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ACKNOWLEDGEMENTS This work is dedicated in memory of K. N. V. S. Kishore Babu, a dear friend who inspired KBM to pursue Astrophysics. We are grateful to Peter Behroozi, Mark Brodwin, Philip Hopkins, Kartheik Iyer, Allison Kirkpatrick, Viraj Pandya, Vicente Rodriguez-Gomez, Brett Salmon, and Raymond Simons for helpful comments and discussions during different manifestations of this work. Special thanks to Daniel Shanaberger for carefully reading the manuscript, thanks to Gillen Brown, Cody Ciaschi, Bandon Decker, Benjamin Floyd, Ripon Saha, and Madalyn Weston for their valuable inputs during the University of Missouri Kansas City (UMKC) Galaxy Evolution Group discussions. DHM and KBM acknowledge support from the Missouri Consortium of National Aeronautics and Space Administration (NASA)’s National Space Grant College and Fellowship Program, and funding from the University of Missouri Research Board. RSS acknowledges support from the Downsbrough family, and from the Simons Foundation. SL acknowledges support from the National Research Foundation of Korea grant, no. 2017R1A3A3001362, funded by the Korea government. This work is based on observations taken by the CANDELS Multi-Cycle Treasury Program with the NASA/ESA HST , which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. Support for Program number HST-GO-12060 was provided by NASA through a grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555. This publication makes use of the SDSS. Funding for the creation and distribution of the SDSS Archive has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the US Department of Energy, the NASA, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. This publication also made use of NASA’s Astrophysics Data System Bibliographic Services, topcat (Tools for OPerations on Catalogues And Tables, Taylor 2005 ), the core python package for the astronomy community (Astropy 1.2.1; Astropy Collaboration et al. 2013 ).


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physical motion of a single massive star-forming region. 3.1. Two Morphological Classes of FIRE-CMZs The FIRE galaxies presented here cover a range of morphologies in their centers because these simulations were not designed to match the detailed structural properties of our Galactic center. Within the sample of seven galaxies, we see two distinct classes of central morphology in their fiery cores (gas and star formation distributions within R ≈ 1.5 kpc): (a) “Asymmetric/Bursty” (m12b, m12c, m12f, and m12w): Large, asymmetric gas clouds and star-forming complexes are seen. Star formation is concentrated in intense starbursts whose feedback dramatically shapes the local gas environment (see m12b, upper subfigure of Figure 1 ). Two simulations falling in this category (m12b and m12f) simultaneously match the MW CMZ gas and SFR measurements. The two others (m12c and m12w) do not simultaneously have SFR and dense gas tracers within the central 145 pc at any point in this time window. (b) “Smooth” (m12i, m12m, and m12r): Gas and star formation is smoothly distributed within the galactic centers, with clear feeding of gas into center, and a cirrus of star formation (see m12m, lower subfigure of Figure 1 ). Feedback events do not dramatically alter the local gas environment, as the feedback is relatively dispersed across their centers. Interestingly, none of the galaxies here exhibit the ring structures, presumed to be long-lived, seen by studies of the central regions of other spiral galaxies and the MW CMZ (Kormendy Kennicutt 2004 ; Molinari et al. 2011 ). We note the lack of strong bars in the centers of any of these FIRE galaxies at this time (m12m did develop a strong bar around z ≈ 0.2, but it does not survive to z = 0; Debattista et al. 2019 ); without the presence of bars in these simulations at z = 0, we cannot speak to the dynamical importance of bars in producing central galactic environments similar to the MW CMZ. Other work has highlighted the potential impacts of bars in funneling gas core-ward and forming rings (Sormani et al. 2015 , 2018 ). However, bar-induced effects would likely push these simulated galactic centers toward more asymmetric states, supporting the idea that bursty, rather than steady-state, star formation is necessary to explain MW CMZ observations. Specifically, Sormani et al. ( 2018 ) showed that the gas flow in a barred potential naturally becomes turbulent and asymmetric, even in the absence of any type of stellar feedback. And so, we leave it to future work to investigate the gas flows driven in FIRE galaxies by the bars that form at higher redshift. The structures in the centers of the FIRE galaxies appear to be fairly transient in nature, existing for ≲10 Myr (similar to the GMC lifetimes seen in these simulations more broadly by Benincasa et al. 2020 ). We should be clear: spiral structures do exist core-ward in these simulations (see the clear presence of spiral arms in both m12b and m12m in Figure 1 ). In


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The authors would like to thank the anonymous referee for useful comments and their colleagues of the SL2S and SLACS collaborations for many useful discussions as well as stimulating exchanges within the Euclid Strong Lensing Science Working Group. R.G. also acknowledges E. Bertin for valuable discussions about massive model fitting. This work was supported by the Centre National d’Études Spatiales (CNES) and the Programme National Cosmologie et Galaxies (PNCG). This research is based on observations obtained with MegaPrime/MegaCam, a joint project of CFHT and CEA/DAPNIA, a joint project of CFHT, Taiwan, Korea, Canada and France, at the Canada-France-Hawaii Telescope (CFHT) which is operated by the National Research Council (NRC) of Canada, the Institut National des Sciences de l’Univers of the Centre National de la Recherche Scientifique (CNRS) of France, and the University of Hawaii. This work is based in part on data products produced at TERAPIX and the Canadian Astronomy Data Centre.


O O O O B-Wavelength O O O B-Wavelength O O O O O O O O O O O O B-Dataset I-Dataset I-Dataset I-Dataset O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O O O B-Dataset I-Dataset I-Dataset O O O B-Wavelength O B-Wavelength O O O O O O O B-Instrument B-Wavelength I-Wavelength O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Wavelength O O O B-CelestialRegion I-CelestialRegion O O O O O O O O O O O B-Wavelength I-Wavelength O O O B-Survey O B-Survey I-Survey O O B-Survey O O O O O B-Model O O O O O O O O O O O O O O O B-Wavelength O B-Survey O O O O O B-Wavelength I-Wavelength O O B-Survey O O O O O O B-Survey I-Survey O O O O O O O O O O O O O O O O O O O O O O O B-Survey O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Survey O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Survey B-Wavelength O O O O O O B-Survey I-Survey O O O O O B-Survey I-Survey O B-Survey O O O O O O B-Survey O O O O O O O O O O O O O O O O O O O O O O B-Wavelength I-Wavelength O O O O O O O O O O O O O O O O O O O O B-Dataset I-Dataset I-Dataset O O O O O O O O O O O O O B-Dataset I-Dataset I-Dataset I-Dataset O O O O O O O B-Model O O O B-CelestialRegion O O O O B-Wavelength I-Wavelength O O O O O O O O O O B-Survey O B-Survey I-Survey O O B-Survey O O O O O O O O O O B-Wavelength I-Wavelength O O O O O O O O O O O O B-Survey O O B-Survey I-Survey O O O O O O O O O O O O O O O O O O O O O O O O O B-Dataset I-Dataset I-Dataset O O O O O O B-CelestialRegion B-CelestialRegion I-CelestialRegion B-CelestialRegion B-CelestialRegion O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O

either the use _ f160w or use _ f850lp flags, as noted in each subsequent case. We further compare to the CANDELS GOODS-S photometric catalog released by Guo et al. ( 2013 ), adopting flags equal to zero for non-contaminated sources. The CANDELS GOODS-S catalog presents the multiwavelength (UV to mid-IR) photometry, with source detection performed in the WFC3 H F160W mosaic using a “hot” and “cold” detection methodology. We first present the number counts in Section 3.1 and then cross-match all catalogs within a 05 radius to compare aperture photometry in Section 3.2 . Finally, in Section 3.3 , we show several example SEDs to showcase the high quality of the photometry. 3.1. Number Counts The number densities of galaxies that satisfy the use _ f160w criterion in the GOODS-S field are shown in Figure 13 as a function of the total H F160W magnitude for both HLF (black), 3D- HST (red), and CANDELS (blue). The error bars are Poisson. Though completely independent data reductions, the three data sets are fairly similar in terms of F160W coverage; HLF covers 207 arcmin 2 in H F160W , whereas CANDELS covers 173 arcmin 2 , and 3D- HST covers 171 arcmin 2 . It is therefore not surprising that the number counts are consistent. The deficit of sources with ABmag in CANDELS relative to the two other fields is likely the result of using a deeper multi-band combined detection image. The further excess of objects at the faint end in HLF results from a combination of effects. In part, this population of faint sources will arise due to the more aggressive source detection settings adopted. But in some cases, it is clear that the HLF F160W imaging is deeper than the earlier 3D- HST version (i.e., explaining why both 3D- HST and HLF have more faint number counts), but HLF appears to further reveal an exciting new population of extremely faint sources. Figure 14 shows 15 × 15 postage stamps (, , H F160W ) of 48 ultra-faint sources with magnitudes between 28 and 29 ABmag rank-ordered by S/N that are identified in the HLF photometric catalog but do not have a match within a radius of 04 in the 3D- HST photometric catalog. Figure 13. Number counts of galaxies with Poisson errors in the GOODS-S field as a function H F160W total magnitude, with no correction for incompleteness. The agreement between HLF (black), 3D- HST (red) and CANDELS (blue) is excellent. Figure 14. Example postage stamps (, , H F160W ) of 48 ultra-faint sources between 28 and 29 ABmag identified in HLF but not 3D- HST . The objects are rank-ordered by S/N, with the outline of the segmentation map shown as a white contour for reference. The depths in the HLF GOODS-S mosaic vary significantly within the field (i.e., HUDF, CANDELS deep, wide, ERS, etc.). Such a heterogeneous weight map implies that the single number count histogram shown in Figure 13 is simply the superposition of the histograms at different depths. In order to better understand the improvement, we separate the weight map into four


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we only observe in the vertical direction, meteor ablation height generally correlates very well with the vector velocity of a meteor. The average ablation height for the three groups of meteors is 105, 103, and 99 km, respectively. Meteors observed on the back-side of the Earth are largely prograde meteors with a heliocentric velocity larger than 30 km s −1 . Their detection rate is very low because of small atmospheric entry velocity. 6 DISCUSSION AND CONCLUSION The most important aspect of the current report is the conclusion that the majority of the meteors observed by the Arecibo radar are in nearly circular orbits. This points to the importance of Poynting–Robertson and solar wind drags. The nature of the two drags is similar, with the latter being about 30 per cent of the former (Gustafson 1994 ). In order for the drag forces to be effective, meteoroid mass is typically smaller than 10 −11 kg. On the other hand, when β, the ratio of radiation pressure to solar gravitation force, is larger than 0.5, a particle will be blown out of the solar system. β = 0.5 corresponds to a particle size of 1 μm (Dermott et al. 2001 ) and a mass of 10 −14 kg, which is the lower boundary of the meteors observed at Arecibo as discussed above. Thus, Arecibo meteors appear to have the ‘Goldilocks’ size, small enough for the drags to force them into nearly circular orbits and being pulled inward to the Sun but large enough not to be blown away by the radiation pressure. The simulations by Borin et al ( 2017 ) show that it takes several million years to reduce the semimajor axis from 3.4 to 1 au for a 10 μm in radius particle. The evolution time for the smaller Arecibo particles is expected to be much shorter. As the Earth lies relatively close to the Sun, the majority of the meteoroids have their initial orbits outside that of the Earth whether they are of asteroidal or cometary origin. For meteoroids whose size is right for drag forces to be important, their orbits will evolve into nearly circular ones at 1 au if not rejected by gravitational scattering by outer planets. Meteoroids with a diameter of 1–9 μm originating from as far as the Kuiper belt have a 20 per cent chance of reaching 1 au and once a meteoroid's orbit reaches within that of Jupiter, it will largely spiral onto the Sun (Liou, Zook Dermott 1996 ). As all meteoroids in the mass range of 10 −14 –10 −11 kg generated outside Earth's orbit have a good chance to be on nearly circular orbits when they reach the inner solar system, their population is thus expected to be uncharacteristically large at 1 au. Further, meteoroids in the ecliptic plane have a higher probability ejected by outer planets than those on inclined planes. Although the distribution shown in Fig. 4 is a convolution of instrument aspect sensitivity and the true distribution, it indicates at least that a large number of meteoroids exist in inclined orbits. The ‘inclined’ population discussed by Divine ( 1993 ) has low


O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation O O O B-Wavelength O O O O O O O O O O O B-Citation I-Citation O B-Citation I-Citation I-Citation O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation O O O O O O O O O B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O O O O O O O O O O O O O O B-Citation I-Citation I-Citation O B-Citation I-Citation O B-Citation O B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O B-Citation I-Citation I-Citation O B-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O O B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O O B-Wavelength O O O O O O O B-Citation I-Citation I-Citation O B-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation O B-Wavelength O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Observatory O B-Observatory O O O O B-Wavelength O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation O O O O B-CelestialObject I-CelestialObject B-Citation I-Citation I-Citation I-Citation O O O O O O O O O B-Wavelength I-Wavelength O O B-Wavelength I-Wavelength O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Wavelength O O B-Wavelength I-Wavelength O O O B-Wavelength I-Wavelength I-Wavelength B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Wavelength O O O O O O O O B-Wavelength O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation O B-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O

respectively (Kouveliotou et al. 1993 ; Mészáros 2006 ). While long GRBs are most likely associated with the collapse of massive stellar objects (Woosley 1993 ; MacFadyen Woosley 1999 ), the short ones are usually attributed to the merger of two neutron stars (Eichler et al. 1989 ; Paczynski 1990 ), or one neutron star with one black hole (Lattimer Schramm 1976 ; Lee Ramirez-Ruiz 2007 ). They are also potential sources responsible for the emission of ultrahigh-energy cosmic rays (Milgrom Usov 1995 ; Waxman 1995a , 1995b ; Dermer Humi 2001 ; Chakraborti et al. 2011 ), high-energy neutrinos (Waxman Bahcall 1997 ; Waxman 2000 ; Abbasi et al. 2012 ; Murase et al. 2013 ), and gravitational waves (Kobayashi Mészáros 2003 ; Corsi Mészáros 2009 ; Chen Holz 2013 ; Lasky Glampedakis 2016 ; Abbott et al. 2017 ). The study of GRBs has been prolific over the past decades (Zhang Meszaros 2004 ; Nakar 2007 ; Kouveliotou et al. 2012 ; D’Avanzo 2015 ). GRBs have opened a unique window into connecting the fields of stellar, interstellar, and intergalactic astronomy as well as cosmology. With the new generation of space-borne experiments, such as the Swift and Fermi observatories, high-energy observations of GRBs have entered a new era in the study of extragalactic background light (EBL; Ackermann et al. 2012 ; Desai et al. 2017 ) and the search for Lorentz invariance violation (LIV; Ellis et al. 2019 ; Lang et al. 2019 ; Wei 2019 ). As one example, GRB 130427A (Ackermann et al. 2014 ) was able to produce photons of up to 94 GeV in observations (126 GeV at the source after redshift correction, which accounts for the attenuation due to EBL). It is unknown up to what energy the spectrum extends, as present-day observations are limited by effective area in the case of space-based instruments. Ground-based detectors may prove to be an important compensator for GRB observation when γ -ray energy is above tens of GeV (Grupen Buvat 2011 ; Abeysekara et al. 2012 ; de Angelis 2016 ). Two types of experimental techniques are available for this purpose: imaging air Cerenkov telescopes (IACTs) and extended air shower (EAS) array. IACTs have better observational sensitivity to stable point sources due to their higher angular resolution and better background rejection power (Poggiani 2017 ). However, IACTs have a limited field of view, which makes them useful for follow-up observations. Moreover, the requirement of a safe slew time makes follow-up of real-time alerts difficult. Even when slew time is not an issue, a small field of view requires accurate pointing in the GRB direction. However, the error boxes provided by wide-field-of-view GRB monitors are usually large at alert time. Another disadvantage is that IACTs operate in good weather on moonless nights, which makes the duty cycle of IACTs only 10%–15% (Hinton 2009 ; Bigongiari 2016 ). A small field of view and low duty cycle make IACTs an inconvenient and inefficient instrument for GRB observation. While EAS arrays


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Acknowledgments This research was partially supported by an iCore grant and an ISF grant. We thank Scott Tremaine, Kenta Hotokezaka and Sivan Ginzburg for useful discussions.


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ACKNOWLEDGEMENTS We thank Andrei Beloborodov, Yuri Lyubarsky, Illya Plotnikov, and Indrek Vurm for helpful information and conversations. We thank the reviewer, Jonathan Katz, for helpful comments. BDM and LS acknowledge funding from the National Science Foundation (grant number 80NSSCK1104). BM is supported by NASA through the NASA Hubble Fellowship grant #HST-HF2-51412.001-A awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., for NASA, under contract NAS5-26555.


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Support for the German contribution to GBM was provided by the Bundesministerium für Bildung und Forschung (BMBF) via the Deutsches Zentrum für Luft und Raumfahrt (DLR) under grant No. 50 QV 0301. The USRA coauthors gratefully acknowledge NASA funding through contract NNM13AA43C. The UAH coauthors gratefully acknowledge NASA funding from cooperative agreement NNM11AA01A. E.B. and C.M. are supported by an appointment to the NASA Postdoctoral Program, administered by the Universities Space Research Association under contract with NASA. D.K., C.A.W.H., and C.M.H. gratefully acknowledge NASA funding through the Fermi -GBM project.


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OSIRIS was built by a consortium of the Max-Planck-Institut für Sonnensystemforschung, Göttingen, Germany, CISAS–University of Padova, Italy, the Laboratoire d’Astrophysique de Marseille, France, the Instituto de Astrofísica de Andalucia, CSIC, Granada, Spain, the Research and Scientific Support Department of the European Space Agency, Noordwijk, The Netherlands, the Instituto Nacional de Técnica Aeroespacial, Madrid, Spain, the Universidad Politéchnica de Madrid, Spain, the Department of Physics and Astronomy of Uppsala University, Sweden, and the Institut für Datentechnik und Kommunikationsnetze der Technischen Universität Braunschweig, Germany. The support of the national funding agencies of Germany (DLR), France (CNES), Italy (ASI), Spain (MEC), Sweden (SNSB), and the ESA Technical Directorate is gratefully acknowledged. We thank the referee, J.-Y. Li, for his comments and suggestions that helped to improve this manuscript.


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This work is supported in China by NSFC (No. 11375052, No. 10975046, No. 10120130794, No. 11165013), the Chinese Ministry of Science and Technology, the Chinese Academy of Sciences, and the Key Laboratory of Particle Astrophysics, CAS, and in Italy by the Istituto Nazionale di Fisica Nucleare (INFN). We also acknowledge essential support from W.Y. Chen, G. Yang, X. F. Yuan, C. Y. Zhao, R. Assiro, B. Biondo, S. Bricola, F. Budano, A. Corvaglia, B. D’Aquino, R. Esposito, A. Innocente, A. Mangano, E. Pastori, C. Pinto, E. Reali, F. Taurino, and A. Zerbini in the installation, debugging, and maintenance of the detector.


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We thank the referee for his or her comments that improved the manuscript, and the Kepler /K2 science teams for providing such excellent data. Funding for the Kepler /K2 mission is provided by NASA’s Science Mission Directorate. The K2 data presented in this paper were obtained from the Mikulski Archive for Space Telescopes (MAST). Support for MAST for non-HST data is provided by the NASA Office of Space Science via grant NNX09AF08G and by other grants and contracts. The research leading to these results has received funding from the European Research Council (ERC) under the European Union’s Horizon 2020 research and innovation programme (grant agreement N o 670519: MAMSIE). This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France; the SAO/NASA Astrophysics Data System; and the VizieR catalogue access tool, CDS, Strasbourg, France.


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We thank the anonymous referee for a careful review that improved this paper. This work was supported by the National Aeronautics and Space Administration (NASA) through Chandra Award Numbers (GO3-14120 A, GO3-14120B) issued by the Chandra X-ray Observatory Center (CXC), which is operated by the Smithsonian Astrophysical Observatory (SAO) for and on behalf of NASA under contract NAS8-03060. Support for program number 13116 was provided by NASA through a grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. The MOJAVE program is supported under NASA-Fermi grant NNX12A087G. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. Facilities: VLA, HST (WFC3), CXO (ACIS).


O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O B-Model O B-Model O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObject B-CelestialObject I-CelestialObject O O O O O O O O B-CelestialRegion I-CelestialRegion I-CelestialRegion O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject B-CelestialObject I-CelestialObject O O O O O O O B-CelestialObject I-CelestialObject I-CelestialObject O O O O O O O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject I-CelestialObject I-CelestialObject O O O O O O O O O O O B-CelestialRegion I-CelestialRegion I-CelestialRegion I-CelestialRegion O O O O O O O O O O O O B-CelestialObject O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Model B-Citation I-Citation I-Citation O O O O B-Survey O O O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject B-CelestialObject I-CelestialObject O O O O O B-CelestialObject I-CelestialObject I-CelestialObject O O O O O O B-CelestialObject I-CelestialObject O O O O O B-CelestialObject O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject O O O O B-CelestialRegion I-CelestialRegion I-CelestialRegion O O O O O O O O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O B-Survey I-Survey I-Survey I-Survey I-Survey I-Survey O B-CelestialRegion O O O O O O O O O O O O

quality factors. A selection of these factors from the original sources is summarized below. 4.1. Selected Galaxy Information When measuring rotation curve data, there are many systematic issues that affect the accuracy of the information. In Oh et al. ( 2015 ), the authors gave a comprehensive discussion of each of the galaxies in the entire sample. Although both CG and MOND capture the essence of the rotation curves across the sample well, there are occurrences where it is useful to have the relevant information about the galaxy to understand why a fit may not be as precise. We have compiled some of the relevant information below for convenience. CVN (UGCA 292) is a gas rich dwarf galaxy in the Canes Venatici constellation. This galaxy has been shown to have a high gas mass fraction and low metallicity suggesting that it is highly unevolved with an extremely LSB (van Zee 2000 ). From a fitting point of view, this galaxy is unique with its extremely low rotational measured velocities. The models shown, however, provide a sufficient fit despite this property. DDO 133 (UGC 7698) is a large, faint galaxy in the Canes Venatici cloud whose extremely LSB shows high stellar evolution (Zhang et al. 2012 ). Due to its lack of tidal motions, this galaxy is assumed to be an isolated system, which is typically beneficial for obtaining a rotation curve. However, this rotation curve was studied by one of us in the past, and can be shown to have a lower than quoted inclination (O’Brien 2012 ). Hence, we use the inclination that was previously studied and published to alleviate any discrepancies. DDO 154 (UGC 8024) is an asymmetrical dwarf galaxy and is a member of the Coma Benenices star constellation. It has been proposed that it is a physical companion of the M64 group. The galaxy is very dim, presenting a low star formation rate and contains a very large amount of atomic hydrogen (Drexler 2006 ). The galaxy is also hypothesized to have a large dark matter to baryon ratio and hence is extremely important to this work. This galaxy has been studied in CG (Mannheim O’Brien 2012 ) after the initial THINGS survey was completed by Walter et al. ( 2008 ). It can be seen in this most recent survey that the alternative gravity models can circumvent the need for the quoted high dark matter to baryon ratio. DDO 168 (UGC 8320) is a galaxy in the Canes Venatici group and is a probable pair with UGC 8308 and sometimes included in the M51 group (Karachentsev et al. 2003 ). The galaxy has a large number of dispersed H ii regions, resulting in an unusual shape to the rotation curve. The essence of the rotation curve is captured in the fits, but due to the overall structure, one cannot expect a high-resolution fit. DDO 210 is located in the Aquarius star constellation. It was one of the original galaxies discovered by van den Bergh ( 1959 ) of the David Dunlap Observatory catalog (DDO) survey. This aquarius dwarf is extremely faint, and has a very LSB, though it contains


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We thank University of Memphis students, Peter Matheny and Adam Chastain, for help with data analysis. This work was inspired and motivated by discussions at the series of Coronal Loops Workshops. The Atmospheric Imaging Assembly on the Solar Dynamics Observatory is part of NASA’s Living with a Star program. CHIANTI is a collaborative project involving the NRL (USA), the Universities of Florence (Italy) and Cambridge (UK), and George Mason University (USA). Solar physics research at the University of Memphis is supported by a Hinode subcontract from NASA/SAO.


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Resources supporting this work were provided by the NASA High-End Computing (HEC) Program through the NASA Advanced Supercomputing (NAS) Division at Ames Research Center. Results were partially obtained using the analysis software Pynbody ( https://github.com/pynbody/pynbody ). We thank the Fisk-Vanderbilt Masters-to-PhD Bridge program for the funding and support of this research. We also wish to thank the anonymous referee for their thorough reading of the manucript and their useful comments which greatly improved its quality. J.B. acknowledges generous support from the Helen Gurley Brown Trust. M.V. acknowledges support from NASA award ATP NNX10AC84G.


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in Section 2.3.2 . 3.3. Gravitationally Unstable Disk Figure 3 (right) shows the simulated observations for a gravitationally unstable disk with Σ ∝ R −0.5 and T ∝ R −0.75 . The disk mass is 0.24 M ⊙ , leading to a disk-to-star-mass ratio of 0.49. Even with this relatively massive disk, the combined system mass lies within the limits discussed in Section 2.3.1 . The simulated unsharp masking observation reproduces the large-scale spiral arms, the dark crescents, and the bright inner ellipse, while the deprojected image shows two large-scale spirals analogous to Figure 1 (right). 4. Discussion The combination of our hydrodynamic modeling and simulated observations allows us to put strong constraints on the disk structure from which the submillimeter continuum emission originates. Throughout all of our models, we find that a steeply declining surface mass density beyond ≈300 au is key to producing a close match to the masked image. Figure 4 shows the scaled radial intensity profile of our successful self-gravitating disk and the radial intensity profile of the successful companion simulations compared to the observational data from Pérez et al. ( 2016 ) and the model of Andrews et al. ( 2009 ). The steep decline beyond ≈300 au in our simulation manifests itself in the form of a steep emission profile at large radii that turns over at approximately the right radius, matching the observed profile well. The modest scaling factor for the self-gravitating disk (∼1.5) is well within uncertainties associated with dust-to-gas ratios and/or grain opacities for circumstellar disks. Figure 4. Radial emission profiles for the successful gravitationally unstable disk simulation (red), which has been decreased by a factor of ∼1.5 (see the text) and the successful companion disk simulation (blue). Also shown are data from Pérez et al. ( 2016 ; black points) and the best-fit model from Andrews et al. ( 2009 ; black dashes). The steep decline beyond ≈300 au is key to matching the observed morphology in the simulated images. We note that the resulting Toomre profile dips marginally below 1 for a limited radial range in our successful self-gravitating disk. However, this disk shows no sign of fragmentation even though we evolve it for many dynamical times for this radial range. On the other hand, a companion beyond the spirals (≳300 au) is unlikely to have formed by core accretion. This suggests that, should the companion hypothesis be correct, this would hint that the Elias 2–27 disk may have formed a fragment by gravitational instability in the past. We also stress the importance of applying an unsharp mask to our simulated observations. While such a mask is primarily applied to increase the contrast of the image, subtle changes are introduced to the resulting image that allow us to remove models from consideration. Changes in the surface density profile affect the radial locations at which there is an excess or deficit with respect to a Gaussian mask. This


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We are grateful to the anonymous reviewer for constructive comments that helped to improve the manuscript. This work has extensively used NASA Astrophysics Data System, and the SIMBAD database, operated at CDS, Strasbourg, France. This publication made use of data products provided by the NASA/IPAC Infrared Science Archive. We are grateful to Dr. J Maíz Apellániz for kindly providing us with the optical spectrum of Cyg OB2 12 and to Dr. M Mapelli for a useful discussion on the dynamics of multiple systems. L.M.O. acknowledges support by DLR grant 50 OR 1302. A.S. is supported by the Deutsche Forschungsgemeinschaft (DFG) under grant HA 1455/26. D.P.H. was supported by NASA grant GO5-16009A. Facility: CXO (HETG/ACIS) .


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We are thankful for fruitful discussions and comments provided by Simone Ferraro and Emmanuel Schaan. The Hyper Suprime-Cam (HSC) Collaboration includes the astronomical communities of Japan and Taiwan and Princeton University. The HSC instrumentation and software were developed by the National Astronomical Observatory of Japan (NAOJ), the Kavli Institute for the Physics and Mathematics of the Universe (Kavli IPMU), the University of Tokyo, the High Energy Accelerator Research Organization (KEK), the Academia Sinica Institute for Astronomy and Astrophysics in Taiwan (ASIAA), and Princeton University. Funding was contributed by the FIRST program from Japanese Cabinet Office, the Ministry of Education, Culture, Sports, Science and Technology (MEXT), the Japan Society for the Promotion of Science (JSPS), Japan Science and Technology Agency (JST), the Toray Science Foundation, NAOJ, Kavli IPMU, KEK, ASIAA, and Princeton University. The Pan-STARRS1 Surveys (PS1), which are used for the photometry and astrometry calibration, have been made possible through contributions of the Institute for Astronomy, the University of Hawaii, the Pan-STARRS Project Office, the Max-Planck Society and its participating institutes, the Max Planck Institute for Astronomy, Heidelberg and the Max Planck Institute for Extraterrestrial Physics, Garching, Johns Hopkins University, Durham University, the University of Edinburgh, Queen’s University Belfast, the Harvard-Smithsonian Center for Astrophysics, the Las Cumbres Observatory Global Telescope Network Incorporated, the National Central University of Taiwan, the Space Telescope Science Institute, the National Aeronautics and Space Administration under grant No. NNX08AR22G issued through the Planetary Science Division of the NASA Science Mission Directorate, the National Science Foundation under grant No. AST-1238877, the University of Maryland, and Eotvos Lorand University (ELTE). This paper makes use of software developed for the Large Synoptic Survey Telescope. We thank the LSST Project for making their code available as free software at http://dm.lsst.org . This work is based in part on data collected at the Subaru Telescope and retrieved from the HSC data archive system, which is operated by the Subaru Telescope and Astronomy Data Center at the National Astronomical Observatory of Japan. The P olarbear project is funded by the National Science Foundation under grant Nos. AST-0618398 and AST-1212230. The James Ax Observatory operates in the Parque Astronómico Atacama in Northern Chile under the auspices of the Comisión Nacional de Investigación Científica y Tecnológica de Chile (CONICYT). The James Ax Observatory would not be possible without the support of CONICYT in Chile. All silicon-wafer-based technology for P olarbear was fabricated at the UCB Nanolab. This research used resources of the Central Computing System, owned and operated by the Computing Research Center at KEK, the HPCI system (Project ID:hp150132), and the National Energy Research Scientific Computing Center (NERSC), a DOE Office of Science User Facility supported by the Office of Science of the U.S. Department of Energy under contract No. DE-AC02-05CH11231. This work was supported in part by the World Premier International Research Center Initiative (WPI Initiative), MEXT, Japan, and MEXT Grant-in-Aid for Scientific Research on Innovative Areas (JP15H05887, JP15H05891, JP15H05892, JP15H05893). In Japan, this work was supported by MEXT KAKENHI grant Nos. 21111002 and JP18H05539, JSPS KAKENHI grant Nos. JP26220709, JP24111715, JP26800125, and the JSPS Core-to-Core Program. In Italy, this work was supported by the RADIOFOREGROUNDS grant of the European Union’s Horizon 2020 research and innovation program (COMPET-05-2015, grant agreement No. 687312), as well as by the INDARK INFN Initiative and the COSMOS network of the Italian Space Agency (cosmosnet.it). Support from the Ax Center for Experimental Cosmology is gratefully acknowledged. The KEK authors acknowledge support from KEK Cryogenics Science Center. The APC group acknowledges travel support from Labex UNIVEARTHS. The Melbourne group acknowledges support from the University of Melbourne and an Australian Research Council's Future Fellowship (FT150100074). Y.C. acknowledges the support from JSPS KAKENHI grant Nos. 18K13558 and 18H04347. T.N. acknowledges the support from the Ministry of Science and Technology (MOST), Taiwan, R.O.C., through the MOST research project grants (No. 107-2112-M-002-002-MY3). H.M. acknowledges the support from JSPS KAKENHI grant No. JP18H04350. M.O. acknowledges the support from JSPS KAKENHI grant No. JP18K03693. R.T. acknowledges the support from Grant-in-Aid for Scientific Research from the JSPS Promotion of Science (17H01131). Numerical computations were in part carried out on Cray XC50 at the Center for Computational Astrophysics, National Astronomical Observatory of Japan. A.K. acknowledges the support by the JSPS Leading Initiative for Excellent Young Researchers (LEADER) and by JSPS KAKENHI grant No. JP16K21744. M.A. acknowledges support from CONICYT UC Berkeley-Chile Seed grant (CLAS fund) No. 77047, Fondecyt project 1130777 and 1171811, the DFI postgraduate scholarship program, and the DFI Postgraduate Competitive Fund for Support in the Attendance to Scientific Events. G.F. acknowledges support from the European Research Council under the European Union’s Seventh Framework Programme (FP/2007-2013)/ERC grant agreement No. 616170. F.M. acknowledges the support by the JSPS Fellowship (grant No. JP17F17025). J.S.S. is supported by the NSF Graduate Research Fellowship.


O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Formula I-Formula I-Formula I-Formula O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialRegion I-CelestialRegion O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O B-Formula I-Formula I-Formula I-Formula I-Formula O O O O O O O O O O O O O O O O O O O O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O

factor. This equation has the solution , where ‘en’ corresponds to the entry time of a perturbation mode into the horizon. Then, at the time at which a dark matter halo becomes non-linear (i.e. δ D ∼ 1), the density contrast of baryonic matter also approaches one. In the non-linear phase of structure formation, assuming a top-hat spherical model, each shell collapses without crossing the other shells and at the turnaround time structures start collapsing and finally virialize (Padmanabhan 1993 ). If we ignore the dissipation of baryonic matter at the first stage of virialization, then baryonic matter particles will have almost the same amount of kinetic and potential energy as dark matter particles and will virialize at the same time as dark matter fluid. Then, in the gravitational potential of dark matter, we expect that baryonic matter extends to the same scale as the dark matter structure. We can now calculate the density of baryonic matter in the dark matter halo from the following relation: where h ν 0 is the energy of CMB photons, m p is the mass of a proton and the ‘zero’ sub/superscript corresponds to the value of parameters at the present time. The important issue concerning the amount of baryonic matter within the gravitational potential of dark matter is the possibility of formation of baryonic structure. Formation of baryonic structure after the virialization of a halo is possible if plasma can cool and fall to the centre of the gravitational potential of the dark matter. The essential criterion is comparing the dynamical time-scale with the cooling time-scale of the baryonic matter. The dynamical time-scale for a test particle in the potential of dark matter is given by and, on the other hand, the cooling time-scale is where σ T is the Thomson cross-section and α is the fine-structure constant. For t cool t grav , we have a sufficient condition for the formation of baryonic structures. Substituting temperature from the virial condition, we get the simple result where baryonic and dark matter are enclosed within the same radius R . Integrating from equation ( 32 ) and equating it with equation ( 35 ), we obtain the following relation between the size of structure and the mass of baryonic matter in terms of η B : where η 10,B is the definition of η B normalized to 10 −10 . For a Galactic halo with mass of order of M D ≃ 10 12 M ⊙ , for η 10,B > 10 −5 , a structure with the mass of a globular cluster can form within this dark halo. Another lower value for the baryonic asymmetry could be η 10,B ≃ 10 −12 , where only one star with M min ∼ 0.1 M ⊙ condenses to form a single star within a 10 12 solar mass halo. We can think about another extreme case, having only one star within the cosmological horizon. In this case, we can write the η B parameter in terms of the minimum mass for a star and the mass of matter within the horizon (i.e. M H ), where m X and m p are the mass of a dark matter particle and the proton mass. Substituting the


O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Model O O O O O O O O O O O O O O O O O B-Model O O O O O O O O O O O O O O O O O B-Model O O O O O O O O O O O O O O O O O O O O O B-Model O O O O O O O O O O O O O O O O O O O O O O O O B-Model O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Model O O O O O O O B-Model I-Model O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Model O O O O O O O O B-Model O O O O O O O O O B-Model O O O O O O O O O B-Model O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Model O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Model I-Model O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Model O O O O O O O O B-Model I-Model O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Model O O O O O O

but it also reorganises the field topology, reducing the pinching of field lines close to the protostellar object, which also hinders the braking mechanism. Magnetic flux freezing and magnetic braking play a central role in the formation and development of the structures during the collapse. While in iMHD strongly amplified fields launch powerful magnetically supported outflows, these latter are much weaker or even disappear in niMHD. Misalignment between the initial rotation axis and the magnetic field direction does not appear to affect the niMHD results, showing that the physical flux dissipation due to ambipolar diffusion dominates the effects of initial configuration or numerical diffusion. For iMHD models, the additional mixing due to a tilted configuration also produces a diffusion plateau, similar to the ambipolar calculations. The disks that form in iMHD, characterised by strong magnetic support and an inflated shape (due to toroidal field lines loaded by infalling matter) strongly differ from the flat disks that form in niMHD. Formation and long-term evolution of disks in ideal MHD, however, are of dubious validity. Furthermore, the excessive magnetic braking generates unphysical counter-rotation inside the outflow or the magnetic tower. Interchange instabilities develop at the interface between the core and the disk, producing a redistribution of the flux and displacement of the core and disrupting the top-bottom symmetry in the system. Mesh resolution is found to strongly affect the simulation results in iMHD. None of these effects is observed in the niMHD simulations if the resolution is high enough. Disks with Keplerian velocity profiles around the protostar form in all our niMHD simulations for all different magnetisations and inclination angles. Their size and mass, however, is significantly reduced (by a factor ) in the more magnetised case () because of the increased braking. Such a magnetisation value seems to be typical of most molecular clouds. Aligned and misaligned initial configurations have no consequence on the disk properties in niMHD and yield disks with very similar properties. Magnetic flux diffusivity due to ambipolar diffusion thus appears to play a dominant role during the first protostellar collapse and the formation of the first Larson core and its surrounding disk, and it appears to dominate processes such as magnetic field and rotation axis orientations. Flux diffusion during the collapse allows the formation of quasi-Keplerian disks, solving the “disk formation crisis” found in ideal MHD. The mass, size, and magnetic properties (plasma ) of these disks, however, strongly depend on the initial magnetisation ( parameter) of the cloud, and in all cases the disks are significantly smaller and less massive than those found in pure hydrodynamics calculations. Characterisation of this plasma is of major importance to characterise viscous transport in disks, notably in the MRI. To complete our study, we will


O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Survey O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation O O O O O O O O O O O O O O B-CelestialObject O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Instrument B-Wavelength I-Wavelength O O O O O O B-Wavelength I-Wavelength O O O O O O O O O O O O O O O B-Wavelength I-Wavelength B-Wavelength I-Wavelength O O O O O O O O O O O B-Instrument O O O O O O O O O O O O O O B-TextGarbage O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O B-CelestialObject O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O B-Telescope B-Instrument O O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O B-Telescope I-Telescope I-Telescope I-Telescope B-Telescope O B-EntityOfFutureInterest I-EntityOfFutureInterest I-EntityOfFutureInterest I-EntityOfFutureInterest O O O O B-Instrument B-Wavelength O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation O O O O O O O O O O O O O B-Survey O O O O O O O B-Wavelength I-Wavelength O O O O O O B-Wavelength I-Wavelength O O O O O O B-Wavelength I-Wavelength O O B-Wavelength O O O O O O B-Citation I-Citation I-Citation O O O O O O O O O O O O O O

For Ion2 , the observed ratio is 15 making Ion2 an outlier in the oxygen abundance vs. the [O iii ]/[O ii ] ratio relation and the EW(H) (i.e., sSFR) vs. the [O iii ]/[O ii ] ratio (Fig. 2, Stasińska et al. 2015 ); the Ion2 ratio is higher by 0.6 dex for the observed EW(H) (Table 1 ) and the ratio is higher by 1 dex than galaxies in the SDSS sample with similar metallicities. The Ion2 [O iii ]/[O ii ] ratio is also higher than expected for the derived stellar mass, SFR, and sSFR ( Nakajima Ouchi 2014 ). Therefore, we conclude that the extreme [O iii ]/[O ii ] ratio is due to unusual physical conditions (density-bounded nebula), which imply a low column density of neutral gas, and so favor leakage of ionizing photons ( Nakajima Ouchi 2014 ). 3.3. Physical properties derived from the photometry While the emission lines highlight Ion2 ISM physical properties, we also use the photometry to confirm findings such as the large [O iii ] emission and derive other physical parameters. From Fig. 8 , the jump in the -band is evident, while the continuum level in the adjacent bands, , and IRAC Channel 1 is 24.15. The magnitude in the HAWKI -band corresponds to an additional flux of erg/s/cm, which corresponds to an observed equivalent width of 6700 Å (1600 Å rest-frame). The total flux derived for H+[O iii ] from the MOSFIRE spectrum (see Table 1 ) is fully compatible with the flux estimated from the -band excess (after correcting for slit losses for the three lines, see above). The derived + [OIII] equivalent width is one of the largest ever observed at . The measured line ratios among the H+ [O iii ] suggest the [O iii ] rest-frame equivalent width is 1100 Å, with a similar value to the few extreme cases reported in the literature at ( Atek et al. 2011 ; Amorín et al. 2014b ; Maseda et al. 2014 ). The [O iii ]/H ratio is also higher than the value reported in Holden et al. (2014) for LBGs at , while the ratio is higher in this sample than in the local Universe ( Brinchmann et al. 2004 ). At large equivalent widths of H+[O iii ] have also been reported, measured through photometric excess detected in the Spitzer /IRAC channels (e.g., Finkelstein et al. 2013 ; Smit et al. 2014 ; Oesch et al. 2015 ). While these cases can be spectroscopically investigated only with future facilities, such as the James Webb Space Telescope (JWST) and Extremely Large Telescope class, in our case the MOSFIRE near-infrared spectroscopy has allowed us to access the details of the lines. Fig. 7 Comparison of the metallicity and ionization parameter () between the mean sample properties from Nakajima Ouchi (2014) and Ion2 with 1 uncertainties. Ion2 is compared to local galaxies from the SDSS survey, a galaxy sample at , a Lyman break galaxy (LBG) sample at , a Lyman alpha emitter (LAE) sample at , local Lyman break analogues (LBAs), LyC emitters, and Green Peas (GPs). See Nakajima Ouchi (2014) and references therein for a complete description of these samples. We derive the SFR


O B-CelestialObject O O O O O O O O O O O O O B-Model I-Model O O B-Model I-Model O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObject O O O B-CelestialObject O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject B-CelestialObject O O O O O O O O O O O B-CelestialObject O O O B-Wavelength O B-Wavelength O O O B-Citation I-Citation I-Citation I-Citation I-Citation O O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O B-CelestialObject O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObject O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O B-CelestialObject O O O O O O O O O O O O O O O O O O O B-Survey B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O B-Software B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O O O O B-Wavelength O B-Wavelength O O O O O O O O O B-CelestialObject O O O O O O O O O O O O O O O O O O O O O O O B-Wavelength O O O O O O O O O B-CelestialObject O O O O O O O O O O B-CelestialObject I-CelestialObject O O O O O O O B-Wavelength O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O B-Wavelength O O O O O O O O O O O O O O O O O O O O B-CelestialObject O O O O O O O B-Survey I-Survey O O O O O B-Organization O O B-URL B-Citation I-Citation I-Citation I-Citation O B-Citation O B-Citation I-Citation I-Citation I-Citation O O O B-Wavelength O B-Wavelength I-Wavelength O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Wavelength I-Wavelength O O O O O O O O B-Survey O O O O B-ObservationalTechniques I-ObservationalTechniques O O O O O O O O O O

of AT2019DSG, the S/N is too low (3), hence the polarization no longer follows a Gaussian distribution, but a Rice distribution instead (Plaszczynski et al. 2014 ). To obtain a robust estimate of the upper limit, we used the formulation by Plaszczynski et al. ( 2014 ). More specifically, for a 68% confidence level (corresponding to 1 σ ), the upper limit can be estimated as:We thus obtained an upper limit of 2.7% at the 1 σ level for the second epoch of AT2019DSG. It appears that AT2019DSG exhibits a polarization of ∼2% at the later epoch. However, we also need to take into account the interstellar polarization (ISP), which comes from the line-of-sight dust contributed by the Milky Way (MWG) and the host galaxy. By checking the extinction map, the line-of-sight MWG extinction— E ( B – V )—is ∼0.088 from Schlafly Finkbeiner ( 2011 ) and 0.102 from Schlegel et al. ( 1998 ). According to the empirical upper limits from Serkowski et al. ( 1975 ), this can translate into P as high as ∼1%. Another way to measure line-of-sight MWG dust polarization is to use field stars. In this regard, we also measure the polarization of a bright field star ∼1′ west of AT2019DSG. The field star polarization is at 0.5% level, consistent with the empirical upper limits from Serkowski et al. ( 1975 ). In addition to line-of-sight MWG dust polarization, there can also be polarization coming from the host. To constrain it, we used multiband photometry from Pan-STARRS (Chambers et al. 2016 ) and fit the host galaxy spectral energy distribution (SED) with LePhare (Arnouts et al. 1999 ; Ilbert et al. 2006 ). This gives an E ( B – V ) of 0.05, relatively small compared to the line-of-sight MWG extinction, hence the host galaxy plays an insignificant role in contributing to the polarization. This suggests that the high polarization (∼9%) in the optical in the earlier epoch cannot be explained by the MWG dust extinction alone. This is different from the case of Swift J164449.3+573451, where there was no polarization in the radio wavelength, and Wiersema et al. ( 2012 ) attributed the high linear polarization in the near-infrared to the dust induced polarization by the host galaxy. Table 2 Polarization Measurements We also include a light curve of AT2019DSG from the alerts issued by the public ZTF survey, aggregrated by the community broker ANTARES 8 8 https://antares.noao.edu (Saha et al. 2014 , 2016 ; Narayan et al. 2018 ) in both g and r bands, as shown in Figure 3 . With the light curve in hand, we can also check whether the TDE faded significantly between the first and second epoch, and whether the polarization signal is diluted (hence reduced) by the light from host galaxy. As can be seen in Figure 3 , the light curve indicates that the brightness of the TDE decreased from 17.87 ± 0.06 mag to 18.46 ± 0.10 mag in the g band between the two epochs. We note that the ZTF measurement is based on difference imaging, meaning that the photometry is solely from the light of


O O O B-Person I-Person B-Person I-Person B-Person I-Person B-Person I-Person B-Person I-Person B-Person I-Person B-Person I-Person I-Person I-Person B-Person I-Person B-Person I-Person B-Person I-Person B-Person I-Person O O O O O O O O O O O O O O O B-Person I-Person O O O B-Instrument O O O B-Person O O O B-Organization O B-Fellowship I-Fellowship I-Fellowship I-Fellowship I-Fellowship I-Fellowship O O O B-Organization I-Organization I-Organization O O O O O B-Observatory I-Observatory I-Observatory O B-Organization O O B-Grant B-Person O O O O B-Organization I-Organization I-Organization O B-Grant O O B-Organization O B-Grant I-Grant I-Grant O B-Grant O B-Grant O O O B-Organization I-Organization I-Organization I-Organization B-Organization O O O O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization B-Organization O O B-Organization O B-Grant O O O O O O O B-ComputingFacility I-ComputingFacility O B-Organization O O O O O O O O O O O O O B-Software O O O B-Software B-Citation I-Citation I-Citation I-Citation I-Citation O B-Software B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O B-Software B-Citation I-Citation I-Citation I-Citation O O B-Software B-Citation I-Citation I-Citation O O O O O O O B-Archive I-Archive I-Archive I-Archive

The authors thank Michael Cooper, Marla Geha, Coral Wheeler, Jose Oñorbe, Ferah Munshi, Erik Tollerud, Frank van den Bosch, Andrew Wetzel, Massimo Ricotti, Phil Hopkins, Justin Read, and the anonymous referee for valuable comments that have improved the manuscript. We also thank Aaron Dutton for providing the NIHAO data. Support for SGK was provided by NASA through Einstein Postdoctoral Fellowship grant number PF5-160136 awarded by the Chandra X-ray Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS8-03060. MBK acknowledges support from the National Science Foundation (grant AST-1517226) and from NASA through HST theory grants (programmes AR-12836 and AR-13888) awarded by the Space Telescope Science Institute (STScI), which is operated by the Association of Universities for Research in Astronomy (AURA), Inc., under NASA contract NAS5-26555. We also acknowledge the support of the Greenplanet cluster at UCI, where much of the analysis was performed. This work also made use of matplotlib (Hunter 2007 ), astropy (Astropy Collaboration et al. 2013 ), numpy (van der Walt, Colbert Varoquaux 2011 ), scipy (Jones et al. 01 ), and ipython (Perez Granger 2007 ). This research has made use of NASA's Astrophysics Data System.


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Acknowledgements We thank the staff of the GMRT who have made these observations possible. The GMRT is run by the National Centre for Radio Astrophysics of the Tata Institute of Fundamental Research. NK acknowledges support from the Department of Science and Technology via a Swarnajayanti Fellowship (DST/SJF/PSA-01/2012-13). We thank an anonymous referee for a very helpful report on an earlier draft that significantly improved the quality of this paper.


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We thank the anonymous referee for a careful review and valuable comments. We appreciate Y. Fujita and Y. Ohira for helpful discussion. This work is partially supported by the Grant-in-Aid for Scientific Research, No. 25400227 from the MEXT of Japan (KA).


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We wish to thank Marc De Rosa for contributing the PFSS library to SolarSoft, and for helpful reponses to questions about the software. We also wish to acknowledge Figen Oktem for useful discussions regarding regularization in optimization problems, and Leon Ofman for similarly useful discussions of coronal MHD simulations.


O O O O O O O O O O O O O O O O B-Instrument O O O O O O B-CelestialObjectRegion O O O O O O O O O O O O O O O O O O B-Instrument O O O O O O O O O O O O O O O O B-Instrument O O O O O O O O O O O O B-Instrument O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Instrument O B-Instrument O O O O O B-Wavelength O O O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObjectRegion O O O O O O O O O O O O O O O O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Formula I-Formula I-Formula O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Model O O O O O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Instrument O O O O O O O O O O O O O O O O O O O O O O O O O B-Instrument O O O O O O O O O O O O O O O O O O O O

range. Figure 2. Normalized (see Equations ( 1 )–( 3 )) power spectral densities observed by MMS1 at MHD scales in the quasi-perpendicular MSH during the time interval 07:00–08:00 UT on 2017 January 8. The parallel and perpendicular magnetic-field components measured by FGM are shown in red and dark blue, respectively. The parallel and perpendicular electric-field components measured by EDP are shown in orange and light blue, respectively. The density measured by FPI is displayed in green. The associated black solid lines represent the slope of the observed spectra. The vertical dashed line displays the electron gyrofrequency, while the dotted and dashed–dotted lines show the electron gyroradius and inertial length, respectively, estimated using Taylor’s hypothesis. The insert shows the anisotropy A at different frequencies. 3. Kinetic-scale Turbulence In order to study magnetic fluctuations in the sub-ion range, the burst-mode magnetic waveforms from FGM and SCM instruments are merged using an FIR filter (see, e.g., Breuillard et al. 2016 ), and B 0 and n 0 are defined as low-pass, filtered fluctuations at 0.1 Hz (i.e., below the frequency range of interest). Spatial and temporal scales are again determined as described in Section 1 ; however, as we will see below, intense whistler turbulence is detected in the quasi-perpendicular MSH, which might violate Taylor’s hypothesis. This is discussed in the next two sections. 3.1. Quasi-parallel Magnetosheath In this case, we have v bulk / v A > 3 and v bulk / v ms > 2 (see Table 1 , case 3), thus the Taylor hypothesis should not be significantly violated, and the frequency shift and flattening of the spectrum should be negligible (see Section 1 ). Figure 3 shows that at sub-ion scales (i.e., below the largest electron scale being f ρe here) the anisotropy of magnetic fluctuations is much smaller (here A ≈ 1) than in the MHD range, as noted in the previous section. Both perpendicular and parallel magnetic spectra follow the classical power law, f −2.8 , which is expected by KAWs and whistler turbulence. In this range, both parallel and perpendicular electric fluctuations follow an f −0.8 power law, i.e., , as expected from theoretical models (see Section 4 ). However, if the density spectrum also roughly follows an f −2.8 power law, , as in the MHD range, which is not expected from any of the models cited above. We note here that the density spectrum is affected by Poisson noise (for more details, see Gershman et al. 2015 ) above 10 Hz (not shown) and the flattening of the spectrum is not physical. Figure 3. Normalized (see Equations ( 1 )–( 3 )) power spectral densities observed by MMS1 at kinetic scales in the quasi-parallel MSH during the time interval 00:16–00:18 UT on 2017 January 18. The parallel and perpendicular magnetic-field components measured by FGM are shown in red and dark blue, respectively. The parallel and perpendicular electric-field components measured by EDP are shown in


O O O O B-Organization I-Organization I-Organization I-Organization I-Organization B-Organization O O O O O B-Organization I-Organization O O O O O O B-Location B-Collaboration I-Collaboration I-Collaboration I-Collaboration O O B-Grant O O B-Collaboration I-Collaboration I-Collaboration I-Collaboration B-Collaboration O O B-Grant O O O O O O O O B-Organization O O B-Grant I-Grant O O B-Grant O O B-Organization O

We gratefully acknowledge the Nederlandse Organisatie voor Wetenschappelijk Onderzoek (NWO) for the support of the FELIX Laboratory. This work is supported by the European MCSA ITN network “EUROPAH” (grant No. 722346) and the NWO Dutch Astrochemistry Network (DAN-II, grant No. 648.000.030). For the computational work, we acknowledge support by NWO under the “Rekentijd” program (grant No. 17603) and the SurfSARA staff.


O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject I-CelestialObject O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObject O O O O O O O O O O O O O B-CelestialRegion O O O O O O B-CelestialObject I-CelestialObject O O B-CelestialObject I-CelestialObject O O O O O O O O B-CelestialObject I-CelestialObject O O O O O O O O O O O O O O O O O B-CelestialRegion O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O B-Telescope O O O O O O O O O O O O O O O O O O O O B-CelestialRegion I-CelestialRegion I-CelestialRegion O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O B-CelestialRegion O O O O O O O O O O B-CelestialRegion O O O B-CelestialObject I-CelestialObject O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject O O O O B-Wavelength I-Wavelength I-Wavelength O B-Wavelength I-Wavelength I-Wavelength O O O O O O O O O O O O B-CelestialObject I-CelestialObject O O O O O O O O O O O B-CelestialObject I-CelestialObject O O O O O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject O O O O B-CelestialObject I-CelestialObject O O O O O O O O O O O O O O O B-Telescope B-Instrument O O O O O O O O O O O O B-CelestialRegion O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O

in Figure 2 . The black solid curves in panel (a) are the same isochrones as used in Figure 2 , but for a distance of 29 kpc. For the Sgr Stream, solid blue lines are used in panel (a) instead of dashed blue lines to represent the Sgr bright arm, and the distance of 23 kpc was adopted for these isochrones. In panel (b), we also show the Law Majewski ( 2010 ) N -body model particles in right- (leading) or left-facing (trailing) triangles. These particles are color coded following Law Majewski ( 2010 ), such that different colors represent the time at which a given debris particle became unbound from the Sgr dSph: magenta for 1.5–3 Gyr ago; and cyan for 3–5 Gyr ago. The gray cross is the same as in Figure 2 . An expectation value of 3 MW halo stars is predicted within the area spanned by the gray cross. The ORPH-F2 field is located on the main Sgr Stream (or the bright arm) at Λ ⊙ = 2264, and we expect Sgr Stream stars to be present in this field. We have overlaid the same age+metallicity isochrones used in the ORPH-F1 field for identifying Sgr stream stars, but at a distance of 22 kpc based on Figure 4 of Belokurov et al. ( 2014 ). Among the stars consistent with these isochrones, seven are easily identified as a conspicuous clump in the PM diagram. These are plotted in blue circles in Figure 4 . For comparison, we also plot the Law Majewski ( 2010 ) N -body model particles that are within 15 in both right ascension and declination from the center of our HST field. The observed PM clump is very close to the model particles as expected, and similar to the findings for Sgr FIELD 4 in Sohn et al. ( 2015 ) at Λ ⊙ = 2511. Beyond these, no other significant excess has been found in the PM diagram. 3.3. ORPH-F3 Figure 5 shows the CMD and PM diagram for the ORPH-F3 field. If present, Orphan Stream stars should be at a distance of 24 kpc. We used the fiducial isochrones to identify Orphan Stream stars, and found four stars consistent with the isochrones within their photometric errors. However, all of these stars are far away from the PM prediction of Newberg et al. ( 2010 ). There is only one star (22) that is close enough to the model prediction to be considered a member of the Orphan Stream, but its color ( F 606 W − F 775 W ≃ 0.8) is too red for its brightness to be considered an Orphan Stream star. We conclude that there is no star associated with the Orphan Stream in this field. Given that this field is about 12 away in the orthogonal direction from the densest part of the Orphan Stream, and given that the Orphan Stream is a thin stellar stream, it is unsurprising to find no star in a small HST ACS field. Figure 5. Selection of stars associated with stellar streams in the ORPH-F3 field. The panels and symbols are similar to those in Figure 2 . The black solid curves in panel (a) are the same isochrones as used in Figure 2 , but for a distance of 24 kpc. The dashed green curve is a fiducial isochrone for a population with an age of 12 Gyr and a