adsabs/FOCAL huggingface dataset:
UGC 4284 is classified as Scd(s) in the RC3 (de Vaucouleurs et al. 1991), and as SABc by Paturel et al. (2003). Kandalian & Kalloghlian (1998), Chapelon, Contini & Davoust (1999), and Cabrera-Lavers & Garzón (2004) found evidence of a central isophotal distortion that they explain postulating the existence of a bar with diameter of 38 arcsec. Other authors did not confirm the presence of a bar in UGC 4284, nevertheless they encountered non-circular and off-plane gas motion in the inner region of UGC 4284 that they discuss as originated by a warped gaseous disc (Coccato et al. 2004; Józsa 2007). Bearing in mind the purpose of the present study, it is important to determine which kind of baryonic components do actually exist for UGC 4284, and given the incongruity of the above-reported results we required to perform 2D isophotal analysis of the surface brightness of this galaxy, using galfit (Peng et al. 2002, 2010) and considering at least three principal components: disc, bulge, and bar. UGC 4284 has been observed with Hubble Space Telescope (HST) in the filters F555W (WFPC2 V), F814W (WFPC2 wide I), and F606W (WFPC2 wide V), none the less the filters F555W and F814W show only a small part of the periphery of the galaxy. In contrast the filter F606W exhibits a larger portion of the disc of UGC 4284, so we use that filter to decompose the disc of UGC 4284. The results of the 2D exponential disc fit are presented in Table 1. In the case of UGC 4284, we also examined the possibility of a bulge and a bar components given the inconsistency in the morphological classification of this galaxy. We use the Spitzer IRAC 4.5 μm image because UGC 4284 is indiscernible in the 2MASS J, H, K bands, and in all the SDSS filters are too faint to allow an obvious separation from the background. Neither the Sérsic component nor the Ferrers (Ferrers 1877) component produces acceptable fit parameters to reveal the presence of a bulge and a bar, and for this reason we conclude that UGC 4284 has only an exponential disc component in total agreement with the RC3 classification.
Citation: Peng et al.; 2010
Citation function: Uses
The idea of ionized material arising from solar flares and causing geomagnetic storms appears to date from around the 1930s (Hale 1931; Dellinger 1937; Newton 1939, 1943; Chapman 1950) and the idea persisted into the space age, where interplanetary shocks were attributed to being caused by flares (e.g., Sonnet et al. 1964; Gosling et al. 1968). By the 1970s, however, some workers had begun to express doubts that solar flares could be the source for interplanetary shocks, geomagnetic storms, and, after their discovery (Tousey 1973), CMEs observed by coronagraphs. Evidence against the flare being the source included: the absence of a 1:1 relationship between geomagnetic storms and flares (e.g., DeMastus et al. 1973; Gosling et al. 1974; MacQueen 1980; Joselyn & McIntosh 1981); the energy required to launch the mass ejection was much larger than that of the flare (e.g., MacQueen 1980; Webb et al. 1980); the location of the flare being at only one footpoint of the larger expanding structure (Harrison & Simnett 1984; Simnett & Harrison 1984, 1985); the lack of correlation between the time of the projected onset of the CME and that of the flare (Harrison et al. 1985, 1990; Harrison & Sime 1989); and a lack of lateral movement of the expanding bright front as its leading edge expanded outwards (Sime et al. 1984). These and other measurements were presented in a review paper by Kahler (1992) and another by Gosling (1993a), the latter now well known as the “solar flare myth” paper. In that paper, flares are described as being either connected to CMEs via a common primary physical process (labeled in his Figure 16 as “evolving solar magnetic fields”), or as being caused by the CME launch. By around 1996 Gosling’s “modern paradigm” was widely accepted as the correct general physical model for CMEs. This picture has changed little in the last two decades, apart from perhaps the general acceptance that a magnetic flux rope is the driving core of a CME, although this too is an idea dating from before the space age (e.g., Morrison 1954; Cocconi et al. 1958; Gold 1959).
Citation: DeMastus et al. 1973
Citation function: Background
Both for the hydrodynamical and the synthetic simulations, a blind analysis was performed on the simulated images. First, surface-brightness profiles were extracted by computing the azimuthal median in concentric annuli. The profiles were then deprojected using a non-parametric onion-peeling technique (Kriss, Cioffi & Canizares 1983; McLaughlin 1999) and converted into radial density profiles. For the synthetic simulations, we found a very good agreement between the true and the reconstructed density profiles, in all of the situations considered here. Even when adding as many as 103 randomly distributed clumps in the simulated image, our method is able to recover the underlying density profile with an accuracy of a few per cent at all radii. The addition of significant triaxiality to the halo did not affect our results (in agreement with the findings of Buote & Humphrey 2012), and similar conclusions were reached for the inclusion of a filamentary structure. In the case of the hydrodynamical simulations, we show in Fig. 1 the average surface-brightness profiles (in arbitrary units) obtained using our azimuthal median method and with the standard azimuthal mean. We can see that the azimuthal median results in a much more regular profile, showing the expected smooth radial decline. In Fig. 2 we show the comparison between the average true 3D density profile in the sample of Vazza et al. (2011b) and the average profile obtained with our method through deprojection. On the same figure, we also show the same comparison in the case of one particularly perturbed cluster, E2 (see Vazza et al. 2011b). As shown in the bottom panel, the deviations are typically at the level of 10 per cent or less, even in the case of the very perturbed system E2, and do not show any particular trend with radius. A similar result is found for the integrated gas mass. This demonstrates the ability of our method to recover the true density profiles from X-ray observations in the presence of inhomogeneities.
Citation: Buote & Humphrey 2012
Citation function: Similarities
Figure 4(a) shows a time–latitude diagram of the subsurface toroidal flux density calculated by integrating the toroidal field over the range of 0.7–1.0R
⊙. The toroidal field is antisymmetric because the dipolar parity always prevails in our model. The latitudinal migration pattern of the toroidal field is different from that of FTD models, which usually show concentrated flux at high latitudes due to the radial shear in the polar portion of the tachocline. For each cycle, a weak toroidal field starts from high latitudes along with another overlapped branch of the previous cycle near the equator. With the equatorward migration of the toroidal field, its flux density increases as the cycle evolutes. This scenario is consistent with that of the solar cycle and the extended solar cycle (Howard & Labonte 1980; Wilson et al. 1988; McIntosh et al. 2014). Below 40° latitude, the toroidal flux density is larger than 2.5 × 1023 Mx Radian−1. Considering the spatial scale of a typical active region to be a few degrees, these fluxes are comparable with that in active regions. According to Cameron & Jiang (2019), the toroidal field at lower latitudes corresponds to the expected amount of flux residing in active regions. The weaker toroidal field at higher latitudes corresponds to ephemeral regions, which are small bipolar magnetic regions that contain a maximum total flux of the order of 1020 Mx (Harvey & Martin 1973). In Figure 4, the vertical dashed line represents cycle minimum (t
min) when the toroidal flux of the last cycle disappears, and vertical dotted line represents cycle maximum (t
max) when the toroidal flux of the new cycle reaches a maximum. The interval between t
min and t
max is 4 yr, which is close to the average period of the rising phase for solar cycles (Hathaway 2015; Jiang et al. 2018). The solid and dashed curves of Figure 5 show the time evolution of unsigned toroidal flux, 



∫5090°∫0.7R⊙R⊙r∣Bϕ∣drdθ

, of each hemisphere for odd and even cycles, respectively. The cycle period is 10.96 yr based on the interval between successive maxima of the toroidal flux. This period is consistent with the average period of observed solar activity (Hathaway 2015). If the cycle period is calculated as the interval between the start and end of the toroidal flux of a cycle, then it is 18.1 yr. This result is consistent with approximately the 18–22 yr period of the extended cycle (Wilson et al. 1988). The maximum toroidal flux produced in one hemisphere is about 2.6 × 1023 Mx, which matches the observed magnetic flux at the solar surface (Cameron & Schüssler 2015).
Citation: McIntosh et al. 2014
Citation function: Similarities
ZTF J092311.41+423634.16 (ZTF J0923+4236), the fifth-most variable white dwarf within 200 pc according to our Gaia-ZTF ranking parameter, demonstrates multiple transit-like phenomena of varying durations and depths. The highest-resolved of these features at MJD ≈ 58450 closely resembles a mirror image of the dips observed for ZTF J0139+5245 (Vanderbosch et al. 2020): this object displays a gradual ingress and sharp egress, also contrasting the transit shapes of WD 1145+017 (Figure 6) and the K dwarf KIC 12557548 (Rappaport et al. 2012; van Lieshout et al. 2016). Unlike those in ZTF J0139+5245, the observed transits in ZTF J0923+4236 appear more incongruous and complex, repeating on irregular intervals with varying shapes. The dip at MJD ≈  58450 appears to have a secondary feature before the onset of the larger dip. This more complicated structure, which appears to be transiting on the order of days, implies ZTF J0923+4236 may be at an evolutionary stage in between WD 1145+017 (roughly 4.5 hr orbital period) and ZTF J0139+5245 (roughly 107 day orbital period). We followed up ZTF J0923+4236 with high-speed photometry on MJD  =  59138 and 59139, using the McDonald 2.1 m telescope to assess short-term variability, observing small amplitude variations on timescales of ≈1 hr (Appendix C). These observations, however, were constrained to less than two hours in length, due to the observing conditions. McDonald 2.1 m observations on MJD  =  59168 spanning 3 hr in length do not show this same possible short-term variability to within the 4〈A〉 significance threshold (Appendix C). We observe broad Balmer lines in a relatively low-resolution follow-up spectrum from HET (see Figure 7), and classify ZTF J0923+4236 as a DA white dwarf. The lack of emission lines disfavors cataclysmic variable activity as the cause of photometric variability, leaving transiting planetary debris as a likely explanation. Future monitoring is suggested to place stronger constraints on the source of these transit-like features and any metal pollution or circumstellar gas around this white dwarf.
Citation: van Lieshout et al. 2016
Citation function: Compare/Contrast
Mapping out the distribution of dark matter in our Galaxy rests as a paramount task with potentially far reaching implications for astroparticle physics and cosmology. This is important to understand galaxy formation and to feed searches aimed at unveiling the very nature of dark matter. In particular, direct and indirect dark matter searches rely heavily on the findings of numerical simulations. It is therefore essential to extract the Galactic dark matter distribution directly from observations. In the outer Milky Way (at Galactocentric radii greater than ∼20 kpc), where baryons contribute little to the total mass budget, the gravitational potential traces closely the dark matter component and the total mass enclosed can be constrained using convenient tracers (e.g., Sakamoto et al. 2003; Dehnen et al. 2006; Xue et al. 2008; Bhattacharjee et al. 2014; Kafle et al. 2014), although with important degeneracies in the tracer population modeling. By contrast, in the inner Galaxy (i.e., in the inner ∼20 kpc) the baryonic contribution is very significant and its morphology rather uncertain, which makes the evidence for dark matter difficult to establish and the extraction of its distribution a delicate undertaking (Iocco et al. 2015). This has been addressed by many authors with different methods (e.g., Dehnen & Binney 1998; Sofue et al. 2009; Catena & Ullio 2010; Weber & de Boer 2010; Iocco et al. 2011; Bovy & Rix 2013; Nesti & Salucci 2013; Loebman et al. 2014), all of which do, however, make explicit assumptions about the underlying dark matter profile: typically, a multi-parameter profile is fitted to the observations together with a given baryonic component. The class of “local” methods to measure the dark matter density in the solar neighborhood (e.g., Salucci et al. 2010; Garbari et al. 2012; Smith et al. 2012; Zhang et al. 2013; Read 2014) avoids this bias, yet such methods are not easily applicable elsewhere in the Galaxy. An approach free of profile assumptions has been developed and successfully tested in external galaxies (Persic et al. 1996; Salucci et al. 2007), but never applied to our own Galaxy given the sizeable uncertainties of both kinematic data and baryonic modeling.
Citation: Kafle et al. 2014
Citation function: Background; Background
We use the state-of-the-art open-source stellar-evolution code Modules for Experiments in Stellar Astrophysics (mesa, Paxton et al. 2011, 2013, 2015, 2018, 2019) in order to investigate a broader range of rotational velocities and metallicities. Moreover, since very massive stars have been observed in the Tarantula Nebula (e.g. Doran et al. 2013) in the Large Magellanic Cloud (LMC), we consider stellar models with initial masses up to nearly 160 M⊙ (a factor of 1.3 larger than in Starburst99). Finally, we account for the fact that the majority of stars are born in binary systems (Sana et al. 2012, 2013). The observational evidence for such large binary fraction led to a paradigm change in our understanding of stellar evolution. As one (or both) of the stars in a binary system fill their Roche lobes, a phase of mass transfer takes place. Material is exchanged between the companions through the first Lagrangian point or lost by the system. The outermost layers of the donor are stripped off and eventually accreted on to the secondary star. This significantly alters the masses and spectroscopic appearances of the stars and generates evolutionary sequences otherwise unattainable in a single-star scenario (e.g. De Marco & Izzard 2017), influencing the mass-loss and the rotation rates of stars. In some cases, the interaction and mass exchange can be unstable, leading to merger events. Although numerous uncertainties still exist regarding the modelling of binary systems, it is becoming increasingly clear that more realistic estimates of stellar feedback cannot ignore the impact of stellar multiplicity. This applies also to radiative feedback as interacting binaries enhance the production of hydrogen- and helium-ionizing photons (e.g. Eldridge et al. 2017; Götberg, de Mink & Groh 2017) and harden the spectra of a stellar population (Götberg et al. 2019). Of particular interest are stripped helium stars, i.e. massive helium stars produced by binary interaction which emit the majority of their light at wavelengths shorter than the Lyman limit (Stanway, Eldridge & Becker 2016; Götberg et al. 2019) on a time-scale beyond few Myr.
Citation: Sana et al. 2012
Citation function: Uses
However, recent work by Brook & Di Cintio (2015a, 2015b) used results from simulations in which stellar feedback processes alter the DM content of dwarf galaxies to show that they can simultaneously match the densities and velocities of observed dwarfs. In this scenario, feedback from stars and supernovae creates bursty star formation histories in dwarf galaxies that fluctuate the gravitational potential well at the center of the dwarf (Pontzen & Governato 2012, 2014; Teyssier et al. 2013; Chan et al. 2015; Oñorbe et al. 2015; Dutton et al. 2016). DM core creation leads to a better match between theory and observed rotation curves (Katz et al. 2017; Santos-Santos et al. 2017). Feedback is particularly effective in dwarf galaxies with halo masses of a few 1010 





 (Governato et al. 2012; Di Cintio et al. 2014b), where it can transform an initially steep inner DM density profile into a flatter “cored” profile. At lower halo masses, there is less star formation, leading to less energy injection and lower core formation efficiencies (Peñarrubia et al. 2012; Maxwell et al. 2015). At higher masses, the deeper potential wells of galaxies make core formation increasingly difficult (Di Cintio et al. 2014b; Pontzen & Governato 2014), at least if an additional source of feedback is neglected, such as active galactic nuceli (AGN) (Martizzi et al. 2013). In this model with baryonic feedback, it is possible to assign dwarf galaxies to relatively massive halos, despite the low rotational velocities measured from their spatially resolved stellar kinematics. This is because baryonic feedback can push DM out of the central regions, lowering the enclosed mass at the radii that stellar kinematics probe (but without affecting the total halo mass8

8
Modulo a slight reduction in halo mass caused by the loss of baryons or preventive feedback (Munshi et al. 2013; Sawala et al. 2013).
). Hence, the densities are lowered, and the apparent velocities of the galaxies, reconciling the observations with theory.
Citation: Dutton et al. 2016
Citation function: Background
In summary, we have found quantitative evidence of thermal conduction suppression in a hot flare loop by coronal seismology of the slow-mode waves. This result suggests that the flare loop should cool much slower than expected from the classical Spitzer conductive cooling. Our studied event is indeed of such long-duration events (LDEs) which are the flares with a slower-than-expected decay rate in soft X-ray and EUV radiation (e.g., Forbes & Acton 1996; Takahashi & Watanabe 2000; Qiu et al. 2012). To explain the LDEs and flare loop-top sources, some previous studies have suggested the mechanism of continuous heating (Warren 2006; Liu et al. 2013; Sun et al. 2013) or conduction suppression (McTiernan et al. 1993; Jiang et al. 2006; Li et al. 2012). Our study confirms the effect of the latter mechanism in a more direct way. Laboratory experiments and numerical studies showed that the actual conductivity is smaller (by at least a factor of two) than that given by Spitzer when l ≲ 30λ where l is the temperature gradient scale length and λ the mean free path of thermal electrons (Bell et al. 1981; Luciani et al. 1983). We estimate l/λ ∼ 30 in this case, suggesting that the nonlocal conduction may account for the observed conduction suppression (Matte & Virmont 1982; Rosner et al. 1986). By studying the evolution of flare loop-top sources, Jiang et al. (2006) suggested that plasma waves or turbulence may play an important role in suppressing the conduction during the decay phase of flares. The mechanism is similar to that used for interpreting the significant reduction of thermal conductivity in galaxy-cluster cooling flows by a tangled magnetic field (Chandran & Cowley 1998). Finally, we conclude that the result of conduction suppression may also shed light on the coronal heating problem (see the review by Klimchuk 2015), because weak thermal conductivity implies smaller conductive losses, and an extended lifetime of individual nanoflares, increasing the average coronal temperature for the same heating rate.
Citation: Qiu et al. 2012
Citation function: Compare/Contrast
The linear relationship that we observe between breakpoint energy and core temperature is in line with previous measurements (e.g. McComas et al. 1992; Štverák et al. 2009), for both the halo and strahl. According to Scudder & Olbert (1979), a linear trend in the halo relation also follows under the assumption that binary Coulomb collisions dominate electron dynamics in the solar wind. However, in order to align with available experimental data, Scudder & Olbert (1979) set a scaling factor of Ebp/kBTc = 7, which differs from our scaling factor of Ebp/kBTc = 5.5 ± 0.1. With a scaling factor of Ebp/kBTc = 7, Scudder & Olbert (1979) predict that a transformation of thermal electrons into the suprathermal population occurs as the solar wind flows out from the Sun. Findings by Štverák et al. (2009), on the other hand, show that the (nh + ns)/nc ratio remains roughly constant with heliocentric distance in the slow wind, suggesting a lack of interchange between the thermal and suprathermal populations. However Štverák et al. (2009) observes some variability in the (nh + ns)/nc ratio in the fast wind, which they attribute to either statistical effects due to a lack of samples or a possible “interplay” between thermal and suprathermal electrons. Scudder & Olbert (1979) also predict that the halo Ebp/kBTc ratio remains constant with heliocentric distance, whereas Štverák et al. (2009) find that the halo Ebp/kBTc ratio decreases with heliocentric distance. These findings by Štverák et al. (2009), along with the discrepancy between our calculated ratio of Ebp/kBTc = 5.5 ± 0.1 and the prediction of Ebp/kBTc = 7, suggest that the model of Scudder & Olbert (1979) requires a minor update to either the theory or to the input parameters. The discrepancy, however, may also be indicative of other processes, such as wave-particle scattering (e.g. Gary et al. 1994), that possibly modifies the ratio between breakpoint energy and core temperature while preserving its linear relationship.
Citation: Štverák et al. (2009)
Citation function: Compare/Contrast; Compare/Contrast