adsabs/FOCAL huggingface dataset:
T Tauri stars in general are known to flare (e.g., Gahm 1990; Guenther & Ball 1999). A potentially better analog of the type of variable Hα emission expected from chromospheric emission and flaring on PTFO 8-8695 is the WTTS V410 Tau, with v sin i = 77.7 km s−1 (e.g., Carroll et al. 2012) compared to the measured v sin i = 80.6 ± 8.1 km s−1 for PTFO 8-8695 (van Eyken et al. 2012). The Hα emission equivalent width (3 Å with a typical value ∼1–2 Å ) on V410 Tau (e.g., Hatzes 1995; Fernández et al. 2004; Mekkaden et al. 2005) is weaker than seen on PTFO 8-8695, although V410 Tau has an earlier spectral type that raises the continuum level without necessarily affecting the strength of the chromospheric emission. V410 Tau has been observed to flare in a number of studies. Outside of flares, the Hα line of V410 Tau is fairly symmetric, relatively narrow, and is similar in shape to the chromospheric Hα profiles for PTFO 8-8695 seen in Figures 3, 5, and 6 (Hatzes 1995; Fernández et al. 2004; Skelly et al. 2010). The Hα line of V410 Tau can grow much stronger and broader during a flare, and also show asymmetries; however, the observed asymmetries seen during flares do not show excess emission with apparent peaks shifted out to greater than ±200 km s−1 (Hatzes 1995; Rice et al. 2011) as seen here in PTFO 8-8695. The typical pattern in a flare is for the line to very rapidly (timescale of a few minutes) strengthen and broaden with only a slight asymmetry developing. The strength and width of the line then decay exponentially with a timescale of ∼1 hr for strong flares (e.g., Fernández et al. 2004). This is not the temporal behavior observed in PTF08-8695. There is at least one additional piece of evidence against the flaring interpretation for the excess Hα emission seen in PTFO 8-8695. Whenever V410 Tau shows flare emission in Hα, significant He i 5876 Å emission also appears. This He i line is covered in the echelle formats of both our McDonald and Kitt Peak data. We have searched both data sets for evidence of this emission, including co-adding the spectra when the Hα emission appears stationary (UT 9:44 to 11:17 for McDonald; UT 6:45 to 8:25 for Kitt Peak) to increase the signal-to noise. No evidence of He i emission is seen. Lastly, if the observed excess Hα emission seen in Figures 2 and 7 were the result of a stellar flare, it would be a remarkable coincidence that the flare-induced asymmetry just happened to appear at and move with the same velocity position in the line profile as that expected for the planetary companion. In particular, the motion shown in Figures 2 and 4 where the excess emission first appears strongly on one side of the line profile and then moves to the other side has not to our knowledge been observed in the Hα emission of flare stars. Flares have been observed on PTFO 8-8695 (van Eyken et al. 2012; Ciardi et al. 2015), and while flares on this star likely will produce changes in the strength and shape of the Hα emission line, we believe all the points described above argue strongly against a purely stellar origin.
Citation: Rice et al. 2011
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The Hα imaging for the galaxies discussed in this paper was generated from two different sources: NGC 3628 was observed on May 8, 1991 using the ESO New Technology Telescope (NTT, red arm of the ESO Multi-Mode Instrument (EMMI)) for 900 s in Hα and 300 s in R as part of ESO program 047.01-003. The detector was a Ford 20482 single CCD (ESO CCD #24). Both images were taken with subarcsecond seeing. The data and all relevant calibration files taken during the run and a few days before and after the observing run were retrieved from the ESO archive and re-reduced by us using IRAF in the usual manner. L.A. Cosmic (van Dokkum 2001) was used to clear the image of cosmic rays to produce the final Hα and final R image. The continuum subtracted image was then produced by aligning Hα and R band, homogenizing the PSF, and subtracting the appropriately scaled R image from the Hα image (e.g., Skillman et al. 1997). To determine the scaling factor, we measured the apparent fluxes for several stars in the field on both the Hα and R images. Scaling, subtraction, and correction to reach an optimal correction of the continuum light of the galaxy was performed with our own IRAF scripts. The NGC 4522 Hα data were taken directly from the work of Koopmann et al. (2001) and the Galaxy On Line Database Milano Network (GOLDMine, Gavazzi et al. 2003). The images were astrometrically calibrated to the Digitized Sky Survey (DSS) system. Flux calibration of the continuum subtracted Hα images was performed by transferring the total Sloan Digital Sky Survey (SDSS) r′ band flux of each galaxy to the its continuum image. With the derived scaling factor (see above), the flux scale can then be directly transferred to the continuum corrected Hα image. A conversion from magnitudes to flux was performed using the fact that SDSS is calibrated in AB magnitudes, so that the zero-point flux density of each filter is 3631 Jy (with 1 Jy = 1 Jansky = 10-26 W Hz-1 m-2 = 10-23 erg s-1 Hz-1 cm-2)1.
Citation: van Dokkum 2001
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Understanding the influence of environment is contingent on being able to identify and quantify galaxy environments. Common environmental measures include the projected number density of galaxies out to the Nth nearest neighbour, the halo mass of a host group or cluster, or the projected separation from the centre of a group or cluster. Star formation and morphology of galaxies correlate well with these environment proxies, with galaxies in high densities regions (or alternatively, high halo mass or small group/cluster-centric radius) being preferentially red, passive, and early type (Dressler 1980; Goto et al. 2003; Poggianti et al. 2008; Kimm et al. 2009; Li, Yee & Ellingson 2009; Wetzel, Tinker & Conroy 2012; Wilman & Erwin 2012; Fasano et al. 2015; Haines et al. 2015). An alternative way to parametrize the environment of a host group or cluster is to classify the degree to which a system is dynamically relaxed. A relaxed, dynamically old group or cluster should be characterized by a central galaxy which is the brightest (most massive) member by a significant margin (e.g. Khosroshahi, Ponman & Jones 2007; Dariush et al. 2010; Smith et al. 2010) and is located near the minimum of the potential well (e.g. George et al. 2012; Zitrin et al. 2012, however also see Skibba et al. 2011), satellite galaxies which are distributed in velocity space according to a Gaussian profile (e.g. Yahil & Vidal 1977; Bird & Beers 1993; Hou et al. 2009; Martínez & Zandivarez 2012), and diffuse X-ray emission which is symmetric about the group/cluster centre (e.g. Rasia, Meneghetti & Ettori 2013; Weißmann et al. 2013; Parekh et al. 2015). The dynamical state of clusters is related to the age of the halo and the time since infall for member galaxies, which simulations have shown is an important quantity in determining the degree to which galaxy properties are affected by environment (e.g. Wetzel et al. 2013; Oman & Hudson 2016; Joshi, Wadsley & Parker 2017). Unrelaxed groups and clusters are systems which formed more recently or which have recently experienced a significant merger event, and in either case it would be expected that the time-since-infall on to the current halo for member galaxies will be relatively short. Therefore galaxies in unrelaxed groups may have properties which have been less influenced by environment compared to galaxies in more relaxed systems.
Citation: Smith et al. 2010
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Using the RADEX3 non local thermodynamic equilibrium (LTE) modeling program (van der Tak et al. 2007) with collisional rate coefficients from Wiesenfeld & Faure (2013), we modeled the relation between the gas kinetic temperature and the measured average of para-H2CO 0.5 × [(322–221 + 321–220)/303–202] ratios, adopting a 2.73 K background temperature, an average observational linewidth of 4.0 km s−1, and column densities N(para-H2CO) = 2.7 × 1012 and 3.7 × 1012 cm−2 for N113 and N159W, respectively. The results are shown in Fig. 5. The values of the para-H2CO column density were obtained with APEX data (beam size ~30″; Tang et al. 2017b), which cover similar regions. Different column densities of para-H2CO only weakly affect derived kinetic temperatures (see Fig. 3 in Tang et al. 2017b or Fig. 4 in Tang et al. 2018a; this was also shown in Fig. 13 and discussed in Sect. 4.3.1 of Mangum & Wootten 1993) as long as all lines are optically thin. Considering that the relation between the gas temperature and the para-H2CO line ratio may vary at different spatial densities (see Fig. 2 in Tang et al. 2017b), we modeled it at spatial densities 104, 105, and 106 cm−3 in Fig. 5. It appears that Tkin at n(H2) = 105 cm−3 is consistently lower than values at 104 and 106 cm−3 by ≲23% and ≲34%, respectively, for Tkin ≲ 100 K. Local thermodynamic equilibrium (LTE) is a good approximation for the H2CO level populations under optically thin and high-density conditions (Mangum & Wootten 1993; Tang et al. 2017a,b, 2018b). Following the method applied by Tang et al. (2017b) in their Eq. (2), we plot the relation between the LTE kinetic temperature, TLTE, and the para-H2CO (3–2) line ratio in Fig. 5. Apparently, TLTE agrees well with Tnon-LTE at volume densities n(H2) ~ 105 cm−3 as long as Tkin ≲ 100 K. Previous observations show that para-H2CO (3–2) is sensitive to gas temperature at density 105 cm−3 (Ginsburg et al. 2016; Immer et al. 2016; Tang et al. 2017b). The spatial density measured with para-H2CO (303–202) and C18O (2–1) in N113 and N159W is n(H2) ~ 105 cm−3 on a size of ~30″ (Tang et al. 2017b). Therefore, here we adopt 105 cm−3 as an averaged spatial gas density in the N113 and N159W regions.
Citation: Tang et al. 2017; b
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Citation: Tang et al. 2017b
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Citation: Tang et al. 2017b
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Citation: Tang et al. (2017b)
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Citation: Tang et al. 2017b
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Citation: Tang et al. 2017b
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Citation: Tang et al. 2017b
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Our case studies focus on several regions of the parameter space of explosive transients that may be related to specific source categories. From Sect. 3, different types of transient emissions from highly magnetized pulsars (also magnetars) can be affected by secondary acceleration. As mentioned previously, magnetars have been identified in many studies as promising candidates for the acceleration of cosmic rays and the production of secondary high-energy neutrinos, for instance Blasi et al. (2000), Fang et al. (2012, 2013), Lemoine et al. (2015), Kotera et al. (2015). The case of newborn magnetars with millisecond periods illustrates a non-relativistic source class, whereas the case of magnetar giant flares involves relativistic outflows. Moreover, tidal disruption events, low-luminosity gamma-ray bursts and blazar flares are examples of relativistic outflows, whose properties partially overlap in the parameter space of explosive transients. In these overlapping regions, they can be similarly affected by secondary acceleration. Therefore, we chose to describe the case of jetted tidal disruptions, while keeping in mind that this case study can be used as a benchmark example for low-luminosity gamma-ray bursts and blazar flares. We note that beyond standard scenarios involving gamma-ray bursts (e.g., Waxman & Bahcall 1997; Murase & Nagataki 2006; Murase et al. 2008; Mészáros 2015) and active galactic nuclei (e.g., Bednarek & Protheroe 1999; Atoyan & Dermer 2001; Halzen & Hooper 2005; Dermer et al. 2014; Petropoulou et al. 2016; Murase et al. 2018; Gao et al. 2019), jetted tidal disruptions have also been proposed as candidate sources for the production of high-energy cosmic rays and neutrinos (Wang et al. 2011; Senno et al. 2016; Dai & Fang 2017; Lunardini & Winter 2017; Wang & Liu 2016; Zhang et al. 2017; Biehl et al. 2018; Guépin et al. 2018). These case studies are associated with different types of photon fields that we simply model by hard or soft broken power laws and we can thus assess their impact on the high-energy neutrino spectrum. From Sect. 3, we can see that all these source categories are affected by strong secondary synchrotron losses and should be affected differently by secondary acceleration.
Citation: Kotera et al. (2015)
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We detect a deep (τsil  ∼  2.3) absorption feature due to silicate grains at around 10 μm. We compared the optical depth of the silicate feature in LEDA 1712304 with those in other AGNs in a wide range of IR luminosities (1010 L⊙   LIR   1013 L⊙). The spectra of the AGNs to be compared are taken from those of the IR galaxies observed by Spitzer/IRS (Stierwalt et al. 2013; Imanishi et al. 2007; Imanishi 2009; Roussel et al. 2006) with the threshold that the equivalent width of the PAH 6.2 μm feature is smaller than 0.27 μm (Stierwalt et al. 2013). In addition, we also take AGNs with low IR luminosities (LIR   1011 L⊙) from Wu et al. (2010). Figure 3 shows the relation between the IR luminosity and the optical depth of the silicate feature for LEDA 1712304 and the AGNs. We estimated the optical depths of the AGNs of Stierwalt et al. (2013), Wu et al. (2010), and Roussel et al. (2006) by ourselves in the same manner as was performed for LEDA 1712304 with the method defined by Imanishi et al. (2007), the spectra of which were obtained from the NASA/IPAC IR Science Archive (IRSA). On the other hand, we adopted the values given in each reference paper for the other AGNs. The blue dotted lines in Fig. 3 show the differences between our estimates and those by Stierwalt et al. (2013), which confirms that the differences between the different methods are not large enough to change a global relation. Figure 3 shows that galaxies with low IR luminosities (LIR   1011 L⊙) show significantly shallow silicate absorption features, as already pointed out by Stierwalt et al. (2013). Imanishi (2009) suggested that the number of heavily obscured AGNs that have deep silicate absorption features (τsil >  2) increases with the IR luminosities of the host galaxies. Therefore, LEDA 1712304 may be a rare galaxy from the aspect of having both deep absorption feature τsil ∼ 2.3 and low IR luminosity (5 ± 1)×1010 L⊙. Such galaxies have hardly been observed; an exception is NGC 1377 (Roussel et al. 2006), as can be seen in Fig. 3. NGC 1377 is a lenticular galaxy (de Vaucouleurs et al. 1991), the stellar mass of which is 109.3 ± 0.1 M⊙ (Skibba et al. 2011). The IR spectrum of NGC 1377 shows a featureless continuum except the silicate feature due to circumnuclear dust (Imanishi 2006; Roussel et al. 2006).
Citation: Imanishi et al. 2007
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Citation: Imanishi et al. (2007)
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The interaction of H2 and HD with atomic hydrogen is among the most widely investigated and important processes in elementary chemical reactions. The H + H2, H + D2, and H + HD reactions serve as benchmarks for experimental and theoretical investigations of bimolecular processes (Marinero, Rettner & Zare 1984; Zhang & Miller 1989; D’Mello, Manolopoulos & Wyatt 1991; Fernández-Alonso & Zare 2002; Harich et al. 2002; Aoiz, Bañares & Herrero 2005; Yang 2007; Gao et al. 2015; Karandashev et al. 2017; Yuan et al. 2018a,b, 2020; Goswami et al. 2020) and continue to attract much attention in the quest for unraveling quantum effects such as the geometric phase (GP) in chemical reactions (Kendrick, Hazra & Balakrishnan 2015; Hazra, Kendrick & Balakrishnan 2016; Croft et al. 2017; Kendrick 2018; Yuan et al. 2018a,b; Kendrick 2019). These elementary reactions are also of considerable interest in early universe chemistry models in determining H2 and HD column densities and the relative abundance of H/D in the interstellar medium (Flower 1999, 2000; Flower & Roueff 1999; Neufeld et al. 2006; Wrathmall, Gusdorf & Flower 2007; Gay et al. 2011; Nolte et al. 2011; Desrousseaux et al. 2018; Walker, Porter & Stancil 2018; Neufeld et al. 2019; Zhou et al. 2020). The H + HD ↔ D + H2 chemical reaction is especially important in this context as it cycles the heavier isotope between molecular form (HD) and purely atomic form. The HD molecule by virtue of its small dipole moment also serves as a tracer of H2 through its j = 1 → 0 rotational transition at 112 μm (Neufeld et al. 2006; Desrousseaux et al. 2018). Observations (Howat et al. 2002; Yuan et al. 2012) using the Infrared Spectrograph on the Spitzer Space Telescope have reported HD emissions from excited rovibrational level v = 1,  j = 5, as well as the pure rotational R(3) and R(4) lines, but relevant reaction rate coefficients are still limited in terms of initial HD rovibrational levels and gas temperature.
Citation: Kendrick, Hazra & Balakrishnan 2015
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Now with the recent availability of high quality full 6D phase-space information for large numbers of sources, much effort has been made to decrease the uncertainties in the Milky Way mass estimate. Recent works using a tracer mass estimator with 6D phase-space information include Sohn et al. (2018, globular clusters), Watkins et al. (2019, globular clusters), and Fritz et al. (2020, satellites). The most recent work using the spherical Jeans equation by Zhai et al. (2018) is very similar to our current investigation in method and data (LAMOST K giants) although only line-of-sight velocities were included, whereas we additionally make use of proper motions from Gaia to obtain the stellar tangential velocities. Using Bayesian analysis to fit a distribution function to full 6D phase-space data (globular clusters, satellites, and halo stars) has been a recent popular choice among many works (Eadie & Jurić 2019; Posti & Helmi 2019; Vasiliev 2019; Deason et al. 2021; Correa Magnus & Vasiliev 2022; Shen et al. 2022; Slizewski et al. 2022; Wang, Hammer & Yang 2022) and a similar distribution function analysis using 5D phase-space data from Gaia (Hattori, Valluri & Vasiliev 2021). In addition to fitting the observational data with a distribution function, several works have incorporated into the fitting a comparison of the observed data with Milky Way-type galaxies from cosmological simulations (Callingham et al. 2019; Li et al. 2020). Newly discovered high velocity stars with full 6D phase-space information have been used to estimate the mass of the Milky Way (Hattori et al. 2018; Monari et al. 2018; Deason et al. 2019; Grand et al. 2019; Koppelman & Helmi 2021; Necib & Lin 2022). Vasiliev, Belokurov & Erkal (2021) and Craig et al. (2021) have estimated the Milky Way mass by fitting models for the Sagittarius and Magellanic Steams, respectively. Several recent studies have estimated the Milky Way mass using measurements of the rotation curve (de Salas et al. 2019; Eilers et al. 2019; Ablimit et al. 2020; Cautun et al. 2020; Karukes et al. 2020; Jiao et al. 2021). Other works have used 6D satellite phenomenology, characterizing simulated Milky Way-type satellite populations and comparing to the observations of satellites in the Milky Way, to estimate the mass of the Milky Way (Patel et al. 2018; Villanueva-Domingo et al. 2021; Rodriguez Wimberly et al. 2022). Zaritsky et al. (2020) apply the timing argument to distant Milky Way halo stars to derive a lower limit to the Milky Way mass.
Citation: Hattori, Valluri & Vasiliev 2021
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Citation: Wang, Hammer & Yang 2022
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Solar filaments, or prominences as they are called when observed above the solar limb, can be observed in a stable state for many days or weeks. Sometimes, they suddenly start to ascend as a whole (full eruptions) (Joshi & Srivastava 2011; Holman & Foord 2015) or within limited sections of their length (partial eruptions) (Gibson & Fan 2006; Kliem et al. 2014). The ascending of a filament can go on high into the corona (successful eruptions) and gives rise to a coronal mass ejection (CME) or can stop at some greater height in the corona (confined or failed eruptions) (Ji et al. 2003; Török & Kliem 2005; Alexander, Liu & Gilbert 2006; Kuridze et al. 2013; Kushwaha et al. 2015). Occasionally two-step eruptions are observed. A filament after the first jump decelerates and stops at a greater height as in failed eruptions, but after a rather short period of time it starts to rise again and develops into the successful eruption with a CME formation. Byrne et al. (2014) observed on 2011 March 8 at the solar limb the erupting loop system that stayed in a metastable intermediate position for an hour and then proceeded and formed the core of a CME. Gosain et al. (2016) analysed observations of the eruption of a long quiescent filament on 2011 October 22 observed from three viewpoints by space observatories. A two-ribbon flare and the onset of a CME appeared 15 h after the filament disappearance on the disc. The filament was not observed at the high metastable position but some coronal structures that can be attributed to a corresponding flux rope were recognized. A clear example of the two-step filament eruption on 2015 March 14–15 was reported by Wang et al. (2016) and Chandra et al. (2017). In this event, a part of a large filament separated from the main body of the filament at the height of ≈30 Mm and rose upwards to the height of ≈80 Mm, where it stayed for 12 h clearly visible in chromospheric and coronal spectral lines. Finally, it erupted and produced a halo CME.
Citation: Ji et al. 2003
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In configuration 1S, the field close to the separatrix was sheared, producing steady interchange reconnection modulated by quasi-periodic reconnection bursts. We can roughly estimate the free energy release rate of the steady component from the energy injected before the onset of reconnection in the first 200 time units of the simulation. By t = 200, around 2 units of free magnetic energy are injected into the closed field; see Figure 16(a). Accounting for the ramp up of the driver and scaling the values, this corresponds to an energy injection rate of ≈5.6 × 1023 erg s−1 at the maximum driving speed. During the quasi-steady phase, this injection is balanced by losses to numerical diffusion and equates to roughly the free energy available for heating the plasma. Even after accounting for the unrealistically fast driving speed (see below), this energy release rate compares well with the observed values of 1023–1024 erg s−1 for bright points (Golub et al. 1974; Priest et al. 1994). The energy released by the bursts was a small fraction of the stored free magnetic energy—≈0.5 × Es = 4.9 × 1025 erg occurring with a period of ≈240 × ts = 8 minutes—while the outflow speeds reached typical values of ≈0.05 × Vs ≈ 60 km s−1 along the outer spine. The energy released in each burst corresponds to ≈18% of the energy released over the same period by the steady component. Many bright points exhibit quasi-periodic intensity increases, with periods ranging from a few minutes to a couple of hours (Kariyappa & Varghese 2008; Tian et al. 2008; Zhang et al. 2012). Our results demonstrate that some of this periodicity can be explained by the natural modulation of the interchange reconnection that occurs as minority-polarity elements are moved by surface motions. The predicted outflow speeds, and certainly the periods of the reconnection cycles, are likely too fast because the driving speed (12.5 km s−1) employed in our simulations is too high. However, configuration 1F demonstrated that the cycle period is mainly set by the displacement of the minority polarity. We speculate that, at more typical photospheric speeds (≈1.5 km s−1; e.g., Brandt et al. 1988), the reconnection cycle period would increase by a factor of 12.5/1.5 × 8 minutes ≈67 minutes, corresponding to the longest observed oscillations in brightness. Without a full treatment of the thermodynamics, however, it is not clear whether the repetitive, low-intensity reconnection jets in this case would be observable.
Citation: Golub et al. 1974
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