adsabs/FOCAL huggingface dataset:
In our model, we compute the tidal response of the oceans and the solid-Earth to luni-solar semi-diurnal forcing, both combined with mimetic continental drift driven by plate tectonics. We focus on the dependence of dissipation on the Earth’s spin rate. We combine two analytical approaches that describe long-wavelength barotropic tidal flows over shallow spherical and hemispherical shells. The spherical shell describes a global ocean that we assume had existed in the earliest eons of the lifetime of the Earth (Motoyama et al. 2020). The existence of an early ocean is supported by evidence from the analysis of detrital zircon around 4.4 Ga (Wilde et al. 2001), from the interaction between the ocean and continental crust 4 billion years ago (Mojzsis et al. 1996), and from records of the oxygen isotope composition of seawater (Peck et al. 2001; Johnson & Wing 2020). The “globality” of this ocean is justified by the analysis of continental crust growth curves based on geochemical evidence in zircon crystallization ages (Dhuime et al. 2012; Hawkesworth et al. 2020). In compliance with these curves, we consider that a hemispherical oceanic shell has taken over in the most recent times. In our model, the center of this hemispheric continental cap follows the evolution of the paleogeographic center. In doing so, we emphasize on the role of “continentality” in the tidal response, while avoiding the under-sampling of geometric scenarios due to theoretical limitations (Hansen 1982; Tyler 2021) or due to uncertainties in plate tectonic models (Matthews et al. 2016; Daher et al. 2021). To compute this evolution, we adopt the recently developed paleogeographic reconstructions that cover the past billion years (Merdith et al. 2021). A postprocessing of these reconstructions allows us to produce the latitudinal evolution of the center of the continental cap captured in Fig. 1. The tidal frequencies at which oceanic resonances are excited and the amplitudes of these resonances vary with the surface position of the hemispherical ocean (Fig. B.1). Super-continental formations and breakups thus have their mark on the predicted lunar recession rate.
Citation: Peck et al. 2001
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Quiescent prominences are objects formed by relatively cool plasma in the hot corona. Their cool material occurs mostly in the dipped magnetic field lines. In quiescent prominences, which can persist from several hours to several days, magnetic dips form quasi-vertical structures called threads. Quiescent prominences are often observed as part of magnetic structures composed of three coronal patterns: the prominence itself, surrounded by a low-density cavity and a dense helmet streamer overlying the cavity (Engvold 1989). The disruption of the helmet streamer often signifies the beginning of a coronal mass ejection (CME) that can reflect the three-part structure of the helmet streamer (when observed in the white light): the bright leading shell surrounding a dark cavity in which bright prominence material occurs (Crifo et al. 1983; Hundhausen 1999). As the masses of CMEs are most often inferred from the visible-light observations, the masses given in literature are usually only the masses of the leading shell. However, Burkepile et al. (2004) also saw highly structured material most probably originating from a prominence in 63% of CMEs associated with eruptive prominences at the limb observed in visible light using the broadband filter (λλ 5000−5350 Å). Although such observations are called as white light in literature, it would be better to consider them as observations in visible light, because naturally integral intensities in the visible spectral range are observed including both continuum and absorption lines, while white light is just a continuum caused by Thomson scattering on free electrons without the lines. Thus, the presence of absorption spectral lines in visible light complicates very much the estimation of prominence mass, but even if the influence of absorption lines was eliminated, continuum alone only allows us to estimate the mass of ionised material, thus mass of the prominence would be underestimated. Low (1996) stated that the contribution of an erupting prominence to the total mass of CME is usually one order of magnitude less than that of the shell, but in some cases these contributions can be comparable. Low et al. (2003) proposed theoretically an importance of the prominence mass for the deposit of magnetic energy for driving a CME. Thus, developing methods for the estimation of the total mass of prominences can help to explain the connection between prominences and CMEs and provide an important quantity for more accurate estimation of total mass ejected by a CME.
Citation: Burkepile et al. (2004)
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Another possible explanation for the peculiar solar composition is that some of the dust in the pre-solar nebula was radiatively cleansed by luminous hot stars in the solar neighbourhood before the formation of the Sun and its planets. This dust-cleansing scenario is supported by the finding that the open cluster M67 seems to have a chemical composition closer to the solar composition than most solar twins (Önehag et al. 2011, hereafter O11 and Önehag, Gustafsson & Korn 2014, hereafter O14). They suggested that the proto-solar nebula was dust-cleansed by massive stars, similar to what happened for the proto-cluster cloud of M67, while the majority of solar twins in the field would presumably have formed in less massive clusters where no nearby high-mass star (≥15 M⊙) was formed. M67 offers the possibility for studying the solar-type stars in a dense cluster environment. This cluster has about solar metallicity ([Fe/H] in the range −0.04 to +0.03, e.g. Hobbs & Thorburn 1991; Yong, Carney & Teixera de Almeida 2005; Randich et al. 2006; Pasquini et al. 2008). The age of M67 is also comparable with that of the Sun: 3.5–4.8 Gyr (Yadav et al. 2008). Pasquini et al. (2008) listed 10 solar-twin candidates in M67, including M67-1194 which hosts a hot Jupiter with a period of 6.9 d and a minimum mass of 0.34 MJup (Brucalassi et al. 2014). M67-1194 was studied by O11 and re-visited by O14. They found that unlike nearby solar twins, which were systematically analysed by Meléndez et al. (2009) and Ramírez, Meléndez & Asplund (2009), the chemical composition of M67-1194 is more solar-like and therefore suggested that the Sun may have been formed in a similar cluster environment, perhaps even in M67. While this scenario is plausible when considering chemical abundances and ages, the dynamics are problematic: the Sun has an orbit close to the Galactic plane, while M67 is presently some 450 pc above the plane. The probability that the Sun, if formed in M67 with an orbit similar to the present cluster orbit, was scattered or diffused out of the cluster into the present solar orbit, is found to be quite low, if the existence of the outer planetary system is taken into account (Pichardo et al. 2012). Gustafsson et al. (2016) suggested that the high-altitude metal-rich clusters (such as M67) were formed in orbits close to the Galactic plane and later scattered to higher orbits by interaction with giant molecular clouds and spiral arms. Thus, it is possible, though not very probable, that the Sun formed in such a cluster before scattering occurred. Currently, only one solar twin in M67 (M67-1194) was spectroscopically analysed with high precision (∼0.02 dex). It is thus crucial to analyse the chemical composition of additional solar-type stars in M67 of very high precision.
Citation: Randich et al. 2006
Citation function: Motivation
Figure 9 summarizes our constraints on the molecular gas mass fraction in the analyzed samples of galaxies. In order to convert CO luminosity to molecular gas mass, we consider two different assumptions for αCO. First, we assume a constant value of αCO = 3.6 M⊙ (K km s−1 pc2)−1 adopted by some previous studies (e.g., Daddi et al. 2010b; Decarli et al. 2014; Walter et al. 2016). We then also consider a metallicity-dependent conversion factor, evaluated by assuming a redshift and stellar mass–metallicity relation (Genzel et al. 2015). Previous studies investigating optically and FIR-selected galaxy samples have estimated the relationship between gas mass fraction, stellar mass, redshift, and SFR offset from the main sequence (e.g., Genzel et al. 2015; Scoville et al. 2017). While the PHIBSS project estimated molecular gas masses by measuring the CO(3–2) line emission (Tacconi et al. 2013; Genzel et al. 2015), Scoville et al. (2016, 2017) used the flux on the Rayleigh–Jeans tail of the dust continuum emission to estimate the total gas masses. We here assume that the samples of galaxies plotted, although not complete to any degree due to their preselection for having a spectroscopic redshift, may be somewhat representative of the star-forming main sequence (Figure 9). Their star formation rates are consistent with scatter around the main sequence and appear to include as many galaxies above and below the main sequence estimated by Speagle et al. (2014). Although our CO detections (both blind and with previous spectroscopic redshifts) are indicative of gas fractions compatible with or above (GN19) expectations for main-sequence galaxies, the individual CO nondetections and the stacked signal appear to be systematically lower than the predicted averages, suggesting lower gas mass fractions than might be expected (Figure 9). We can quantify the apparent deficit in stacked signal relative to expectations for the 





 sample by calculating expected gas masses for the individual stacked galaxies predicted as a function of their redshift, stellar masses, and star formation rates. The expected sample average molecular gas mass is 





 adopting the best-fit relation by Genzel et al. (2015) and 





 according to the relation by Scoville et al. (2017). The constant CO luminosity conversion factor above would imply a ratio between expected and observed stacked CO luminosity of 4.8 ± 2.4 and 6.3 ± 3.1, according to the relations by Genzel et al. (2015) and Scoville et al. (2017), respectively. Applying instead the metallicity-dependent CO conversion factor suggested by Genzel et al. (2015) to individual galaxies would somewhat reduce the tension, implying ratios of 3.0 ± 1.7 and 3.8 ± 2.1, according to the relations by Genzel et al. (2015) and Scoville et al. (2017), respectively. While the constraints for low stellar mass galaxies may be compatible with an evolving CO conversion factor due to low metallicity, this is unlikely to resolve the apparent conflict at the high-mass end and may point to lower-than-expected gas masses.
Citation: Tacconi et al. 2013
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Another important test for galaxy formation models, is whether they place the right amount of stellar mass into discs and bulges. Moffett et al. (2016) measured the SMF separating galaxies into different morphological types and also into discs/bulges. With this, they derived the fractional contribution from bulges/discs to the total stellar mass in bins of stellar mass. This is quite a difficult measurement to do in observations as light profile fitting is required, which can be robustly done in very disc- and bulge-dominated galaxies, but when both components contribute similarly, the measurement is less robust (Robotham et al. 2017). In Fig. 11, we compare Shark with these measurements for two measurements of bulge mass.9 The first one is considering all the mass in the central concentration (regardless of whether it was formed due to mergers or disc instabilities; solid line), and the second one assumes that the bulge mass formed through disc instabilities (either through the starburst triggered by the gas being fueled to the centre or the stars that are transferred from the disc to the bulge) is part of the disc (dashed line). The bulge mass formed via mergers include both the stars formed via merger-driven starbursts and stars that were accreted by the bulge as a result of a merger (but that were formed in the disc of the primary and/or in the secondary galaxies). The latter is done as pseudobulges in Moffett et al. (2016) were added up to the disc rather than the bulge, and those are thought to form through secular processes taking place in the discs of galaxies (Kormendy & Kennicutt 2004). The effect of assigning the bulge mass formed via disc instabilities to the disc has the effect of shifting the transition from disc- to bulge-dominated stellar budget to higher stellar masses, much closer to the Moffett et al. (2016) observations. In Shark, we find that the formation of elliptical galaxies (i.e. spheroid dominated, massive galaxies) is dominated by galaxy mergers rather than disc instabilities, as disc instabilities only increase the bulge contribution by ≈10 per cent at $10^{11}\, \rm M_{\odot }$. This is qualitatively similar to the finding in large cosmological hydrodynamical simulations that elliptical galaxies form primarily via galaxy mergers (Wellons et al. 2016; Clauwens et al. 2018; Lagos et al. 2018a).
Citation: Robotham et al. 2017
Citation function: Motivation
Set II: we now focus on the observational constraints on the model parameters after the inclusion of the local measurement of H0 by Riess et al. (2018) with the previous data sets (CMB, Pantheon, and CC) in order to see how the parameters could be improved with the inclusion of this data point. Since for this present UM, the estimation of H0 from CMB alone is compatible with the local estimation of H0 by Riess et al. (2018), thus, we can safely add both the data sets to see whether we could have something interesting. Following this, we perform another couple of tests after the inclusion of R18. The observational results on the model parameters are summarized in Table 4. However, comparing the observational constraints reported in Table 3 (without R18 data) and Table 4 (with R18), one can see that the inclusion of R18 data (Riess et al. 2018) does not seem to improve the constraints on the model parameters. In fact, the estimation of the Hubble constant H0 remains almost similar to what we found in Table 3. In order to be more elaborate in this issue, we have compared the observational constraints on the model parameters before and after the inclusion of R18 to other data sets. In Figs 7 (CMB versus CMB+R18), 8 (CMB+CC versus CMB+CC+R18), 9 (CMB+Pantheon versus CMB+Pantheon+R18), and 10 (CMB+Pantheon+CC versus CMB+Pantheon+CC+R18), we have shown the comparisons which prove our claim. One can further point out that the strong correlation between the parameters μ and H0 as observed in Fig. 5 still remains after the inclusion of R18 [see specifically the (μ, H0) planes in Figs 7–10]. The physical nature of μ does not alter at all. That means the correlation between H0 and μ is still existing after the inclusion of R18 to the previous data sets, such as CMB, Pantheon, and CC. In addition to that since μ ≲ 0.9 according to all the observational data sets, thus, the transition from past decelerating era to current accelerating era occurs to be around z ≲ 0.6, similar to what we have found with previous data sets (Table 3).
Citation: Riess et al. (2018)
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Citation: Riess et al. 2018
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Citation: Riess et al. (2018)
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In Figure 8, we also find small differences in the [Na, Al, O/Fe] abundance ratios between the stars in the bright and faint RC groups, although it is not as clear as in the case of [Na, Al, O/H] abundances. In particular, unlike Figure 7, stars in the bRC group are more enhanced in [Na/Fe] but appear to be more depleted in [Al/Fe] and [O/Fe] than those in the fRC group. The mean differences are 0.053 ± 0.021 dex, 0.032 ± 0.018 dex, and 0.071 ± 0.045 dex in [Na/Fe], [Al/Fe], and [O/Fe], respectively, which are marginally significant at p-values of 0.22, 0.18, and 0.23. When the relative fraction of RC stars is taken into account (27%; see Section 4), the difference in [Na/Fe] between the genuine RC stars would correspond to Δ[Na/Fe] ∼ 0.20 dex, which is comparable to that expected from our chemical evolution model for the bulge stars (Δ[Na/Fe] = 0.2 ∼ 0.3 dex; Kim & Lee 2018; Lee et al. 2019).10

10
The previous study by Lee et al. (2019) noted a clear separation of the two groups according to Na abundance among bright RGB stars in the outer bulge. The apparent lack of such a distinct difference between the two groups in this study may be due to a larger uncertainty on abundances of relatively faint sample stars.
 The overall chemical patterns, however, are not identical to those observed in typical GCs, where the later-generation stars are more enhanced in [Na, Al/Fe] and more depleted in [O, Mg/Fe] than the first-generation stars at a given metallicity, although the trend of [Na, Al O/Fe] between the two RCs is less clear. Figure 9 shows the comparison of stars in this study with stars in metal-rich GCs ([Fe/H] > −1.0) on the Na–O diagram. The stars used in this study have a different distribution from stars in GCs. Although the bRC group is slightly more enhanced in [Na/Fe] and more depleted in [O/Fe] than the fRC group, the [Na/Fe] variation of RC stars is smaller than that of GC stars. This discrepancy might imply the different chemical evolution between stars in the bulge and typical GCs. We note, however, that even though we employ only metal-rich GC stars for the comparison, the majority of stars are still far more metal-poor ([Fe/H]  −0.5) than stars in the bulge. Because the relatively small [Na/Fe] variation is expected from the chemical evolution model for metal-rich bulge stars and the O-depletion is indistinct in some metal-rich GCs, such as NGC 6121 and 47 Tuc (see Kim & Lee 2018; Lee et al. 2019), the direct comparison of bulge stars with similarly metal-rich GCs on the Na–O plane would require further spectroscopic observations for such GCs in the bulge.
Citation: Lee et al. 2019
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Citation: Lee et al. 2019
Citation function: Similarities
Gaia DR2 contains unprecedented information about the Galaxy, which should allow us to infer its current structure, its equilibrium state, its evolution, modes of mass growth over time, dark matter distribution (and perhaps nature), to cite a few of the questions of modern Galactic astrophysics. As an example, it has been known for several decades that the Galactic disc contains large-scale non-axisymmetric features, including a central boxy/peanut-shaped bar (Okuda et al. 1977; Maihara et al. 1978; Weiland et al. 1994; Dwek et al. 1995; Binney et al. 1997; Babusiaux & Gilmore 2005; López-Corredoira et al. 2005; Rattenbury et al. 2007; Cao et al. 2013) and its possible in-plane extension (Hammersley et al. 2000; Benjamin et al. 2005; Cabrera-Lavers et al. 2007; Wegg et al. 2015), a warp (Burke 1957; Kerr 1957; Westerhout 1957; Weaver 1974; Djorgovski & Sosin 1989; Evans et al. 1998; Gyuk et al. 1999; Drimmel & Spergel 2001; López-Corredoira et al. 2002; Momany et al. 2006; Robin et al. 2008; Reylé et al. 2009; Amôres et al. 2017), and spiral arms (Georgelin & Georgelin 1976; Taylor & Cordes 1993; Drimmel 2000; Bissantz & Gerhard 2002; Churchwell et al. 2009; Vallée 2014; Reid et al. 2014; Hachisuka et al. 2015; Hou & Han 2015). However, full knowledge of these asymmetric structures, that is, of their spatial extent, pattern speeds, and number (in case of spiral arms) is still lacking. Since asymmetries constitute the driver of the secular evolution in galaxy discs (see e.g. Minchev et al. 2012; Fouvry et al. 2015; Halle et al. 2015; Aumer et al. 2017 and Kormendy 2013, for a review) by redistributing angular momentum between the inner and outer disc and between its baryonic and dark matter content (Debattista & Sellwood 2000; Bournaud & Combes 2002; Athanassoula 2003; Martinez-Valpuesta et al. 2006; Combes 2011), quantifying their characteristics is fundamental for understanding to what extent the Milky Way has “simply” evolved secularly in the last ~9 Gyr (Hammer et al. 2007; Martig et al. 2014), or whether some more complex evolutionary scenarios need to be invoked.
Citation: Westerhout 1957
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In order to understand the origin of the termination of the powerful jet in Mrk 231, consider the very low accretion rates in these FRI NLRGs, three to four orders of magnitude less than a quasar (Chiaberge et al. 1999, 2002; Hardcastle et al. 2009). However, there are the occasional FRI broad-line galaxies such as the Seyfert 1 galaxy, 3C 120, which has a very prominent FRI morphology on a 400 kpc scale (Walker et al. 1987). Thus, the low accretion rate of FRI NLRGs is not the full explanation of the discrepancy with Mrk 231, but is likely related. One difference between 3C 120 and Mrk 231 is that Mrk 231 has a low-ionization BAL wind (Lipari et al. 1994; Smith et al. 1995). In low radio states, it has also displayed evidence of a high-ionization X-ray-absorbing wind (Feruglio et al. 2015; Reynolds et al. 2017). There are also extreme amounts of intrinsic optical absorption in the galaxy itself from dusty gas (Lipari et al. 1994; Smith et al. 1995). All three circumstances point to a very dense nuclear environment through which the jet must propagate. Furthermore, in Reynolds et al. (2009), it was argued that the density of the free–free absorbing screen at K1 is consistent with the jet being stopped by the BAL wind at K1. By contrast, the extremely low, undetectable accretion of gas in FRI NLRGs is consistent with a very low-density nuclear environment. The launching of the jet in Mrk 231 does not seem to be stopped by the dense nuclear environment, but its propagation does seem to be thwarted. The case of 3C 120 does not seem to be accommodated by this discussion. It is noted that there is no evidence of either a BAL wind or a high-ionization X-ray-absorbing wind in 3C 120 (Oke & Zimmerman 1979; Ballentyne et al. 2004). Thus, it might be the case that there is not an extremely dense nuclear environment near the source of the jet in 3C 120; hence, jet propagation is not hindered. We cannot rule out the possibility that the BAL wind in Mrk 231 also diminishes the power of the jet-launching mechanism as well as providing a drag on its propagation.
Citation: Feruglio et al. 2015
Citation function: Background