adsabs/WIESP2022-NER huggingface dataset:
Paper: 2016ApJ...827L..18O (fulltext)
agrees well with the assumption. In addition to the reactions presented above, the following unimolecular hydrogen molecule elimination (UHME) process may occur during the photolysis of methanol at low temperatures (Gerakines et al. 1996 ; Hama et al. 2009 ): The reaction products of UHME process may be incorporated into AAs. If these reactions occur in a statistical manner, the frequencies of reactions ( 4 b ) and ( 5 a ) are twice as large as those of their counterparts (reactions ( 4 a ) and ( 5 b ), respectively) because the number of C–H bonds in CH 2 DOH is larger than that of the C–D bond, which may affect the D/H ratio of the reaction products. The production of C–H bonds is also possible from other pathways that do not use methanol as a reactant. Briggs et al. ( 1992 ) reported the formation of various COMs such as glycolic acid (HOCH 2 COOH), glycerol (HOCH 2 CH(OH)CH 2 OH), and Gly after the photolysis of solid mixtures (CO:H 2 O:NH 3 = 5:5:1) at 10 K. Although their gas composition is different from that in our study, similar products can be produced in the present experiment, which has the potential to decrease the deuteration level of Gly. Although the formation of other AAs was not reported by Briggs et al., others reported the formation of various AAs upon the hydrolysis of the organic residues even in the absence of methanol (Muñoz Caro et al. 2002 ), which should have the potential to drastically decrease their deuteration levels. 4.2. Effect of Hydrolysis on Deuteration Levels of Amino Acids We demonstrated that the acid hydrolysis of the organic residues generally results in significant increases in AA abundances, which is consistent with that demonstrated in previous studies (Nuevo et al. 2008 ; Evans et al. 2012 ). On the other hand, there has been no report on the variation in the deuteration levels of AAs after the acid hydrolysis of the organic residues. The present result suggests that the five AAs can be divided into two groups depending on whether the deuteration level decreases upon hydrolysis: those whose D/H is significantly decreased ( α -Ala and β -Ala; referred to as the Ala group) and those whose D/H did not change (Gly, Sar, and Ser; referred to as the Gly group). This indicates that the two groups have different formation mechanisms upon hydrolysis. We propose one possible scenario to explain the observed D/H ratio: we assume that these five AAs are produced by reactions of small radicals during warming of the irradiated ice and that most AAs formed are incorporated into macromolecules via relatively weak bonds such as peptide bonds. Upon hydrolysis, these weak bonds are broken and the unit AAs are released. If this is the only pathway for the formation of AAs, the deuteration levels of the AAs would not vary between the free and bound samples, as is the case for Gly-group AAs. On the other hand, for Ala-group AAs, if there are other formation pathways that do not use deuterated radicals and/or molecules
Paper: 2019MNRAS.483..873O (acknowledgments)
ACKNOWLEDGEMENTS The authors wish to acknowledge the DJEI/DES/SFI/HEA Irish Centre for High-End Computing (ICHEC) for the provision of computational facilities and support. This work used the BATS-R-US tools developed at the University of Michigan Center for Space Environment Modeling and made available through the NASA Community Coordinated Modeling Center. DÓF wishes to acknowledge funding received from the Trinity College Postgraduate Award through the School of Physics. AAV acknowledges funding received from the Irish Research Council Laureate Awards 2017/2018. SJ acknowledges the support of the German Science Foundation (DFG) Research Unit FOR2544 ‘Blue Planets around Red Stars’, project JE 701/3-1 and DFG priority program SPP 1992 ‘Exploring the Diversity of Extrasolar Planets’ (RE 1664/18). We thank Jackie Villadsen and Joe Llama for their useful discussion on topics of stellar radio emission. The authors would like to thank our referee, Dr Jorge Zuluaga, for his valuable input on this work. Software: bats-r-us (Powell et al. 1999 ), radiowinds (Ó Fionnagáin 2018 ), scipy (Jones et al. 2001 ), matplotlib (Hunter 2007 ), and tecplot and seaborn (Waskom et al. 2018 ).
Paper: 2016A&A...585A...6G (fulltext)
of the next shock after one pulsation period. Any dust formed at 1 may thus not accumulate and grow over several pulsation periods in the shocked photosphere. The situation is different at larger radii, where the post-shock gas temperature drastically decreases. 3.2.1. Alumina The alumina clusters chiefly form in the post-shock gas at 1 , and their abundances in the first oscillation are illustrated in Fig. 6 . The formation of AlO involves the dimerisation of AlO through termolecular recombination while AlO forms from the reaction of Al with OH in the first place. The dimerisation of AlO is also described by termolecular recombination to form the AlO dimer, (AlO), which has a binding energy of eV ( Li Cheng 2012 ), and can thus withstand the high temperatures of the post-shock gas. This dimer is considered as the seed that will further condense into alumina grains. Over one oscillation resulting from the first shock at 1 , only a small quantity of alumina dimers forms, reaching an abundance of at . A large fraction of Al stays in the form of molecular rings (e.g., AlO and AlO). Because the dimerisation of AlO and AlO is termolecular and thus dependent on the background pressure, cluster formation occurs over a reduced range of radii in the inner wind and the formation of AlO cluster ceases after 2 . The grain size distributions for alumina in one oscillation at 1 is presented in Fig. 7 as a function of pulsation phase. Because of the high post-shock gas temperatures after the shock passage and the extremely low abundance of alumina dimers, no grains form at . Very small grains start forming after , but coagulation starts to be efficient for dust growth in the phase range , for which gas temperatures span the range K and the alumina cluster abundance reaches a maximum. The distribution peaks at grains in the size range Å, despite the small fraction of AlO clusters available. However, the post-shock gas densities are very high at 1 , as seen from Table 2 , and the clusters coagulate and condense with high efficiency and lead to the formation of AlO grains of small size over one oscillation and pulsation period. Under the assumption that small alumina grains may not survive the passage of periodic shocks, the results show that there exists a steady population of small size (m) alumina grains located in the shocked photosphere at 1 . If we consider that grain growth may also proceed through surface deposition when the grain radius has reached several tens of Å (at and after), the grain size will increase, but the growth will be time-limited and proceed only between . We now consider the extreme case where alumina grains survive the periodic shocks at each shock passage, while molecules are destroyed and reform. We may then investigate grain growth over several pulsation periods. For a gas drift velocity of 0.5 km s, the gas layer at 1 takes 12 pulsations to reach 2 . We thus consider the build-up of dust grains over a number of 12 pulsations for
Paper: 2016MNRAS.458.2307K (fulltext)
detail. As mentioned before, we identified two mild-Am stars, HD 33204 and HD 46304. We show abundance distributions of these stars in Fig. 10 . A typical Am star exhibits overabundances of iron-peak elements and some heavy elements (Zn, Sr, Zr and Ba), but Ca and Sc abundances of these stars are underabundant (Gray Corbally 2009 ). As can be seen in Fig. 10 , the mild-Am stars in our study have nearly solar abundances of Ca and Sc elements. Only HD 33204 shows overabundances in some heavy elements typical for Am star. In the case of HD 46304, most of lines are blended due to high rotation velocity of the star. These blended lines cause difficulties in abundance calculations. The abundance differences between HD 46304 and a typical Am star can be caused by this effect. In the spectral classification process some stars were defined as metal-poor, mostly taking into account Mg, Mn and Fe lines (see Table 2 ). For these stars, the average abundance of Mg (7.57 dex) was found to be close to the solar abundance. However, the average abundances of Mn (5.13 dex) and Fe (7.27 dex) are slightly lower than the solar abundances. For some of our targets (HD 26298, HD 33204, HD 106103, HD 110379, and HD 126516) chemical abundances were already obtained before the present study. We compared our atmospheric parameters and abundance results with the literature values. Abundances of HD 33204 and HD 10613 were derived by Varenne Monier ( 1999 ) and Fossati et al. ( 2008 ), respectively. In their work, similar methods for abundance analysis were used but different atomic data bases were adopted. For HD 33204 higher abundances of Sc, Mg and Y were obtained in the present study. Our result is consistent with the Am: type peculiarity of this star. For HD 106103 only the Y abundance is different. Abundances of HD 26298, HD 110379 and HD 126516 were derived by Bruntt et al. ( 2008 ) using the equivalent width method and the different atomic data base. In the case of these three stars, the marked differences were detected only for V and Ba abundances. Figure 10. Chemical abundances of Am: stars compared with the solar values (Asplund et al. 2009 ). HD 33204 is represented by stars and HD 46304 by triangles. Additionally, we checked the possible correlations of the element abundances with atmospheric parameters, i.e. T eff , log g and ξ. Similar as in Niemczura et al. ( 2015 ), no correlations were found. Moreover, the element abundances do not depend on v sin i . The same result was found by Fossati et al. ( 2008 ). However, Takeda et al. ( 2008 ) reported negative correlations between v sin i and C, O and Ca elements. Finally, we checked the relations between abundances of iron and other elements. Strong positive correlations were found for Mg, Si, Ca, Sc, Ti Cr, Ni, Y and Ba. Similar correlations between Fe abundances and iron-peak elements were found by Niemczura et al. ( 2015 ). On the other hand, a negative correlation was obtained between Fe and O abundances. 6
Paper: 2016MNRAS.463L..17A (acknowledgments)
NA and SV thank Denis Erkal for stimulating discussions. We thank the referee, Ray Carlberg, for a constructive report.
Paper: 2020MNRAS.491.4783G (acknowledgments)
ACKNOWLEDGEMENTS This work was supported by the Science Technology Facilities Council (STFC) (Research Grants ST/L000628/1 and ST/M503599/1). We thank S. Falle for the use of the mg hydrodynamics code used to calculate the simulations in this work and S. van Loo for adding SILO output to it. The calculations used in this paper were performed on the DiRAC1 Facility at Leeds, which is jointly funded by STFC, the Large Facilities Capital Fund of the Department of Business, Innovation and Skills, and the University of Leeds. The 3D volumetric renderings were created using the VisIt visualisation and data analysation software (Childs et al. 2012 ). The data associated with this paper are available from the University of Leeds data repository ( https://doi.org/10.5518/573 ).
Paper: 2021ApJ...908..151T (fulltext)
; Tiwari et al. 2019 ), and the hottest loops can live up to multiple hours (López Fuentes et al. 2007 ; Klimchuk et al. 2010 ). Most of the longest lived loops are probably a result of interaction of multiple loop strands, which sequentially heat making the loop bundle live much longer than a loop strand (e.g., Warren et al. 2002 ). Loop 23 in Table 1 shows this behavior of sequential heating (see the animation of Figure 1 ). The coordinates of the center of the box outlining the two footpoints for each loop are also given in Table 1 . For easy identification, the coordinates in bold font indicate that the surrounding box has unipolar magnetic flux; all others have mixed-polarity magnetic flux. We quantified the distribution of magnetic field strength and polarity at each foot of each loop via histograms of LOS magnetograms within the selected footpoint region, in the way presented in Figure 5 for three examples. We use a 6′′ × 6′′ boxed area centered at each foot to obtain LOS magnetogram histograms at each foot. The 6′′ × 6′′ size of the box ensures that the loop foot is completely covered within the box. We have three categories of loops: (1) loops with unipolar magnetic flux in both feet, (2) loops with mixed-polarity magnetic flux in both feet, and (3) loops having one foot in mixed-polarity flux and the other foot in unipolar magnetic flux. Of the 23 loops that we examined, ∼40% (9/23) have both feet in unipolar magnetic flux, ∼4% of loops (1/23) have both feet in mixed-polarity flux, and ∼56% (13/23) have one foot in unipolar and one in mixed-polarity flux. We have considered the presence of mixed-polarity only when the LOS magnetogram at the loop foot contains values ≥20 G of positive or negative minority polarity flux. The random noise level in the LOS magnetograms is about 7 G (Couvidat et al. 2016 ), fairly well below our selected lower limit. We would have counted footpoint 2 of our only category 2 loop as unipolar if there was no minority-polarity pixel with a negative B z value larger than 20 G, i.e., if there were no pixel in the second bin (in 20–40 G range, bin size is 20 G) on the negative/minority polarity side of the zero line in this footpoint box’s histogram in Figure 5 . This case is the most marginal one in our sample—in all other mixed-polarity cases there are several to many minority-polarity pixels with the magnitude of B z larger than 20 G. Because there is only one loop in the category of both loop-feet having mixed-polarity flux, in our discussions we often count it in the category of loops having at least one foot in mixed-polarity flux region. In Figure 2 we show three example loops, one from each of the three categories. In the top row of Figure 2 , we show an example loop that has both feet in unipolar magnetic flux (category 1); the peak time in the hot 94 image comes at 05:02:59 and the associated LOS HMI magnetogram is obtained at 05:02:52. The light curve for this loop in Figure 3 (a) also shows another
Paper: 2020A&A...641A.124L (fulltext)
[CII] redshift ( Carilli et al. 2013 ). The rest-frame frequencies of HF J = 1–0 and H 2 O (2 20 –2 11 ) are 1232.476 and 1228.789 GHz respectively. In contrast, the line width of the H 2 O (2 20 –2 11 ) in the QSO is, within the uncertainties, exactly the same as for the other lines. Neither the QSO or the SMG show clear evidence for an outflow in the H 2 O line, although there is a possible excess in the line profile of the QSO seen ∼400 km s −1 blueward of the systemic H 2 O redshift, consistent with the velocities of the weak [CII] outflow ( Carilli et al. 2013 ). 4. Discussion 4.1. HF: Molecular gas irradiated by AGN and young stars The detection of HF emission in the QSO is unusual. In most of the sources observed thus far, the HF J = 1–0 line is observed as an absorption line with little or no emission ( van der Werf et al. 2010 ). Only a handful of sources are known to have HF purely in emission, without any obvious associated absorption. These sources are the Orion Bar ( van der Tak et al. 2012 ; Kavak et al. 2019 ), Mrk 231 ( van der Werf et al. 2010 ), and a few nearby galaxies ( Pereira-Santaella et al. 2013 ; Lu et al. 2017b ). To increase the number of galaxies observed with HF in emission, we extracted the HF emission line fluxes from local galaxies observed with Herschel in the sample of Yang et al. (2013) , which all host AGN. We find that the galaxies that host AGN and the QSO in BR 1202-0725 have an approximately constant ratio of L HF / L IR irrespective of L IR (Fig. 3 ). We note, however, that the exact slope of the relation only considering the local AGN and QSO is dependent on the far-infrared luminosity used. Other estimates of the infrared luminosity of the QSO in the literature would lower L IR / L HF (cf. Salomé et al. 2012 ; Wagg et al. 2014 ; Lu et al. 2017a ). Fig. 3. Relationship between the infrared luminosity and the ratio of the HF luminosity and infrared luminosity of BR 1202-0725, a sample of local AGN ( Yang et al. 2013 ), and the Orion Bar ( Salgado et al. 2016 ; Nagy et al. 2017 ). All sources are labeled. We show a least-squares fit to the local AGN, QSO, and the Orion Bar (black line with the dotted line extension to the Orion Bar), which implies the ratio is constant. We note that fitting the AGN and QSO only yields a similar fit but with a slightly steeper slope. For comparison, we show the ratio L HF / L IR for the entire IR luminosity of the Orion Bar and the luminosity if only 2% of the total IR luminosity is within the Herschel beam used to measure the HF flux (see text for details). The Orion Bar is the only known galactic source or sight-line where HF appears purely in emission (although we note that the absorption of HF and the nearby H 2 O (2 20 –2 11 ) could mask any emission in other galactic and extra-galactic sources; see e.g., Monje et al. 2014 ; Neufeld et al. 2010 ; Sonnentrucker et al. 2015 ). If we compare L HF / L IR and L IR for the Orion Bar ( Salgado et al. 2016 ; Nagy et al.
Paper: 2021ApJ...908..200W (fulltext)
afterglow emission takes place when the Lorentz factor of the core of the blast slows down to Γ ∼ 1/ θ v which falls into mildly relativistic domain if θ v is relatively large. Since we wish to follow the blast emission up to, and shortly after, its peak, a treatment that follows the blast deceleration down to a trans-relativistic speed is required. In this work we use a semi-analytical model following the shock jump conditions derived by an approximate trans-relativistic EoS (Uhm 2011 ). This model works well at both high and modest Lorentz factors and can smoothly transform to a Sedov–Taylor solution when the shock is no longer relativistic. Our model assumes one-dimensional radial solution for each blast segment. Therefore, we neglect sideways (lateral) expansion in our model. The sideways expansion of the blast is important only at later stages of the blast evolution and it affects the light curve predictions after its peak (Gill Granot 2018 ). However, the key factor to constrain inclination angle, of interest in this study, is the arising slope before the peak flux, where the jet is still relativistic. So the sideways expansion will not influence our result. To describe the physical quantities more clearly, here we consider two reference frames. One is burst rest frame (hereafter BR frame) which is at rest with the central engine. The other is the shock co-moving frame (hereafter CM frame) which is at rest with respect to the gas just downstream of the forward shock. Primed variables (i.e., denoted with ′) are measured in CM frame, while non-primed ones are measured in BR frame. Under the approximation of neglecting sideways expansion, we can regard the jet evolution at each radial direction as being independent. It is therefore convenient to describe the evolution at each direction by using isotropic equivalence quantities. Consider an isotropic burst of total energy E iso = Γ 0 M ej c 2 , where M ej is the mass of ejecta and Γ 0 is its initial Lorentz factor. The outflow spreads in a cold uniform environment of number density n and forms a relativistic shock. We assume that the blastwave is expanding adiabatically without further energy injection (for considerations of energy injection in the blast see Lamb et al. 2020 ). Energy loss due to radiation is negligible compared with E iso . We define the radius R of the shock by the distance from burst center to the forward shock. All swept-up mass is compressed in the downstream of forward shock, moving with Lorentz factor Γ measured in BR frame, and its isotropic kinetic energy is E k . We also assume that all ejecta are moving together with shocked ambient at the same Lorentz factor behind the contact discontinuity, and its total energy is E ej = Γ M ej c 2 . The trans-relativistic shock jump conditions is (Uhm 2011 ; Ryan et al. 2019 ) where , , and are the number density, internal energy, and pressure of shocked region, respectively. is the dimensionless speed. c is the speed of light.
Paper: 2019ApJ...870...83J (fulltext)
enhanced, so Tri II does not follow the strange K–Mg anticorrelation in NGC 2419. Note that the halo sample here is different than in Figures 2 and 4 because our usual halo compilation does not have K abundances (Frebel 2010 ). Adapted from Venn et al. ( 2017 ). 4.4. Iron-peak Elements: Cr, Mn, Co, Ni, Zn The Fe-peak abundances were determined with equivalent widths, except for Mn, which is synthesized owing to hyperfine structure (Kurucz Bell 1995 ). In Gru I we can constrain Cr, Mn, Co, and Ni, finding that both stars have essentially identical abundances of these elements. Though the Cr and Mn abundances are similar to those in metal-poor stars in other UFDs or in the Milky Way halo, the Co and Ni abundances are somewhat higher. However, this difference is not very significant, especially for Co, which is derived from only a few bluer lines. One Zn line is marginally detected in GruI-032 with an abundance consistent with the halo trend, though with large uncertainty. In Tri II, we can detect Cr and Ni in TriII-40 and provide upper limits in TriII-46. Mn, Co, and Zn are unconstrained, as they only have strong lines blueward of 5000 Å. The upper limits for Cr and Ni in TriII-46 are uninteresting. For TriII-40, we detect a normal [Cr/Fe] ratio, but Ni appears significantly enhanced ([Ni/Fe] = 0.57 ± 0.16), in agreement with Kirby et al. ( 2017 ) and Venn et al. ( 2017 ). 4.5. Neutron-capture Elements: Sr, Ba, Eu Strontium is detected only in Gru I, as the strong Sr ii lines at 4077 and 4215 Å are out of the range of the Tri II (GRACES) spectra. The abundance of both lines is determined with spectrum synthesis. The Sr abundances in these two stars are very similar, [Sr/Fe] ≈ −2, which is much lower than what is found in most halo stars but similar to most UFDs (Figure 4 ). Barium is measured with four different lines in the Gru I stars including hyperfine structure and isotope splitting (McWilliam 1998 ). We use solar isotope ratios (Sneden et al. 2008 ), but given the low overall abundance, changing this to r - or s -process ratios does not significantly affect our abundances. GruI-032 has a low [Ba/Fe] ≈ −1.6, but GruI-038 has a much higher Ba abundance [Ba/Fe] ≈ −1.0. This is formally only 1.6 σ different, but differential comparison of the line strengths (e.g., the λ 6497 line in Figure 1 ) suggests that the difference is real. We discuss this more in Section 5.1 , but both Ba abundances are low and similar to those in most UFDs. Ba is not detected in either Tri II star, so instead we place 5 σ upper limits. The Ba limit for TriII-40 is [Ba/Fe] −1.25, suggesting a low Ba abundance similar to other UFDs. Kirby et al. ( 2017 ) determined [Sr/Fe] = −1.5 and [Ba/Fe] = −2.4 from their HIRES spectrum of this star, consistent with our upper limit and showing that TriII-40 clearly has very low neutron-capture element abundances. The Ba limit for TriII-46 is only [Ba/Fe] ≲ −0.2, but this is still at the lower envelope of the halo trend (Figure 4 ).
Paper: 2017ApJ...836..239V (acknowledgments)
We thank the anonymous referee for insightful remarks that helped us to improve the paper. B.V. acknowledges the support from an Australian Research Council Discovery Early Career Researcher Award (PD0028506). This work was partially supported by grants ARC FT130101593, and HST /GO 13767, 12905, and 12572. Facilities : HST (ACS), HST (WFC3).