O O O O O B-Telescope I-Telescope O O B-Organization B-Organization I-Organization I-Organization O B-Observatory O O O O O O B-Database I-Database I-Database B-Database O O O O O B-Organization I-Organization I-Organization B-Organization I-Organization I-Organization I-Organization O O O O B-Organization I-Organization I-Organization I-Organization I-Organization B-Person O O O O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization O O B-Organization I-Organization I-Organization I-Organization I-Organization B-Person O O O B-Grant I-Grant I-Grant I-Grant O O O B-Observatory O O O O O O O B-Person I-Person O O O O O O B-Person I-Person O O O O O O O O O O O O O O O O O O O O O O O O O

Based on observations with the 100-m telescope of the MPIfR (Max-Planck-Institut für Radioastronomie) at Effelsberg. This research 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. A.P. is a member of the International Max Planck Research School (IMPRS) for Astronomy and Astrophysics at the Universities of Bonn and Cologne. C.C.-G. acknowledges support by UNAM-DGAPA-PAPIIT grant number IA101214. We thank the Effelsberg operators for the useful help. We acknowledge Roberto Fanti for useful discussions and suggestions and Jochen Heidt for providing us with spectral redshifts of some of the sources. We thank the anonymous referee for careful revising of the paper and useful comments.


O O O O O O O O O O O O O O O O O O B-Grant O O B-Organization I-Organization I-Organization O O O O O O O B-Organization I-Organization I-Organization B-Organization I-Organization O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization O O O O O O O O O O O B-Software I-Software

The results of the research presented in this paper have been performed within the scope of the contract ASI/2007/I/082/06/0 with the Italian Space Agency. The authors are particularly thankful to the University of Rome ‘La Sapienza’ and the University of Bologna-Forlì Radio Science teams, for their precious support in the validation and testing of the orbit14 software.


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 B-Telescope I-Telescope I-Telescope I-Telescope I-Telescope I-Telescope B-Telescope B-Telescope O O O O O B-Organization I-Organization I-Organization I-Organization I-Organization B-Organization O B-Organization O B-Grant O O B-Organization I-Organization I-Organization B-Organization O B-Organization O B-Grant I-Grant I-Grant O O O O O O B-Organization I-Organization I-Organization B-Organization I-Organization I-Organization I-Organization O O O O B-Organization I-Organization I-Organization I-Organization I-Organization B-Organization I-Organization O O

Acknowledgements We thank an anonymous referee for several very useful comments that improved our paper. We are grateful to the SOFIA engineering and operations teams for their support which enabled the observations presented here. The research reported here is based largely on observations made with the NASA/DLR Stratospheric Observatory for Infrared Astronomy (SOFIA). SOFIA is jointly operated by the Universities Space Research Association, Inc. (USRA), under NASA contract NAS2-97001, and the Deutsches SOFIA Institut (DSI) under DLR contract 50 OK 0901. This work was performed at the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration. USA Government sponsorship acknowledged.


O O B-Person I-Person I-Person I-Person O B-Person I-Person O O O O O O O O O O O O O O O B-Grant

We thank Victor De la Luz and Ernesto Aguilar for reading the manuscript and providing important comments. This research is partly supported by grant CIC-UMSNH-4.9.


O O O O O B-Person I-Person O O O O O O O O O O O B-Person I-Person B-Person I-Person O B-Person I-Person O O O O O O O O B-Instrument O O B-Person O O O B-Grant I-Grant I-Grant I-Grant O B-Grant I-Grant I-Grant I-Grant I-Grant I-Grant I-Grant O O B-Telescope O B-Person O O O O B-Telescope O

We thank the referee and Karin Muglach for helpful comments that led to improvements in the paper, and Brian Dennis, Richard Schwartz, and Kim Tolbert for their continuing help and support with the RHESSI data analysis. G.H. was supported by NASA Heliophysics Guest Investigator and Living with a Star TR T Grants and the RHESSI project. A.F. was supported by the RHESSI project.


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 B-Telescope 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 B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O B-Telescope I-Telescope I-Telescope 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 B-CelestialRegion I-CelestialRegion I-CelestialRegion I-CelestialRegion 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-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 O B-CelestialRegion I-CelestialRegion I-CelestialRegion O O O O O O O O O B-Citation I-Citation O B-Formula I-Formula I-Formula I-Formula O O O O O B-Formula I-Formula I-Formula I-Formula 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 B-Collaboration 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 O O O O O O O O O O O O O O B-Wavelength I-Wavelength B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O B-Dataset I-Dataset I-Dataset I-Dataset I-Dataset I-Dataset B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject O O O B-Collaboration I-Collaboration I-Collaboration I-Collaboration I-Collaboration I-Collaboration 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 B-Wavelength O O O B-CelestialObject I-CelestialObject B-CelestialObject B-CelestialObject I-CelestialObject O B-CelestialObject I-CelestialObject 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 B-Wavelength O O B-Wavelength I-Wavelength O O O O O O O O O O O O O O O B-Dataset I-Dataset 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 B-Telescope O O O B-Telescope O O O B-Collaboration 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O

the FD energy scale (determined calorimetrically; Abu-Zayyad et al. 2013b ; Tsunesada et al. 2017 ). The systematic uncertainty on the TA energy scale is 21% (Abbasi et al. 2016 ) and its energy and angular resolutions are 15%–20% and 10–15, respectively, depending on the event geometry and energy (Abbasi et al. 2014 ). In this Letter we use data collected by the TA SD array in a 9 yr period from 2008 May to 2017 May with reconstructed energies above 43 EeV, zenith angles less than 55°, and declinations δ > −10° using the same quality cuts as in Abbasi et al. ( 2014 ). This data set comprises 284 events. We neglect the finite angular and energy resolution of TA events, and consider the detector fully efficient, i.e. with a flat response for all showers with energies and zenith angles in the considered range, so that its directional exposure ω TA equals the geometrical one for δ > −10°, which varies with declination but not with right ascension (Sommers 2001 ):where φ TA = +393 is the detector latitude and θ m = 55° is the maximum zenith angle accepted. The energy threshold of E min = 43 EeV used in this analysis corresponds to the Auger energy threshold of 39 EeV, at which the most significant correlation with SBG was found. Here we took into account the 10.4% difference between the energy scales of the two experiments as estimated by a comparison of energy spectra around 5 EeV (Verzi et al. 2017 ; Abu-Zayyad et al. 2018 ). 2.2. Source Catalog Following the Auger analysis (Aab et al. 2018 ), we select the candidate sources from a sample of 63 SBGs outside the Local Group compiled by the Fermi -Large Area Telescope (LAT) Collaboration (Ackermann et al. 2012 ) for the gamma-ray emission search. 36 36 Only four of those objects were actually successfully detected in gamma-rays in that work: NGC 253, M82, NGC 4945, and NGC 1068. Imposing the cut of flux greater than 0.3 Jy at 1.4 GHz leaves 23 objects in the catalog of candidate sources. Their UHECR fluxes were assumed to be proportional to their radio fluxes at 1.4 GHz. These objects are listed in Table 1 . Table 1 Selected Source Candidates from the SBG Catalog used in this Analysis (the same as in Aab et al. 2018 ) Note. The last column shows the relative source contribution weighted with the TA directional exposure ω TA . In the Auger analysis, the effect of energy losses by UHECRs during their propagation was found to be negligible in the SBG model, as most of the anisotropic flux originates from sources within a few Mpc; in this Letter, we neglected the losses for simplicity. 2.3. Test Statistic and Flux Model Let be the unit vector representing a direction in the sky, pointing away from the observer. Given two flux models describing a null hypothesis and an alternative hypothesis, respectively, and the directional exposure of an experiment, the test statistic (hereinafter ) is defined as twice the log-likelihood ratioand being the reconstructed arrival direction of the i -th observed


O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 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 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

in Eq. (12) , each locally reflected vector I , and hence each local reflection matrix R 1 , is rotated such that the reference plane is no longer the local meridian plane, but the planetary scattering plane. The local rotation angle β depends on the local viewing angle θ and the location of surface area d O on the planet. The geometric albedo A G of the planet with radius r at distance d is given by where π F 0 (0°, λ ) is the reflected total flux at wavelength λ and phase angle α equal to 0°. In P Y M IE DAP, the integration in Eq. (12) is replaced by a summation over locally reflected Stokes vectors. In order to do so, we divide the planetary disk on the sky in pixels, and compute the locally reflected Stokes vector at the center of each pixel. A pixel contributes to the disk signal when its center is within the disk radius. The integration is then given by with N the number of illuminated and visible pixels on the planetary disk, and with subscript n indicating that μ, μ 0 , φ − φ 0 , and β depend on the location of the pixel on the planet. In addition, μ 0 , φ − φ 0 , and β at a given location of the pixel on the planet depend on the planet phase angle α ( Appendix B provides relations that can be used to derive these angles). In the summation, d O n refers to the area of the pixel as measured on the surface of the planet. Although not explicitly indicated in Eqs. (14) and (12) , R 1 will usually also depend on the location (of a pixel) on the planet. Typical horizontally inhomogeneities would be the surface coverage and altitude, and the atmospheric composition and structure. The obvious horizontal variations on Earth are of course the oceans and the continents, and in the atmosphere, the clouds. Horizontal inhomogeneities can be taken into account by using different local reflection vectors R 1 across the planet (in that case, R 1 in Eq. (14) would include a subscript n ). Given a model planet and a planetary phase angle α , the steps to evaluate Eq. (14) are the following: Divide the planet into pixels small enough to be able to assume that the planet properties across each pixel are horizontally homogeneous, and to be able to follow the limb and the terminator of the planet. Calculate for (the center of) each pixel the angles μ (i.e., cos θ ), φ , and the rotational angle β . These angles are independent of the location of the sun or star with respect to the planet. Calculate for the given phase angle , and for (the center of) each pixel, μ 0 (i.e., cos θ 0 ) and φ 0 . These angles depend on the location of the sun or star with respect to the planet. Calculate for (the center of) each pixel the column vector R 1 of the locally reflected light using the appropriate database file (see Sect. 5 ). Perform the summation described by Eq. (14) . The pixels can be defined on the planet, for example, using a latitude and longitude grid, in which case μ d O , the projected area of the pixel (see Eq. (12) ), that is, the pixel area as “seen”


O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O

choose a unique for each particle upon emission. We assume the field changes on a length scale , where λ is a model parameter, and at the beginning of each time step we calculate the timescale for the particle to experience a significant change in :where z is a random number between 0 and 1. If is less than the time for a particle to leave a zone or end the grid time step, the magnetic “scattering” action is selected, the particle position and energy are updated, and a new field direction is chosen at random. The re-orientation preserves energy and the magnitude of the momentum in the cmf. This is reasonable because these discrete scatters are standing in for smoother and more gradual field gradients. However, it also leads, on average, to energy losses in the center-of-explosion frame, which decrease thermalization efficiency. This can be thought of as the particles transferring their energy to the bulk kinetic energy of the ejecta rather than the thermal background. It is a small effect. The propagation of a particle along field lines depends on total velocity and Particles are emitted on the grid with a random direction defined by the unit vector , so μ is . In the case of radial fields, μ evolves to preserve the magnetic adiabatic invariant . This evolution encourages outward motion. Inward-moving particles encounter ever-stronger fields, and decrease their inward velocity in response, eventually “mirroring” off the field and beginning to travel outward. Particles streaming outward move into weaker magnetic fields, so their pitch angle increases and they stream out even faster. We include magnetic beaming and mirroring only for radial fields, where changes to the pitch angle facilitate particle escape and can be meaningful for thermalization. Particles in toroidal fields are confined to one region of velocity space; changes in their pitch angle cannot promote outward motion and so cannot affect thermalization. While particles in random fields could experience beaming and mirroring, they may also undergo pitch angle scattering—interactions with small fluctuations in the field that can change pitch angle, and would act to isotropize μ , counter to beaming and mirroring. We assume the two effects balance out, and hold pitch angle constant for particles in random and toroidal fields. A.3.2. Fission Fragment Transport As discussed in Section 2.2 , at late times, fission fragments may have Larmor radii comparable to the magnetic field coherence length. This could affect the transport, especially for disordered fields, where field lines in close proximity may have very different orientations. In such a system, the guiding center approximation breaks down, and the motion of the particle must be resolved. We estimate the importance of this effect by modeling fission fragment transport in a tangled field as a random walk of the fragment itself (as opposed to its guiding center). In this simple scheme, fission fragments travel in straight lines and


O O O O O O O O O O O O O O O O O O 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 B-Formula I-Formula I-Formula 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 B-CelestialObject B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O B-Telescope O B-Telescope 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 B-CelestialObject O O B-Telescope I-Telescope I-Telescope O B-Telescope O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O B-Wavelength O O B-Telescope B-CelestialRegion I-CelestialRegion O O O O O O O O B-CelestialRegion I-CelestialRegion O O O O O O O O O B-Survey O O O B-Telescope O O O O O O O O O B-Telescope O O O O O B-Dataset O O O O O B-CelestialRegion I-CelestialRegion O O O B-Wavelength O O O O O O B-Telescope O B-Telescope O O O O B-Survey O O O O O O O O O O O O O O O O O B-CelestialObject O O B-Wavelength O O O O O O O B-CelestialRegion I-CelestialRegion O O O O O O O O O O O B-Telescope O O O O O B-Survey O O O O O O O O O O O O O O B-Tag O B-Telescope 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 B-Software 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-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O B-Telescope B-CelestialObject O O B-Telescope O O O O O O O O O B-CelestialObject O O O O O B-Wavelength I-Wavelength O O B-Wavelength I-Wavelength B-Wavelength I-Wavelength O B-Wavelength I-Wavelength 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 B-Tag O B-Telescope O O O O O O O O O O O O B-Dataset O O B-Telescope O O O O B-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O B-Telescope O B-Proposal I-Proposal O B-Person O O O B-CelestialRegion I-CelestialRegion O O O O O O O B-Dataset O O O O O O B-Telescope O O O O O O O O O O B-CelestialRegion I-CelestialRegion O O O O B-Telescope O O O O O O O O B-Dataset O O O O O O O O

the SED fitting used to determine spectral types for companion stars in HMXB candidates. In Section 5 we discuss our results in the context of previous studies, and in Section 6 we provide a brief summary of our results. We assume a Galactic column density, , and a photon index, Γ = 1.7 (Stiele et al. 2011 ), to convert count rates into absorbed energy flux. We assume a distance of 776 kpc to M31 (Dalcanton et al. 2012 ) for luminosity calculations. 2. Data In this study we employ data from NuSTAR , Chandra , and HST . We now describe each data set in more detail below. For an overview of the area observed by each telescope, see Figure 1 . Figure 1. Near-UV image of M31 from the Galaxy Evolution Explorer ( GALEX ; left; Gil de Paz et al. 2009 ) and three-color X-ray image of Chandra Field A (right), see Section 2.2 for more information on Field A data reduction. Magenta regions outline the area observed by PHAT. Green outlines the NuSTAR observed region, and cyan outlines the area observed by Chandra , with solid lines indicating ChandraPHAT observations and dashed lines indicating Field A observations. In the UV image, the 64 sources observed by NuSTAR and Chandra that fall within the PHAT footprint are marked with white crosses, and a 5′ circle outlines the densest inner bulge region of M31. In the X-ray image, all sources detected by wavdetect within Field A are marked with white ellipses (including sources that do not match NuSTAR sources or are outside the PHAT footprint, and thus are not presented in our sample in Table 1 ). 2.1. <italic>NuSTAR</italic> Data NuSTAR source catalogs and source classifications come from Wik et al. (2018, in preparation). Observations were taken between February and October 2015 covering the area outlined in green in Figure 1 with an average exposure time of ∼400 ks, and they were reduced using the nupipeline software. Sources previously observed with Chandra were used for astrometric alignment. For detailed information on data reduction, source detection, and source classification, see Wik et al. (2018, in preparation), who present the entire NuSTAR M31 survey. The NuSTAR observations cover the nucleus and inner disk regions of M31 at an energy range of 4–25 keV using the 4–6 keV, 6–12 keV, and 12–25 keV energy bands. The completeness of the NuSTAR observations starts to fall off at a luminosity of ∼3 × 10 36 erg s −1 and reaches zero at ∼2 × 10 36 erg s −1 . 2.2. <italic>Chandra</italic> Data Chandra data used in this study are comprised of two data sets: the ChandraPHAT data, a Chandra Large Project survey by B. Williams et al. (2018, submitted), and one additional Chandra field (obsid 18046, P.I. Hornschemeier), hereby referenced as Field A, that was reduced for this paper. The ChandraPHAT data set is comprised of seven Chandra pointings, each with a depth of about 50 ks. The Field A data consist of one Chandra pointing with a depth of 25 ks. The ChandraPHAT observations were taken in 2015 October and the


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

use for power beaming are in the millimeter band, so are outside the microwave range the Allen Telescope Array observed. Similarly, power beaming using near-IR frequencies would be undetectable by the Boquete observatory. Therefore, the Harp et al. ( 2016 ) and Schuetz et al. ( 2016 ) limited observation times and wavelength coverage are not sufficient to produce firm conclusions on power beaming from KIC 8462852. Most applications would be seen in the radio—if transmissions were oriented in our direction at the proper time and at the frequencies observed—but are not. More extensive observations should be made in more systematic studies of power beaming leakage, including observing at higher radio frequencies and for longer times at visible wavelengths. 5. CONCLUSIONS AND IMPLICATIONS We have listed several classes of power beaming applications/missions, quantifying the principal observable parameters. Applying this reasoning to the recent observations of KIC 8462852, we conclude that if power beaming were in use at the time observed, generally pointed in our direction, and at frequencies between 1 and 10 GHz, that most power beaming applications would have been detectable. The non-detection provides a weak (owing to the caveat of the probabilities listed) rejection of the popular hypothesis that the system is inhabited by an advanced spacefaring extraterrestrial civilization. As discussed above, the beaming power levels are high and transient and easily dwarf any ETI civilization’s diffuse leakage to space (Sullivan et al. 1978 ). Power beaming described here is larger than that necessary for beaming systems for communication: for a 1000 ly range beacon (Benford et al. 2010 ). SETI programs could explore a different part of parameter space by observations suitable to finding leakage from power beams. Such beams would be visible over large interstellar distances. This implies a new approach to the SETI search: instead of focusing on narrowband beacon transmissions generated by another civilization, look for more powerful beams with much wider bandwidth. This requires a new approach for their discovery by telescopes on Earth. Past SETI observations have been in the 1–10 GHz microwave band. For our atmosphere, future observations should look in bands where with lower oxygen and water vapor allow transmission: windows at 35 GHz, 70–115 GHz, 130–170 GHz, and 200–320 GHz. And, of course, such transient sources require longer observing times. A promising avenue is to revisit past observations of transient events, of which there are many, to look for patterns and identify as possible regions of the sky to emphasize. ETIs would know their power beams could be observed. They could put a message on the power beam and broadcast it for our receipt at little additional energy or cost. By observing leakage from power beams, we may well find a message embedded on the beam. That message may use optimized power-efficient designs such as spread spectrum and


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LS is grateful for the support from SNSF grant SCOPES IZ73Z0 152581, GNSF grant FR/339/6-350/14. This work was supported in part by DOE grant DEFG 03-99EP41093. NASA's Astrophysics Data System Bibliographic Service and the arXiv e-print service were used for this work. Additionally, we wish to acknowledge gnuplot , a free open-source plotting utility which was used to create all of our figures and perform the function fitting central to this work. We wish to acknowledge the many useful conversations with Bharat Ratra while performing this work. Lastly, we would like to thank the referee for their useful comments.


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

determines the age, metallicity, and dust attenuation of the stellar populations. Consequently, the light-weighted ages and metallicities derived above are driven by spectral features that are less degenerate than the spectral shape alone (see Figure 1 and Table 1 ). Table 1 Prominent Stellar Features Note. The first column gives the feature and wavelength. The second and third columns give the age range that the features are found in, in a 0.4 Z ⊙ bpass and starburst 99 model, respectively. The two main types of FUV stellar features are strong, broad stellar wind P-Cygni features and weak stellar photospheric absorption lines. Both types of features can be contaminated by neighboring ISM absorption and require high S/N and moderate spectral resolutions to resolve. In the following two subsections, we discuss both types of spectral features individually, and we illustrate how the individual features relate to the inferred stellar population ages and metallicities. The purpose of these subsections is not to advocate for determining the stellar ages and metallicities using single features, but rather to demonstrate that the stellar properties inferred from the full spectral fits are entirely consistent with the trends of stellar spectral features. 4.1. Stellar Wind Features The most notable stellar features in the FUV are the broad blueshifted absorption and redshifted emission profiles (called P-Cygni profiles; orange labels in Figure 1 ). These P-Cygni profiles arise from strong winds that are radiatively driven off of stellar photospheres (Castor et al. 1975 ; Lamers Cassinelli 1999 ). The terminal velocity, ionization structure, and mass-outflow rates sensitively depend on the stellar luminosity (Castor et al. 1975 ; Lamers Leitherer 1993 ; Lamers et al. 1995 ; Leitherer et al. 1995 ; Puls et al. 1996 ; Kudritzki Puls 2000 ). The terminal velocity describes the maximal velocity extent of the absorption component, while the mass-loss rate determines the depth of the profile. In turn, these establish the shape of the P-Cygni absorption and emission. Whether a given ion is observed as a P-Cygni profile in the wind depends on the ionization structure of the stellar wind. The peak ionization stages for stellar winds typically are the C v , N iv , O iv , and Si v states (Lamers Cassinelli 1999 ), none of which have resonant transitions in the rest-frame FUV. Alternatively, the presence of adjacent ionization stages with P-Cygni profiles (N v , C iv , or Si iv ) provides information on the stellar temperature of the most luminous stars and, by inference, the stellar population age. The age of a stellar population can be inferred from the strength of the observed P-Cygni transitions: O v and N v have the highest ionization states and are strongest in stars with lifetimes of 2–3 Myr, while C iv and Si iv are lower ionization and peak in stars with lifetimes near 5 Myr. Stellar temperature, or age, is not the sole determinant of the stellar wind


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

are required to grid the raw data generated by FAST in real-time with the existing methods, which could greatly increase the cost of computation for this survey. In the past decade, various studies on optimization of gridding performance have been carried out. For example, Winkel et al. ( 2016b ) have presented a versatile gridding module ( Cygrid : A fast Cython-powered convolution-based gridding module for python ) for radio astronomical data reduction, featuring a single-level lookup table based on the c ++ standard template library (STL) for gridding parallelization, as well as the Hierarchical Equal Area Latitude Pixelation ( healpix) -based spherical tessellation for fast neighbour searching. And the performance of the Cygrid code can be significantly improved when using multiple CPU cores. However, utilizing the STL vector, which features dynamical arrangements of storage allocations with a relatively small initial storage volume, to construct a lookup table for input raw data is not always practical. For example, if the raw samples are unevenly distributed, with more data points in the central regions and less on all sides, the required storage needs to be expanded in the middle part of each healpix -tessellated ring frequently. In this case, the STL vectors have to make massive data copies of the data frequently, which would consume broad bandwidth, thus resulting in degradation of gridding performance. In addition, the main workload of gridding can be considered as a combination of a large number of single-instruction, multidata stream operations. For such tasks, the high parallelization level, as well as throughput of GPUs brings a possibility for better performance. Thus, Romein ( 2012 ) has developed a work-distribution scheme for radio data gridding with GPUs, which could significantly reduce the memory access time of computing device, and map observed samples on to a grid with high efficiency. Merry ( 2016 ) has further optimized this algorithm with thread coarsening strategy, making one thread handle multiple samples simultaneously. However, the method proposed by Romein ( 2012 ) and Merry ( 2016 ) has been deeply customized for radio telescope arrays, strongly relying on the spatial coherence of the interference matrix, thus cannot be easily adapted for single-dish telescopes. Accordingly, in order to meet the requirements of large data volume from the latest single-dish radio telescopes, and to overcome the shortcomings of GPU-based gridding methods mainly designed for radio telescope array, inspired by Winkel et al. ( 2016b ), and based upon our previous work (Luo et al. 2018 ), here we present a convolutional gridding framework with hybrid computing environments, HCGrid , for radio astronomy, which can greatly improve the performance of the gridding process. Key features of our framework can be summarized as follows: We have designed a gridding module with high performance for large single-dish radio telescopes. We


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 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 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 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 B-Citation I-Citation I-Citation 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 B-CelestialObject I-CelestialObject O O O O B-Telescope 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-Model 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 B-Model I-Model O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject 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 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-Telescope O O O O O O B-Wavelength O B-Wavelength O O O B-Telescope B-Telescope O B-Telescope O O O O B-Model O O O O B-Model 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 B-Citation I-Citation I-Citation I-Citation O O O O O O O O B-CelestialObject I-CelestialObject

for a triangle of interferometric baselines: For a symmetrical Gaussian intensity distributions on the sky, the visibility phase is expected to be zero and so is the closure phase. Any deviation from a zero closure phase will indicate the source deviation from a Gaussian or point-like structure, and this observable can in principle be used to compare the model and observed emission shape without reconstructing the radio maps. We have calculated the theoretical visibility closure phases. For the model with , which shows the crescent on the N side of the BH (see Fig. 8 ), the closure phases are positive; ,19,11,11, where the four values correspond to different time moments of the observation. The evolution is caused by the rotation of the Earth, and it is probing slightly different values. A typical observation duration is two to three hours. This is about three times shorter than the dynamical time scale of the source (8.5 h) with its BH mass, . For , for which the crescent is on the S side of the BH (Fig. 8 ), the closure phases are negative: , , , . In both cases, the values are consistent with the observed value: (Akiyama, priv. comm.). In summary, for the fiducial model RH100 (for both viewing angles of and ), visibility amplitudes and closure phases are roughly consistent with the preliminary observations of the M 87 core obtained by the EHT (visibility amplitudes and closure phases on a single VLBI triangle). Fig. 9 Theoretical visibility amplitudes along E-W ( left ) and N-S ( right ) baselines, computed for model RH100 with (solid lines) and 160(dashed lines). 4. Discussions Deriving the appearance of a jet in the direct vicinity of a SMBH is not straightforward. The jet formation mechanism, as well as particle acceleration in jets, is generally not understood well. Moreover, one has to take spacetime curvature into account, which affects the plasma dynamics and light propagation. Using GRMHD simulations of a weakly magnetized accretion flow, a jet appears naturally, and we calculate the appearance of the M 87 jet base at radio and millimeter wavelengths. For the electron heating, we assume that the electrons are weakly coupled to protons in the accretion disk and strongly coupled in the jet – a simple, but crucial concept that we have already used successfully to explain the appearance of the SMBH in the center of the Milky Way. Below we discuss our results in a context of observational constraints. We also discuss the model limitations. Fig. 10 Comparisons of the visibility amplitudes from our models and those from the EHT observations. The visibility amplitudes, computed at mm ( GHz), along three baselines (SMT-CARMA, SMT-HAWAII, and CARMA-HAWAII) are shown for model RH100 with (solid lines) and 160 (dashed lines). The distance between the VLBI stations is expressed in units of ( km). The observed visibility amplitudes (points) and the tracks are taken from Doeleman et al. (2012) . 4.1. Mass-accretion rate The accretion rate onto M 87


O 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 B-Person O O O B-Person I-Person I-Person B-Person I-Person B-Person I-Person B-Person I-Person I-Person O B-Person I-Person O O O O O O O O O B-Grant I-Grant I-Grant I-Grant O O B-Organization I-Organization I-Organization I-Organization B-Organization O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization O O O O O O O O O B-Telescope I-Telescope I-Telescope O O O O O B-Archive I-Archive O O B-Organization I-Organization I-Organization O O O O O B-Telescope O O O O O O B-Organization I-Organization I-Organization I-Organization O O O B-Organization I-Organization I-Organization O O O B-Location B-Location B-Location B-Location B-Location O O B-Organization I-Organization I-Organization

The authors are grateful to an anonymous referee for a careful reading of the manuscript and several stimulating comments that helped to improve this work. A.F.L. and S.M. are grateful to Drs. J. Bouvier, A. Frasca, A. Klutsch, A. C. Lanzafame, and E. Moraux for interesting discussions. The authors acknowledge support through the PRIN-INAF 2012 funding scheme of the National Institute for Astrophysics (INAF) of the Italian Ministry of Education, University and Research. This work is based on data obtained by the CoRoT Space Experiment and made available through the CoRoT Archive at the IAS Operation Center whose availability is gratefully acknowledged. CoRoT is a space mission of the French Space Agency CNES in association with French CNRS Laboratories and participation of Austria, Belgium, Brazil, Germany, Spain, and the European Space Agency.


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

between 323 and 608 MHz (black solid-line). Radio flux density at 323 MHz across the source extension (red dashed-line). The data-points correspond to the regions displayed in Fig. 1d . Error-bars on the -axes correspond to the region size. Table 2 Galaxy cluster properties. Notes. (a) Coordinates of the BCG. (b) Rescaled with spectral index from 323 MHz. (c) k-corrected with spectral index 0.7. 4. Discussion A prediction from simple DSA theory is a relation between the Mach number of the shock and the injection spectral index of the accelerated electrons: Therefore, a spectral index implies an unphysical Mach number. In the case of this specific relic the injection spectral index appears extremely flat. If we assume the outer region as representative for the injection index we find: . This value seems to point towards a stronger Mach number shock than what is usually found for radio relics. However, we note that Eq. ( 2 )has been shown not to hold in at least the well-studied case of the “toothbrush” cluster ( van Weeren et al. 2016 ). Also the integrated spectral index (extracted from -harmonised images at 323 and 608 MHz) appears quite flat () with respect to the average spectral index of radio relics (; Feretti et al. 2012 ). As already noted for another radio relic with a flat integrated spectral index (Abell 2256; van Weeren et al. 2012b ; Trasatti et al. 2015 ), if the stationary shock conditions do not hold. As a consequence, a possible way to reconcile the flat spectrum with shock acceleration is that the relic must have been produced very recently (0.1 Gyr ago). This means that the time for energy losses is longer than the age of the shock at some observed energies and that the integrated spectrum of the relic will still evolve, getting steeper. We note that using the total flux density at 148, 323, 608, and 1400 MHz ( mJy; from NVSS), we obtain a steeper integrated spectral index of . However, given the different -coverages, these maps can be sensitive to different spatial scales. The X-ray luminosity extracted from ROSAT data is erg/s. From this luminosity, we derived the associated cluster mass: ( Pratt et al. 2009 ). The cluster is amongst the least massive clusters hosting a radio relic discovered so far (see Fig. 7b in de Gasperin et al. 2014 ). Only two other less-massive systems are known to host radio relics: Abell 3365 (, , from X-ray; van Weeren et al. 2011 ) and Abell 3376 (, , from SZ; Kale et al. 2012 ). Both galaxy clusters host a double radio-relic system. In de Gasperin et al. (2014) , we discussed the presence of a scaling relation between radio-relic power and hosting-cluster mass. Given that relation, Abell 3527-bis is almost an order of magnitude over-luminous for the hosting cluster mass. Nonetheless, Abell 3527-bis provides a rare example of low-mass cluster hosting a radio relic at . Nuza et al. (2012) investigated the predicted abundance of radio relics in current and future radio surveys. They estimated 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 O O O O O O O O O O O O B-Software O O O B-Software O O O O O O O O O O O O O O O O B-Software 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 B-Tag O O O B-Software B-Software 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 B-Citation I-Citation O B-Citation I-Citation O B-Citation I-Citation 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 O O O 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 I-Formula I-Formula

panel: p rich ). For the helium burning problem, discussed in this study, we find global agreement between the two solvers, whereas there are larger differences in the small-scale features. Figure 1. Density colour map of the in initial state for the floating torus. Figure 2. Density colour map of the floating torus at time 0.2 sec (left-hand panel: vulcan ; right-hand panel: rich) . Figure 3. Density colour map of the floating torus at time 0.3 sec (left-hand panel: vulcan ; right-hand panel: rich) . In particular, although the two solvers successfully suppress numerical transition to detonation in very different ways, the final conclusions are the same. Finally, some caution must be taken due to the issue of possible turbulence. We recognize that our simulations cannot follow fully developed turbulence since the spatial resolution is too coarse and also the missing third dimension limits the correct cascading of turbulent scales. It is well known that fully developed turbulence may enhance burning and as a result may trigger detonations. However, since we reach detonations in our 2D simulations, the presence of possible turbulence cannot alter our final conclusions. 2.1 The <sc>vulcan</sc> methodology The code v2d (Vulcan-2D; Livne 1993 ) consists of an ALE method with a second-order advection scheme. In the problems discussed here, the grid moves rather slowly and the simulations are nearly Eulerian. For astrophysical combustion problems, a ‘multi-phase’ method has been added, which allows to handle separately fuel and ashes as different substances. The advection algorithm can use either a ‘slopes method’ (a la Van-Leer scheme) or VOF (Volume of Fluid) method. VOF is actually a family of schemes in which sharp boundaries between phases are constructed for advecting phases with minimal numerical diffusion (Hirt 1981 ; Ubbnik 1999 ; Gopala 2008 ). In v2d , we implement the basic idea during the advection step, by constructing, in each mixed cell, a local planar interface using the known partial volume fractions. For two phases, the partial volumes and the orientation of the surface provide enough information for interface reconstruction. The orientation, or a vector normal to the interface, is calculated from the phases gradients, based on cell centres values. In cases of more than two phases, the algorithm is much more complicated, not uniquely defined and less accurate. Once we have the local linear interface, we use it to cut the boundaries of the cell and to compute exact phase fluxes between cells. The VOF method is rather accurate as long as the boundaries between the phases are smooth. When small-scale fragmentation develops it is harder to define accurate normals and as a result the method becomes less reliable. The VOF method is more appropriate to the issue of this work and was implemented in all our helium burning simulations. The phases are defined according to a ‘critical temperature,’ which is in our case T c = 3 × 10 8 K.


O O O B-CelestialObject 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 B-CelestialObject O B-CelestialObject B-Citation I-Citation I-Citation I-Citation O O O B-Wavelength O O O O O O O O O B-CelestialObject I-CelestialObject O O O O O B-CelestialObject I-CelestialObject O O O O O O O O O O O O O B-CelestialObject I-CelestialObject O B-CelestialObject B-CelestialObject I-CelestialObject 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-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 B-CelestialObject I-CelestialObject 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 B-CelestialObject I-CelestialObject O O O O O O B-Wavelength O O B-CelestialObject O B-CelestialObject O O O O O O O O O O B-CelestialObject O O B-CelestialObject I-CelestialObject O O O O O O O O O O B-Wavelength O B-Citation I-Citation I-Citation I-Citation O O O O O O O B-CelestialObject I-CelestialObject O O O O O O B-CelestialObject O B-CelestialObject 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 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-Model B-Model O O O O 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 B-Citation I-Citation I-Citation I-Citation 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 B-CelestialObject O B-CelestialObject 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 O O B-CelestialObject I-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

for the planets d and e (along with the other wavelengths), despite the lack of non-equilibrium chemistry, non-solar metallicity compositions, and local variation of the cloud thickness in the models, suggest that these ingredients are not critically needed here. Our analysis rather demonstrates that a higher content in dust and reduced CIA of H resulting from a low surface gravity can explain the deviation of the planet’s SED with respect to those of standard field L and T dwarfs. This does however not mean that such additional ingredients do not play a role, especially for HR8799b and c. Barman et al. (2011a) noted that the -band spectrum of the high-amplitude variable brown dwarf T1 dwarf 2MASS J21392676+0220226 fitted well the SED of HR8799 b. The spectrophotometric properties of the two objects representing the best the properties of HR8799 b and c, 2MASS J13243553+6358281 and SDSS J151643.01+305344.4 , can be modeled by blended spectra of L-T dwarfs, e.g. as unresolved L/T transition binary at field dwarf ages ( Burgasser et al. 2010 ; Geißler et al. 2011 ). Beside its possible young age, we can speculate that SDSS J151643.01+305344.4 is a good template because the integrated flux of the possible binaries mimics layered patchy atmospheres. Conversely, Heinze et al. (2015) find that 2MASS J13243553+6358281 experiences high amplitude variability in the red-optical. The planets e and d could also be variable objets. The best fitting template to HR8799d’s SED – PSO J318.5338-22.8603 – was recently found to experience photometric variability in the near-infrared ( Biller et al. 2015 ). In addition, the red late-L dwarf WISE 1738, which reproduces well the properties of HR8799d and e also has a highly variable near-infrared spectrum ( Mace et al. 2013 ). In that sense, the good fit of the planet spectra with field dwarfs templates reddened by sub-micron iron, corundum, or silicate grain opacities in Sect. 3 is in line with the conclusions of Yang et al. (2015) who find that the observed variability of L5 dwarfs can be explained by the presence of spatially varying high-altitude haze layers above the condensate clouds. Such a layer may exist into the planet atmospheres and could be added to the current self-consistent cloud models (e.g. BT-SETTL, DRIFT-PHOENIX) to bring the and radii estimate in agreement with the expectations from evolutionary models. 5.2. Masses and diversity among the HR8799 system Marois et al. (2008) , Marois et al. (2010) , and Currie et al. (2011b) used the bolometric luminosity of the planets to estimate masses of 5, 7, 7, and 7 for HR8799b,c,d and e, respectively if the system is 30 Myr old. We retrieve the same luminosity – and then mass estimate – for HR8799d and e using the bolometric correction estimated for the dusty L7 dwarf WISE 0047 by Gizis et al. (2015) . The radii and log derived from the best atmospheric model fit yield masses marginally consistent with these values. This is most likely due to the


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This work has made use of the HPC resources of CINES (Jade and Occigen supercomputer) under the time allocations 2013047012, 2014047012 and 2015047012 made by GENCI. This work is partially supported by the Spin(e) grants ANR-13-BS05-0005 ( http://cosmicorigin.org ) of the French Agence Nationale de la Recherche and by the ILP LABEX (under reference ANR-10-LABX-63 and ANR-11-IDEX-0004-02). Part of the analysis of the simulation was performed on the DiRAC facility jointly funded by STFC, BIS and the University of Oxford. NEC is supported by a Beecroft Postdoctoral Research Fellowship. We thank S. Rouberol for running smoothly the Horizon cluster for us. NEC thanks Rachel Mandelbaum for useful discussion regarding the grid-locking effect.


O O O O O O O O O O O O O O O O O B-Wavelength I-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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 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-Telescope B-Instrument O O O O O O O O O B-TextGarbage I-TextGarbage I-TextGarbage I-TextGarbage I-TextGarbage I-TextGarbage O B-TextGarbage 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 O O O O O O O O O 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 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 B-Software O

clarity only scales are shown. Left: total fraction of fake sources recovered as a function of input H 160 -band magnitude. Recovery of faint sources is improved with the wavelet subtraction method. Bright sources are lost as they are subtracted in the wavelet decomposition process; to ensure that we can detect bright as well as faint sources, we therefore use both the original and wavelet-subtracted images for source detection. Middle: recovered−input magnitudes of fake sources as a function of input magnitudes for sources detected in the original and wavelet-subtracted images, but all photometry measured in the original image. The accuracy of photometry is noticeably affected when the objects are detected in images subtracted to scale 3. Right: same as the middle panel, but with photometry measured in the wavelet-subtracted images. Here, we recover systematically fainter magnitudes, as part of the light of some of the galaxies is subtracted during the wavelet subtraction procedure. We also note that detection in the wavelet-subtracted images is hampered for bright galaxies as they get removed when we perform the subtraction to smaller scales. Although we do not expect to find any high-redshift galaxies in this regime, we nonetheless choose to combine catalogs from the wavelet-subtracted and original images in order to ensure that we include potential bright or extended sources. 4. Source Selection We detect galaxies in these fields using SExtractor in dual image mode, where one detection image is used to measure photometry in each of the seven available HST filters. In order to maximize our sensitivity to faint high-redshift galaxies at a range of different colors, we use 11 different detection images: each of the four HST /WFC3 filters , plus seven stacked combinations of these images: , , , , , , and . We carry out this process with detection in both the original PSF-matched images and the wavelet-subtracted ones, for a total of 22 detection images. In all cases, photometry is measured from the original, PSF-matched images in all seven filters, as well as from the wavelet-subtracted images. We use the wavelet-subtracted photometry where it is within 1 σ of the photometry from the original image in every band, because the signal-to-noise ratio is higher, especially in the central regions of the clusters. Where the photometry differs (i.e., where some part of the Galaxy has been subtracted in the wavelet decomposition process), we use the photometry from the original image. We also visually inspect every galaxy to ensure that no part of the Galaxy has been subtracted; we find that all partially subtracted galaxies have photometry that differs by between the original and wavelet-subtracted images, so the original image photometry is used. We then combine our 22 catalogs into a single master catalog, selecting unique sources within a matching radius of 02. We measure photometry in the small Kron aperture flux, with the SExtractor parameter


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ACKNOWLEDGEMENTS We would like to thank the anonymous referee and the journal’s editorial team for their work in improving the text. In addition, the authors would like to thank Simon Glover, Thomas Greif, Sacha Hony, Ben Keller, Eric Pellegrini, Daniel Rahner, and Stefan Reissl for their useful comments and discussions during the preparation of this work. This work was granted access to HPC resources of CINES under the allocation x2014047023 made by GENCI (Grand Equipement National de Calcul Intensif). The authors acknowledge support by the High Performance and Cloud Computing Group at the Zentrum für Datenverarbeitung of the University of Tübingen, the state of Baden-Württemberg through bwHPC and the German Research Foundation (DFG) through grant no INST 37/935-1 FUGG. This work has been funded by the European Research Council under the European Community’s Seventh Framework Programme (FP7/2007-2013). SG and RSK have received funding from grant agreement no. 339177 ( starlight ) of this programme. RSK further acknowledges support from the Deutsche Forschungsgemeinschaft in the Collaborative Research Centre SFB 881 ‘The Milky Way System’ (subprojects B1, B2, and B8) and in the Priority Program SPP 1573 ‘Physics of the Interstellar Medium’ (grant nos KL 1358/18.1 and KL 1358/19.2). PH has received funding from grant agreement no. 306483. JR acknowledges support from the ORAGE project from the Agence Nationale de la Recherche, grant ANR-14-CE33-0016-03.


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-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 B-Software I-Software B-Identifier I-Identifier O B-Citation I-Citation I-Citation I-Citation O B-Citation 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 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 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 O B-Citation I-Citation I-Citation O O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula 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 B-Formula I-Formula I-Formula 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 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-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 B-Wavelength O O O O O B-Citation I-Citation I-Citation I-Citation O B-Citation O B-Citation O B-Wavelength 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-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

abundance of polycyclic aromatic hydrocarbons (PAHs; Geers et al. 2009 ). Because of these drawbacks and the chemical complexity, we restrict our model to EUV and X-ray dominated heating. 3 Numerical model All simulations of this paper are two-dimensional axisymmetric, including the toroidal velocity and magnetic field component (i.e., 2.5 dimensional). They were carried out with the PLUTO code (version 4.3) ( Mignone et al. 2007 , 2012 ), which allows solving the MHD equations on a spherical grid, including ohmic resistivity. Throughout this section, we present the basis of the numerical model, that is, the X-ray and EUV heating approach, disk structure, magnetic field, and dynamic treatment of ohmic diffusion. 3.1 MHD equations Relevant equations are the conservation of mass, the conservation of momentum, the energy equation, and the induction equation, where P is the thermal gas pressure, ρ the gas density, B the magnetic flux density vector, J the electric current density vector, η the ohmic diffusion coefficient, v the gas velocity vector, and the gravitational potential. To update the cells, the HLLD Riemann solver ( Miyoshi Kusano 2005 ) was used. To ensure the condition ∇ ⋅ B = 0, the divergence cleaning method ( Dedner et al. 2002 ) was applied. All simulations were evolved with an adiabatic equation of state, using an isentropic exponent of κ = 5∕3. Fig. 1 Ionization parameter–temperature relation based on the data in Picogna et al. (2019) . 3.2 Diagnostics The relativestrength of the magnetic field can be expressed with the dimensionless plasma parameter β : which is the ratio of thermal pressure over magnetic pressure. In order to evaluate the nonideal MHD effects (ohmic diffusion) on possible MRI activity, the notion of an ohmic Elsasser number Λ Ω ( Turner et al. 2007 ) is useful: Here v A is the Alfvén velocity and η the ohmic diffusion coefficient. Generally, MRI is considered to be suppressed for Λ Ω ≲ 1. The integrated mass-loss rate caused by the wind is computed by the expression and the disk region is excluded from the range in θ . Here, v R denotes the radial velocity component of the flow. 3.3 EUVand X-ray heating Following the approach of Owen et al. (2010 , 2012 , 2013 ), X-ray heating was provided by using a fitted ionization parameter - temperature relation T = f ( ξ ), taken from Picogna et al. (2019) . In addition to the previously studied models, the temperature curves depend on the local column density, as shown in Fig. 1 . Knowing the dependence of the temperature on the local variables, we traced rays from the inner simulation boundary radially outward up to column densities of 5 × 10 22 cm −2 . Within this range, the temperatures were set corresponding to the locally computed ionization parameter using the fitted relation f ( ξ ). With intermediate column densities, the temperature is linearly interpolated. The temperature adjustment was only invoked when the temperature computed by the ionization model was


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 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 B-Citation I-Citation I-Citation I-Citation 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 O B-Citation I-Citation I-Citation I-Citation 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 B-Citation I-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 O O O O 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 O O B-Wavelength I-Wavelength 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 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 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 B-CelestialObject I-CelestialObject 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 B-Telescope O O O O O O O B-Instrument O O B-Instrument O O B-CelestialObject I-CelestialObject O O O O O O O O O O O B-Instrument O O B-Telescope I-Telescope O O O B-Observatory I-Observatory I-Observatory I-Observatory O B-Location O O O O O O O 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-Citation I-Citation I-Citation I-Citation I-Citation I-Citation O O O O O O O O O O B-Wavelength I-Wavelength O O O O O B-Formula I-Formula I-Formula O O O O O B-Wavelength 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 B-CelestialObject I-CelestialObject O O O O O O O O B-Wavelength I-Wavelength O O B-CelestialObject I-CelestialObject O O O O O O O O O O O O O O O O O O B-Software B-Instrument O O O B-TextGarbage I-TextGarbage O O

have a Ly α luminosity of . It was the first SMG detected in CO, revealing it to contain a massive reservoir of molecular gas ( M ⊙ , Frayer et al. 1998 ). This gas reservoir was later shown to cover a large volume and contain multiple, merging galaxies (Ivison et al. 2010 ), one of which is a dusty, broad-absorption-line (BAL) quasar (Villar-Martín et al. 1999a , 1999b ; Vernet Cimatti 2001 ; Frayer et al. 2018 , e.g., Genzel et al. 2003 ). The existence of a large reservoir of diffuse molecular gas has been recently inferred by observations of the CH + cation (E. Falgarone et al. 2019, in preparation), following the detections against several other SMGs (Falgarone et al. 2017 ). Inspired by these earlier results, we observed the system with KCWI tuned to the Ly α nebula and directed to cover it and its neighboring galaxies and quasar. To our surprise, the data reveal a new, luminous Ly α source—the primary focus of this paper, which is structured as follows. We describe our KCWI observation and data reduction process in Section 2 . In Section 3 , we present the detection of the ≳100 kpc Ly α nebula powered by this system and the discovery of a new dark cloud illuminated by the nearby BAL QSO. Finally, we discuss the Ly α nebula and the physical interpretations of the newly discovered dark galaxy in Section 4 , and give a brief summary in Section 5 . Throughout this paper, we assume a flat cosmological model with and , which implies the physical scale is ≈7.85 kpc per arcsec at the redshift of SMM J02399−0136. We follow the naming convention presented in Ivison et al. ( 2010 ): L1 marks the BAL quasar; L2 is a low-surface-brightness companion to the east of L1; L2SW is a massive, extremely luminous and dust-obscured starburst galaxy, east–south–east of L1 and south–west of L2; L1N is a faint, compact northern component, visible in post-SM2 HST imaging. 2. Observations and Data Reduction 2.1. KCWI Observations The KCWI observations of SMM J02399 were carried out on UT 2017 October 21 (seeing ∼15) using KCWI on the Keck II telescope of the W. M. Keck Observatory in Hawaii. We used the BM1 grating and the medium slicer (slice width ∼07), which yields an IFU datacube with FoV of 20′′ × 168 (pixel scale of 03 × 07) centered on J023951.88−013558.0, the quasar optical position given by Ivison et al. ( 1998 ). The grating was tilted to give a central wavelength of 4620 Å and provides a spectral resolution, R ≈ 4000. The good wavelength coverage is ∼4230–5010 Å. The total on-source exposure time is four 10-minute exposures, each dithered by ∼06. For sky observations, we used an “offset-target-field” to construct the sky datacube. The offset-target has a different redshift and is located ∼2 degrees from SMM J02399. It is a compact point source at the Ly α wavelength of SMM J02399. To convert the spectral images and calibration frames (arcs, flats, bias) to a calibrated datacube, we used the IDL-based KCWI data reduction pipeline. 11 11 Available at


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Acknowledgements We thank Brad Foley for his constructive review, which greatly helped to improve this paper. D.H. has been supported through the NWA StartImpuls and N.T. has been supported by the Helmholtz Association (project VH−NG−1017) and by the German research foundation (DFG) through the priority programs 1922 “Exploring the diversity of extrasolar planets” (TO 704/3-1) and 1883 “Building a habitable Earth” (TO 704/2-1).


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This work is part of the research project entitled ‘Dynamical Analysis, Advanced Orbit Propagation and Simulation of Complex Space Systems’ (ESP2013-41634-P) supported by the Spanish Ministry of Economy and Competitiveness. The author has been funded by a ‘La Caixa’ Doctoral Fellowship for he is deeply grateful to Obra Social ‘La Caixa’. The present paper would not have been possible without the selfless help, support, and advice from Prof. Jesús Peláez. An anonymous reviewer made valuable contributions to the final version of the manuscript.


O O O O O O O O O O B-Instrument 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-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 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 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 O O O O O O B-Organization I-Organization I-Organization I-Organization B-Organization O O O O O O O O B-Organization O O O O O O O O O O O O O O O O O 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 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 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-Instrument O O O O O O O O O O O O O O O O O O B-Organization O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 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

especially when high consistency is required for large-scale surveys like DESI. Further discussion of the VPH grating physics is given by Barden et al. ( 1998 ). 2.1. Requirements and Specifications The DESI VPH gratings have physical dimensions of 180 × 160 mm 2 × 16 mm (thickness) with varying clear aperture areas. The incident field of view (FOV) of the grating is ±75 from the grating central axis, illustrated in Figure 2 , and is illuminated by a collimated beam of various pupil sizes. All DESI VPH gratings are optimized such that the +1 diffracted order efficiency is maximized. There were no specifications on off-axis FOV efficiency. The performance requirements and specifications for the DESI VPH gratings are detailed in Table 1 . These are driven by both science requirements and optical design limitations. Efficiency requirements listed are averaged over polarization and over the clear aperture. Since the actual measurements of the diffraction efficiency will include transmission losses from the VPH gelatin, glue, antireflection (AR) coating, and the substrate glasses, the “adjusted” efficiencies are used to compare results. A total of 36 full-size gratings (12 per channel) were produced by Kaiser Optical Systems Inc. (KOSI). A preliminary set of measurements was conducted by KOSI on a few points across the clear aperture (Figure 4 ), albeit in a less comprehensive manner compared to our measurements presented here. Based on our measurements, the top performing gratings in each channel with the best efficiencies are chosen for DESI integration. Table 1 DESI VPH Grating Specifications and Requirements for Unpolarized Incidence; Wavefront Specifications at 632 nm Notes. a Incidence angles are to be within an allowed installation range of ±10 from the specified α 0 (Edelstein et al. 2018 ). b Including glue and gelatin losses, not including AR coating and glass transmission losses. c Adjusted for glue, gelatin, AR coating, and glass transmission losses. We compare measurements to these values. 2.2. Preliminary Evaluation of Test Gratings Prior to the full production of the DESI VPH gratings, preliminary performance studies were conducted on full-size test VPH gratings for each bandpass, also produced by KOSI. Measurements were made in all three bandpasss to verify that the gratings met desired design specifications and also to prototype a set of reliable and efficient test procedures discussed in this paper. Efficiency performance at the specified incidence angle and the peak-to-valley (P-V) WFE for power and irregularity met the requirements (DESI Collaboration et al. 2016b ). Performance results of the test gratings were used to constrain the properties of the final DESI VPH gratings discussed in this paper. 3. Measurement Rig We built an automated optical test bench to measure the +1 order diffraction efficiency of the gratings. The design was motivated by the desire to conduct a comprehensive yet efficient performance


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Part of the calculations were performed with the high-performance computing facility MESOCENTRE, hosted by OCA. We are grateful to Nicolas Prantzos for useful suggestions that improved the content of this article. We thank the anonymous referee for her/his constructive comments. T. Masseron and B. Plez were particularly kind in providing us with some molecular data. We appreciate that D. Romano, C. I. Johnson, and C. Kobayashi kindly shared their model data. This work was partly supported (Š.M.) by the grant from the Research Council of Lithuania (LAT-16019).


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ACKNOWLEDGEMENTS The authors thank the anonymous referee for a careful reading and for useful and detailed comments which improved the quality of the paper. PHB thanks the Institut Universitaire de France for its support. DRF acknowledges support from STFC (ST/L00075X/1), including provision of local computing resources. This work has been supported by the Agence Nationale de la Recherche (ANR-HYDRIDES), contract ANR-12-BS05-0011-01.


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Acknowledgements This research was supported in part by the Polish Ministry of Science and Higher Education under grant 0149/DIA/2013/42 within the Diamond Grant Programme for years 2013–2017 and by the Polish National Science Centre under grant 2013/10/A/ST9/00023. MV acknowledges support from HST-AR-13890.001, NSF awards AST-0908346, AST-1515001, NASA-ATP award NNX15AK79G.


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ACKNOWLEDGEMENTS We thank M. Begelman, A. Beloborodov, E. Quataert, K. Parfrey, L. Sironi, A. Spitkovsky, and V. Zhdankin for fruitful discussions. This work was supported by DOE grant DE-SC0008409, NASA grants NNX17AK57G and NNX16AB28G, NSF grant AST-1411879, and NASA through Einstein Postdoctoral Fellowship grant PF7-180165 awarded to AP by the Chandra X-ray Center, operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS803060. The simulations presented in this paper used computational resources of the NASA/Ames HEC Program and XSEDE/Stampede2 (allocation PHY140041) (Towns et al. 2014 ). DAU gratefully acknowledges the hospitality of the Inst. for Advanced Study and support from the Ambrose Monell Foundation.


O O O O B-Instrument O O O B-Wavelength O B-Wavelength 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 B-Wavelength 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 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 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 B-Wavelength 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 B-CelestialObject 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 B-Wavelength O O O O O O O O B-CelestialObject O O O O O O O O B-CelestialObject O B-CelestialObject O O O O O O O O B-Wavelength 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 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 O O O B-Wavelength 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-Wavelength 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 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-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 O B-CelestialObject 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 B-CelestialObject 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-Wavelength O B-Wavelength I-Wavelength O O O O O O O O O O O B-Wavelength I-Wavelength O O O O B-CelestialObject O O O O O O O O O O O O B-Instrument O O O O B-ObservationalTechniques I-ObservationalTechniques O O O O O O O O O O

broadband SEDs, 19 with GROND and 3 with optical and NIR data from different instruments. For the rest of the GRBs at , we collected measurements from the literature if available; these measurements are listed in Table 3 . 4. Results The results, i.e., the best-fit models and parameters, from fitting the combined NIR, optical, and X-ray SEDs with the local extinction curves are summarized in the top portion of Table 2 and a plot containing all the SEDs is shown in Fig. 7 . As mentioned above, the SEDs of the 10 GRBs that show evidence for a small to medium amount of dust ( mag) were also fitted with the SN extinction curve. These SEDs are additionally plotted in Fig. 8 , in comparison to the best results from the local extinction curves (zoomed into the optical and NIR regime; the full plot is displayed in Fig. C.1 ). The corresponding parameters are given in the bottom portion of Table 2 . With these measurements we increase the number of determined values for GRBs at by a factor of 2–3. As previously observed for GRBs at lower redshift, the local extinction curves provide a good fit to the data, with the featureless SMC extinction curve usually best describing the observed magnitudes ( Greiner et al. 2011 ; Zafar et al. 2011a ). We find that from the 22 modeled GRBs at , 16 are best fit with the SMC and only 6 are best fit with the LMC or MW extinction curve. In contrast to other samples with GRBs at mostly , however, we find that all of the GRBs are only behind small to medium dust column densities ( mag) within their host galaxies line of sight, which is also true for the values collected from the literature (see Table 3 ). The distribution of our best-fit values compared to those from the unbiased samples from Greiner et al. (2011) and Covino et al. (2013) and the data from Zafar et al. (2011a) , based on the GRB sample from Fynbo et al. (2009) , is shown in Fig. 9 . While all of these samples include a % of GRBs with mag (and 3–5% with mag), only % of our GRBs at redshift have a visual extinction of mag and the remaining % have mag. A two-sample Kolmogorov-Smirnov test returns a rejection probability of % for the null hypothesis that the values from our GRBs at and those from Greiner et al. (2011) are drawn from the same sample. Considering the 10 SEDs that were also fitted with the SN extinction curve, we find that, besides GRB 090205 and 090516, the goodness of fit is comparable to the result from the local extinction curves, while for two of the bursts, GRB 140428A and 151027B, the SN extinction curve provides a better fit to the data. In case of GRB 140428A, the somewhat flat SED between the and band can be identified with the flattening between (rest-frame) 3000 and 1700 Å of the SN extinction curve () or as a small (rest-frame) 2175 Å feature when fitting the MW extinction curve (). This result is consistent with that from fitting the GROND SED to determine a photometric redshift, where we also find that the SN extinction curve is


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

for a realistic simulated pulsar. Apart from different f - and -scales, the corresponding plot for the semicoherent mismatch looks very similar. The mismatch agrees well with its metric approximation for mismatch m ≤ 0.5, which is typical: in Appendix B , we show that the highest sensitivity at given computing cost for an elliptical search is obtained with an average mismatch (see Table B1 ). Figure 3. Comparison of the coherent metric approximation to the actual mismatch, for parameters of a simulated eccentric orbit binary pulsar in 4FGL J0523.3−2527. Blue contours show the actual mismatch and red contours the metric approximation, at m = 0.25 and 0.5. The sky position parameters { n x , n y } are not shown in Figure 3 because we assume that for spider pulsars they are known to high precision from optical observations. A partially informed search for binary pulsars in elliptic orbits, without exact information about the sky position and constraints on the orbital parameters, is computationally impossible. There are too many parameter space dimensions—even for circular orbits with reasonable parameter ranges, the grid has too many points. To make a search possible, one needs tight constraints derived from optical/X-ray observations of the pulsar’s companion star. In the next section, we will discuss constraints and the search design for the probable eccentric orbit binary gamma-ray pulsar in 4FGL J0523.3−2527 (Strader et al. 2014 ; Saz Parkinson et al. 2016 ). 4.2. Search Design for Low-eccentricity Binary In this section, we discuss how to reduce the search parameter space using orbital constraints for the gamma-ray source 4FGL J0523.3−2527, presumed to be a pulsar in an eccentric binary orbit. The source is named 3FGL J0523.3−2528, 2FGL J0523.3−2530, or 1FGL J0523.5−2529 in previous LAT source catalogs. This is similar to the circular example of Section 3.2 . The gamma-ray source itself was investigated by Saz Parkinson et al. ( 2016 ) and ranked ninth highest in a list of most significant 3FGL unassociated sources predicted to be pulsars. It shows typical pulsar-like properties: the photon flux is stable over time, and the spectrum is fit by an exponential cutoff power law. The source is not in the Galactic disk, which increases the odds that it hosts an MSP. Earlier optical observations identified a likely companion and indicate an orbit with small, but not negligible, eccentricity of (Strader et al. 2014 ). In contrast to the previous paragraph, this suggests that the pulsar is a YP, because binary MSPs tend to have rather circular orbits (Phinney 1992 ). The frequency and spin-down search ranges are chosen following the logic of the previous search design (Section 3.2 ). For YPs we search f ∈ [0, 44] Hz and . For MSPs we search f ∈ [44, 1500] Hz and . The f -dimension is efficiently searched using FFTs with bandwidth f BW = 8 Hz, and the -dimension is covered by a uniformly spaced lattice. The sky position search range of the probable


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open field lines. The plots of null height (Figure 4 (a), red line) and free magnetic energy (Figure 4 (b), red line) also indicate that the system has almost relaxed back to its equilibrium state. The free magnetic energy (b) sharply drops as soon as the photospheric driving stops. This is because energy is being propagated out of the open top boundary along the jet spire, but is no longer being injected through the lower boundary. The null height (a) continues to increase until around t = 70, showing that the jet dome has continued to expand. The decrease in null height after this point is more rapid than that in PAD09 (black line) as there is no longer photospheric driving injecting energy and twist into the field. We now compare our PAD10 simulation with that of Pariat et al. ( 2010 ), as well as PAD09 and PAD09R. In Pariat et al. ( 2009 ) they see a peak and sudden drop in free magnetic energy when the eruption occurs after kinking breaks the symmetry of the model. In Pariat et al. ( 2010 ), the model is initially set up to be non-symmetric, with the background magnetic field inclined at 10° to the vertical. They continue to rotate the positive magnetic polarity after the first eruption, resulting in a series of eruptive events, hence a series of peaks and drops in free energy. The free magnetic energy in our model can be seen in Figure 4 (b) (blue line). We do see a series of peaks in our model, but they are not sharply defined, as there is not as rapid a decrease in energy after each peak. Again, this is due to the quasi-static nature of the magnetofrictional method. We cannot accurately follow the dynamics of an eruption, so the result is that not as much energy and twist is released before it starts building up again due to the photospheric driving. The first peak occurs at a similar time to that of our PAD09 simulation. The peak free magnetic energy built up before the first release is 1.45 × 10 25 erg and 1.29 × 10 25 erg for the PAD09 and PAD10 simulations, respectively. The energy buildup is greater in PAD09, as changes in connectivity do not occur until after the onset of the kink and the breaking of the structure’s symmetry. In PAD10, the symmetry is already broken in the initial setup, so changes in connectivity can occur as soon as the simulation begins and magnetic energy can already be released. Figure 5 (a) shows in the xz -plane at y = 9 Mm for PAD09 (left), PAD10 (center), and PAD09R (right). An animated version of these panels is available online. In PAD09 and PAD09R, current is initially all contained within the twisted dome of the jet and along the narrow, vertical jet spire, indicating that no twist has been released and almost all non-potentiality is contained within the dome. After the onset of the kink, the symmetry of the system is broken and we see a broader jet spire, as in Figure 5 (a). In PAD10, the spire broadens more rapidly, as connectivity changes occur with background magnetic field right from the start