tvir ∼ 107 K) filling the space between the galaxies in the clusters, known as the ICM, has been observed in its X-ray emission for a long time (Reichert et al. 1981; Jones & Forman 1984; Branduardi-Raymont et al. 1985; Oukbir, Bartlett & Blanchard 1997; Diego et al. 2003a; Diego, Silk & Sliwa 2003b; Cavagnolo et al. 2009; Hurier et al. 2015). However, the hot gas (Tvir > 106 K) present in the form of circumgalactic medium (CGM) in massive galaxies (Mh ∼ 1012–1013 h−1 M⊙; Birnboim & Dekel 2003; Kereš et al. 2005; Singh et al. 2015) is less explored in X-rays due to its fainter X-ray emission. Some of the recent observations and studies (Grcevich & Putman 2009; Anderson & Bregman 2011; Dai et al. 2012; Putman, Peek & Joung 2012; Anderson, Bregman & Dai 2013; Bogdán et al. 2013a,b; Gatto et al. 2013) indicate that the CGM can account for a good fraction of the baryons in these galaxies. The X-ray emission from the CGM, therefore, is a promising tool to put strong constraints on the distribution and energetics of the gas (Singh et al. 2016) with eROSITA. At energies above 2 keV, the extragalactic point sources like AGN completely dominate the X-ray sky (Lehmann et al. 2001; Kim et al. 2007). Even below 2 keV, where the X-ray emission from the hot gas in the ICM and CGM is significant, the major contribution to the observed X-ray sky comes from AGNs (Sołtan 2007). Therefore, studying the X-ray emission from the AGN is crucial to understand the origin and evolution of the AGN as well as to extract the X-ray signal from the subdominant components. We thus also compute the angular power spectrum of the unresolved AGNs that are expected to contribute to the diffuse X-ray background of eROSITA and contaminate the angular power spectrum due the ICM/CGM in the 0.5–2 keV X-ray band.
Citation: Grcevich & Putman 2009
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We explore another mechanism that can possibly result in torque reversals in the neutron star on long time-scales. It has been known for a long time that radio pulsars (which are usually isolated systems) show unexplained stochastic deviations in their spin-down behaviour (known as ‘timing noise’) on varied time-scales of a few hundred days to a few tens of years. This manifestation is akin to the random-walk behaviour in spin frequency observed in wind-fed accretion powered pulsars like Vela X-1. In an interesting study of timing irregularities of a sample of 366 pulsars, Hobbs, Lyne & Kramer (2010) found some radio pulsars showing quasiperiodic structures in their long-term timing residuals. From power spectrum analysis, significant periodicities ranging from about 1.4 to 10 yr were found in PSR B1540−06, PSR B1642−03, PSR B1818−04, PSR B1826−17, PSR B1828−11, and PSR B2148+63 (Hobbs et al. 2010). Interestingly, we also find quasiperiodic variations in the long-term spin evolution of Vela X-1 on time-scales of about 5.9 yr that is comparable to those inferred for radio pulsars showing quasiperiodic changes in their timing residuals. The underlying phenomena causing quasiperiodic structures in timing noise of radio pulsars is elusive. However, it has been suggested that these changes are driven by changes in the magnetosphere of the neutron star (Lyne et al. 2010). In this ‘state-switching model’, the magnetosphere of the neutron star is suggested to harbour two or more magnetospheric states which can be stable on time-scales of years but the pulsar can switch abruptly between these states driven by changes in the parameters regulating the spin-down (Lyne et al. 2010). This can possibly happen in Vela X-1 where the coupling between the dynamic magnetosphere and the neutron star can change in a quasiperiodic fashion. Interestingly, the disc-fed pulsar LMC X-4 has been found showing a near cyclic spin period evolution on time-scales of about 6.8 yr (Molkov et al. 2016) which is within a factor of 1.2 of the inferred time-scale in wind-fed pulsar Vela X-1. Recent observations of transient pulsar V0332+53 suggests switching of coupling between the accretion disc and the neutron star magnetosphere in a disc-fed pulsar (Doroshenko et al. 2017).
Citation: Hobbs et al. 2010
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Since the star-formation histories of galaxies (ETGs included, e.g. De Lucia et al. 2006; Maraston et al. 2009) can be stochastic and include multiple bursts, we also verify the full-spectrum fitting capabilities to retrieve more complex SFHs. In particular, we take an 11 Gyr old composite stellar population with an exponentially delayed SF (τ = 0.3 Gyr) as the main SF episode (this age is compatible with the age of the Universe at z ~ 0.15, which is the median redshift of our sample, see Sect. 3). We then define more complex SFHs by combining this single CSP with a burst of SF at different ages (5, 6, 7 Gyr) and with different mass contributions (3, 5, 10 %). In all cases, we consider a solar metallicity for the main SF episode and, according to the results of Maraston et al. (2009), a subsolar metallicity (Z = 0.004) for the later one. We do not mask any spectral feature of the input spectra, we assume AV  =  0.1 mag for the two components and apply a velocity dispersion of 200 km s-1. We show the results for a S/N of 80, which matches the typical S/N of the SDSS median stacked spectra analyzed in the following (see Sect. 3). Useful information can be derived from the comparison between the output SFH obtained from these input simulated spectra and the one provided when the single CSP alone is taken as input SFH. Fig. 5 shows that the single CSP alone is well recovered by the full spectrum fitting. In particular, ~80% of the stellar mass is retrieved within ~1 Gyr from the SFH peak. When a burst is added to this major episode of SF, the full-spectrum fitting is able to recognize the presence of a more complex SFH, as indicated by the tail appearing at smaller ages, and the total mass percentage of the later burst is retrieved within 1 Gyr from the expected age. However, we note that the main episode of SF is spread on a time interval longer than expected, and 50% of the stellar mass is retrieved around ~1 Gyr from the SFR peak. We also find that, in this case, the mean properties of the global stellar population are well retrieved, with a percentage accuracy larger than 10% starting from S/N ~ 15 for age, ~7 for metallicity, ~20 for AV, ~8 for σ and that the metallicities of the two SF episodes are separately recovered. These S/Ns are well below those typical of the stacked spectra analyzed in the following sections.
Citation: Maraston et al. 2009
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Citation: Maraston et al. (2009)
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The ReTOF results from irradiated phosphine and deuterated methane ice provide crucial information regarding the mechanism of formation for methylphosphanes by analyzing the intensities of various isotopologues. Figure 8 shows the possible formation routes that would lead to each of the three observed isotopologues of methylphosphine (CH3PH2). To obtain m/z = 50 (CHD2PH2), CD4 has to decompose via the loss of molecular hydrogen or two deuterium atoms to form carbene (CD2), which has been observed in previous irradiated ice studies (Holtom et al. 2005; Bennett & Kaiser 2007), and then insert into a phosphorus–hydrogen bond of phosphine (reaction (5)). If the carbene is formed in its first excited singlet state (a1A1), the insertion is barrierless (Gordon et al. 1987). For m/z = 51 (CD3PH2), methane and phosphine each lost a hydrogen or deuterium atom, and the resulting methyl (CD3) (Kaiser et al. 1997) and phosphino (PH2) radicals recombined barrierlessly (reaction (6)). Finally, the formation of m/z = 52 (CD3PHD) mirrors that for CHD2PH2 but in this case phosphine lost two hydrogen atoms or molecular hydrogen to create the phosphinidene (PH) radical and then inserted into a carbon–deuterium bond of methane (reaction (7)). Phosphinidene, like imidogen (NH) (Fueno et al. 1983), is expected to insert barrierlessly in its first excited singlet state (a1Δ):
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Therefore, our results provide compelling evidence that methane decomposes not only to the methyl radical, but also to carbene. Likewise, phosphine was found to fragment to the phosphino radical and also to phosphinidene. The ratio of ion intensities for m/z = 50:51:52 is 2:10:1, indicating that radical recombination was the preferred formation pathway with CD3PH2 as the most abundant isotopologue. This could either be a result of the methyl and phosphino radicals reacting quickly or that more of these radicals were produced than carbene and phosphinidene.
Citation: Gordon et al. 1987
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In addition to the galaxy colors (e.g., Baldry et al. 2004; Bell et al. 2004; Borch et al. 2006; Xue et al. 2010; Salim 2014; Lee et al. 2015; Wang et al. 2017), there are other indicators characterizing the star formation nature of galaxies, e.g., the morphology (Strateva et al. 2001; Barro et al. 2013, 2014; Pan et al. 2013) and some spectral features such as the Balmer absorption line (e.g., Kuntschner et al. 2002; Kim et al. 2013; Kim & Yoon 2017) and the 4000 Å break (e.g., Kauffmann et al. 2003; Lambas et al. 2012; Rowlands et al. 2018; Angthopo et al. 2019). Using these indicators, a number of studies have found that the color or SFR transition of galaxies is associated with the consumption of their gas content (e.g., Kereš et al. 2005; Dekel & Birnboim 2006; Kruijssen 2015; Nelson et al. 2018). If this “quenching” scenario is correct, the timescales of the galaxy color/SFR transition and their gas depletion must be consistent. Some authors have studied galaxy quenching by quantifying the evolution of SMFs (or luminosity functions) of SFGs and QGs at high redshifts (e.g., Fritz et al. 2014; Rowlands et al. 2018). The others, with the help of some infrared or submillimeter surveys (e.g., Hershel and ALMA), have managed to study the cosmic evolution of the gas content as well as the gas-depletion rate of galaxies with different SFRs (e.g., Geach et al. 2011; Lagos et al. 2011; Tacconi et al. 2018; Liu et al. 2019; Castignani et al. 2020; Magnelli et al. 2020). These studies have revealed that both the galaxy quenching and gas-depletion timescales increase with decreasing redshifts. The typical quenching timescale of GV galaxies can be as long as several billion years in the local universe (Rowlands et al. 2018; Correa et al. 2019; Phillipps et al. 2019). Methods such as fitting the spectral energy distributions (SEDs) of galaxies (e.g., Belfiore et al. 2018; Phillipps et al. 2019; Zick et al. 2018) and using cosmological simulations (e.g., Feldmann et al. 2017; Nelson et al. 2018; Correa et al. 2019; Donnari et al. 2019) have been widely used to derive the star formation histories of galaxies and their lifetimes.
Citation: Kereš et al. 2005
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