O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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-Software 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 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 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 B-Citation I-Citation O O O O O O B-Software O O O O O O O O O O O O O O O O O O O O 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-Software O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O O O O O O O O B-Software O O O O O O O O O O O O O O O O B-Tag O O O O O B-Software O O O O O O O O O O O O O O O O O O O O O O O

properties. This means that the stellar mass (or the dust luminosity when there is no stellar population included in the model) is not a free parameter even though it is technically possible to treat it as such. This would greatly expand the size of the parameter space by adding an extra dimension, while slowing the computation of the grid and the analysis by a similar amount, and degrade the accuracy of the estimation of the physical properties. Optionally, CIGALE can also handle fluxes for which only upper limits have been determined. We adopt the method presented in Appendix A2 of Sawicki (2012) . This affects the aforementioned computing steps in several ways. First, the computation of the χ 2 is divided between regular quantities (first three terms corresponding to Eq. ( 14 )) and those with only an upper limit (last three terms): with “erf” being the error function, f ul the flux upper limit, and the indices i , j , and k indicating respectively the fluxes, extensive properties, and intensive properties. See Eq. (A6)–(A10) of Sawicki (2012) for a full derivation 6 . The main difficulty is to determine α . For this we have to numerically solve ∂ χ 2 /∂ α = 0 for every model, which is equivalent to solving Eq. (A11) from Sawicki (2012) : We do so using the scipy.optimize.root root-finding method. Once the χ 2 are computed, the subsequent steps no longer depend on the presence or absence of upper limits. We have to note that objects in a catalogue do not all need to be fitted with the same set of data. CIGALE will automatically only consider the available data for a given object. The lack of certain data for some targets has a direct impact on some of the aforementioned computation steps. For steps 2 and 3, we simply do not include the data that have been marked as invalid in the input file (value lower than 0 or set to “nan”, by convention) in the computation of α and χ 2 . Another feature of this module is the possibility to assess whether or not physical properties can actually be estimated in a reliable way through the analysis of a mock catalogue. The idea is to compare the physical properties of the mock catalogue, which are known exactly, to the estimates from the analysis of the likelihood distribution. To build the mock catalogue we consider the best fit for each object. We then modify each quantity by adding a value taken from a Gaussian distribution with the same standard deviation as the uncertainty on the observation. This mock catalogue is then analysed in the exact same way as the original observations. Physical properties for which the exact and estimated values are similar can be estimated reliably. Applications of the pdf_analysis module to a sample of representative galaxies is presented in Sect. 5.3 . 5. Examples of <monospace>CIGALE</monospace> use cases As seen above, CIGALE has been designed to be flexible and versatile so that it may have various applications: estimation of the physical properties of an object


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B-CelestialObject I-CelestialObject 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 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 O O O O O O O O O O O B-Model I-Model I-Model I-Model I-Model I-Model I-Model B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O B-ObservationalTechniques I-ObservationalTechniques 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 B-ObservationalTechniques I-ObservationalTechniques 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 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 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 O 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 O O O O O O O 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-Software I-Software I-Software I-Software I-Software I-Software I-Software I-Software I-Software B-Software B-Citation I-Citation I-Citation I-Citation O B-Citation O B-Citation O B-Citation O B-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 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

HD 269902 (bottom). In HD 269902, most of the power in the WWZ is contained in the lower frequencies where the stochastic low-frequency background dominates—unsurprisingly given that the star has the second largest value of α out of the five FYPS (0.024 ± 0.001). Only by altering the scaling of the image of the WWZ are we able to see the ridge associated with the lowest-frequency peak, and it is only detected with low S/N. In both stars, the lowest-frequency peak in the RPS is associated with a broad band of power in the WWZ, similar to the WWZ of HD 269953, with a secondary band appearing at day −1 in HD 269840, at the approximate location of the highest peak in the RPS. 4. Discussion 4.1. Asteroseismic Modeling The discovery of fast pulsations in YSGs has very interesting implications for the study of the interiors of evolved massive stars. The 25 M ⊙ solar metallicity Geneva model (Ekström et al. 2012 ) has a main-sequence lifetime of approximately 7 Myr. It then crosses the H-R diagram in under 1 Myr, and half a Myr later, has evolved bluewards once more to become a Wolf–Rayet star (Massey et al. 2017 ). Due to the incredibly short lifetime of YSGs on both crossings of the H-R diagram, theoretical uncertainty has long stymied our understanding of massive star evolution (Kippenhahn Weigert 1990 ). Pulsation frequencies extracted from long-baseline lightcurves assembled from space-based observations are perhaps the most precise measurements we can make; typical values of for frequencies listed in Tables A1 – A5 are . In better-studied stars, such precision allows for the diagnosis of incredibly complicated physics and is truly the benchmark of testing stellar evolution theory (Aerts 2019 ). No asteroseismic models for YSGs that reliably converge exist (Jeffery Saio 2016 ), so the era of precision YSG asteroseismology is not yet upon us. However, we can compare the frequencies observed in the FYPS with the characteristic Lamb and Brunt–Väisälä angular frequencies in a model YSG: where l is the harmonic degree, c s is the sound speed, r is the radius of the local shell, g is the gravity, and the remaining variables are defined as in Aerts et al. ( 2010 ). Modes with angular frequencies and are mostly restored by pressure ( p -modes), and modes with and are mostly restored by buoyancy ( g -modes). While rudimentary, such a comparison would illustrate the approximate regions of the stellar structure the observed pulsations will allow us to probe. To do this, we use version 12778 of Modules for Experiments in Stellar Astrophysics (MESA, Paxton et al. 2011 , 2013 , 2015 , 2018 , 2019 ) to evolve a grid of nonrotating stellar models with initial masses between 15 and 30 M ⊙ , with a spacing of 1 M ⊙ , initial metallicity set to Z = 0.006, and initial helium mass fraction set to . We first evolve the models through the entire pre-main-sequence stage until the entirety of the star’s luminosity comes from nuclear burning and then introduce more


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

by the interaction of the supernova (SN) ejecta with the circumstellar medium (e.g. Ofek et al. 2007 ; Quimby et al. 2011c ; Chevalier Irwin 2011 ), (ii) pair-instability driven ( Gal-Yam et al. 2009 ; Gal-Yam Leonard 2009 ), or (iii) neutron-star driven ( Kasen Bildsten 2010 ; Dessart et al. 2012 ; Kotera et al. 2013 ; Metzger et al. 2014 ; Murase et al. 2015 ; Suzuki Maeda 2017 ). Given the variety of observed spectra, it is plausible that different processes are at play in different objects (e.g. Gal-Yam 2012a ; Nicholl et al. 2014 ). Interestingly, scenarios (i) and (iii) predict bright associated γ -ray emission in the GeV to TeV range ( Murase et al. 2011 , 2014 , 2015 ; Katz et al. 2012 ; Kotera et al. 2013 ). The search for such γ -ray emission with the Fermi Large Area Telescope (LAT) data is the scope of this paper. In the most conventional model (scenario i) SLSNe are powered by the interaction between the SN ejecta and a massive, optically thick circumstellar medium ( Smith McCray 2007 ; Smith et al. 2008 ; Miller et al. 2009 ; Benetti et al. 2014 ). SN2003ma and SN2006gy for example seem to be explained well by this phenomenology ( Smith McCray 2007 ; Ofek et al. 2007 ; Smith et al. 2010 ). Several authors ( Murase et al. 2011 , 2014 ; Katz et al. 2012 ) have demonstrated that a collisionless shock could then be formed and would host efficient cosmic-ray acceleration leading to non-thermal emission from radio-submillimeter to γ -rays. In the GeV range, this emission can escape from the system without severe attenuation, for specific shock velocities (about 4500–5000 km s −1 ) ( Murase et al. 2015 ), and at late times after the shock breakout. Ackermann et al. (2015) searched for this specific radiation with the LAT at the location of core-collapse SNe (Types IIn and Ib), with standard luminosity, spanning typical time windows of a few months to a year starting from the optical luminosity peak. No detection was reported and model-independent flux upper limits were derived. A fast-rotating central neutron star releasing its rotational energy into the SN ejecta could also drive SLSNe (model iii; e.g. Kasen Bildsten 2010 ). The rotation period has to be close to milliseconds to transfer significant energy to the ejecta. The strength of the initial dipole magnetic field of the star sets the timescale over which the energy is injected (a stronger field leads to faster decline). Magnetars ( B ~ 10 15 G) have thus been proposed as central engines for SLSNe ( Kasen Bildsten 2010 ; Dessart et al. 2012 ). Pulsars with millisecond periods at birth and milder dipole magnetic fields B ~ 10 13 G would also lead to bright peaks as well as a high-luminosity plateau lasting for several months to years ( Kotera et al. 2013 ; Murase et al. 2015 ). These authors further calculated that the young neutron-star wind nebula would present a bright X-ray and γ -ray peak, respectively, through synchrotron radiation and inverse Compton (IC) scattering,


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Acknowledgements H.R. thanks a CAPES fellowship and the CAPES PDSE programme (88881.132145/2016-01). J.M. thanks support by FAPESP (2012/24392-2 and 2014/18100-4) and CNPq (Productivity fellowship). We also thank Dr. Akito Tajitsu for his important support and initial data reduction during the observations with the Subaru Telescope. This study is based (in part) on data collected at Subaru Telescope, which is operated by the National Astronomical Observatory of Japan.


O O O O O O O B-Model I-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 B-Formula I-Formula I-Formula I-Formula 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 B-CelestialObject O B-CelestialObject O O O O O O O O B-CelestialObjectRegion I-CelestialObjectRegion O O O O O O O O O O O O O O O O O O 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 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 B-Telescope B-Dataset 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 B-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 O O O O O O O O B-CelestialObject I-CelestialObject B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O O O O O B-Citation I-Citation I-Citation I-Citation O B-Citation O O O O O O O O O O O O O B-Telescope B-Dataset 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 B-Telescope B-Dataset O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 B-Telescope B-Dataset 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 B-Telescope I-Telescope O O O O O 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 O O 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 O O O O O O O

radius from the particle data of the N -body simulation. We also determine the circular frequency, Ω, by averaging the radial accelerations of particles in an annulus with a width of 50 pc at each radius. In Fig. 8 , we show Ω + κ/ m ( m =2, 3, 4, and 5) and Ω as functions of R . The horizontal line in Fig. 8 indicates the pattern speed of the bar. The radii at which the line intersects with the frequency curves correspond to the guiding-centre radii for the 2:1, 3:1, 4:1, 5:1, and the corotation resonances. In our model of the MW, the Sun, at a distance of ∼8 kpc from the Galactic Centre, is close to the 3:1 OLR. It is then not surprising that this resonance is dominant in the local stellar population. Figure 8. The orbital frequencies as functions of the galactic radius, R . Ω was determined from the velocity curve and Ω + κ/ m ( m =2, 3, 4, and 5) as expected within the epicycle approximation are plotted. The vertical axis is normalized by Ω 8kpc . The horizontal line indicates the bar’s pattern speed and the vertical lines represent the guiding radii for the respective resonant orbits. 4.2 The MW’s bar pattern speed When projecting the Gaia DR2 data into the R – v φ plane, a series of ridges appear over a wide range of R (Antoja et al. 2018 ; Kawata et al. 2018 ; Ramos et al. 2018 ). Although the origin of these structures is still in debate, previous studies (e.g. Monari et al. 2017b , 2019a ; Fragkoudi et al. 2019 ; Barros et al. 2020 ) have discussed that these ridges may originate from resonances, while other studies have suggested perturbations by a satellite galaxy such as the Sagittarius dwarf (Khanna et al. 2019 ; Laporte et al. 2019 ) or by winding transient spiral structure (Hunt et al. 2018 , 2019 ). In this section, we study the relation between the observed ridges in Gaia DR2 and resonances in our simulated MW model: by projecting our simulated and classified stars into the R – v φ plane, we can directly compare to the same projection of the Gaia DR2 data. If the stars classified by resonance type align with the observed ridges, then this provides compelling evidence that these ridges are indeed caused by resonantly trapped stars. In Fig. 9 , we show the distributions of stars in the R – v φ plane as it arises in our simulated MW model (left-hand column) and in the Gaia DR2 (right-hand column). The top panels show 2D histograms of particles within 2 kpc from the Sun [, in our model]. This panel reveals a selection effect caused by Gaia ’s measurement uncertainties: due to us having selected stars with relative parallax error below 10 per cent the right-hand panel displays systematically fewer stars in approximately concentric rings around the Sun’s position at Visually, this causes the smaller extend of the red area in the right-hand panel. We expected, however, that our analysis is robust with respect to this selection effect: it implies that fewer stars make it into our DR2 projection than into the projection of our


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J.H. greatly appreciates the CNES post-doctoral fellowship program. J.H. and M.D. were supported by the project under the contract 11-BS56-008 (SHOCKS) of the French Agence National de la Recherche (ANR), JD by grant GACR 15-04816S of the Czech Science Foundation, DO by the grant NCN 2012/S/ST9/00022 of Polish National Science Center, and A. Marciniak by grant 2014/13/D/ST9/01818 of Polish National Science Center. We thank the referee, Mikko Kaasalainen, for his thorough review of our manuscript and his constructive comments and suggestions that led to a significant improvement of the text. The computations have been carried out on the “Mesocentre” computers, hosted by the Observatoire de la Côte d’Azur, and on the computational cluster Tiger at the Astronomical Institute of Charles University in Prague ( http://sirrah.troja.mff.cuni.cz/tiger ). Data from Pic du Midi Observatory were partly obtained with the 0.6 m telescope, a facility operated by observatoire Midi-Pyrénées and Association T60, an amateur association. The Joan Oró Telescope (TJO) of the Montsec Astronomical Observatory (OAdM) is owned by the Catalan Government and operated by the Institute for Space Studies of Catalonia (IEEC). We thank Franck Pino (INO-AZ) and Lech Mankiewicz (EU-HOU/Comenius) for the remote access to Ironwood North.


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 O O O O O O O O O O O O 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 B-Wavelength 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 B-Citation I-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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O B-Wavelength O O 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 I-Citation O B-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

supported. The key features of each multiphase flow are summarized as follows. As rotation dominates over turbulent motions (Ta t v rot / v 1) in the massive galaxy, the condensation occurs along a conical helix, spinning up the gas which finally settles on to the equatorial region in a multiphase disc (Section 3 ). The disc is cylindrically stratified: the ionized (UV) gas envelops the neutral (optical/IR) layer which englobes the very thin molecular (radio) structure, dominating 95percent of the condensed mass. There is no major fragmentation observed, with the warm phase experiencing mild perturbations via pressure torques induced by condensation. The BHAR is suppressed by almost 2 dex compared with the cooling rate as most gas has found circularization. Accretion proceeds along a narrow polar funnel (10) or after reaching the inner equator through gas which condensed out of the plasma with low angular momentum. After 50Myr the disc has reached 1kpc extension. Extended molecular and warm discs with similar surface density have been observed in 1/3 of observed massive, typically quiescent, galaxies (Caon etal. 2000 ; Ocaa Flaquer etal. 2010 ; Young etal. 2011 ; Werner etal. 2014 ; Hamer etal. 2016 ). When rotation and turbulence are both inhibited, the hot plasma condenses radially, forming a compact multiphase cloud within the inner 100pc (Section 4 ). The inner multiphase density profiles are steep, following a free-fall slope ( r 3/2 within the SMBH influence region). The BHAR is almost identical to the cooling rate. Adiabatic, compressional heating inhibits the formation of the third molecular phase. The warm 8000K (optical/IR) gas dominates the condensed medium, albeit highly concentrated, with a very thin ionized UV layer. Similarly, a significant number of quiescent massive galaxies display a spherically symmetric warm H core (Caon etal. 2000 ; Werner etal. 2012 , 2014 ) and tend to lack cold gas, in particular when no AGN feedback or other major perturber are detected. In the more frequent, intermediate state the hot halo experiences both rotation and subsonic turbulent motions (160kms 1 ) leading to CCA (Section 5 ). The multiphase condensation occurs in the high density peaks via local runaway cooling. AGN heating prevents a global cooling inflow, facilitating the growth of multiphase structures. Initially, warm filaments condense out of the low-entropy plasma. The denser, perturbed regions further condense into molecular clouds, following a top-down rainfall. The CCA angular momentum transport is driven by pressure torques on 1kpc scale and dominated by chaotic inelastic collisions in the core (ram-pressure drag is secondary). The accretion rate can be modelled via quasi-spherical viscous accretion (Section 5.2 ), with an effective clump collisional viscosity c v c ; the mean free path is c 100pc with large scatter. The core cloud volume filling is a few per cent, rapidly declining beyond 1kpc. As ensemble, the multiphase clouds


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ACKNOWLEDGEMENTS This research has been supported by Arizona State University. We would like to thank Soma De, Tanmay Vachaspati, Daniel C. Jacobs, Adam Beardsley and Samuel Gordon for their insightful conversations. Some of the results in this paper have been derived using the healpix (Górski et al. 2005 ) package. This research made use of astropy , a community-developed core Python package for Astronomy (Astropy Collaboration et al. 2013 ).


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

z LAEs with high-resolution HST narrow-band imaging (Bond et al. 2010 ; Finkelstein et al. 2011 ). Bond et al. ( 2010 ) suggested that Ly α sizes are compact and similar to UV emission, but Finkelstein et al. ( 2011 ) posited that the half light radius of Ly α appears ∼1.6 times larger than the half light radius of UV continuum. These narrow-band HST images of high- z LAEs are very hard to obtain and have low S/Ns. Thus, the Ly α and UV sizes of low- z LAEs are valuable. Because Green Peas are analogs of high- z LAEs, our results suggest that most high- z LAEs likely have larger Ly α than UV sizes. The extended Ly α emission probably indicates gas outflows around galaxies illuminated by Ly α light. One interesting question regards the redshift evolution of Ly α sizes of LAEs. Recently, Wisotzki et al. ( 2016 ) measured Ly α radial profiles of a sample of LAEs at z = 3–6 from VLT/MUSE data. They found that, in 12 LAEs with both Ly α and UV continuum sizes, the Ly α light is considerably more extended than the UV continuum light. Here, we compare the sizes of Green Peas and the MUSE LAEs. Using the Ly α radial profiles in Wisotzki et al. ( 2016 ), we measured the deconvolved Ly α scale radius r e , assuming intrinsic exponential profiles, so that the methods are same when measuring the r e of MUSE LAES and Green Peas. As shown in Figure 5 , the Ly α -to-UV size ratios of Green Peas and MUSE LAEs are very similar. Notice that some MUSE LAEs have extended Ly α halos far beyond the scale radius. For Green Peas, however, we do not have robust data to characterize the Ly α emission beyond a few Kpcs. Figure 5. Left: Comparison of the Ly α- to-UV scale length of Green Peas (green squares) and MUSE z = 3–6 LAEs sample (red dots) (Wisotzki et al. 2016 ). The Ly α scale lengths of MUSE LAEs are measured from the radial profile in Wisotzki et al. ( 2016 ) using the same method as Green Peas. Right: Comparison of the Ly α -to-UV Petrosian 20% (P20) radius of LARS ∼ 0 galaxies and MUSE LAEs. Notice that the Petrosian 20% radius of MUSE sample is measured from the best-fit model of radial profile. The dashed cyan lines show constant ratios of 1:1, 2:1, 5:1, and 10:1. In the right panel of Figure 5 , we compare the Petrosian 20% radius ( R P20 ) of MUSE LAEs (Table 2 in Wisotzki et al. ( 2016 )) to that of the LARS sample (Table 1 in Hayes et al. 2014 ). Compared to the five strong Ly α emitters in LARS sample (marked by stars, LARS01, 02, 05, 07, and 14), the MUSE LAEs only have about two times larger ratios of R P20 (Ly α ) to R P20 (UV). One caveat of the comparison is that the Petrosian radius of MUSE sample is measured from the best-fit model of radial profile, instead of the observed data, which is different from the method used in the LARS sample. This might be the reason that the R P20 (UV) of MUSE LAEs are about 2–4 kpc, approximately a factor of three larger than the R P20 (UV) of the five LARS Ly α emitters. Based on our rough comparison of Green


O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 B-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 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 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 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 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O

this correlation is presently unclear. Figure 8. Periastron distance (top panel; ) and the tidal force acting on the primary star at periastron divided by the star’s surface gravity (bottom panel; see Equation ( 4 )), both as a function of the full amplitude of the photometric HB signal at periastron. In both panels the markers’ radius is linear in the primary star’s (See Table 1 ). 4.5. Higher Multiplicity Systems Many of the HB stars examined here may be members of triple or higher multiplicity systems. Several works (e.g., Duquennoy Mayor 1991 ; Meibom Mathieu 2005 ; Tokovinin et al. 2006 ; Raghavan et al. 2010 ) have found that the fraction of higher multiplicity systems among short period and highly eccentric binaries is very high, exceeding 90% for binaries with days (Tokovinin et al. 2006 ). Indeed, tertiary bodies may excite the orbital eccentricity of HB progenitors via Kozai–Lidov oscillations, producing close periastron passages and allowing them to be detected as HB systems. Quadruple systems composed of two binaries may also be common among HB systems, and these types of systems have been predicted by Pejcha et al. ( 2013 ), who found that HB systems formation is greatly enhanced in quadruple systems relative to triple and binary systems. Despite the probable multiplicity of many HB systems, we find it unlikely that many of our Keplerian RV solutions have been greatly affected by orbital motion induced by a third body. A putative third body exciting the orbital eccentricity would likely have an orbital period, P 3 , more than several times longer than the HB period (depending on its mass, inclination, and eccentricity) in order to allow long term dynamical stability (Kiseleva et al. 1994 ). If a typical HB progenitor system was born with an orbital period days before having its orbital eccentricity excited, we expect for most third bodies, with much larger values of P 3 in most cases. Although the orbital RV amplitude induced by such a third body can be several km s −1 , we expect only a small RV change over the ∼100 days baseline of our observations, and therefore this motion can be safely included as a constant RV offset. We are currently pursuing a follow-up survey on many of these HB systems that will reveal whether any of them have tertiary companions with orbital periods . 4.6. Future Prospects The orbits measured here comprise a relatively large sample obtained with a single instrument, and they can facilitate several follow-up scientific studies in addition to that of tidal circularization and the shape of the e – P distribution (Section 4.3 ). Other examples include: 1. Stars in binary systems with highly eccentric orbits are expected to rotate near a “pseudo-synchronous” rotation period, where the stellar rotation is synchronized with the orbital motion close to periastron which is faster than the mean orbital motion along the entire orbit. The pseudo-synchronous rotation rate depends on eccentricity and orbital period


B-Person O O O B-Organization I-Organization I-Organization B-Grant I-Grant B-Grant I-Grant I-Grant B-Organization I-Organization I-Organization I-Organization B-Organization O O O O O O B-Fellowship I-Fellowship I-Fellowship O O B-Organization I-Organization I-Organization I-Organization I-Organization O O B-Organization I-Organization I-Organization I-Organization O O O B-Organization I-Organization I-Organization I-Organization O O O B-Organization B-Person O B-Person I-Person O B-Person I-Person O O O O O B-Software B-Person I-Person O O O O O B-Software O O B-Person I-Person O O O O O B-Person O B-Person I-Person O O O O B-Telescope I-Telescope I-Telescope O O O O O O O B-CelestialObject I-CelestialObject B-Person I-Person I-Person O O O O B-Telescope I-Telescope I-Telescope B-Person I-Person O O O O O B-Instrument O B-Person I-Person O O O O O O O O B-Telescope O O B-Person I-Person O O O O O O B-Telescope O O O B-Person I-Person I-Person I-Person O O B-Organization O O O O O O O O O O O O O B-Instrument O B-Person O O O B-Observatory I-Observatory O O B-Person I-Person O O O O O B-Software O O O O B-Observatory I-Observatory I-Observatory I-Observatory O O O O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization B-Organization O O O O O O B-Organization I-Organization I-Organization O O O O O O O O O O O O B-Archive I-Archive I-Archive I-Archive I-Archive B-Archive B-Organization O O O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization O B-Organization O B-Grant O O O O O O O B-Organization O O O B-Grant I-Grant I-Grant O O B-Organization I-Organization I-Organization I-Organization O O O O B-Organization I-Organization O B-Organization O B-Grant I-Grant O O O O O O B-Observatory I-Observatory I-Observatory B-Telescope I-Telescope I-Telescope O O O O O O O B-Organization I-Organization I-Organization O B-Telescope I-Telescope I-Telescope I-Telescope O O O B-Observatory I-Observatory O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization O O B-Observatory I-Observatory I-Observatory I-Observatory O O O O O O O O O B-Organization I-Organization I-Organization B-Organization O B-Telescope O B-URL O O O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization B-Organization B-URL O O O O B-Organization O O O O O O O O O O O O O B-Collaboration I-Collaboration I-Collaboration


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tidal forces from the Milky Way, or any other physical process 9 9 For example, relaxation cannot be responsible for the spatial distribution of the different stellar populations in Carina since the relaxation times in dwarf galaxies are much larger than the age of the universe (e.g., Freeman 2008 ). , were not strong enough to move a significant fraction of these stars to the outer region. Otherwise, a large number of intermediate-age stars would be present in the outskirts of Carina, and thus we would not observe such a strong break in the SFH and MDF of Carina between the intermediate and outer regions. Since the radial trends in Carina do not seem to be governed by the tidal influence of the Milky Way, we suggest that this hints that the SFH is not dominated by tidal forces either. In this line, the fact that we find two episodes of star formation in Carina instead of three, as has often been claimed (Hurley-Keller et al. 1998 ; Rizzi et al. 2003 ; Pilkington Gibson 2012 ; McMonigal et al. 2014 ), makes it harder to explain the star-formation episodes in Carina as being the result of periodic close passages to the Milky Way (as done in Pasetto et al. 2011 ), even more so since proper motions suggest that the orbital period of Carina is close to 2 Gyr (Piatek et al. 2003 ; Pasetto et al. 2011 ). Another mechanism proposed to explain the SFH of Carina is gas accretion from the Milky Way (e.g., Lemasle et al. 2012 ; de Boer et al. 2014 ). In this view, the first episode would have ended because of gas depletion, and subsequent star formation would have been stopped for a few gigayears until the dwarf galaxy renewed its gas content thanks to an inflow from the Milky Way. Even though this is a plausible scenario, the in-falling gas should have had the same metallicity as the stars from the end of the first episode to reproduce the age–metallicity relation in Carina. A simpler way to explain why the metallicity of the stars at the end of the first episode coincides with that of the stars at the beginning of the second episode is to assume that the gas that formed the second episode in Carina was enriched within the same galaxy by the first generation of stars. Internal evolution would also be consistent with the mean metallicities and ages at different Carina regions, which according to what we see in Figure 9 are roughly constant throughout the extent of the galaxy. This would not necessarily be the case if the gas that formed those stars had been accreted at different moments and each region had different levels of influence from each inflow. For all these reasons, we conclude that the SFH of Carina must be dominated by internal evolution. Our interpretation is that, after the first episode formed the majority of the stars in Carina, the physical conditions of the gas made the galaxy unable to form stars for a few gigayears, until it contracted to activate star formation again. We claim that the temporal gap in star formation in Carina might


O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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-Software O O O O O O B-Software 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 I-Model I-Model 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 B-Model I-Model 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 B-Formula I-Formula I-Formula I-Formula 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 B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation 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 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 B-Citation I-Citation O O O O O O

of the disc will be In our model, the relative change in disc mass and accretion rate in an encounter is determined by γ and the truncation radius relative to the disc characteristic radius just before the encounter, . If the disc is truncated, its viscous time-scale decreases and the disc starts evolving faster. We ignore the orientation of the discs, so the equation for truncation radius is the average truncation radius over all disc inclinations. 2.4 Numerical implementation We use the fourth-order Hermite N -body code ph4 , which is incorporated into the AMUSE framework. The softening length is and the time-step parameter is η = 10 −2 . The simulations start in virial equilibrium and last for 2 Myr. Collisional radii for each star in the cluster are defined depending on the stellar mass and the theoretical size of the disc evolving in isolation, given by equation ( 1 ). The initial collisional radius for each star is given by the distance at which the most massive star in the cluster can truncate its disc. The collisional radii of all stars are updated every 2000 yr to account for the viscous evolution of the discs. Two stars are considered to be in an encounter if the distance between them is less than the sum of their collisional radii. The encounter distance between the stars is then computed by analytically solving the two-body problem. If the encounter is strong enough as to truncate the discs, disc parameters of both stars are updated as described in Section 2.3 . After each encounter, the collisional radii of both stars are set to 0.49 r enc in order to prevent the N -body code from detecting the same encounter multiple times during the same 2000 yr window. 3 RESULTS 3.1 Initial conditions 3.1.1 Cluster properties Simulations were run for clusters with 1500 stars. The stellar masses are randomly sampled from a Kroupa initial mass distribution (Kroupa 2001 ) with lower mass limit 0.1M ⊙ and upper mass limit 100M ⊙ . The stars are initially distributed in a Plummer sphere (Plummer 1911 ) with Plummer radius 0.5 parsec. All stars are assumed to be single and coeval. All simulations were carried out for . We define three different gas scenarios for the simulations: No gas Gas presence Gas expulsion starting at Each simulation was run for two values of the turbulence parameter: α = 10 −2 and α = 5 × 10 −3 . 3.1.2 Disc properties Based on the observations of low-mass stars carried out by Isella et al. ( 2009 ) and the estimations obtained by Hartmann et al. ( 1998 ), the initial characteristic radius R c of the circumstellar discs was chosen to be with . Circumstellar discs with masses larger than 10 per cent of the mass of their star are the most likely to be gravitationally unstable (Kratter Lodato 2016 ). According to hydrodynamical simulations of collapsing gas clouds, the disc-to-star mass ratios of embedded protostars are large enough for gravitational instabilities to occur (Vorobyov 2011 ), but these instabilities lead to


O O B-Dataset I-Dataset O O O O 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 I-Formula O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula 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 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 B-Citation I-Citation I-Citation I-Citation 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 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O O O O O O O O O O O O O O O O O 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 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 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

from the MPA-JHU DR7 catalogue, and we obtained MBH only for galaxies classified as ellipticals (see Section 3 ). The values are given in Table 8 . MLLINERs cover the range between log (MBH/M ⊙ ) = 7.03–8.57 with a median value of log (MBH/M ⊙ ) = 7.45, while LLINERs show MBH in the range log (MBH/M ⊙ ) = 6.24–8.54 and median value of log (MBH/M ⊙ ) = 8.04. Interestingly, our MLLINERs do not contain the most massive BHs in their centres. It could be surprising that MLLINERs, having on average higher SFRs than LLINERs, show in general lower stellar masses. Different works, both observational (Kauffmann et al. 2003b ; Mateus et al. 2006 ; Leauthaud et al. 2012 ; Pérez et al. 2013 ) and numerical (Shankar et al. 2006 ; Behroozi et al. 2012 ), revealed a stellar mass of ∼6 × 10 10 M ⊙ as critical for the growth rate of stellar populations. In particular, in Pérez et al. ( 2013 ) by studying a 3D spectroscopic sample of 105 local galaxies, the authors found that in galaxies more massive than 5 × 10 10 M ⊙ the inner regions ( 0.5 R 50 ) grew as much as 50–100 per cent faster than in the lower mass galaxies. They found that the peak of relative growth rates of inner and outer galaxy regions correspond to the stellar mass of 6–7 × 10 10 M ⊙ (see their fig. 5), while for lower and higher masses the growth rate decreases and therefore SFRs (LSF) as well. The median stellar mass of our MLLINERs (6.58 × 10 10 M ⊙ ) corresponds perfectly to this region, while for most LLINERs their stellar masses are already higher and correspond to lower values of the relative growth rate (lower LSF). This explains why MLLINERs having in average lower stellar masses in comparison to LLINERs, have higher LSF. In addition, we studied the stellar mass distributions of MLLINERs and LLINERs for the three morphological groups. Table 9 shows the median stellar masses for different morphological types. As can be seen, of three morphological types the highest difference was obtained for early-type galaxies. These galaxies represent a significant fraction of MLLINERs (40 per cent, see Fig. 8 ) and their median mass corresponds exactly to the highest relative growth rate of stellar populations (according to Pérez et al. 2013 ), which then could explain their high SFR values. This is not the case for early-type LLINERs that are characterized by higher stellar masses (lower growth rates) and lower SFRs. Table 9. Median total stellar masses of MLLINERs and LLINERs in relation to morphology (given as logarithm and in M ⊙ ). 5.1.2 Extinction We found that our MLLINERs can be hosted by galaxies with a wide range of extinctions. When using the Av measurements based on the Hα and Hβ emission lines, we find that most of them reside in galaxies with high extinctions. The median Av is 1.65 mag, covering the range 0.49–3.46 mag (see Fig. 7 , top plot). For comparison, the bottom plot in Fig. 7 shows the Av distributions of LLINERs and MLLINERs when taking into account the SDSS MPA-JHU DR7 3


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The authors acknowledge research funding by Deutsche Forschungsgemeinschaft (DFG) under grant SFB 963/1, project A5.


O O O O 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-Archive I-Archive I-Archive O O O O O O B-Archive I-Archive I-Archive O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization B-Organization O B-Location I-Location O B-Archive I-Archive I-Archive I-Archive I-Archive

The authors would like to thank the anonymous referee for helpful comments and suggestions. The scientific results reported in this article are based on data obtained from the Chandra Data Archive. We also made use of the ROSAT data archive of the Max Planck Institute für Extraterrestrische Physik (MPE) at Garching, Germany and NRAO VLA Sky Survey archives.


O O O O O O O O B-Organization I-Organization I-Organization B-Grant O B-Grant I-Grant I-Grant I-Grant I-Grant I-Grant O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization B-Organization B-Person O O O O B-Person I-Person O B-Person I-Person O O O O O O O O O O O O B-Mission O O O O O B-Mission O O O O O O B-Organization B-Database O B-URL O

This work was supported by the grant of Universiti Kebangsaan Malaysia (GUP-2018-086) and Grant-in-Aid for Scientific Research (C) 17K05673 from the Japan Society for the Promotion of Science (JSPS). W.L.T. would like to thank Takuma Nakamura and Bo-Jhou Wang for useful discussions. Special thanks to the dedicated efforts of the entire MMS mission team for data access. MMS data can be downloaded through the NASA CDAWeb at http://cdaweb.gsfc.nasa.gov/ .


O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O B-Wavelength 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 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 B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O B-Wavelength I-Wavelength I-Wavelength O B-Formula I-Formula I-Formula I-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 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 O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O O O O O O O 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 O O O O 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 O O O O O O O O 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 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 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 B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O B-Formula I-Formula I-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 I-Formula I-Formula I-Formula O O O O O

environments. Our results will aid in the preparation of multicomponent ices containing acetaldehyde by allowing for calibrations and determinations of that compound in an ice sample. We also envision applications of our data to the measurement of reaction yields in experiments that might produce acetaldehyde in polar ices. Our IR intensities also can be used to estimate band strengths in spectral regions beyond the mid-IR used in the present work. See Gerakines et al. ( 2005 ) for an example. Astrochemical results based on IR band strengths published before those presented here can now be updated with our new measurements. For example, and as already described, the reference band strength of Schutte et al. ( 1999 ) for amorphous acetaldehyde's CH 3 deformation at 1350 cm −1 is A ′ = 1.5 × 10 −18 cm molecule −1 , which now is seen to be in error by nearly 80 per cent. Molecular abundances or column densities based on that older A ′ will be in error by that much, although published relative band strengths might not change substantially. Finally, we again point out that few aldehydes have been examined with the methods used here. As an example, propionaldehyde (propanal) is an interstellar molecule, and at least two laboratory papers recently have been published involving it (Jonusas, Guillemin Krim 2017 ; Qasim et al. 2019 ), but no band strength measurements are available for the solid compound. The situation is no better for H 2 CO (formaldehyde). Many laboratory papers have been published involving H 2 CO, but only estimates and indirect measurements are available for its IR intensities. Unfortunately, their accuracy is unknown. 5 SUMMARY AND CONCLUSIONS Here we have reported the first direct measurements of band intensities of solid acetaldehyde, removing the need to rely on either calculated gas-phase results or values that cannot be traced to their original source. We also have presented mid-IR spectra of the two crystalline polymorphs previously identified by neutron-diffraction studies, one stable under our conditions and one metastable. The IR spectrum of H 13 C(O) 13 CH 3 has been recorded and used to identify ketene as a radiolytic decomposition product of solid acetaldehyde for the first time. APPENDIX A The procedure used to determine absorption coefficients ( α ′) and band strengths ( A ′) is the same as in our earlier papers, and is essentially the method of Hollenberg Dows ( 1961 ). Briefly, IR spectra were recorded of ice samples having a range of thicknesses. The absorbance at each peak of interest was measured and bands of interest were integrated. Multiplication of each absorbance or band area by ln(10) ≈ 2.303 was then done to convert from a common to a natural logarithmic scale. Representative plots of the results are shown in Fig. A1 . The slope of the upper graph is α ′ and the slope of the lower one is ( A ′ ρN A / M ), where M = molar mass = 44.05 g mol −1 , N A = 6.022 × 10 23 molecules mol −1 , and density ( ρ


O O O O O B-Tag O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O O 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 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 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 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 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 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 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 B-Formula I-Formula I-Formula I-Formula O B-Formula I-Formula I-Formula I-Formula 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 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 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 B-Citation I-Citation I-Citation I-Citation O B-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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O

(Linder 2008 ). 3.4. Growing <italic>ν</italic> Mass Growing neutrino mass models, wherein the mass of the neutrino ( m ν ) increases with time and stops the dynamical evolution of the dark-energy scalar field are invoked to solve the cosmological coincidence problem, i.e., the problem that the present-day matter density and density of Λ are similar order of magnitude despite their different dependence on the scale factor (Fardon et al. 2004 ; Wetterich 2007 ). The combined dark sector (scalar field plus neutrinos) energy density (where a = 1/(1 + z )) is given by where Ω ds = 1 − Ω M is the present-day dark-energy density. The scale factor at which the neutrinos become nonrelativistic, is given by preserving continuity between the early and late-time terms (i.e., setting Equations ( 14 ) and ( 15 ) at a = a t ). The two free parameters are the early dark-energy density Ω e and the neutrino density Ω ν . The normalized Hubble parameter for this model is given by 3.5. Dark-energy Transitions at Low Redshift Mortonson et al. ( 2009 ) proposes models with large fluctuations in the dark-energy equation of state at low redshifts, typically z ≲ 0.02, that induce step-like transitions in the H ( z ). Such changes are hidden from constraints coming from Hubble flow SNe Ia (as even the lowest redshift SNe can typically be at z > z t ), but calibrators offer an additional restriction. Step-like responses in H ( z ) also evade several model-independent constraints, as they often assume smoothness of H ( z ). Phenomenologically, we can write the first Friedmann equation in these models aswhereIn this definition, z t , and δ are the position, width, and amplitude of the transition, respectively. In particular, δ = 0 corresponds to the standard ΛCDM . Hence, this model has three free parameters at a given fixed z t : , and δ . The observed Hubble constant is . In this work, we either assume z t = 0.02 or z t = 0.1; the former case motivated by Mortonson et al. ( 2009 ) and the latter case motivated so we can use all the low- z SNe at z ≈ 0.1. The adopted priors for and are equivalent to the priors shown on Table 1 for H 0 and Ω M . The difference in notation highlights the fact that and cannot be interpreted as the observed Hubble constant and the dark-matter density at redshift zero in these models (see Mortonson et al. 2009 ). 3.6. Cosmographic Expansion Along with the different dark-energy models described above, we also look into a more model-independent method, by expanding the expression for H ( z ) as a Taylor series in a cosmographic approach (for e.g., Arendse et al. 2019a , 2019b ; Feeney et al. 2019 ; Macaulay et al. 2019 ). Expanding H ( z ), we getwhere , 2 and . 4. Results In this section we present the results of fitting the different dark-energy models described in Section 4.1 to the combined calibrator and Hubble flow SNe Ia. We also discuss the impact of the systematics covariance matrix on the inferred value of H 0 in Section 4.2 .


O O O O B-Citation I-Citation I-Citation O B-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 B-ObservationalTechniques O 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-Person 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 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 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 O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O

( 1969 ) and Yoshimura ( 1975a , 1975b ) succeeded in numerically reproducing the latitudinal evolution of the toroidal magnetic field (butterfly diagram). In order to obtain the equatorward migration of the toroidal magnetic field, it is required that the angular velocity should decrease as the radius increases in the solar convection zone ( for the positive α in the northern hemisphere; : angular velocity). However, the developments of helioseismology revealed that the real radial distribution of the angular velocity in the solar convective zone does not satisfy this requirement. In addition to this point, Parker himself discussed the limitation of the α dynamo in the solar convection zone from the viewpoint of magnetic-field buoyancy (Parker 1975 ). Once the magnitude of a magnetic field has been amplified to some level, the magnetic flux tube starts moving upward from deeper to shallower regions. Since this buoyancy convection timescale, even for a very weak magnetic flux tube, is estimated to be very short as compared with the α -dynamo amplification timescale, it reaches the solar surface without being sufficiently amplified by the dynamo. This is called the magnetic buoyancy dilemma intrinsic to the magnetic flux tube. These difficulties have been considered to be obviated by assuming that the dynamo action is operated in a convectively stable thin layer between the radiative and convection zones (Parker 1975 ). This region is associated with a large differential rotation and is called the tachocline . In the tachocline dynamo picture, the solar magnetic field is generated by a strong differential rotation there and may have enough time to be amplified by the α dynamo before it is raised by the magnetic buoyancy (Spiegel Weiss 1980 ; van Ballegooijen 1982 ). As mentioned above, depending on the location of the field generation, the dynamo models are divided into two types: (i) the flux-transport dynamo, which operates in the overshoot layer below the bottom of the convection zone; and (ii) the mean-field distributed dynamo, which works in the bulk of the convection zone, in particular at the subsurface shear layer. On the surface of the Sun, large-scale meridional circulations with a speed of ∼15 m s −1 in the poleward direction have been observed. The combination of the meridional circulation and the α –Ω mean-field model demonstrated its ability to reproduce the basic features of the solar magnetic activity cycle such as the equatorward migration of the sunspots, the cycle period, and the phase difference between the poloidal and toroidal magnetic fields. In the combination with meridional circulations or convective-diffusive motions such as the Babcock–Leighton-type convection, which produces the east–west tilt of the sunspot pairs under the action of the Coriolis force on the magnetic flux tube during its rise through the convection zone, it is expected that the tachocline dynamo can reproduce several properties of the


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We thank the referee and B. Czerny, C. Done, C. Knigge, J. Matthews, H. Netzer, and D. Starkey for discussions. P.H. and G.S. acknowledge the support of the Natural Sciences and Engineering Research Council of Canada (NSERC), funding reference number 2017-05983. K.H. acknowledges support from STFC grant ST/M001296/1. Software: MATLAB.


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 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 O 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 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 B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O B-CelestialObject I-CelestialObject O B-CelestialObject I-CelestialObject B-CelestialObject I-CelestialObject 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 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 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 B-CelestialObject I-CelestialObject O B-CelestialObject I-CelestialObject B-CelestialObject I-CelestialObject 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 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 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 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 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 B-CelestialObject I-CelestialObject O B-CelestialObject I-CelestialObject B-CelestialObject I-CelestialObject O O O O O O O O O O B-CelestialObject I-CelestialObject B-CelestialObject I-CelestialObject 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 O O O O O O O O O O O O B-CelestialObjectRegion I-CelestialObjectRegion I-CelestialObjectRegion O O O O O O O O O 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 O O O O O O O O O O O O O O O O O B-Citation O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject O B-CelestialObject I-CelestialObject O B-CelestialObject I-CelestialObject O B-CelestialObject I-CelestialObject B-CelestialObject I-CelestialObject O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation I-Citation O O O B-Telescope O O O O O O O O O B-CelestialObject

motion in declination direction. The proper motion of SV Cru, (μ αcos δ, Cep , μ δ, Cep ) = (− 6.43 ± 0.04, 1.31 ± 0.05) mas yr −1 , deviates 10.9σ in declination direction from the proper motion of Ruprecht 97, (μ αcos δ, cl , μ δ, cl ) = (− 6.49 ± 0.07, 0.44 ± 0.07). The mean proper motion is estimated by 23 OC members and the deviation of proper motion may be caused by the strict selection of OC members. Considering the good matching of other constraints, we regard SV Cru as an OC Cepheid. The distance and age of Ruprecht 97 are DM = 12.95 ± 0.53 mag and log ( t yr −1 ) = 8.13 ± 0.21, which are consistent with Cantat-Gaudin et al. ( 2020 ). 3.1.10 UBC 231 and WX Pup WX Pup is located at about 25.9 arcmin from the centre of UBC 231. Its proper motion is 2.2σ from the mean proper motion of UBC 231. The parameters of UBC 231 DM = 12.31 ± 0.17 mag and log ( t yr −1 ) = 7.87 ± 0.30 are in agreement with the reference values DM = 12.31 mag and log ( t yr −1 ) = 8.07, given by Cantat-Gaudin et al. ( 2020 ). 3.1.11 UBC 375 and V438 Cyg V438 Cyg is located at approximately 7.0 arcmin from the centre of UBC 375. V438 Cyg satisfies the constraints of the spatial position and parallax. The proper motion of V438 Cyg is (μ αcos δ, Cep , μ δ, Cep ) = (− 3.28 ± 0.07, −4.27 ± 0.07) mas yr −1 , which is 6.4σ away from the mean proper motion of UBC 375 (μ αcos δ, cl , μ δ, cl ) = (− 3.42 ± 0.07, −5.22 ± 0.11) mas yr −1 . The determined OC’s parameters are DM = 11.08 ± 0.21 mag and log ( t yr −1 ) = 8.03 ± 0.30, which are consistent with Cantat-Gaudin et al. ( 2020 ). Since only 32 OC members are adopted to estimate the mean proper motion, the deviation of proper motion may be attributed to the strict selection of OC members. Considering the good matching of other constraints, we still regard V438 Cyg as an OC Cepheid. 3.1.12 FSR 0158 and GQ Vul GQ Vul is located at approximately 2.4 arcmin from the centre of FSR 0158. GQ Vul agrees well with FSR 0158 in the proper motion (0.6σ) and parallax (0.1σ). This new OC Cepheid is ∼5.3 kpc from us and almost at the detection limit of our method based on proper motions and parallaxes. The far away OC Cepheids in the Milky Way’s disk are helpful to testify the probable bias in Cepheid’s PAR. 3.2 Issue about the known OC Cepheid In this subsection, we discuss known OC Cepheids not included in our catalogue. We compared with Turner ( 2010 ) who established an OC Cepheid catalogue based on optical data and found that 11 of the 24 OC Cepheids in that work are in our list. Nine of the other 13 Cepheids are in small clusters or star associations and they are not included in Cantat-Gaudin’s membership analysis catalogue. Three OC Cepheids have proper-motion differences larger than 10σ. The remaining OC Cepheid is BB Sgr in Collinder 394. 3.2.1 Collinder 394 and BB Sgr BB Sgr was treat as an OC Cepheid for a long time from Turner Pedreros ( 1985 ). However, based on Gaia , the membership is questionable. The proper motion of BB


O O O O O O B-Observatory O O O O O O O O O O O B-Telescope O 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 I-CelestialObjectRegion O 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 B-Identifier O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 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 O O O B-Wavelength I-Wavelength B-Wavelength 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

of [2] and [3], which show FERMI spectra because they provided better statistics. We note that the corresponding RHESSI spectra at those times showed very similar results. The X-ray spectra of background-subtracted data (black solid lines) are plotted together with the fitting results for the isothermal component (red dashed lines), and the nonthermal component (blue dashed lines). The background is represented by the gray solid line. The electron temperature, T , emission measure, EM, electron distribution index, δ , and cutoff energy, E C , as well as the chi-squared of the fitting, χ 2 , are listed. Figure 4. (a)–(d) EM-weighted temperature maps during the M1.3 flare. The dotted lines represent an artificial slit along the trajectory of the jet-like feature, along which we extracted the DEM evolution. (e) Time–distance plot for EM-weighted temperatures along the slit. (f) EM distribution for the position and times marked by the colored circles in panel (e). An animation of panels (a)–(d) is available. The video begins at 21:30:01 and ends at 22:14:49. The realtime duration of the animation is 15 s. (An animation of this figure is available.) In order to infer information on the triggering of the event and its magnetic topology, we initially studied the magnetic field of AR 11944 during the flare under study as well as a few days before. We note that due to the rapidly varying photospheric magnetic field of the AR prior to the occurrence of SOL2014-01-13T21:51M1.3, its closeness to the limb, and the small-scale nature of the event, the magnetograms did not provide us with useful information for the magnetic configuration on the date of the flare occurrence. 3.3. Causal Relationship of Precursors and Flare Features In addition to the bright loop and associated chromospheric brightenings observed for P2 and P3, the impulsive flare phase starts with the appearance of an additional loop (marked by the orange dashed line in the inset of Figure 2 (d)) rooted at a chromospheric brightening (labeled (3)). This was immediately followed by the generation of flare ribbons neighboring L1 and the enhanced emission at L2 as a consequence of nonthermal bremsstrahlung (Figure 2 (e)). No X-ray sources were found at the locations where the ribbons formed encompassing L1 (Figures 3 (d), (e)), possibly due to RHESSI ’s limited dynamic range. However, strong nonthermal emission during the impulsive flare phase was registered at L2, in the form of a compact hard x-ray (HXR) source and a flat nonthermal ( δ ∼ 4.2) component in the corresponding spectrum (Figure 3 (d)). This shows that L2 was the main compact flare region. During this phase, a jet-like feature emerged from L2 (see inset in Figure 2 (e) and Figures 4 (d), (e)), reaching its maximum altitude at . The EM distribution of a position along the jet-like feature before, during and after its occurrence is shown in Figure 4 (f) (see position and corresponding times marked in Figure 4 (e)), showing that it


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This work was supported in part by JSPS KAKENHI Grant Number 24540226.


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F.C. acknowledges support from Swedish National Science Council grant 637-2013-7261. A.L. acknowledges support from the ERC Advanced Grant “ FEEDBACK ”. R.H.I., M.A.P.T. and A.A. acknowledge support from the Spanish MINECO through grants AYA2012-38491-C02-02 and AYA2015-63939-C2-1-P.


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We thank the anonymous referee for many comments and suggestions that improved the manuscript. 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. This research has also made use of NASA's Astrophysics Data System (ADS), and the NASA/IPAC Infrared Science Archive, which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. This research made use of Montage, funded by the National Aeronautics and Space Administration's Earth Science Technology Office, Computation Technologies Project, under Cooperative Agreement Number NCC5-626 between NASA and the California Institute of Technology. Montage is maintained by the NASA/IPAC Infrared Science Archive. This publication makes use of data products from the WISE , which is a joint project of the University of California, Los Angeles, and the Jet Propulsion Laboratory/California Institute of Technology, funded by the National Aeronautics and Space Administration. The UKIRT is supported by NASA and operated under an agreement among the University of Hawaii, the University of Arizona, and Lockheed Martin Advanced Technology Center; operations are enabled through the cooperation of the Joint Astronomy Centre of the Science and Technology Facilities Council of the UK. When some of the data reported here were acquired, reported here were acquired, UKIRT was operated by the Joint Astronomy Centre on behalf of the Science and Technology Facilities Council of the UK. Some of the data presented in this paper were obtained from the MAST. STScI is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS5-26555. 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. Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the US Department of Energy, the National Aeronautics and Space Administration, the Japanese Monbukagakusho, the Max Planck Society, and the Higher Education Funding Council for England. The SDSS website is http://www.sdss.org/ . RC was supported by NASA Contract 1391817 and NASA ADAP grant 13-ADAP13-0155. MSY acknowledges support from the NASA ADAP grant NNX10AD64G.


O O 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 O B-Person I-Person O O O O O O O O O O O O O O O O O B-Person O O O O O O O O B-Organization I-Organization I-Organization I-Organization O O O O O B-Organization B-Person O O O O O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization O B-Organization O O O O O O O B-Organization O B-Grant I-Grant B-Organization B-Grant I-Grant I-Grant B-Organization B-Grant I-Grant I-Grant O B-Grant O B-Organization I-Organization O O O O O O O O B-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility

We thank Andrey Kravtsov, Surhud More, Phil Hopkins, Frank van den Bosch, Andrew Hearin, Erik Tollerud, Erwin Lau, and Benedikt Diemer for useful discussions and/or comments on an early draft. We also thank the reviewer for useful comments. A.R.W. acknowledges the hospitality and stimulating environment of the Aspen Center for Physics, supported in part by the NSF. A.R.W. also gratefully acknowledges support from the Moore Center for Theoretical Cosmology and Physics at Caltech. This work was supported in part by NSF grants AST-1412768 1009811, NASA ATP grant NNX11AE07G, NASA Chandra grants GO213004B and TM4-15007X, the Research Corporation, and by the facilities and staff of the Yale University Faculty of Arts and Sciences High Performance Computing Center.


B-Software 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O O O O O O O

GASP2D photometric decomposition of the mock images of the simulated galaxies is shown in Fig. 3 for the lenticular remnant resulting from the merger experiment gS0dE0o100 seen at an inclination . We listed the more relevant best-fitting structural parameters (i.e., the effective radius , Sérsic index , and axial ratio of the bulge, the scale-length and axial ratio of the disk, and the difference between the position angles of bulge and disk ) of the mock images of the simulated lenticular remnants seen at different inclinations in Table A.1 for the giant-giant galaxy mergers, Table A.2 for the giant-dwarf galaxy mergers, and Table A.3 for the giant S0-dwarf E0 galaxy mergers with different orbital parameters. Following Méndez-Abreu et al. (2017) , we adopted and as uncertainties on the axial ratio and position angle of both bulge and disk, respectively. 4.3. Bulge intrinsic shape of the simulated lenticular remnants seen at different inclinations We made use of our statistical method to retrieve from the mock images the probability distribution of the intrinsic axial ratios and of the bulges of the simulated lenticular remnants seen at different inclinations (, and ). We excluded from the analysis the edge-on configurations () because they do not allow us to constrain and due to the unknown orientation of the triaxial bulge in the disk plane. The probability distribution of and for the bulge of the simulated lenticular remnants resulting from the merger experiment gS0dE0o100 and seen at an inclination is shown as an example in Fig. 4 . The face-on configurations of the simulated lenticular remnants ( and ) provided the same result in terms of the probability distribution of and for all the simulated bulges, as expected if no observational and theoretical bias affected the adopted statistical method. Furthermore, we confirmed that the tightest constraints for are given when galaxies are seen face on, whereas remains unconstrained for these galaxies. We also found consistent probability distributions of and for the same simulated lenticular remnant seen at intermediate inclinations (, , and ). Because of this, we were confident of having correctly recovered the bulge intrinsic shape and suggested us to set a limit on the inclination of real galaxies () to robustly apply our statistical method to their bulges (see Sect. 2 ). In general, the probability distribution of and is tighter at with respect to or . Therefore, we considered this inclination as the ideal viewing angle for future analyses of the bulge intrinsic shape in real galaxies. The probability distributions of and of the bulge of the simulated lenticular remnant resulting from the merger experiment gS0dE0o100 seen at different inclinations are shown in Fig. 5 , while the remaining galaxies are in Fig. B.1 . The values of and derived for all the simulated lenticular remnants seen at different inclinations are listed in Table 1 . We realized that the accuracy of the photometric


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This paper is dedicated to the memory of Janusz Kaluzny, founder and long-time leader of CASE, who prematurely passed away in March 2015. We thank an anonymous referee for very insightful remarks that significantly improved the final shape of the paper. A.A.P. acknowledges partial support from the Polish National Science Center (NCN) grant 2015/17/B/ST9/02082. M.R. was supported by the NCN grant DEC-2012/05/B/ST9/03931.


O O O O O O O O O O O O O O O O O O O O O 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-Telescope 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 O O O O O O O O O O O O O O O O O O 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-Telescope 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 O O B-Telescope O O O O O O O O O O O O B-Telescope 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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-Telescope O O O O O O O O O O O O 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-Telescope O O O O O O O O O O O O O B-Observatory B-Telescope B-Wavelength I-Wavelength 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 B-Telescope O O O O O O O O O O O O O O O O O O O O O O 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-Telescope O O O O O O O O O O O O O O O O O B-Formula I-Formula I-Formula 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 B-Telescope O O O O O O 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-Telescope O O O O B-Telescope O O O O O O B-Formula I-Formula I-Formula 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 B-Telescope O O O

caused by the erupting filament plasma and the bright ribbon fronts on the outer border of the growing dimming region that come into the pixel’s FOV. Figure 7. Time evolution of the relative density changes, i.e., actual density divided by its pre-event value, derived from AIA DEM maps for pixels corresponding to the EIS spectra in Figure 6 . The segments in red indicate the phase of the impulsive density drop: its start is defined by the data steeply decreasing below 1, and its end is calculated from the time derivative of the curve (defined by the first positive increase). The green vertical lines indicate the time at which the corresponding pixel was covered by the EIS slit. To relate the evolution of the densities reconstructed from the AIA DEM maps (Figure 7 ) in the 4 pixels indicated in Figure 5 with the corresponding spectra and flows determined from the EIS scan (Figure 6 ), we need to relate the dynamics observed in AIA with the time at which the EIS slit covered the particular pixels. We first discuss the behavior of the top three pixels plotted in Figure 7 , and thereafter the bottom-most one, which is related to different structures as we will show. In the density evolution retrieved from the DEM maps (Figure 7 ), we can see that the sharp density drop due to the dimming occurs first in the leftmost (easternmost) pixel, and thereafter in the two neighboring pixels to the right. This is consistent with the expansion of the dimming region toward the west, i.e., from left to right (see Figure 8 and the associated movie). On the other hand, EIS was in scanning mode, and we may expect to see the largest upflow components in those pixels that have been activated as dimming pixels (i.e., the associated fields were “opened”) closest in time to when the EIS slit was mapping them. Figure 8. Evolution of the western dimming region in SDO /AIA 211 Å filtergrams. We show the same zoomed FOV as for the EIS maps in Figure 5 . The yellow boxes mark pixels for which we show the EIS spectral profiles in Figure 6 . The still figure shows the event at two time steps (annotated in each panel). The animated figure shows the event evolution from 2012 September 27, 23:50 UT to September 30, 01:16 UT. (An animation of this figure is available.) The EIS scan was from west to east, and the strongest upflow component is observed at the pixel ( x = 563′′, y = 44′′), which was covered by the slit at 00:04:16 UT. This is the pixel where we find that the total amount of upflowing plasma is larger than the stationary plasma component. The AIA density evolution shows that in this pixel, the strong drop in the density, indicative of the impulsive dimming onset, starts exactly at that time when it was mapped by the EIS slit. In the neighboring EIS pixel to the left, centered at x = 560′′, the impulsive density drop occurred during 23:58 to 00:06 UT, and the EIS slit covered it during this time range at 00:04:58 UT. This finding is consistent with EIS still showing a