adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019ApJ...880...62M (fulltext)
open field lines. The plots of null height (Figure 4 (a), red line) and free magnetic energy (Figure 4 (b), red line) also indicate that the system has almost relaxed back to its equilibrium state. The free magnetic energy (b) sharply drops as soon as the photospheric driving stops. This is because energy is being propagated out of the open top boundary along the jet spire, but is no longer being injected through the lower boundary. The null height (a) continues to increase until around t = 70, showing that the jet dome has continued to expand. The decrease in null height after this point is more rapid than that in PAD09 (black line) as there is no longer photospheric driving injecting energy and twist into the field. We now compare our PAD10 simulation with that of Pariat et al. ( 2010 ), as well as PAD09 and PAD09R. In Pariat et al. ( 2009 ) they see a peak and sudden drop in free magnetic energy when the eruption occurs after kinking breaks the symmetry of the model. In Pariat et al. ( 2010 ), the model is initially set up to be non-symmetric, with the background magnetic field inclined at 10° to the vertical. They continue to rotate the positive magnetic polarity after the first eruption, resulting in a series of eruptive events, hence a series of peaks and drops in free energy. The free magnetic energy in our model can be seen in Figure 4 (b) (blue line). We do see a series of peaks in our model, but they are not sharply defined, as there is not as rapid a decrease in energy after each peak. Again, this is due to the quasi-static nature of the magnetofrictional method. We cannot accurately follow the dynamics of an eruption, so the result is that not as much energy and twist is released before it starts building up again due to the photospheric driving. The first peak occurs at a similar time to that of our PAD09 simulation. The peak free magnetic energy built up before the first release is 1.45 × 10 25 erg and 1.29 × 10 25 erg for the PAD09 and PAD10 simulations, respectively. The energy buildup is greater in PAD09, as changes in connectivity do not occur until after the onset of the kink and the breaking of the structure’s symmetry. In PAD10, the symmetry is already broken in the initial setup, so changes in connectivity can occur as soon as the simulation begins and magnetic energy can already be released. Figure 5 (a) shows in the xz -plane at y = 9 Mm for PAD09 (left), PAD10 (center), and PAD09R (right). An animated version of these panels is available online. In PAD09 and PAD09R, current is initially all contained within the twisted dome of the jet and along the narrow, vertical jet spire, indicating that no twist has been released and almost all non-potentiality is contained within the dome. After the onset of the kink, the symmetry of the system is broken and we see a broader jet spire, as in Figure 5 (a). In PAD10, the spire broadens more rapidly, as connectivity changes occur with background magnetic field right from the start
Paper: 2020MNRAS.496..206C (fulltext)
variables to their conservative form and back (using the Newton–Raphson-based root-finder of Noble et al. 2006 for the conservatives-to-primitives step, as described next). The detailed Jupyter notebook used to generate needed GRHD equations in MANGA may be found in the NRPy+ github repo 4 or be viewed directly via nbviewer . 5 Before proceeding, we note two important points. Previously, MANGA used only cgs units, but for GR, G = c = 1 units are superior so the code was extended to support them. Second, although our MM code is limited to periodic boundary conditions, no boundary conditions are applied to the metric (as it is static and known), so there is no ‘periodic gravity.’ We simply ensure that the outer boundaries of our numerical domain are sufficiently far away from regions of interest, so that as outer (periodic) boundary conditions are applied to the hydrodynamic variables, they have little to no effect on the simulation. In future work, we may as needed adopt radiation inflow and outflow boundary conditions as described in Chang et al. ( 2020 ) to the hydrodynamic variables on the boundary, to produce arbitrary boundary conditions. We note that this implementation would be similar to the boundary conditions for a ‘sphered cube’ (Burns et al. 2020 ). 3.1 Time-stepping We time integrate equation ( 9 ) as follows: Estimate the Courant-limited time-step Δ t for each cell, as described in Chang et al. ( 2017 ). The time-step can either be an individual time-step in a multistep algorithm (Prust Chang 2019 ) or a global time-step (Chang et al. 2017 ). Estimate the half time-step state of the cell using the second-order van Leer scheme described in Chang et al. ( 2020 ). In brief, we solve the RHS of equation ( 9 ), but with the replacement of Δ t → Δ t /2 and use piecewise continuous reconstruction to get face values. We then solve for the fluxes following the methodology described next and include the source terms on a half-time-step. Drift the mesh generating points by a half-time-step and recompute the half-time-step tessellation to provide second order convergence in time (e.g. provide an estimate for ). Use the half-time-step states and use full linear reconstruction to derive face values to compute the half-time-step fluxes, (described next). Then, apply the full step including the source terms (using the half-step values), e.g. . Update the state of the cell to the full step. As stated in Chang et al. ( 2020 ), the use of the van Leer half-step prediction allows the source terms to be automatically included at second order. It also simplifies the code as the equations are only written once. Moreover, the van Leer method can be easily adapted to multistepping schemes as we have done here. The only change is that the half-time-step estimate must be taken from the cell’s initial state. MANGA was originally designed to reconstruct the conserved values at cell faces. For Newtonian codes, this is a reasonable choice, but modelling
Paper: 2020ApJ...901..136M (fulltext)
−1 , with an average value of 0.2 eV molecule −1 . For crystalline ice the energy needed to convert it into compact amorphous ice is higher. For instance, Leto Baratta ( 2003 ) estimated that this energy ranges between 0.59 and 1.42 eV molecule −1 for pure crystalline water ice and for various projectiles. Famá et al. ( 2010 ) obtained an energy of eV molecule −1 for a pure crystalline water ice impacted by H + ions with an energy of 225 keV. Note that Raut et al. ( 2008 ) suggested the existence of a stopping power threshold of 4 eV Å −1 to trigger the destruction of the open porosity of the ice. However, the presence of such a threshold in compaction at low stopping power was not confirmed for the high-energy cosmic rays encountered in interstellar space (Dartois et al. 2013 ). These results are not surprising since every molecule is coordinated with the neighbors in the lattice. The photodissociation of a single water molecule, breaking and reformation of bonds, and formation of radicals or other species such as H 2 O 2 affect the unit cell of the crystalline structure, and hence a larger number of molecules are affected (Leto Baratta 2003 ). Figure 8 shows the impact of the energy needed to compact ice on the depth down to which cometary ice inside the nucleus is fully converted into compact amorphous ice. For an energy varying between 10 −2 and 10 eV molecule −1 , this depth varies between ∼28 and ∼8 m. The depth down to which the ice structure is fully changed by irradiation is not strongly dependent on the energy needed to compact ice but is rather controlled by the sharp drop of the deposited energy by GCRs inside the nucleus. Figure 8. Maximal depth until which the cometary ice is completely converted into compact amorphous ice as a function of the energy required to convert the ice into amorphous ice. One important implication of this result is that, regardless of the water ice structure inside the cometary nucleus, its outer layer (the first ∼10–20 m) consists of compact amorphous ice and is not necessarily representative of the whole nucleus. This has strong implications for dynamically new comets (i.e., a comet that is approaching the Sun and outgasses for the first time). For instance, recent observations by Meech et al. ( 2009 ) showed that the activity at large heliocentric distances of comets passing through the solar system for the first time may be caused by the annealing of amorphous water ice. According to our results, such observations of the altered outer layers of a cometary nucleus cannot be taken as a proof that comets as a whole are composed of amorphous ice and cannot be used to infer the temperature and region where comets have formed in the protosolar nebula. 5. Discussion and Conclusion In this paper we analyzed the effects of GCRs on the isotopic composition, chemical composition, and ice structure inside a cometary nucleus. We did not consider surface effects such as sputtering, which leads to erosion of the
Paper: 2016MNRAS.459L..71N (acknowledgments)
We thank Seshadri Nadathur, András Kovács and Juan García-Bellido for correspondence and discussion. Some of the results in this Letter were derived using the healpix (Górski et al. 2005 ) package. ABL thanks CNES for financial support through its post-doctoral programme. This work was supported in part by the Kavli Institute for Cosmological Physics at the University of Chicago through grant NSF PHY-1125897 and an endowment from the Kavli Foundation and its founder Fred Kavli. OL acknowledges support from a European Research Council Advanced Grant FP7/291329, which also supported ABL and KN.
Paper: 2017MNRAS.471.3398A (acknowledgments)
Acknowledgements We would like to thank the anonymous referee for many helpful comments. This work is supported by the Canadian Natural Sciences and Engineering Research Council. TAAS also gratefully acknowledges support from Western University. This research made use of the SIMBAD data base, operated at CDS, Strasbourg, France.
Paper: 2020ApJ...898L..45O (fulltext)
for the water production rate of both asteroids. In Section 5 , we calculate the fraction of ice coverage required to achieve this upper limit, discuss our results, and present our conclusions. 2. Observations and Data Analysis 2.1. Herschel HIFI Observations Themis was observed with HIFI on Herschel on 2012 January 30 (see Table 1 ). The asteroid was traveling inbound (2.92 au) and close to perihelion (2.75 au) when the observation was carried out. Cybele was observed with HIFI on 2012 December 21 (see Table 1 ) when the asteroid was 3.52 au from the Sun and inbound, having already passed aphelion (3.8 au). It is important to highlight that with the low escape velocity from these asteroids, any escaping water vapor would have been visible to Herschel independent of its source location, as the footprint of the telescope was much larger than the asteroids themselves. Table 1 Observations and Observing Conditions for the Previously Unpublished Measurements Used in This Paper Note. Columns: target (the two asteroids), observatory and instruments used, ObsID (observation ID reference number), date and time in UT, rh (heliocentric distance), Δ (Earth to asteroid distance), α (phase angle). The line emission from the fundamental ortho-H 2 O 1 10 –1 01 line at 556.936 GHz was searched for in the lower sideband of the HIFI band 1b mixer. The observations were performed in the frequency-switching observing mode using both the wide-band spectrometer (WBS) and the high-resolution spectrometer (HRS). The spectral resolution of the WBS is 1 MHz (0.54 km s −1 at the frequency of the observed line), while the HRS was used in its high-resolution mode with a resolution of 120 kHz (0.065 km s −1 ). Initial processing of the HIFI data set was carried out using the standard HIFI processing pipeline of the Herschel Interactive Processing Environment v15.0.1 (HIPE; Ott 2010 ), followed by processing with the GILDAS CLASS software. In CLASS, vertical and horizontal polarizations were folded and averaged. The frequency-switching observing mode introduces quite a strong baseline ripple, which was removed through a polynomial baseline correction performed on the final spectra. The averaged and baseline-corrected HRS spectra of the two asteroids are shown in Figures 1 (a) and (b). The pointing offset of both the vertical and the horizontal polarization spectra is 35 in band 1b with respect to the target position. The half-power beam width is 381 at 557 GHz. The main beam brightness temperature scale was computed using a forward efficiency of 0.96 and a beam efficiency of 0.62. The output line area rms values are provided in Table 2 . Figure 1. Averaged HIFI spectra and thermophysical modeling output. Shown are the HRS-averaged spectra for (a) asteroid Themis and (b) asteroid Cybele with TPM outputs (c) for Themis (left) and Cybele (right). These outputs represent the Herschel view with orientation “ecliptic sky” at the specific HIFI observing start epochs. The plots in
Paper: 2016MNRAS.463.4459B (fulltext)
forbidden line emission such as [O i ] lines (Appenzeller, Oestreicher Jankovics 1984 ; Hartigan, Edwards Ghandour 1995 ). Our spectra show that the EW of the [O I] (6300 Å) line increases by a factor of 10 after the eclipse (Figs 6 and 7 ). The observed EW increase is in agreement with measurements by Petrov et al. ( 2015 ) and Takami et al. ( 2016 ). Petrov et al. ( 2015 ) attribute the change in the forbidden line's EW to a drop in the photospheric continuum, as reported by the resolved photometry of RW Aur A by Antipin et al. ( 2015 ). Petrov et al. ( 2015 ) further imply that the obscuring body must be blocking the inner parts of the system while the outflow remains visible. The He i (6678 Å) and Hα emission lines remain clearly visible in our spectra after RW Aur enters the eclipse. This again indicates that at least part of the accretion flow is still visible. The He i EW remains constant (Fig. 8 ), implying the He-emitting regions are obscured in the same way as the photosphere. The Hα EW decreases after the eclipse (Fig. 7 ). Our EW measurements show lower values for the line during the dimming compared to the values reported by both Petrov et al. ( 2015 ) and Takami et al. ( 2016 ). We attribute this to the fact that our spectroscopy does not resolve the A and B components, hence we get an additional flux in the continuum from RW Aur B, resulting in a lower EW value. The Hα line profiles show a suppressed blueshifted peak after the start of the dimming (Fig. 5 ). This can also be seen in the Hα profiles from Takami et al. ( 2016 ) from both the 2010 and 2014 dimmings. This blueshifted absorption can be interpreted as the presence of an absorbing material moving towards the observer's line of sight, i.e. an outflow. 4.3 Origin of the eclipse In Section 1 , we introduce two possible explanations for the long dimming event observed in RW Aur. The obscuring screen could either be a tidal arm in the outer disc, as suggested by Rodriguez et al. ( 2013 , 2016 ), or a wind emanating in the inner disc, suggested by Petrov et al. ( 2015 ). Our observations provide new evidence to support the occultation by a wind. In particular, we show that the blueshifted portion of the Hα line is suppressed since the beginning of the eclipse, indicating the presence of outflowing and thus blueshifted gas with velocities of 100–200 km s −1 in the occulting screen. This scenario is further supported by the increase in the mid-IR fluxes of the system in eclipse, reported by Shenavrin et al. ( 2015 ) and visible in archive NEOWISE data in the W 2 band (Fig. 4 ). Shenavrin et al. ( 2015 ) fit a blackbody to the spectral energy distribution of RW Aur in its faint state and find that the excess IR emission is plausibly explained by the presence of hot dust with temperature around 1000 K along the line of sight. Thus, the occulting body is outflowing and contains hot dust, which fits with a scenario where a wind emanating from the inner disc eclipses the
Paper: 2015MNRAS.450L..71F (acknowledgments)
We thank Peter Goldreich, Phil Hopkins, Christine Moran, Evan Kirby, and Sterl Phinney for useful discussions. We acknowledge stimulating workshops at Sky House where these ideas germinated. JF acknowledges support from NSF under grant no. AST-1205732 and through a Lee DuBridge Fellowship at Caltech. CDO acknowledges support from NSG through grants PHY-1151197 and AST-1205732.
Paper: 2017MNRAS.471..317R (acknowledgments)
Acknowledgements The authors thank the reviewer for a set of helpful comments, Richard Booth, Charles Gammie, Sebastien Fromang, and Ji-Ming Shi for important feedback on an earlier draft of the paper, and Kacper Kornet for advice on numerical issues. This research is partially funded by STFC grant ST/L000636/1. Many of the simulations were run on the DiRAC Complexity system, operated by the University of Leicester IT Services, which forms part of the STFC DiRAC HPC Facility ( www.dirac.ac.uk ). This equipment is funded by BIS National E-Infrastructure capital grant ST/K000373/1 and STFC DiRAC Operations grant ST/K0003259/1. DiRAC is part of the UK National E-Infrastructure.
Paper: 2016MNRAS.458..105K (acknowledgments)
We thank the anonymous referee for his/her careful reading of the manuscript and the valuable feedbacks. We extend our gratitude to Jacob A. Arnold and Kristin A. Woodley regarding their help in observations. We thank Blesson Mathew and Nicola Pastorello for the careful reading of the manuscript. We also acknowledge the members of SAGES group, especially Christopher Usher, Vincenzo Pota and Joachim Janz, for the support and enlightening discussions. SSK thanks the Swinburne University for the SUPRA fellowship. DAF thanks the ARC for support via DP-130100388. This work was supported by NSF grant AST-1211995. This paper was based in part on data collected at Subaru Telescope, which is operated by the National Astronomical Observatory of Japan. This paper uses data products produced by the OIR Telescope Data Center, supported by the Smithsonian Astrophysical Observatory. This work is based on observations made with the NASA/ESA Hubble Space Telescope , and obtained from the Hubble Legacy Archive, which is a collaboration between the Space Telescope Science Institute (STScI/NASA), the Space Telescope European Coordinating Facility (ST-ECF/ESA) and the Canadian Astronomy Data Centre (CADC/NRC/CSA). Some of the data presented herein were obtained at the W. M. Keck Observatory, operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration, and made possible by the generous financial support of the W. M. Keck Foundation. We wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community. We are most fortunate to have the opportunity to conduct observations from this mountain. This research has made use of the NASA/IPAC Extragalactic Database (NED) operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. The analysis pipeline used to reduce the DEIMOS data was developed at UC Berkeley with support from NSF grant AST-0071048. We acknowledge the usage of the HyperLeda data base ( http://leda.univ-lyon1.fr ).
Paper: 2019MNRAS.488.4942P (fulltext)
are dominated by dark matter halo once one goes sufficiently far out in the galaxy. What is relevant is that the halo is dense and compact and hence the dark matter dominates even at inner radii as in LSBs (e.g. Ghosh Jog 2014 ) or in dwarf galaxies like DDO 168 (Ghosh Jog 2018 ). It is this feature that influences the development of a slow bar due to dynamical friction. DDO 168 has a high central dark matter density, ρ 0 = (39.81 ± 6.37) × 10 −3 and a scale radius, r s = 2.81 ± 0.83 kpc for an isothermal dark matter halo (Oh et al. 2015 ). It has been found that the dark matter halo dominates the observed dynamics of this galaxy starting from the central region (Ghosh Jog 2018 ). Thus, the detection of a slow bar in DDO 168 supports the proposition that the dense, compact dark matter halo slows down the bar through dynamical friction. We note that all the three galaxies that so far have been seen to host an H i bar (NGC 2915, NGC 3741, and DD0 168), also have the property that the bar is slow. We further note that the H i gas is very extended (several times the optical radius) in all three galaxies, and in NGC 3741 and DDO 168 the halo is dense and compact (Ghosh Jog 2018 ). Here, it should be emphasized that now, there is evidence that slow bars are more common than what was thought before. For example, using an innovative and simple technique, Font et al. ( 2017 ) determined the corotation radius in 68 galaxies. Using the photometric data, they further determine the bar length in these galaxies and consecutively estimate the to find that almost one-third of their sample galaxies host ‘slow bars’ (i.e. ). This commonality of slow bars in galaxies can be thought of due to more effective braking by the dark matter halo than what was believed before. However, the status of the pattern speeds in dark matter dominated galaxies, especially LSBs, is not well studied as compared to the HSBs (Font et al. 2017 ; Peters Kuzio de Naray 2019 ). As shown in our works for DDO 168 and NGC 3741, the dense and compact halo might play an import role in slowing down the H i bars. Thus, though the braking of bars by dark matter halo is more common, yet the dark matter dominance from the central region might have an enhanced effect in determining the slowdown of the bar. Further, Font et al. ( 2014b ) have shown by simulation that, due to the dynamical braking by the dark matter halo, the corotation radius moves outwards, at the same time, the bar can also grow in length significantly due to accretion. This implies, even for a slow bar the ratio, can settle well below the nominal value of 1.4, thus declaring the bar to be a fast one. Hence, all the fast bars as decided by criteria might not be fast ones. However, we clarify that the traditional definition of a slow bar (i.e. ) does not change our conclusion on the existence of the slow bar in DDO 168, albeit it puts questions on its rarity. 4 CONCLUSION We have identified an H i bar in the dwarf galaxy DDO
Paper: 2015ApJ...808...83Z (acknowledgments)
We would like to thank the NASA CDAWEB for providing the public WIND 3DP, MFI, and SWE data. We also acknowledge Dr. Lief Svalgaard and Dr. Lan Jian for providing the sector boundary list and the lists of ICMEs and CIRs. This work is jointly supported by the National Basic Research Program of China (Grant No. 2012CB825601), the National Natural Science Foundation of China (41231068, 41174152, 41374188, 41304146), and the Specialized Research Fund for State Key Laboratories.