adsabs/FOCAL huggingface dataset:
The most prominent metallicity structure in the Milky Way disk is the decrease of metallicity with increasing Galactic radius—the radial metallicity gradient. All of the major tracers reveal radial gradients, typically with slopes between −0.03 and −0.09 





. The radial metallicity gradient can be explained by an inside-out galaxy formation where the disk grows via a radially dependent gas infall rate and star formation rate (e.g., Matteucci & Françoi 1989). But why is there such a wide range of radial gradient slopes measured? There are several possibilities. (1) Measurement uncertainty. Measurement errors will cause some variations, but it is unlikely to produce a factor of 3 difference in slope. Homogeneity in observing procedures and data analysis may be more important given the variations in abundance that can be derived for the same source (for further discussion see Rudolph et al. 2006; Henry et al. 2010; Balser et al. 2011). (2) Dynamical evolution. Radial gradients calculated with stellar tracers may be affected by radial migration, where stars are scattered into different orbits (Sellwood & Binney 2002). Radial migration should flatten the radial metallicity gradient, but there is some evidence that this may not be a large factor for stars in the Milky Way (Di Matteo et al. 2013; Kubryk et al. 2013; Bovy et al. 2014). (3) Temporal evolution. Many tracers have a wide range of age and therefore are probing the Milky Way disk at different times. For example, the radial gradient has been observed to flatten with time when using Open clusters (e.g., Friel et al. 2002). But accurate ages are crucial to separate out these temporal effects. PN studies indicate a flattening of the radial gradient with time (Maciel et al. 2003), a steepening with time (Stanghellini & Haywood 2010), or no temporal variation (Henry et al. 2010). These vastly different conclusions reveal the difficulty in deriving accurate PN ages and distances. (4) Sample Evolution. Hayden et al. (2013) measure a flattening of the radial gradient away from the Galactic mid-plane. Therefore, including sources out of the Galactic plane, which may be from an older population, may alter the derived radial gradient slope (see Minchev et al. 2014). Azimuthal abundance variations have been reported (e.g., Pedicelli et al. 2009; Balser et al. 2011, Section 3.3). If real, these will complicate any analysis of radial gradients. Here we detect metallicity gradient slopes over different azimuth ranges that span the values observed in the literature. The usual assumption is that the disk is well mixed at a given radius but this may not be true.
Citation: Balser et al. 2011
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There are two possibilities resulting in the different magnetic field strengths inferred from the polarimetric and molecular-line observations: (1) the rotational-to-gravitational energy βrot is overestimated, and (2) there are additional contributions in the polarized intensity from other mechanisms, such as dust scattering. In our MHD simulations, βrot is adopted to be 0.4% based on the observational estimates of the core mass of ∼1 M and the angular speed of the core rotation of 4 × 10−14 s−1. The angular speed was estimated based on the global velocity gradient along the major axis of the dense core observed with single-dish telescopes (Saito et al. 1999; Yen et al. 2011; Kurono et al. 2013). Numerical simulations of dense cores including synthetic observations show that the specific angular momentum derived from the synthetic images of the dense cores can be a factor of 8–10 higher than their actual specific angular momentum computed by the sum of the angular momenta contributed by the individual gas parcels in the dense cores (Dib et al. 2010). In addition, if there are filamentary structures in the dense core in B335, which could not be resolved with the single-dish observations, infalling motions along the filamentary structures could also contribute to the observed velocity gradient, leading to an overestimated angular speed of the core rotation (Tobin et al. 2012). We have also performed our simulations with a lower βrot, and we find that the rotational velocity on a 100 au scale in the simulations decreases with decreasing βrot. Thus, the discrepancy in the magnetic field strengths inferred from the field structures and the gas kinematics can be reconciled, if the core rotation in B335 is overestimated by a factor of a few in the observations, and these results would suggest a weak magnetic field of initial λ of 9.6 in B335. Further observations combining single dishes and interferometers to have a high spatial dynamical range and to map the velocity structures of the entire dense core in B335 at a high angular resolution are needed to study coherent velocity features and provide a better estimate of the core rotation.
Citation: Yen et al. 2011
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Citation: Saito et al. 1999
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The LCs analyses were made using the phoebe v.0.29d software (Prša & Zwitter 2005) that is based on the 2003 version of the Wilson-Devinney (W-D) code (Wilson & Devinney 1971; Wilson 1979, 1990). In the absence of spectroscopic mass ratios, the “q-search” method (for details see Liakos & Niarchos 2012) was applied in modes 2 (detached system), 4 (semi-detached system with the primary component filling its Roche lobe) and 5 (conventional semi-detached binary) to find feasible (“photometric”) estimates of the mass ratio. The step of q change during the search was 0.05–0.1 starting from q = 0.05−0.1. The effective temperatures of the primaries (T1) were given the values derived from the spectral classification (see Sect. 2) and were kept fixed during the analysis, while the temperatures of the secondaries T2 were adjusted. The Albedos, A1 and A2, and gravity darkening coefficients, g1 and g2, were set to generally adopted values for the given spectral types of the components (Rucinski 1969; von Zeipel 1924; Lucy 1967). The (linear) limb darkening coefficients, x1 and x2, were taken from the tables of van Hamme (1993). The dimensionless potentials Ω1 and Ω2, the fractional luminosity of the primary component L1, and the orbital inclination of the system i were set as adjustable parameters. At this point it should to be noted that since the Kepler’s photometer has a spectral response range between approximately 410–910 nm with a peak at ~ 588 nm, the R filter (Bessell photometric system-range between 550–870 nm and with a transmittance peak at 597 nm) was selected as the best representative for the filter depended parameters (i.e. x and L). Moreover, there is evidence of maxima brightness changes in all of the systems, therefore parameters of photospheric spots on the surface of the secondary were also adjusted. The selection of the magnetically active component was based on the effective temperatures of the members of the systems. In all cases the secondaries are clearly cooler than the primaries (i.e. large minima difference), therefore, they host a convective envelope that better suits a magnetically active star. In addition, the hotter stars are candidates for δ Sct type pulsations and it is rather rare to present magnetic activity also. For all EBs, except for KIC 111, the third light parameter (l3) was also adjusted because the systems are candidates for triplicity (see Sect. 1). However, during the iterations it resulted in unrealistic values, therefore, it was omitted in the final analysis. Finally, all systems were found to have the minimum ∑ res2 in mode 5. KIC 066 and KIC 111 have a minimum at q = 0.3, while KIC 105 and KIC 106 at q = 0.15. In Fig. 7 the respective q-search plots are shown.
Citation: von Zeipel 1924
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Here, for the first time, we have identified highly dynamic non-potential activity on QS-like supergranular network scales. These events overlie mixed polarity network flux elements near the spatial resolution of HMI, and are the first non-potential structures to be unassociated with strong concentrations of bipolar magnetic flux. One event (2011b August 05) shows eruptive activity in the form of jets, which is similar to larger-scale micro-sigmoids (Raouafi et al. 2010) and even QS bright point sigmoids (Chesny et al. 2015). The existence of flaring non-potential fields in the QS-like mixed network immediately shows that supergranular-scale magnetic fields can support processes similar to sigmoid formation (Chesny et al. 2013). Strong non-potential field arcades have been observed in hot X-ray sigmoids on scales of hundreds of arseconds (Moore et al. 2001; Gibson et al. 2002; Liu et al. 2010), micro-sigmoids in soft X-ray imaging on scales of ∼50″ (Mandrini et al. 2005; Raouafi et al. 2010) and small-scale AR EUV (Zheng et al. 2012, 2013), and now at EUV temperatures in QS-like mixed network fields with lengths down to ∼10″. This range of lengths over a range of temperatures and magnetic field topologies points directly to self-similar mechanisms influencing plasma and magnetic field dynamics at a range of scales. Our findings suggest that the processes driving some large-scale eruptions (i.e., flux emergence, helicity build-up, and flux cancellation leading to non-potential field heating (Chen et al. 2014)), can also manifest in a range of configurations on sub-network size scales in QS-like magnetic field configurations. These QS flaring non-potential fields are similar to their large-scale counterparts, but not as strict in their evolution. The diversity in the observed topologies may scale with the diversity of magnetic configurations that exist in the supergranulation network. QS non-potential fields evolve in multi-polar environments, and are not restricted to strongly bipolar dominated regions as in larger-scale, higher temperature events. Despite this, one of the presented events (2011a August 05) results in a post-flare potential loop arcade, which is similar to some observed AR sigmoid fields (Moore et al. 2001).
Citation: Liu et al. 2010
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Citation: Mandrini et al. 2005
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Subpulse drifting has been a topic of intensive research with the phenomenon expected to be present in 40–50 per cent of the pulsar population. There are around 120 pulsars known at present to exhibit some form of periodic modulations in their single-pulse sequence (Weltevrede et al. 2006, 2007; Basu et al. 2016). The drifting effects are very diverse and associations are seen with other phenomena like mode changing and nulling in some cases (Wright & Fowler 1981; Deich et al. 1986; Hankins & Wolszczan 1987; Vivekanand & Joshi 1997; Redman et al. 2005; Basu & Mitra 2018b). However, to gain a deeper understanding of the phenomenon at large, more comprehensive studies involving classification of the drifting population and identification of underlying traits are essential. The first such attempt was undertaken by Rankin (1986) who estimated the drifting behaviour within the empirical core–cone model of the emission beam. The pulsar profile, which is formed after averaging several thousand single pulses, has a highly stable structure and is made up of one or more (typically 5) components. The components can be classified into two distinct categories: the central core component and the adjacent conal components. A number of detailed studies suggest that the radio emission beam consists of a central core emission surrounded by conal emission arranged in nested rings (Rankin 1990, 1993; Mitra & Deshpande 1999). The differences in the observed profiles are a result of the different line of sight (LOS) traverses across the emission beam. Rankin (1986) suggested subpulse drifting to be primarily a conal phenomenon, and the drifting to arise due to the circulation of the conal emission around the magnetic axis. This would result in an association between the drifting properties (particularly the phase variation of the drifting feature across the pulse) and the profile classification. The most prominent drifting with clear drift bands and large phase variations was associated with the conal single (Sd) and barely resolved conal double (D) profile classes. This corresponded to the LOS traversing the emission beam towards the outer edge. Progressively more interior LOS traverses of the emission beam result in well-resolved conal double (D), conal Triple (cT), and conal Quadruple (cQ) profile shapes. In the above classification scheme, such profile classes were expected to show primarily longitude stationary drift with very little phase variations. The core-dominated profiles were associated with central LOS traverses of the emission beam. The principal profile classes were categorized as core single (St), Triple (T) with a central core component and pair of conal outriders, and Multiple (M) with central core and two pairs of conal components. Subpulse drifting was expected to be phase stationary and only seen in the conal components of the T and M class profiles. However, the core components sometimes show longer periodic structures that were not classified as subpulse drifting.
Citation: Rankin 1990
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Surveys of radio bursts in decimetric wavelengths is presented in papers by Isliker & Benz (1994) and Jirička et al. (2001), within 1–3 GHz and 0.8–2.0 GHz frequency ranges, respectively. Some of these bursts are still not well understood. This is a case of the slowly positively drifting bursts (SPDBs). They appear in groups or as single bursts, with a duration of an individual burst from 1 to several seconds and their frequency drift is lower than about 100 MHz s−1 (Jirička et al. 2001). The SPDBs seem to be similar to the reverse type III bursts (Aschwanden 2002) but their frequency drift is much smaller. The majority of observed SPDBs are connected to solar flares (Jirička et al. 2001), and they appear many times at the very beginning of the flares (Benz & Simnett 1986; Kotrč et al. 1999; Kaltman et al. 2000; Karlický et al. 2018). Kaltman et al. (2000) reported on several SPDBs observed during three solar flares in the 0.8–2 GHz frequency range. They found frequency drifts of the observed SPDBs to be within the 20–180 MHz s−1 range. Kotrč et al. (1999) studied one of those flares. By combining the radio and spectral plus imaging Hα observations, they explained the observed SPDBs as radio emission generated by downwards propagating shock waves. Based on numerical simulations of the formation of thermal fronts in solar flares, Karlický (2015) proposed that SPDBs observed in the 1–2 GHz range could be a signature of a thermal front. Furthermore, Karlický et al. (2018) reported the observation of an SPDB (1.3–2.0 GHz) observed during the impulsive phase of an eruptive flare. They found time coincidence between the SPDB occurrence, an appearance of an ultraviolet (UV)/EUV multithermal plasma blob moving down along the dark Hα loop at approximately 280 km s−1, and the observed change of Hα profile at the footpoint of that dark loop. Combining these observations they concluded that observed SPDB was likely generated by the thermal front formed in front of the falling EUV blob.
Citation: Jirička et al. 2001
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Citation: Jirička et al. 2001
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Citation: Jirička et al. (2001)
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Forbush decreases (FD) are the results of influence of solar wind (SW) large-scale disturbances on the background cosmic ray (CR) flux. They often demonstrate relatively fast CR intensity decrease, accompanied by large and extremely variable CR anisotropy, which is followed by a slower recovery (Forbush 1937; Lockwood 1971; Iucci et al. 1979a; Belov 2009). Most FDs have a sporadic character and are induced by interplanetary disturbances such as Interplanetary Coronal Mass Ejections (ICMEs) caused by coronal mass ejections (CMEs) (Cane 2000; Gopalswamy 2010a; Richardson & Cane 2011a). FDs caused by high-speed streams (HSS) from coronal holes (CHs) have a recurrent character (Iucci et al. 1979b; Richardson 2004; Singh & Badruddin 2007b). The influence of different types of SW disturbances on galactic CR has been well-documented in scientific literature (e.g. Dumbovic et al. 2012; Chertok et al. 2013; Kumar & Baddruddin 2014a,b; Melkumyan et al., 2018, 2019). It is known that most of CMEs, and almost all energetic CMEs, are originated from the active regions being accompanied by solar flares which intensity correlated with the CME kinetic energy (e.g. Yashiro & Gopalswamy 2009). Comparison between the impact on the heliosphere of active-region (AR) and non-AR CMEs is presented in Gopalswamy et al. (2010b). The authors reveal that: (i) active regions produce almost all CMEs with above-average energy, for example geoeffective CMEs, CMEs associated with solar energetic particles, shock-driving CMEs; (ii) the quiescent filament regions produce the other type of CMEs which are not related to the sunspots and occur mostly at high latitudes during the maximum phase of solar activity. One of the factors responsible for the variability of FDs is different origin of related interplanetary transients. In Marii et al. (2020, 2021) the influences of different parts of ICMEs (turbulent sheath and magnetic obstacle) on the associated FDs are examined according to the CME origin. The events are separated into three subsets: AR CMEs, disappearing filament CMEs, and stealthy CMEs. CR variations caused by turbulent sheath and magnetic obstacle parts of the corresponding ICME are compared for the three subsets and two different linear relations are found: for AR ICMEs on one side and for filament/stealthy ICMEs on the other side.
Citation: Iucci et al. 1979b
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Given the potential huge information content of these observables advocated by some recent works (Patton et al. 2017), the predictions from first principles developed in this paper could be successfully applied to forthcoming data along with the standard power spectrum based analysis and could bring additional information beyond ΛCDM parameters like massive neutrinos (Liu & Madhavacheril 2019) or dark energy (Codis et al. 2016a). Let us stress that implementing nulling in weak-lensing analysis is central in order to avoid extracting biased information from the small scales that lack a full theoretical understanding (including due to the effect of baryon physics that needs to be modelled in weak-lensing surveys Hildebrandt et al. 2018; Yoon et al. 2019). Not only this general nulling technique should be used for one-point statistics but could also be applied to standard power spectrum analysis (and more generally to the full two-point PDF) in order to disentangle the effects of the different physical scales. However, more realistic effects have to be accounted for before the here mentioned formalism could be directly applied to real data. In particular, we have not investigated the precise impact of the galaxy redshift distribution ns(z) (for which one needs to go from a set of discrete source planes to a source distribution), photometric redshift errors or shape noise, which are left for future works. Promising extensions include an application of the formalism to (i) compensated filters such as for aperture mass which require the joint modelling of the field at two different scales, (ii) two-point statistics in order to model cosmic variance, and (iii) the joint analysis of multiple redshift bins. All of these ideas are within reach of LDT as was shown in the case of the three-dimensional matter density in Bernardeau, Codis & Pichon (2015) and Codis et al. (2016b) for, respectively, the multiscale and two-point statistics.
Citation: Hildebrandt et al. 2018
Citation function: Future Work
We justify our use of the Two-Zone approximation as follows: we anticipate that the magnetic structure within the outflow would be perpendicular to the plane of the host galaxy and thus perpendicular to the magnetic field orientation within the Zone B ISM region. Indeed, such perpendicular magnetic structure in outflows is seen in simulation work where, e.g. the action of a CR-driven dynamo yields a perpendicular magnetic field configuration compared to the host galactic plane (Kulpa-Dybeł et al. 2011), or by the advection of the magnetic fields by the flows themselves (Bertone et al. 2005), by magnetic amplification via the CR streaming instability (Uhlig et al. 2012) along the outflow. This magnetic structure would also be consistent with polarised radio synchrotron emission above and below the planes of galaxies known to host outflows in the nearby Universe, with the polarisation direction aligned with the orientation of the outflow cone (see, e.g. Hummel et al. 1988; Sukumar & Allen 1990, 1991; Hummel et al. 1991; Brandenburg et al. 1993; Chyży et al. 2006; Soida et al. 2011; Mora & Krause 2013). We argue that the principle mechanism for CRs to permeate the Zone A/Zone B interface would be via diffusion. With magnetic field lines aligned in a direction parallel to the inter-zone boundary, diffusion across the interface would be severely hampered – the cross-boundary diffusion coefficient would effectively be perpendicular the the local magnetic field lines, and so would be around two orders of magnitude smaller than that along the field directions (e.g. Shalchi et al. 2004, 2006; Hussein & Shalchi 2014), and substantially less than the effective ISM diffusion coefficient. The detailed substructure of the magnetic fields in these interfacing regions is not yet fully understood (Veilleux et al. 2005), but we argue that our prescription is consistent with existing work on relevant scales and that adopting an alternative model for CR transport across this boundary at this point would not imply an interpretation that is any more physical than that adopted here. We acknowledge that, in future studies, it will be critical to assess the magnetic fields in these interfacing regions across a range of length-scales to properly determine the permeability of the Zone A/Zone B interface to diffusing CRs.
Citation: Hummel et al. 1991
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