adsabs/WIESP2022-NER huggingface dataset:
Paper: 2019MNRAS.490.3098T (acknowledgments)
ACKNOWLEDGEMENTS Calculation of this work is performed of the XC30 and XC50 system at the Center for Computational Astrophysics (CfCA), National Astronomical Observatory of Japan (NAOJ). This work supported by JSPS KAKENHI Grant Number JP 19J13373 (RT), Society for the Promotion of Science Grant-in-Aid for Scientific Research (A) (17H01102 KO; 16H02170 TT), Scientific Research (C) (16K05309 KO; 18K03710 KO), and Scientific Research on Innovative Areas (18H04592KO; 18H05463 TT). This research is also supported by the Ministry of Education, Culture, Sports, Science and Technology of Japan as ‘Priority Issue on Post-K computer’ (Elucidation of the Fundamental Laws and Evolution of the Universe) and JICFuS. RT acknowledges the support by JSPS Overseas Challenge Program for Young Researchers. CD acknowledges the Science and Technology Facilities Council (STFC) through grant ST/P000541/1, and visitor support from Kavli IPMU.
Paper: 2018A&A...614A..25L (fulltext)
2009 ). 3 Test model A simple 1-D test model is now defined. This allows numerous simulated data sets to be created so that the p -values computed according to Sects. 2.1 and 2.2 can be compared and any differeces established with high statistical confidence. 3.1 Hubble expansion The null hypothesis is that the local Universe is undergoing an isotropic, cold Hubble expansion with Hubble parameter h 0 = 70 km s −1 Mpc −1 . Moreover throughout the local Universe there exists a population of perfect standard candles of absolute bolometric magnitude M = −19.0. Finally, we suppose that the redshifts z of these standard candles are measured exactly, but that their apparent bolometric magnitudes m have normally-distributed measurement errors with σ m = 0.3 mag. 3.2 Ideal samples In order to test p -values when the null hypothesis is correct, we need data sets consistent with the above description of the Hubble flow. We define the local Universe as extending to z max = 0.2, so that to a good approximation the geometry is Euclidean. If we suppose the standard candles to be uniformly distributed in space, a sample { z i } is created with the formula where the x i are independent random numbers ∈ (0, 1). The distances of these n standard candles are and their measured apparent bolometric magnitudes are where the z G are independent gaussian variates from . The n pairs comprise a data set D from which the posterior probability density Λ( h | D ) of the Hubble parameter h can be inferred. Moreover, the global GOF between Λ( h | D ) and D can be tested with the p -values of Sects. 2.1 and 2.2 . 3.3 Imperfect samples In order to test p -values when the null hypothesis is false, we need to modify either the model or the data. If we continue to make inferences based on the Hubble flow defined in Sect. 3.1 , we can introduce a violation of the null hypothesis by contaminating the data sample with an unrecognized second class of standard candles with a different absolute bolometric magnitude. Thus the data set D now comprises n 1 pairs with given by Eq. ( 15 ) and n 2 pairs with given by and where n = n 1 + n 2 . Thus, the sample is now contaminated with a second population of standard candles that are brighter by Δ M magnitudes. The investigator is unaware of this, and so carries out a Bayesian analysis assuming an ideal sample. 4 Comparison of <italic>p</italic>-values The Bayesian p -values defined in Sect. 2 are now computed for simulated data samples derived as described in Sect. 3 . Fig. 1 Comparison of p -values. For J = 200 ideal samples D j , is plotted against log p B . 4.1 Null hypothesis correct In order to compare p -values in this case, J independent ideal samples D j are created as described in Sect. 3.2 , and then p -values for each D j computed as described in Sects. 2.1 and 2.2 . With a constant prior π ( h ) in the interval ( h 1 , h 2 ) with h 1 = 64 and h 2 = 76 km s −1 Mpc −1 , the posterior density Λ( h | D j ) is obtained from Eq. ( 1 )
Paper: 2016A&A...589A..87B (fulltext)
The saturation of the MRI and the outcome of a disc subject to these non-ideal effects is still the subject of intense debate. Stratified models including Ohmic and ambipolar diffusion only have shown that the disc could become essentially laminar (i.e. not turbulent) in the midplane, with a strong magnetically driven jet at the disc surface ( Bai Stone 2013 ; Simon et al. 2013 ; Gressel et al. 2015 ). For radii 1 between 1 and 10 au (where Hall effect dominates diffusive processes) a large-scale midplane stress is found by adding the Hall effect, with a strong field polarity dependence and no turbulence ( Lesur et al. 2014 ; Bai 2014 ), whereas the outer disc is found to be essentially insensitive to the Hall effect ( Bai 2015 ), though some polarity dependence could persist ( Simon et al. 2015 ). The results of Lesur et al. (2014) and Bai (2014) and, in particular, the presence of a large-scale “laminar” stress are both important and potentially problematic results. In effect, these simulations were done in the stratified shearing box framework ( Hawley et al. 1995 ) with a very limited radial extension (typically a few scale heights), so the Lesur et al. (2014) and Bai (2014) results indicate that the MRI is indeed active in these simulations, but it tends to produce structures that are larger than the simulation box. As a matter of fact, unstratified shearing box models have already indicated that the Hall-dominated MRI is prone to produce large-scale structures, such as zonal fields and flows ( Kunz Lesur 2013 , hereafter KL13) in wide enough boxes. All of these elements tend to point toward the fact that the Hall effect could be an efficient mechanism for triggering self-organisation in protoplanetary discs. In this paper, we explore this hypothesis using full 3D cylindrical models of accretion discs dominated by the Hall effect. Our aim is not to reproduce the exact structure of a disc with its detailed ionisation equilibrium but to test the hypothesis of self-organisation in a global model, free of the artefacts found in local, i.e. shearing-box simulations. We start by describing the physical framework in detail and giving and some remarkable facts concerning Hall dynamics. We then depict our numerical model and present the tests performed to validate it along with a reproduction and discussion of the multifluid models of O’Keeffe Downes (2014) . Next, we move to the core of the paper with the presentation and characterisation of several Hall-dominated models, putting emphasis on the physical mechanisms driving self-organisation. We then explore some potential treats to these mechanisms such as diffusive effects and more complex field configurations. We conclude by discussing the link between these results and recent observations along with extensions of this work to more realistic numerical setups. 2. Framework 2.1. Physical model We want to model a thin disc of partially ionised gas orbiting a central mass and threaded by a
Paper: 2017MNRAS.471.3152P (acknowledgments)
Acknowledgements We thank Andrew Baker, Laure Ciesla, Annalisa De Cia, Loretta Dunne, Haley Gomez, Leslie Hunt, Suzanne Madden, Joshua Peek, Molly Peeples, Paola Santini, Raffaella Schneider, Jason Spyromilio, Phil Wiseman and Svitlana Zhukovska for stimulating discussions and/or providing data. We thank the referee for useful suggestions. We further acknowledge the Kavli Institute for Theoretical Physics ‘Cold Universe’ workshop where some of this work was completed. This research was supported in part by the National Science Foundation under Grant No. NSF PHY11-25915. RSS thanks the Downsbrough family for their generous support, and acknowledges support from the Simons Foundation through a Simons Investigator award.
Paper: 2017MNRAS.469L..63R (fulltext)
have half-light radii of 60pc, would be barely resolved with NIRCam. 3.2.2 The main epoch of GC formation Each individual GCP will be totally dark (in UV) before experiencing its main (1G) burst of star formation, quickly reaching its peak SFR and brightness, then followed by rapid decline, possibly punctuated by a series of secondary (2G) bursts. Then rapidly turning passive and fading well below observability. A uniform GC formation rate between z 8 and 3 has been assumed above, but it may well be that the GC formation rate is initially very low, then peaks at a specific redshift and then declines, with few clusters still forming at lower redshifts. Having explored this redshift range, observations should then be able to identify such a golden age of GC formation, therefore establishing the relative timing of GC and galaxy formation, and providing, incidentally, a straight check for local GC ages as inferred from isochrones. 3.2.3 Clustering GC are not uniformly distributed in the local Universe, but clusters around their parent galaxies. This must be the case also at high redshift, while they are forming, though the parent galaxy may not be in place yet. Thus, one expects candidate GCPs to be highly clustered, lying within groups of 100kpc, or 20arcsec size (see Table 1 ). With sufficient statistics, it should be possible to directly establish the connection between the peak of GC formation and the stage of development of the host galaxy: clusters of GCPs as signposts of incipient massive galaxy formation. 4 DISCUSSION AND SUMMARY GCPs identified via BB photometry will need a spectroscopic confirmation, and the most obvious and timely follow up will be with JWST itself, either with the NIRCam grizm capability or with NIRSpec. In the longer run, 3040m class telescopes with their adaptive optics capabilities will offer unparalleled resolution. However, even before spectroscopic observations, some confirmation may come for those GCP candidates detected in more than one band, with SED fits providing age, mass and redshift constraints. Of course, more information could be gathered for those objects detected with NB filters, but they could be at most just a few. Designing a specific observing plan is far beyond the scope of this Letter. From the estimates given here it is relatively easy to infer what kind of information could be obtained for given invested observing time. On the other hand, deep multiband imaging over many fields, either primary or parallel, will certainly be done with JWST . This exercise has demonstrated that GCPs can be detected, with the number of detections critically depending on the mass of the progenitors, i.e. on the mass budget factor f MB . If GCPs were just marginally more massive than their present-day progeny then only the most massive ones could be detected, i.e. would have BB fluxes at formation above the thresholds given in Table 2 , scaled to the actual integration time. If this were the case, than the massive
Paper: 2019ApJ...885...34W (fulltext)
features may be seen within a “mossy” region that is oriented roughly along the axis of the supergranular network elements. In this case, the presence of small loops is hardly surprising, given that the underlying network is interspersed with minority-polarity flux elements. Figure 12. Moss-like region associated with supergranular magnetic network located just north of NOAA 12321, 2015 April 18. The circled area has radius 30′′ and is centered at (265′′, 391′′). Top panel: sharpened 17.1 nm image recorded at 14:29:47 UTC. Second panel: sharpened 19.3 nm image recorded at 14:29:42 UTC. Third panel: sharpened 21.1 nm image taken at 14:29:47 UTC. Bottom panel: HMI magnetogram recorded at 14:29:49 UTC and saturated at ±100 G. In this case, the looplike fine structure seen inside the moss is consistent with the presence of positive-polarity flux elements scattered in and around the negative-polarity network. Finally, Figure 13 shows an inverted-Y structure embedded in an area of moss inside NOAA 12335, as observed at 06:14 UTC on 2015 May 5. The base of the structure is rooted in strong negative-polarity plage, but the bright moss occupies a much wider area, extending through the weaker-field region to the south. The dark areas in the 17.1 nm image in the top panel represent fibrilar and filament material located near large- and small-scale polarity inversion lines. Figure 13. Moss region inside NOAA 12335, 2015 May 5. The circled area has radius 12′′ and is centered at (−376′′, −152′′). Top panel: 17.1 nm image taken at 06:14:35 UTC. Second panel: sharpened version of the same image. Third panel: sharpened 19.3 nm image taken at 06:14:30 UTC. Bottom panel: HMI magnetogram recorded at 06:14:51 and saturated at ±100 G. The inverted-Y structure is located in strong negative-polarity plage, but the underlying moss extends southward into areas of weaker photospheric field. These examples suggest that a strong underlying photospheric field is neither a necessary nor a sufficient condition for the presence of bright moss. 5. Summary and Discussion We have presented illustrative examples of inverted-Y structures embedded inside Fe ix 17.1 nm moss, where the underlying photospheric field was purely unipolar according to simultaneous HMI magnetograms. (Additional examples may be seen in Figures 2, 4, and 5 of Wang 2016 .) The triangular or dome-shaped bases of these structures had horizontal dimensions of ∼3′′–6′′ or ∼2–4 Mm. In the standard topological interpretation, the base contains minority-polarity flux, with a null point located at its apex. The lifetimes of the looplike footpoint features ranged from less than a minute to the order of an hour, with the smaller features generally being shorter-lived than the larger ones. As an alternative possibility that does not require the presence of minority-polarity flux, the triangular structures might represent canopies overlying individual granular cells, with thermal pressure dominating above the
Paper: 2015ApJ...809..100B (acknowledgments)
We thank Benjamin Bromley, Eric Charles, Scott Kenyon, Christopher Kochanek, Sebastien Muller, and Dan Schwartz for valuable comments and discussions. We thank the referee for providing comments that have led to a tighter discussion. A. B. is supported by the Department of Energy Office of Science, NASA the Smithsonian Astrophysical Observatory with financial support from the NCN grant DEC-2011/01/M/ST9/01891. M.J.G. is supported by the Smithsonian Institution.
Paper: 2015ApJ...802...18M (acknowledgments)
We thank the referee, Dennis Zaritsky, for a thoughtful review that led us to broaden the scope of this paper. We also thank Matt Bershady and Rob Swaters for conversations about the DiskMass project. This work is based in part on observations made with the Spitzer Space Telescope , which is operated by the Jet Propulsion Laboratory, California Institute of Technology under a contract with NASA. Support for this work was provided by NASA through an award issued by JPL/Caltech. Other aspects of this work were supported in part by NASA ADAP grant NNX11AF89G. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.
Paper: 2019ApJ...880...21M (fulltext)
signatures, it will also be important to better explore problems such as disk formation, the formation of disk outflows, and the termination of accretion by means of multidimensional simulations in the future. We note that the fallback accretion energy is more likely to be activated in hydrogen-rich progenitors from the perspective of stellar evolution. It is hard to achieve the conditions required to form magnetars in hydrogen-rich progenitors (e.g., Heger et al. 2005 ) because of efficient angular momentum transport from the core to the hydrogen-rich envelope, in particular if dynamo action is considered. For fallback-powered SNe, one needs to form an accretion disk near the compact remnant to have the accretion disk wind powering the SNe. Such an accretion disk can be formed even if the initial rotational period of the core is small because shells farther out in the envelope have higher specific angular momentum and can easily reach Keplerian velocities before being accreted onto the central compact object. Even if the angular momentum in the envelope is not sufficiently large, asymmetric explosions may impart sufficient angular momentum to the outer envelope by imparting some tangential velocity onto the fallback material, as found in a more energetic fallback supernova model (C. Chan et al. 2019, in preparation), which is enough to make an accretion disk as it falls back. Fallback-powered SNe may also be affected by the interaction between SN ejecta and a dense circumstellar medium (CSM) surrounding them. Although the 40 progenitor we studied does not experience much mass loss during its evolution, more massive progenitors are expected to undergo significant mass loss before the explosions due to the pulsational pair instability (Woosley 2017 ). After forming a dense CSM due to the pulsational pair instability, the progenitors are expected to collapse by forming BHs and may result in the fallback-powered SNe. Therefore, the existence of a dense CSM could be common in the fallback-powered SNe. The possibility of CSM interaction in fallback SNe needs to be studied further in the future. Difference in CSM configurations with the same central engine can result in a variety of LCs (e.g., Fischer et al. 2018 ). Finally, fallback-powered SNe may be common among the explosions of the first stars. It has been predicted that the first stars have a top-heavy initial mass function (e.g., Hirano et al. 2015 ), and many of them are likely to be massive enough to form a BH during their explosions. Pair-instability SNe from the first stars (Heger Woosley 2002 ) are generally considered as potential SNe that could become bright enough to be observed even from this early epoch of cosmic evolution in the future transient surveys, especially in near-infrared (e.g., Scannapieco et al. 2005 ). We have shown that the fallback-powered SN from the 40 first star can be as luminous as pair-instability SNe. The mass range of pair-instability SNe that can be as
Paper: 2020MNRAS.495L.144H (fulltext)
the IC file, which sets two haloes + star system simultaneously. The NGC 1052 system and the NGC 1052-DF2 system are always initially placed 400 kpc apart, with a fixed relative velocity of 250 km s –1 towards each other, which is consistent with the gravitational acceleration of free-falling from infinity and checked later to satisfy the velocity projection requirement. In order to control the NGC 1052 stellar distribution precisely to match the deprojection of the cored n = 3.4 Sérsic profile, we use a static baryonic potential kept at the centre of NGC 1052 instead of living particles. The NGC 1052 halo and NGC 1052-DF2 halo are simulated by 5.54 million and 0.06 million halo particles, respectively, and 0.1 million star particles simulate the luminous NGC 1052-DF2 profile. [Such initial configuration and simulation set-up is shared by the simulation for an alternative dark matter model; see the last paragraph. The physics for the pure CDM model is later tested by higher resolution simulation (Yang, Yu An 2020 ).] The gravitational softening length is 114 pc for halo particle and 14.4 pc for star particle. The bulk of the parameter scan is 2D, namely on the impact parameter b and the stellar Plummer radius r P (equivalent to its central density ρ c ). 3 RESULTS AND DISCUSSION In Fig. 3 , we plot the information extracted from the snapshots, which can be consistent with the observation after the encounter; here and later, d is the true 3D distance between the centres of the two galaxies. First, in the left-hand panel, we can see that with a scan of impact parameter at a spacing of 0.2 kpc, all simulations within this b range give small enclosed DM mass within the 7.6 kpc, which can be consistent with the observation, at the level of M / L 8. This is a consequence of the initial radial distribution, since before encounter the enclosed mass at the corresponding radius has a similar ratio. We can see that generally the smaller the impact parameter, the larger the stretch, and the smaller the enclosed mass in the 7.6 kpc. In the lower subpanel, we check the visual velocity relative to the central NGC 1052, and in a region, it can be consistent with the observed value. Figure 3. The combination of the two effects explains the observed lack of DM in NGC 1052-DF2. Left-hand panel: DM enclosed within the 7.6 kpc from several simulation snapshots after the encounter, in which we scan the impact parameter from 3.0 to 4.0 kpc. In the bottom subpanel, we also check the radial velocity relative to NGC 1052 and confirm that the observed 293 km s –1 can be reproduced within several snapshots. The two groups of curves are for NGC 1052-DF2 passing visually in front of or behind the central NGC 1052, respectively. They correspond to the two extremes in their final velocity angles with the line of sight, when considering the non-zero impact parameter and the bending of trajectory. Right-hand panel: The probability histogram by randomly selecting 10 star
Paper: 2018MNRAS.479.4355K (acknowledgments)
ACKNOWLEDGEMENTS KEKC and MPK acknowledge support from the UK’s Science and Technology Facilities Council (grant agreement No. ST/M001008/1). JEG is supported by the Royal Society. MJM acknowledges the support of the National Science Centre, Poland through the POLONEZ grant agreement No. 2015/19/P/ST9/04010; this project has received funding from the European Union’s Horizon 2020 research and innovation programme under the Marie Skłodowska-Curie grant agreement No. 665778. KK acknowledges support from the Swedish Research Council and the Knut and Alice Wallenberg Foundation. NB acknowledges support from the European Research Council Advanced Investigator Program, COSMICISM (ERC-2012-ADG-20120216, PI R.J.Ivison).
Paper: 2017A&A...608A..74O (acknowledgments)
We thank the anonymous referee for valuable comments. We also thank Bertram Bitsch, Hidekazu Tanaka and Takanori Sasaki for useful comments and suggestions. Numerical computations were conducted on the general-purpose PC farm at the Center for Computational Astrophysics, CfCA, of the National Astronomical Observatory of Japan. This work was supported by JSPS KAKENHI Grant Number 16H07415.