adsabs/FOCAL huggingface dataset:
As an additional constraint we have to make sure that our MCMax models are consistent with the observed CO snowline location. Therefore, we analyse the mid-plane temperature profile Tmid-plane(r) for all our models that match the observed SED, which we plot in Fig. 4. As mentioned above, models of the same series have the same temperature structure. We find that all of them have mid-plane temperatures between ∼20and25 K at the observed snowline radius of $R_{\rm sl} \approx 90$ au (Qi et al. 2015). These are well within the range of values generally assumed and observed for the freeze-out temperature of CO: The freeze-out temperature can vary between ∼20 and ∼30 K depending on whether CO is binding to pure CO ice or a mixture with water ice (Collings et al. 2004), which is, in turn, also dependent on the chemical history of the ice (Garrod & Pauly 2011). In general, the CO freeze-out temperature is not known unambiguously and might vary from system to system (Hersant et al. 2009; Qi et al. 2015): Qi et al. (2013a) found a freeze-out temperature of CO of 17 K from their modelling of TW Hya, whereas Jørgensen et al. (2015) obtain temperatures of about 30 K in their study of embedded protostars. Qi et al. (2011) assume a freeze-out temperature for CO of T ≈ 19 K (pure CO ice) for HD 163296. However, in Qi et al. (2015), they perform a new analysis with higher resolution observational data and use a temperature in the mid-plane at Rsl of T ≈ 25 K (mixed CO/H2O ice). Thus, all our models have temperatures in the disc mid-plane at the location of the snowline radius that are well within the plausible range. Our exploration of self-consistent models confirms that all the freeze-out temperatures assumed in these previous literature references fall within the plausible range of temperatures for HD 163296. If the freeze-out temperature of CO was known unambiguously, this would, in combination with an observationally determined CO snowline location, be a powerful model discriminant and we might be able to exclude models based on this constraint. Since, however, it is unclear what exactly the relevant freeze-out temperature is, we find that all the models can match the observed snowline location of 90 au. This weak model discrimination also means that it is impossible to predict the CO snowline radius from SED model fitting alone or even from fitting the molecular line emission together with the SED (Qi et al. 2011): given the uncertainty in the sublimation temperature of CO, our viable SED fits imply predicted radii in the range ∼40–135 au as denoted by the red asterisks in Fig. 4. In general, we find that the location of the CO snowline radius does not further discriminate between models in comparison to the criterion given by the SED; however, it sets the radial location inwards of which no freeze-out is taking place in our models, and which is therefore important for the interpretation of the C18O emission. It is important to note that an SED fit does not determine the CO snowline location and that, vice versa, a CO snowline observation does not discriminate amongst possible SED models.
Citation: Collings et al. 2004
Citation function: Uses
We futher examine the distributions of the radial number fraction of satellites with different colors in different LSE in Figure 5 to study the role central galaxies play in SCA. It can be seen, in general, that the intersections of these lines for each subsample are located between 





 to 





, and with a mean value of 





. Satellites (especially red satellites) prefer to reside in the inner part of groups and the number fraction decreases with the increasing radius. Such evolutionary behavior is consistent with “strangulation” wherein cold gas that could be used for star formation is depleted as satellites are accreted (e.g., Kauffmann et al. 2004; Kang et al. 2005; Luo et al. 2016). The shapes of each curve (the maximum value, the decrease gradient) not only depend on the color of central and satellite galaxies but also depend on the LSE. In the inner part of halos, the fraction of satellites is lower in the more dense LSEs. In the outer parts, it is the opposite, namely, in knots there is a higher fraction of satellites. It must be pointed out is that, in a given environment and for a given central color, the fraction of blue satellites is always higher than red satellites in outer regions. For wall environments, about 25% ∼ 30% satellites reside at the low values of 





, and the fraction decreases very quickly. Additionally, in the inner parts, the fraction of satellites of red centrals (red dashed line in the top left panel of Figure 5) is ∼5% higher than those in blue centrals (blue dashed line in the top left panel of Figure 5). The radial distributions of satellites in filaments are almost absent with the different colors of both central and satellite galaxies. However, in a knot environment, blue centrals have ∼3% more satellites than red centrals at low values of projected radius. These radial satellite distributions with different colors in different LSE indicate that, red central galaxies have merged more satellites than blue central galaxies in the knot LSE, which means satellites with a red central galaxy may be more affected by the interaction with central galaxies and represent a better SCA, especially in the inner region.
Citation: Kang et al. 2005
Citation function: Similarities
The GMC associated with G305 is one of the most massive and luminous clouds in the Milky Way (Fig. 1). It is located in the Galactic plane at l ~ 305°, b ~ 0° and at a kinematic distance of 4 kpc (derived from a combinationof radio and Hα observationsby Clark & Porter (2004); Davies et al. (2012) measured its spectrophotometric distance to be 3.8 ± 0.6 kpc and most recently Borissova et al. (2019) measured the Gaia DR2 average distance to be 3.7 ± 1.2 kpc); this places it in the Scutum-Crux spiral arm. Given this distance, the complex has a diameter of ~ 30 pc (Clark & Porter 2004) and a molecular mass of ~6 × 105 M⊙ (Hindson et al. 2010). The G305 complex consists of a large central cavity that has been cleared by the winds from massive stars belonging to two visible central clusters (Danks 1 and 2) and the Wolf-Rayet star (WR48a; Clark & Porter 2004; Davies et al. 2012). The cavity is surrounded by a thick layer of molecular gas (traced by CO and NH3 emission; Hindson et al. 2010, 2013). Radio continuum observations by Hindson et al. (2012) have revealed that the cavity is filled with ionized gas and identified six ultra-compact HII (UC HII) regions and also one bright rimmed cloud (BRC) at the periphery of the cavity, indicating molecular gas irradiated by UV radiation (Sugitani & Ogura 1994; Thompson et al. 2004), which may cause implosion (Bertoldi 1989) or evaporation. A number of studies havereported star formation tracers (water and methanol masers, HII regions and massive young stellar objects, MYSOs; Clark & Porter 2004; Lumsden et al. 2013; Urquhart et al. 2014; Green et al. 2009, 2012). Furthermore, Hindson et al. (2010) found the concentration of star formation tracers to be enhanced inside a clump of NH3 bearing molecular gas that faces the ionizing sources, which is consistent with the hypothesis that the star formation has been triggered. Analysis of the stellar clusters in the complex reveals them to have ages of 1.5 Myr for Danks 1 and 3 Myr for Danks 2,with the former possibly being triggered by the latter (Davies et al. 2012). Additionally, a diffuse population of evolved massive stars was also found to exist within the confines of the G305 complex that had formed around the same time as the two clusters (Leistra et al. 2005; Shara et al. 2009; Mauerhan et al. 2011; Davies et al. 2012; Faimali et al. 2012; Borissova et al. 2019).
Citation: Davies et al. (2012)
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Citation: Davies et al. 2012
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Citation: Davies et al. 2012
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Due to their non-axisymmetry, bars can cause the gas to flow inwards (e.g., Athanassoula 1992; Regan et al. 1999; Sheth et al. 2000), thus enhancing star formation in the central regions of their host galaxies (Sérsic & Pastoriza 1967; Hawarden et al. 1986; Ho et al. 1997; Sheth et al. 2005; Coelho & Gadotti 2011; Ellison et al. 2011). As a result, barred galaxies have various stellar structures, such as nuclear rings, nuclear disks, inner rings, outer rings, and boxy peanut bulges (e.g., Combes & Gerin 1985; Athanassoula 1992; Sellwood & Wilkinson 1993; Knapen et al. 2002; Kormendy & Kennicutt 2004; Comerón et al. 2010; Kim et al. 2012; Erwin & Debattista 2013; Buta et al. 2015; Emsellem et al. 2015, and references therein). These are dense regions in terms of stellar density. However, bars also develop sparse regions around themselves, usually producing a “Θ” shaped region. These sparse regions are thought to be produced by the bar-driven secular evolution in a way that the bar sweeps out materials around it. Such sparse regions are called by different names, such as “light deficit”(Kim et al. 2016), “star formation desert” (James & Percival 2015, 2016), and “dark spacer” (Buta 2017). In this respect, some barred galaxies have been found to be accompanied by a very faint disk or almost no disk around the bar (Gadotti & de Souza 2003; Gadotti 2008; Łokas 2021). Using Hα imaging, James et al. (2009) found that star formation is suppressed in the immediate vicinity of the bar (James & Percival 2016, 2018). The bar effectively transports gas into the central region once it is formed. Therefore, the star formation is truncated in the star formation desert. James & Percival (2016, 2018) measured the age of stellar populations and put constraints on the bar formation epoch. With zoom-in cosmological simulations, Donohoe-Keyes et al. (2019) found a truncated star formation and a lack of young stellar populations. However, they stressed that the interpretation can be complicated due to the radial migration of stars. Kim et al. (2016) found that the light deficit is more pronounced among galaxies with a longer and stronger bar. Later, Buta (2017) found that these sparse regions occur not only inside the inner ring, but also between the bar and the outer ring. He also found that the degree of the light deficit is a strong function of the bar strength, A
2. Therefore, the bar-driven secular evolution is recorded in the light deficit and how it changes with redshift is worthy of investigation.
Citation: Erwin & Debattista 2013
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Solar flares occur because of the explosive release of magnetic energy via the reconnection process in the solar atmosphere. Similar energetic events have been observed in other stars, especially M dwarfs like YZ CMi (dM4.5e), EVLac (dM3.5e), and AD Leo (Hawley & Pettersen 1991; Kowalski et al. 2010). The general idea is that for low-mass stars, the flare activities would be more frequent if the stellar masses are smaller. The level of flare effects depends also on the rotation periods (P) of the stars For example, the M-type dwarfs with P  10 days have been found to exhibit far more flare activity than those with P > 10 days, according to the Kepler observation (Lin et al. 2019). As P ∼ 1/





, with t being the stellar age on the basis of the gyrochronology relation (Skumanich 1972; Barnes 2010; van Saders et al. 2016), it is expected that the flare activity should be more frequent and powerful for young and fast-rotating stars (Davenport 2016; Stelzer et al. 2016; Davenport et al. 2019). Such time evolution of the stellar activities is important to the study of habitability of exoplanets orbiting around low-mass stars because of the potential impacts on the atmospheric loss and biological damage (Tarter et al. 2007; West et al. 2008; Segura et al. 2010; Armstrong et al. 2016). The large amount of high-precision photometric light curves produced by the Kepler mission (Borucki et al. 2010) is therefore ideal for statistical studies of stellar flares (Walkowicz et al. 2011). Following the discovery of superflares of solar-type stars by Maehara et al. (2015), there are increasing interests in topics related to exospace weather effects across a wide range of spectral types from A-type (Balona et al. 2015), G-type (Maehara et al. 2012, 2015; Wu et al. 2015), to M-type stars (Hawley & Pettersen 1991; Kowalski et al. 2010; Chang et al. 2017). There are two types of exoplanets of binary stars: the S-type for those in orbit around one of the stellar component, and the P-type for those orbiting around both components (Kaltenegger & Haghighipour 2013; Pilat-Lohinger et al. 2019). Their surface temperatures (and hence the condition of habitability) would be determined by the total radiative fluxes and magnetic activity of the binary systems. From this point of view, the binary systems are no less interesting because of the possible presence of the P-type or S-type habitable exoplanets.
Citation: Davenport et al. 2019
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Now, we must consider a specific distribution of electrons. Ideally, with a well-constructed background-subtracted image, only recently injected coronal material should be visible in excess total brightness; hence, ideally, ne should be identically zero outside the CME. However, reality is not so kind as to clearly delineate between CME and non-CME material. The explosive eruption and ejection of coronal material may rearrange previously quiescent material. If such material is pushed into a line of sight that passes through the CME, then, due to the optically thin nature of the corona, it will be impossible to distinguish newly ejected material from newly arranged material. A further issue that must be confronted when choosing a functional form for ne is the lack of observation or direct measurement on the internal mass distribution of a CME, as well as imprecise knowledge on the morphology of the CME. For example, is a CME balloon-shaped with a circular cross section (Millward et al. 2013) or croissant-shaped with an elliptical cross section (Thernisien et al. 2006, 2009)? Are there one or more cavities in the CME, and if so, where are the cavities located, how wide are they relative to the CME, and how low is the cavity density relative to the edge density? Is the mass distributed symmetrically relative to the central axis of the CME? Does the mass distribution change with time, either by pileup at the CME leading edge or by continuous mass inflow from the lower atmosphere (Bein et al. 2013; DeForest et al. 2013a)? Nevertheless, in spite of such epistemic uncertainty on the internal mass distribution, if the shape of the CME is known—either directly through polarimetry (de Koning & Pizzo 2011) or by assuming a shape (Thernisien et al. 2006; Millward et al. 2013)—and hence the start and end points of the line-of-sight integral are known, one might choose a particularly simple distribution, namely, 





 (Feng et al. 2013b). On the other hand, if both the CME mass distribution and morphology are poorly known, as was the case in the pre-STEREO era, then one might choose an even simpler distribution, that is, a point source,
12





where δ is the Dirac delta function and n0 is a constant. To enable direct comparison of our three-view mass determination with the two-view mass determination of CV09, we employ the point-source distribution.
Citation: Millward et al. 2013
Citation function: Uses; Uses; Uses
Citation: Millward et al. 2013
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The crucial aspect we want to test in this work is whether the CR-IGIMF scenario, once embedded in a realistic galaxy formation context, can match the observed trends. To this purpose, we use mean star formation histories (SFHs) predicted by the GAEA semi-analytic model (Hirschmann et al. 2016) for $z$ = 0 galaxies in four different stellar mass bins, namely M⋆ ∼ 1012, 1011.5, 1010.5, and $10^{9.25} \, {\rm M}_\odot$. These SFHs have been extracted from the reference GAEA realization and from runs implementing variable IMF approaches (F17 and F18). Galaxies more massive than ∼1010.5 are typically characterized by an early peak of SFR, followed by a smooth decay. The epoch, height, and width of the peak depend on the final galaxy stellar mass, with more massive galaxies having an earlier, narrower, and higher peak (see also De Lucia et al. 2006). Lower mass galaxies have almost constant SFHs. For the mean SFH in each mass bin, we compute the mass-weighted global IMF using our CR-IGIMF grid and then the corresponding value of fdg. It is worth noting that in our estimate for fdg we only consider the $\lt 1 \, \mathrm{M}_\odot$ mass range, where the mass-weighted IMF coincides with the present-day mass function. Moreover, observational constraints are closer to luminosity-weighted IMFs. However, for very old stellar populations (i.e. massive ETGs), we do not expect this difference to change our conclusions significantly. In Fig. 3, we compare these fdg for UCR/UMW = 1 (blue stars – for each mass bin, we plot separately the values corresponding to the mean SFHs extracted from the three GAEA realizations considered), with the observed values. We correct the observed values of fdg from LB13 to an aperture of 1 effective radius, whose properties we consider to be comparable to SAM predictions. To this aim, we assume an IMF radial profile with the same shape as that recently derived for one massive ETG by La Barbera et al. (2017). For each mass bin, the profile is rescaled, in the central region, in order to match the fdg value within the SDSS fibre aperture, as detailed in LB13. The correction, which is larger at the highest mass and negligible at the lowest mass,4 brings all fdg values from LB13 in the range from ∼0.72 to ∼0.8, as illustrated in Fig. 3 (see red dots with error bars, corresponding to the 18 stacked spectra of LB13). Remarkably, the aperture correction brings the observed fdg values within the range predicted by the CR-IGIMF.
Citation: De Lucia et al. 2006
Citation function: Similarities
To further test the possible role of mass in driving the high emission-line velocity-widths observed in our targets, in Fig. 14 we plot W80,Hα as a function of the emission-line ratio log ([N ii]/Hα). This emission-line ratio is a tracer of the metallicity of star-forming galaxies (e.g. Alloin et al. 1979; Denicoló, Terlevich & Terlevich 2002; Kewley & Dopita 2002), as well as an indicator for the source of ionizing radiation (e.g. Kewley et al. 2006; Rich et al. 2014). Furthermore, there is an observed relationship between mass and metallicity and therefore an expected relationship between this emission-line ratio and stellar mass (e.g. Lequeux et al. 1979; Tremonti et al. 2004; Maiolino et al. 2008; Stott et al. 2013). We can also make a crude prediction for the relationship between W80,Hα and mass, under the assumption that the line width is a tracer of the stellar velocity dispersion. Therefore, for a given stellar mass, we predict the position galaxies would be located in Fig. 14 by combining: (1) the observed z = 0.7 mass–metallicity relation (following Maiolino et al. 2008) and (2) the observed mass–velocity dispersion relationship for massive galaxies (following Bezanson, Franx & van Dokkum 2015). The majority of the star-forming galaxies appear to broadly follow the expected trend. Furthermore, the measurements from the stacked average emission-line profile (Fig. 11) are in agreement with the rough mass-driven prediction for the average mass of these galaxies [i.e. log (M⋆/M⊙) = 10.3; see Fig. 14]. In contrast, the AGN typically have higher log ([N ii]/Hα) emission-line ratios and higher velocity-widths than the star-forming galaxy sample (also visible in the stacked profiles; Fig. 11) and a positive correlation is observed between these two quantities. Such a positive correlation has been shown to be a tracer of shocks and outflows in the ISM through IFS observations of AGN and star-forming galaxies (e.g. Ho et al. 2014; McElroy et al. 2015). Interestingly, the small number of star-forming galaxies with high line widths (i.e. W80 ≳ 400 km s−1) appear to follow the same relationship as the KASHz AGN, which may indicate that these galaxies also have a contribution from shocks and/or host AGN that were not detected in the X-ray surveys. We clarify that some of the AGN targets have log ([N ii]/Hα) emission-line ratios that could also be photoionized by H ii regions (see Fig. 14) and we will explore these ideas further when exploring the spatially resolved outflow kinematics and emission-line flux ratios of the KASHz AGN in future papers.
Citation: Tremonti et al. 2004
Citation function: Uses
The close proximity in the sky of M 16 and M 17, two of the nearest giant HII regions of our galactic neighborhood lying at a similar distance from the Sun, naturally leads to the question of whether they are physically related, and whether they may share a common origin (Moriguchi et al. 2002; Oliveira 2008; Nishimura et al. 2017). Both giant HII regions are projected on the contour of a giant bubble-shaped structure, outlined in the distribution of HI and CO emission as first noted by Moriguchi et al. (2002). Evidence for triggered star formation in M 16 has been examined in detail by Guarcello et al. (2010), who concluded it had been induced externally, and not by the activity of its associated cluster NGC 6611. This suggests that the formation of M 16 and M 17 could have been triggered by the expansion of the bubble, powered by a previous generation of massive stars near its center, thus representing an example of triggered star formation on the scale of several tens of parsecs (Elmegreen 1998). Given the ages of the giant HII regions, the timescale of expansion of a wind-blown bubble, and the short lifetimes of massive stars, it is to be expected that the most massive members of that previous generation may have exploded as supernovae several Myr ago. The spatial dispersion of the members of the association that must have taken place progressively during its existence, combined with the distance of 2 kpc to the M 16–M 17 complex and the large number of unrelated foreground and background stars in that general direction, would make it extremely difficult to identify even its currently hottest members still remaining on the main sequence. Moriguchi et al. (2002) noted the presence of O and early B stars in the area and proposed that they were part of a massive star population responsible for having caused the bubble, but a review of their properties shows them to be generally too bright to be at the distance of the bubble and the giant HII regions, and instead are more likely members of a foreground population.
Citation: Moriguchi et al. 2002
Citation function: Motivation