adsabs/FOCAL huggingface dataset:
The physics of asymmetric inflow reconnection has been investigated in detail for fully ionized plasmas. One of the principal applications of this work has been Earth's dayside magnetopause. Asymmetric inflow reconnection has also been investigated in the context of Earth's magnetotail and elsewhere in the magnetosphere (Oieroset et al. 2004; Muzamil et al. 2014), the solar atmosphere (Murphy et al. 2012; Nakamura et al. 2012; Su & van Ballegooijen 2013; Su et al. 2013), laboratory experiments (Yamada 2007; Murphy & Sovinec 2008; Yoo et al. 2014), and plasma turbulence (Servidio et al. 2009, 2010). Cassak & Shay (2007) performed a scaling analysis for asymmetric inflow reconnection. They found that the outflow is governed by a hybrid upstream Alfvén speed that is a function of the magnetic field strength and density in both upstream regions (see also Birn et al. 2010) and that the flow stagnation point in the simulation frame and magnetic field null were not colocated. The structure and dynamics of asymmetric reconnection in fully ionized collisionless and two-fluid plasmas have been studied previously in several works (e.g., Swisdak et al. 2003; Cassak & Shay 2008, 2009; Mozer et al. 2008; Murphy & Sovinec 2008; Pritchett 2008; Mozer & Pritchett 2009; Pritchett & Mozer 2009; Malakit et al. 2010, 2013; Aunai et al. 2013a, 2013b). Simulations of the plasmoid instability during reconnection with asymmetric upstream magnetic fields by Murphy et al. (2013) showed that the resultant magnetic islands developed primarily into the weak-field upstream region. Because the reconnection jets impacted the islands obliquely rather than directly, the islands developed net vorticity. In addition to asymmetric inflow reconnection, several groups have investigated asymmetric outflow reconnection (e.g., Oka et al. 2008; Murphy 2010; Murphy et al. 2010) and reconnection with three-dimensional asymmetry (e.g., Al-Hachami & Pontin 2010; Wyper & Jain 2013).
Citation: Yoo et al. 2014
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The LMC is known to have brought a large population of star clusters as it has been accreted onto the Milky Way (Bica et al. 2008); thus, it is possible that DELVE 2 may share a similar history to the thousands of known star clusters in the Magellanic system. The LMC star cluster formation history is believed to be three-staged, including a period of rapid cluster star formation in the early universe (τ ≳ 10 Gyr), followed by a long quiescent period between ∼10 and ∼2–4 Gyr ago and then by a period of rapid star cluster formation extending to the present day, potentially due to the interaction between the LMC and SMC (Harris & Zaritsky 2009; Weisz et al. 2013; Rubele et al. 2018; Ruiz-Lara et al. 2020). One consequence of this period of quiescence in the LMC cluster formation history is the so-called “age gap” in the age distribution of LMC clusters, with a small (N ≲ 20) population of globular clusters with ages comparable to most known Milky Way globular clusters, separated by the gap from a much larger population of less massive young clusters (e.g., Bertelli et al. 1992; Girardi et al. 1995; Olszewski et al. 1996). These two populations of clusters obey an overarching age–metallicity relation, within which the older clusters (τ > 12 Gyr) are significantly more metal-poor (−2.2 ≲ [Fe/H] ≲ −1.2) compared to the younger population of clusters ([Fe/H] ≳ −0.7; Meschin et al. 2014).38

38
We note that Gatto et al. (2020) recently discovered 16 cluster candidates believed to be within the LMC cluster age gap (4 Gyr ≲ τ ≲ 10 Gyr).
 Therefore, although the photometrically derived metallicity and age for DELVE 2 are limited in accuracy by the small number of red giant branch stars available to precisely constrain these properties through synthetic isochrone fitting, it is clear that DELVE 2 is more consistent with an old, metal-poor stellar population and thus the former class of LMC clusters (provided the system is not a dwarf galaxy, as discussed in the previous subsection).
Citation: Rubele et al. 2018
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The experiment was carried out using the VULCAN laser facility at the Rutherford Appleton Laboratory (Danson et al. 1998), using the setup shown in Figure 1(a). A laser pulse of duration ∼1 ns and energy ∼ 70 J, was focused onto a 50 μm thick gold foil, at an incidence angle of 




∼
45
°


 to a peak intensity of ∼1015 W cm−2. The interaction of the nanosecond pulse with the gold foil leads to the generation of ablated plasma that is mainly constituted of thermally distributed gold ions and faster lighter ions with an average energy per nucleon of tens of keV (Tan et al. 1984; Gitomer et al. 1986). The lighter ions, such as protons and carbon ions, originate from contaminant layers (water vapor and hydrocarbon) typically present on the surface of the targets (Gitomer et al. 1986). The laser–target interaction was enclosed in a gas cell filled with pure nitrogen at a controlled pressure of ∼10−1 mbar. Hydrodynamic simulations using the code HYADES  (Larsen & Lane 1994) indicate that the gas becomes fully ionized within 100 ps from the start of the laser irradiation by the X-rays emitted from the target (Dean et al. 1971), resulting in a stationary plasma with an electron density and temperature of ∼3 × 1016 cm−3 and 






T


e


∼
1


  keV, respectively. These values imply an electron Debye length of 






λ


D


∼
1.4

μ
m


 and an ion-acoustic speed of 






C


s


∼
2.2
×


10


5



m



s


−
1




. Moreover, the Coulomb logarithm for electron–electron and ion–ion collisions are of the order of 6 and 11, respectively, indicating a characteristic timescale for collisions of 






τ


ee


∼
36

ns


 and 






τ


ii


∼
600

ns


, respectively. The PPI technique (Borghesi et al. 2002; Sarri et al. 2010) was employed to investigate the interaction of the ablated plasma with the background plasma. The probe proton beam was generated by focusing a second laser pulse of ∼1 ps duration and ∼50 J energy, to an intensity of 




∼


10


19




 W cm−2 onto a thin gold foil (thickness ∼20 μm). In the experimental arrangement shown in the Figure 1(a), the distance between proton source and interaction region was 




l
≃
4

mm


, and the detector was placed at 




L
≃
38

mm


 from the interaction region, giving an intrinsic geometrical magnification of 




M
≈
(
l
+
L
)

/

l
∼
10.5


 (Borghesi et al. 2002). A stack of several layers of dosimetrically calibrated (Kirby et al. 2011) RadioChromic Films (RCF), was used to detect the proton beam. The two lasers were temporally delayed so that a proton with an energy of 13 MeV traverses the interaction region at the time 






t


0


=
180
±
20


 ps after the start of the long-pulse irradiation, where the start of the interaction corresponds to 1/10 of the peak intensity. The error in defining the beginning of the interaction is mainly due to the systematic error in the synchronization of both laser beams. This does not affect the temporal resolution of the PPI technique, which is on the order of picoseconds (Sarri et al. 2010).
Citation: Larsen & Lane 1994
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The simplest two component model (power law + disk) is a phenomenological model often used to describe the spectra of ULXs as an empirical description of a disk plus corona geometry. In the presence of a cool (kT ~ 0.1−0.4 keV) and luminous (L ~ 1039−1040 erg/s) disk, it allows inferring the presence of intermediate-mass black holes (e.g., Makishima et al. 2000). This is not the case for M33 X-8, where the disk component describes the high-energy part of the spectrum well and appears to be hot (kT ~ 1.15 keV), leaving a soft excess that is accounted for by the power law. The overall disk parameters are then inconsistent with a massive black hole, but instead are more typical of an ordinary stellar mass black hole: using the relationship between mass, temperature, and luminosity in a standard disk (see, e.g., Makishima et al. 2000), we derive a mass of ~10 M⊙ for a nonrotating black hole, consistent with the estimation obtained by data from other satellites (e.g., Foschini et al. 2006; Weng et al. 2009; Isobe et al. 2012). Sutton et al. (2013) developed a classification scheme based on a disk+power law fit, to be applied to ULX spectra, according to which the spectral state of an ULX source can be defined by the disk temperature, the power-law slope, and the ratio between the flux contribution of the two spectral components in the 0.3−1 keV band. Our result is consistent with that found by Sutton et al. (2013) using XMM-Newton data, and, according to their classification, it identifies M33 X-8 as a broadened disk source, or in other words, as a source whose spectrum is dominated by emission from a hot disk (see Table 2) and where the additional soft component may be the effect of an unrealistic description of the disk spectrum by the diskbb model. In fact, such hot-disk/soft power-law spectra are difficult to explain in the context of the analogy of ULXs with GBHs: the thermal state of GBHs is indeed characterized by a hot disk, but the presence of a soft power-law-like component in addition to the disk is unusual, and its physical interpretation is not simple: if this component is due to the presence of a Comptonized corona, we do not expect it to be dominant at energies lower than the temperature of the seed photons that come from the disk.
Citation: Sutton et al. (2013)
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Citation: Sutton et al. (2013)
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Using all GRBs showing X-ray plateau phases, Dainotti et al. (2008) discovered a tight correlation between L0 and tb (the Dainotti relation). Subsequently, the Dainotti relation has been used to measure cosmological parameters (Cardone et al. 2009, 2010; Dainotti et al. 2013; Postnikov et al. 2014). Cardone et al. (2010) relied on this correlation to build a Hubble diagram of 66 GRBs and present a preliminary constraint on cosmological parameters. The errors on the cosmological constraints are large, which may be related to the nature of GRB light curves and in part due to the sample selection. In previous works, GRBs with X-ray plateaus were used to derive the Dainotti relation, and then to constrain cosmological parameters. However, a newborn magnetar can be spun down through a combination of electromagnetic dipole and gravitational-wave quadrupole emission (Shapiro & Teukolsky 1983). The X-ray luminosity of GRBs is given by the energy input from the electromagnetic and gravitational waves into the surrounding medium (Dai & Lu 1998; Zhang & Mészáros 2001; Metzger et al. 2011). Therefore, in order to standardize GRBs as standard candles through the Dainotti relation, it is recommended that X-ray plateaus caused by the same physical mechanism are used (electromagnetic dipole radiation or gravitational waves). Similar to supernova cosmology, only SNe Ia from accretion channels can be treated as standard candles. Following the Dainotti correlation (Dainotti et al. 2008) between the plateau luminosity and the end time of the plateau in X-ray afterglows, we further confirm this relation out to redshift z = 5.91. In this paper, we perform a first attempt to standardize long GRBs with X-ray plateaus dominated by electromagnetic dipole radiation as reliable standard candles. Using the same method, Hu et al. (2021) considered two other kinds of GRBs to constrain cosmological parameters, i.e., short GRBs with plateau phases dominated by magnetic dipole radiations and long GRBs with gravitational-wave-dominated plateau phases. It is interesting to note that this correlation also holds for the optical sample (Dainotti et al. 2020b).
Citation: Zhang & Mészáros 2001
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Figure 6 shows the orbital eccentricities as a function of [M/H] for the metal-rich disc sample. The solid lines correspond to the required eccentricity (see Eq. (2)) for different values of ISM radial metallicity gradients: −0.10 dex kpc−1 (black), −0.07 dex kpc−1 (our measured gradient for young stars in Table 1; see also Minchev et al. 2018, red), −0.04 dex kpc−1 (orange), and −0.06 dex kpc−1 (Cepheids analysis from Genovali et al. 2014, green). For the three first cases, we assumed ISM[M/H](R⊙) = 0.0 to estimate Rbirth from the stellar metallicity. However, Genovali et al. (2014) have their own zero point, defined as: [Fe/H] = −0.06 * Rg + 0.57, with a clear shift in the relation compared to the other ones assumed in this work. The impact of the ISM gradient value and the zero-point assumption on the derived Rbirth, and therefore on the required eccentricity to reach the solar vicinity without the need for churning, is clearly observed. As described in Hayden et al. (2020), given the measured [M/H] and eccentricity, stars lying to the left are able to reach the solar neighbourhood through blurring, while the stars to the right of the line are possible candidates to have migrated through churning. This is the case for most of the SMR stars (70% of the SMR stars lie below the line that corresponds to the Cepheids analysis); they are therefore likely to have been brought to the solar neighbourhood by churning, which is in close agreement with previous studies (e.g., Kordopatis et al. 2015a; Wojno et al. 2016). However, it is worth noting that the observed metallicity distribution function in Fig. 2 peaks around 0.2 dex, which is higher than previous reported solar vicinity MDFs (see e.g., Fuhrmann et al. 2017). A possible ignored bias towards more metal-rich objects in the sample selection could be pulling the percentage of possible migrators to higher values. Among the entire distribution, our churned candidates with [M/H] >  + 0.1 comprise around 17% of the sample. If we constrain the number of migrators to only stars with [M/H] >  + 0.25, the global percentage decreases to 8% of the sample.
Citation: Genovali et al. 2014
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Citation: Genovali et al. (2014)
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While the angular extent of IP shocks can be directly investigated using multi-point in situ measurements, the size of coronal shocks can only be indirectly inferred via remote-sensing observations of the electromagnetic emissions associated with them. According to Nelson & Robinson (1975), the average angle subtended at the solar surface by fundamental metric type II radio emission sources is 43Â°. Aurass et al. (1994) found particular cases with larger, double type II source structures covering a separation angle beyond 90Â°. Type II radio sources often show non-radial propagation trajectories (see Mann et al. 2003, and references therein). Wave-like large-scale disturbances propagating over the solar disk in extreme ultraviolet observations (usually referred to as âEIT wavesâ or âEUV wavesâ) are in close empirical correlation with type II radio bursts (Klassen et al. 2000). Most EIT waves are accompanied by CMEs, and observations and MHD modeling suggest that they are driven by the lateral expansion of CMEs, while the ultimate nature of the phenomenon remains under debate and could consist of true waves, pseudo waves (e.g., reconnection fronts), or a combination of both (Patsourakos & Vourlidas 2012; Nitta et al. 2013b, and references therein). According to Patsourakos & Vourlidas (2012), EIT waves can reach distances up to 1.3 R (850 Mm) from the source. Single-case studies reported some EIT waves covering a whole solar hemisphere (Klassen et al. 2000; Kienreich et al. 2009; Thompson & Myers 2009). Connections between EIT waves and SEP events have been often suggested (e.g., Bothmer et al. 1997; Krucker et al. 1999), and recently Rouillard et al. (2012) hypothesized that the EIT wave can be used to track the expansion of a coronal shock responsible for particle acceleration. Other authors question the EIT wave acceleration scenario for SEPs, with many EIT waves being observed at well-connected positions having no associated SEP increase (Miteva et al. 2014).
Citation: Klassen et al. 2000
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Citation: Klassen et al. 2000
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As previously mentioned, the small amplitude of the RV modulation in the NIR as measured by GIANO contemporaneously with REM photometry is intriguing. Assuming that the depth of the spectral lines relative to their adjacent continuum is constant and considering a spot with a contrast Cs at latitude ϕ with a filling factor fs, we expect an RV modulation approximately of Csfsvsinicosϕ (cf. Saar et al. 1997; Desort et al. 2007) that is ≈ 0.75−1.0 km s−1 for a spot at ϕ = 60°, which is remarkably higher than what has been observed. A cool spot at a higher latitude would reduce the amplitude of the NIR wide-band flux modulations; a quenching of the convective shifts or the Zeeman effect also do not appear to be viable explanations because they increase the effect of a cool spot on the RV at NIR wavelengths (Reiners et al. 2013). Nevertheless, the variation of the relative depths of the spectral lines in the NIR cannot be neglected and it is the dominant effect in the cool spot responsible for the large photometric modulation in the infrared. The relative line depths are strong functions of the continuum opacity and of the degree of element ionization both remarkably varying in a cool spot area with respect to the unperturbed photosphere. In general, these effects produce a remarkable increase of the relative depths of the spectral lines in the cool spot. This compensates for the decrease of the continuum intensity in the spot, reducing the distortions of the spectral line profiles and yielding an RV variation in the NIR smaller than expected from the wide-band photometric variation in the case of constant relative line depths. In any case, a quantitative analysis is not warranted by our data, since a larger number of observations would be required. This scenario suggests a need to better investigate this kind of target, since they might go through specific activity phases during which the VIS and NIR RV amplitudes are similar, possibly resulting in false positives. Looking at the curve phase shifts might give crucial information in these cases (see e.g. the recent result by Hatzes et al. 2018 for the K-giant γ Draconis).
Citation: Reiners et al. 2013
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Due to their high core temperatures at birth, neutron stars cool pre-dominantly by neutrino emission at ages ≲ 106 yr (Potekhin, Pons & Page 2015). The most rapid changes of surface temperature occur early, first when the temperature of the outer layers achieves equilibrium with the rapidly cooling core at an age of ≲ 100 yr (Lattimer et al. 1994; Gnedin, Yakovlev & Potekhin 2001; see also Nomoto & Tsuruta 1987) and then when the temperature drops below the critical temperature for core neutrons to become superfluid, which activates the efficient neutrino emission process of Cooper pair formation and breaking (Gusakov et al. 2004; Page et al. 2004). The rapid cooling of the CCO in Cassiopeia A (at a ten-year rate of ≈2.2 ± 0.2 or 2.8 ± 0.3 per cent, depending on whether NH varies between observations) indicates the latter starts to take place at an age of ∼200 yr (Page et al. 2011; Shternin et al. 2011). Neutron star cooling models predict that by an age of several hundred years, the cooling rate will be 1 per cent per decade. From the 1σ temperature uncertainties of our fit results with model parameters linked between observations (see Tables 4 and 7), we estimate upper limits on the ten-year cooling rates of 6 per cent for XMMU J172054.5−372652 and 17 per cent for CXOU J160103.1−513353. We measure a possible increase in temperature of ∼4 ± 2 per cent (accompanied by a decrease in emission area; see Table 5 and 6) for 1WGA J1713.4−3949. We also perform fits of the spectra of XMMU J172054.5−372652 and CXOU J160103.1−513353 which allow the temperature to be different between each observation, and the results are shown in Table 9 and Fig. 8 (analogous results for 1WGA J1713.4−3949 are shown in Fig. 3). We point out that we are concerned in Fig. 8 with relative changes in temperature and not in absolute temperature differences between CCOs since absolute temperatures depend on a variety of factors that are intrinsic to each CCO and may be different among CCOs in our sample, e.g. mass (and hence neutrino cooling rate) and radius (and hence gravitational redshift) and envelope composition and thickness. Nevertheless it is noteworthy that the temperatures of all three older CCOs appear to be higher than those of Cassiopeia A. Unlike for the 340 yr old CCO in Cassiopeia A, we do not see that temperatures of the 600 yr old XMMU J172054.5−372652 and 1000 yr old CXOU J160103.1−513353 are changing, at least within measurement uncertainties.
Citation: Page et al. 2004
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Several molecules show a large discrepancy between the abundances derived from observations and calculated by chemical equilibrium, although it is not as severe as for the molecules discussed above. We refer to PN in O-rich stars and H2S in C-rich stars, which are indicated by hatched rectangles in Fig. 2. For PN in O-rich AGB atmospheres, the disagreement between the observed abundances, (1–2) × 10−8 (Ziurys et al. 2018), and the calculated maximum chemical equilibrium abundance is almost three orders of magnitude. However, uncertainties on the observational and theoretical sides mean that the true level of disagreement is unclear. For example, while Ziurys et al. (2018) derived a PN abundance of 10−8 relative to H2 in IK Tau, De Beck et al. (2013) and Velilla Prieto et al. (2017) derived higher abundances, (3–7) × 10−7, in this source. When we give preference to these latter abundances, the level of disagreement would be even higher. On the other hand, the formation enthalpy of PN is rather uncertain (see Lodders 1999), which directly translates into the calculated chemical equilibrium abundance. In this study we adopted the thermochemical data for PN from Lodders (1999), who gives preference to a formation enthalpy at 298.15 K of 171.5 kJ mol−1, while other compilations such as JANAF use lower values that would result in higher chemical equilibrium abundances for PN. This would reduce the level of disagreement. In the case of H2S in C-rich AGB stars, the calculated maximum chemical equilibrium abundance is 7 × 10−11, while the value derived from observations is ~50 times higher. In this case, the observed abundance is based on the detection of only one line in only one source (see Agúndez et al. 2012), and thus it has to be viewed with some caution. In summary, the main failures of chemical equilibrium to account for the observed abundances of parent molecules in circumstellar envelopes are NH3, HCN, CS, SO2, and possibly PN in M-type stars, H2O and NH3 in S-type stars, and the hydrides H2O, NH3, SiH4, PH3, and perhaps H2S as well in C-type stars. The large discrepancies between the abundances derived from observations and those calculated with chemical equilibrium necessarily imply that nonequilibrium chemical processes must be at work in AGB atmospheres. Any invoked nonequilibrium scenario must account for all these anomalously overabundant molecules, but must also reproduce the remaining molecular abundances that are reasonably well explained by chemical equilibrium. No scenario currently provides a fully satisfactory agreement with observations, although two mechanisms that can drive the chemical composition out of equilibrium have been proposed.
Citation: Ziurys et al. (2018)
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Citation: Ziurys et al. 2018
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