O B-Formula I-Formula I-Formula O B-Formula I-Formula I-Formula 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 B-Formula I-Formula I-Formula I-Formula O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O B-Formula I-Formula I-Formula I-Formula 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 B-Formula I-Formula I-Formula O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O O O O O O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O B-Formula I-Formula I-Formula I-Formula O O O O O O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O B-Formula I-Formula I-Formula I-Formula 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 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 B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O 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 B-Citation I-Citation I-Citation I-Citation O O O O O O B-CelestialObject I-CelestialObject 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 B-CelestialObject I-CelestialObject O O O O O O O O O O O O O B-CelestialObject I-CelestialObject O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O O O O O O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula 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 O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O O O O O O O O O O O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O 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 B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O B-Formula I-Formula I-Formula I-Formula O O O O B-CelestialObject I-CelestialObject 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 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-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O O O O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O O O O O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O

is γβ ≳ 1.3, i.e., β ≳ 0.8. So, this means that the fastest ejecta has a velocity v ≳ 0.8 c . For lower-density medium, the fastest velocity should be even higher. On the other hand, Figure 4 shows that the lower limit to the shock energy at t = 156.4 days is E sh ≳ 2 × 10 50 erg assuming B ≲ 10 −2 , implying that the total ejecta kinetic energy is E ej ≳ 2 × 10 50 erg. This constraint is insensitive to B and/or E 1 . 4. Conclusion and Discussion We have used a simple spherical shock model to explain the nonthermal afterglow emission of GW 170817. We consider only the emission from the forward shock. The shock dynamics derived is available not only to ultrarelativistic and Newtonian shocks, but also to mildly relativistic shocks. The model parameters, except for those assuming e = 10 −1 , are constrained by multiband observations, and the main results are as follows: 1. The merger ejecta profile has k ≈ 6.7 and E 1 ≳ 4 × 10 49 erg for n 10 −1 cm −3 . 2. The fastest ejecta has γ max β max ≳ 1.3, i.e., β max ≳ 0.8, for n 10 −1 cm −3 . The slowest ejecta has γ min β min ≳ 0.7, i.e., β min ≳ 0.6, for n 10 −1 cm −3 . 3. The total ejecta kinetic energy is E ej ≳ 2 × 10 50 erg, if B ≲ 10 −2 . Recently, there have been some works claiming that the rising light curve turns over around t turn ∼ 150 days (Alexander et al. 2018 ; Dobie et al. 2018 ; Nynka et al. 2018 ). This turnover implies that there is no longer any energy injection, and thus the light curve starts to decline. So, the lower panel (for t = 156 days) in Figure 4 shows us the constraint of the total kinetic energy of the ejecta. The H i observation of the host galaxy constrained n 0.04 cm −3 (Hallinan et al. 2017 ). However, since the location of GW 170817 appears to be within the effective radius of the host galaxy (Levan et al. 2017 ), the density of the medium at GW 170817's location may not be very low. So, a reasonable circumburst medium density for GW 170817 could be n ∼ 10 −2 cm −3 . Moreover, GRB afterglow modeling usually gives a postshock magnetic field with B ∼ 10 −5 –10 −3 (e.g., Lemoine et al. 2013 ; Barniol Duran 2014 ; Santana et al. 2014 ). Taking n ∼ 10 −2 cm −3 and B ∼ 10 −5 –10 −3 , we can obtain from the lower panel of Figure 4 that the total kinetic energy is E ej ∼ (0.3–5) × 10 51 erg and that the slowest ejecta has γ min β min ∼ 1–1.5, i.e., β min ∼ 0.7–0.8. For the same range of n and B values, the upper panel of Figure 4 implies that the fastest ejecta has a velocity γ max β max ≳ 1.8, i.e., γ max ≳ 2. The gamma-ray emission from GW 170817 may be produced by the fastest part of the ejecta. Its energy E fast ≈ E γ ≃6 × 10 46 erg (Abbott et al. 2017b ) is many orders of magnitude smaller, compared with the constrained kinetic energy of the ejecta. For n ∼ 10 −2 cm −3 and B ∼ 10 −5 –10 −3 , the lower panel of Figure 4 also shows that E 1 ∼ 10 50.5 –10 52.5 erg, with which we can estimate the velocity of the fastest ejecta, γ max β max = ( E 1 / E γ ) 1/ k ∼


O O O O O O O O O O O O B-CelestialRegion O B-CelestialRegion I-CelestialRegion O O O B-CelestialObject O O B-Wavelength O O O O O O O O O B-Citation I-Citation I-Citation I-Citation I-Citation O B-Wavelength O B-Wavelength O B-Telescope I-Telescope I-Telescope I-Telescope I-Telescope I-Telescope B-Telescope B-Wavelength B-Wavelength O O O O B-Citation I-Citation I-Citation I-Citation O B-Citation O O O B-CelestialRegion I-CelestialRegion O O B-Citation I-Citation I-Citation I-Citation O B-Citation O O B-Citation I-Citation I-Citation I-Citation O O B-CelestialRegion I-CelestialRegion O O O O O O O O O O O B-Wavelength B-Wavelength B-Telescope B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O O O O O B-Survey I-Survey I-Survey I-Survey B-Survey B-Citation I-Citation I-Citation O O B-Survey I-Survey I-Survey I-Survey I-Survey I-Survey B-Survey O B-Person I-Person O B-Survey O B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O B-Survey I-Survey I-Survey I-Survey I-Survey I-Survey I-Survey I-Survey I-Survey I-Survey B-Survey O B-Person I-Person O O O B-Citation I-Citation I-Citation I-Citation I-Citation O O O B-Wavelength O B-Wavelength O O O O O O B-Software I-Software B-Citation I-Citation I-Citation O O O O O O O O O O O O O O B-ObservationalTechniques I-ObservationalTechniques B-ObservationalTechniques O O O O B-Wavelength I-Wavelength O B-Wavelength I-Wavelength I-Wavelength O O O O O O O O B-ObservationalTechniques I-ObservationalTechniques B-ObservationalTechniques O O O O B-Telescope O O O O O O O O O O O O O O B-Telescope B-Wavelength O B-Telescope B-Wavelength O O O O O O O O O O O O O O O O O O O O O O B-CelestialRegion O B-CelestialRegion I-CelestialRegion O O O O O O O O O O O O O O O B-Survey B-Wavelength O O B-Wavelength O B-Wavelength O B-Wavelength O O O O O O O B-Telescope B-Citation I-Citation I-Citation O O O O B-Citation I-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation I-Citation O O O B-Survey I-Survey I-Survey I-Survey B-Survey O O O B-Citation I-Citation I-Citation I-Citation I-Citation O O B-Wavelength I-Wavelength O O O O B-Citation I-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation I-Citation O B-Wavelength B-Wavelength O O O B-CelestialRegion I-CelestialRegion O O O B-Dataset I-Dataset I-Dataset O O B-Survey I-Survey B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O O O O B-CelestialRegion I-CelestialRegion O B-Wavelength O O O O B-Dataset I-Dataset I-Dataset O O B-Survey I-Survey I-Survey I-Survey I-Survey B-Survey B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O O O O O 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 B-Survey I-Survey B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O B-Survey I-Survey I-Survey B-Survey B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O B-Survey B-Citation I-Citation I-Citation O B-Citation O B-Survey I-Survey I-Survey I-Survey B-Survey B-Citation I-Citation I-Citation O B-Survey I-Survey I-Survey B-Survey B-Citation I-Citation I-Citation O O O B-Citation I-Citation I-Citation I-Citation O B-Citation O O O B-CelestialRegion I-CelestialRegion O B-CelestialRegion

ratio 3.5. This process resulted in 726 and 1088 sources in the COSMOS and UDS fields, respectively. The source S2CLSJ021830-053130 with an 850m flux of 50mJy is the lensed candidate discussed by Ikarashi etal. ( 2011 ). 2.2 Radio and mid-infrared The Karl G. Jansky Very Large Array (VLA) 1.4GHz radio data were taken from Schinnerer etal. ( 2007 , 2010 ) for the COSMOS field, and from Ivison etal. ( 2005 , 2007 ) and Arumugam etal. (in preparation) for the UDS field. The catalogues include sources for which 3 detections were obtained. The mid-infrared (mid-IR) Spitzer (Fazio etal. 2004 ; Rieke etal. 2004 ; Werner etal. 2004 ) data are from the Spitzer Extended Deep Survey (SEDS; Ashby etal. 2013 ), the Spitzer Large Area Survey with Hyper-Suprime-Cam (SPLASH, PI: P. Capak), the S-COSMOS project (Sanders etal. 2007 ; Le Floch etal. 2009 ) and the Spitzer Public Legacy Survey of the UKIDSS Ultra Deep Survey (SpUDS; PI: J. Dunlop) 2 described in Caputi etal. ( 2011 ). To obtain the 3.6 and 4.5m photometry, we used the de-confusion code T-PHOT 3 (Merlin etal. 2015 ). This utilizes prior information on the positions and morphologies of objects from a high-resolution image (HRI; in this case the K -band or K s -band images) to construct a model of a given low-resolution image (LRI; in this case the Spitzer imaging) while solving for the fluxes of these objects. The 3 depths of the VLA radio and Spitzer mid-IR imaging in both fields are summarized in Table 1 . Table 1. The 3 depths of the multifrequency data used in the COSMOS and UDS fields. Note. a The two alternative values correspond to the shallower and deeper strips of the UltraVISTA near-IR imaging. 2.3 Optical and near-infrared The optical data in both fields were obtained with Subaru/SuprimeCam (Miyazaki etal. 2002 ), as described in Dye etal. ( 2006 ) and Furusawa etal. ( 2008 ), and from the CanadaFranceHawaii Telescope Legacy Survey (CFHTLS), as described in Bowler etal. ( 2012 ). The deep z -band images are described in Bowler etal. ( 2012 ) and Furusawa etal. ( 2016 ). The near-infrared (near-IR) data in the COSMOS field was obtained from Data Release 2 of the UltraVISTA survey (McCracken etal. 2012 ; Bowler etal. 2014 ), while in the UDS field the near-IR data were provided by Data Release 10 of the UKIRT Infrared Deep Sky Survey (UKIDSS; Lawrence etal. 2007 ; Cirasuolo etal. 2010 ; Fontana etal. 2014 ). In both fields, the optical and near-IR fluxes were measured in 3-arcsec diameter apertures, and the resulting 3 depths of this aperture photometry are summarized in Table 1 . Finally, we used a list of spectroscopic redshifts from 3D- HST (Brammer etal. 2012 ; Skelton etal. 2014 ; Momcheva etal. 2016 ), VIMOS Ultra-Deep Survey (VUDS; Le Fvre etal. 2015 ; Tasca etal. 2017 ), zCOSMOS (Lilly etal. 2007 , 2009 ), MOSFIRE Deep Evolution Field (MOSDEF; Kriek etal. 2015 ), PRIsm MUlti-object Survey (PRIMUS; Coil etal. 2011 ) and from Trump etal. ( 2009 , 2011 ) in the COSMOS field and UDSz


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We thank an anonymous referee for several insightful comments that helped improve the content and presentation of this paper. We also thank J.-P. Macquart, R. M. Shannon, M. Bailes, and H. Knight for several useful discussions. This scientific work makes use of the Murchison Radio-astronomy Observatory operated by CSIRO. We acknowledge the Wajarri Yamatji people as the traditional owners of the Observatory site. NDRB is supported by a Curtin Research Fellowship. Support for the operation of the MWA is provided by the Australian Government Department of Industry and Science and Department of Education (National Collaborative Research Infrastructure Strategy: NCRIS) under a contract to Curtin University administered by Astronomy Australia Limited. We acknowledge the iVEC Petabyte Data Store and the Initiative in Innovative Computing and the CUDA Center for Excellence sponsored by NVIDIA at Harvard University and support from the Centre for All-sky Astrophysics (CAASTRO) funded by grant CE110001020.


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ACKNOWLEDGEMENTS I thank V. Belokurov, D. Erkal, N. W. Evans, N. Hambly, S. Koposov, and the anonymous referee for valuable suggestions. This work relied exclusively on the data from the European Space Agency (ESA) mission Gaia ( https://www.cosmos.esa.int/gaia ), processed by the Gaia Data Processing and Analysis Consortium. This work was supported by the European Research council under the 7th Framework programme (Grant No. 308024).


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

of the cavity λ H decreased due to subtraction of solar wind He + ions from ∼19.5 to ∼15.5 au, whereas the average ionization factor β H,E decreased from 7.6 × 10 −6 to 3.9 × 10 −6 s −1 . These values are still well above the expected range of 3–6 au for λ H and (0.5–1) × 10 −6 s −1 for β H,E . Therefore, the presence of solar wind He + ions in SWAP observations does not fully explain this discrepancy, but it is an important factor that influences the fitted spectra of H + PUIs. Moreover, a reasonable agreement may not be possible using the Vasyliunas Siscoe ( 1976 ) model, which is only a simple approximation and does not account for most of the transport processes. More advanced models (see, e.g., Zank et al. 2018 ) can give results consistent with the observations for more realistic parameters. 6. Summary Helium, the second most abundant species in the solar wind, is emitted from the solar corona almost entirely in the form of double ionized alpha particles. Due to charge exchange with ISN atoms, some of the He 2+ ions acquire electrons and form a new population of singly charged solar wind He + ions. Charge exchange with ISN hydrogen is the dominant source of this population. The rate of this reaction significantly increases over the range of the solar wind speeds. As a result, the ratio of solar wind He + ions to solar wind He 2+ can exceed 10% at the termination shock for regions of fast solar wind ≥600 km s −1 , such as at high heliolatitudes around solar minimum (McComas et al. 2000 ). Therefore, solar wind He+ ions are an important component of the plasma and, as such, should be included in analyses of the processing of the solar wind and acceleration processes at the termination shock. In this study, we show that this contribution is visible in some part of the SWAP observations as a “bump” just below the cutoff energy of H + PUIs. This paper is a first report on the observation of this population in the outer heliosphere. Earlier reports on solar wind He + ions are very sparse (Tucker Gould 1966 ; Kozlovsky 1968 ; Feldman et al. 1974 ; Gloeckler Geiss 1998 ). We select observations with a local maximum just below the H + PUI cutoff energy. About ∼3/5 of them are related to significant changes in the solar wind speed up to 10 days before the date of the observation. Consequently, the observed spectrum was likely modified due to this change. In the remaining ∼2/5 cases, the solar wind speed remains stable throughout 10 days, and for most of them, we find significant improvements in the goodness of the fit for H + PUIs when solar wind He + ions are adequately accounted for. In only one observation is the improvement not statistically significant. We point out that the observations of solar wind He + ions are more likely to be visible during times of faster solar wind, with lower temperatures and high relative abundances of solar wind alpha particles. We notice that in some observations also, other heavy ions such as oxygen and carbon


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 B-Survey O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 B-Survey O O O O O O O O O O O O O O O O O O B-Survey O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Survey O O O O O O O O B-Model 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 B-Survey O O O O O O B-Model O B-Model O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Model O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Survey 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 B-Model O B-Model O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Survey O O O O B-Model O O O O B-Model O

rot / σ profiles tend to peak at a median radius of R ∼ 3 R e . By comparison, the ePN.S FRs tend to peak at smaller fractions of R e , at a median 1.3 R e . This difference between the observed and simulated FRs is not a consequence of the selection functions of the samples. The TNG galaxies are selected according to color, mass and intermediate to major axis ratio p . The selection in p removes centrally elongated galaxies, most of which have intermediate to low λ (1 R e ) (Fig. 5 ). The Atlas3D sample, selected as described in Sect. 3 , has some of the disk galaxies removed which as in MANGA (Fig. 5 ) will be mostly located at high λ (1 R e ). We recall that in the comparison we consistently matched the color selection and mass range of the Atlas3D sample to the selection criteria adopted for the TNG galaxies. Thus the TNG sample should in principle contain a larger number of late type galaxies with strong disks, i.e. high V rot / σ , by comparison with Atlas3D. Figure 13 shows instead that the TNG100 and TNG50 ETG samples lack galaxies with high rotational support at 1 R e . In Sect. 6.1 we discussed that the effective radii used to normalize the radial scales in TNG would be expected to be systematically slightly overestimated compared to the observed R e since they are defined using the total bound stellar mass which is often inaccessible in observations. If taken into account, this effect would increase the gap between simulated and observed samples. On the other hand, since the mass-size relation is roughly reproduced in the simulations (Fig. 7 ), the different steepness of the V rot / σ profiles in observations and simulations implies a different distribution of the angular momentum as a function of radius in the simulated galaxies. This could be due to a too efficient condensation of the gas into stars that does not allow the gas to dissipate and collapse to small enough radii. 7.3. Simulated versus observed profiles – outskirts Figure 14 compares the relation between V rot / σ in the central regions and V rot / σ in the stellar halos for the simulated and observed galaxies. The observed galaxies are from the ePN.S survey ( Pulsoni et al. 2018 , their Fig. 9) and are reported in the top left panel. Their measurements use absorption line data at 1 R e and PN data for the halos, which on average cover 6 R e with a large scatter (minimum 3 R e , maximum 13 R e ). The central and right top panels show V rot / σ (6 R e ) versus V rot / σ (1 R e ) for TNG100 and TNG50 ETGs, respectively. Galaxies close to the dashed 1:1 lines show similar rotational support in the central regions and in the outskirts; galaxies below the lines have increased rotational support in their stellar halos; galaxies above the lines instead have reduced rotation at large radii. Fig. 14. V / σ ratio in the central regions compared with the V / σ at large radii. Top : V / σ (1 R e ) versus V / σ (6 R e ) for the ePN.S galaxies ( left ), TNG100 ( middle ), and TNG50 (


O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 B-CelestialObject O O O O O O O O O O O O O O O O B-CelestialObjectRegion I-CelestialObjectRegion O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-ObservationalTechniques I-ObservationalTechniques O B-ObservationalTechniques 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 B-CelestialObject O O O O O O O O O O O O O O O O O O B-CelestialObject O O O O B-ObservationalTechniques I-ObservationalTechniques O O O B-Survey I-Survey I-Survey I-Survey I-Survey O O B-Mission I-Mission O O O O O O O O O O O O O O O O O O B-ObservationalTechniques O O O O O B-ObservationalTechniques I-ObservationalTechniques O O O O B-ObservationalTechniques I-ObservationalTechniques O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O

average confirm the continuous decline of metallicity. The use of age groups in the inner disc revealed an increase of the metallicity level with age. The older clusters (1.7 Gyr) appear more metal rich by about 0.07 dex than clusters of a young group (0.5 Gyr). This is clearly the opposite of what was commonly looked for in the past, a decline of metallicity with age, but it might be an observational confirmation of radial migration. However, to further validate this result, a homogeneous age scale for the clusters and a somewhat larger sample of older objects are needed. Furthermore, we compared open clusters and Cepheids with recent Galactic chemical and chemodynamical models. The models reproduce the observed increase of metallicity with age and generally agree well with the observations of young and older populations. However, the comparison shows that the models probably underestimate the effect of radial migration, or that there is another not yet included mechanism acting in the outer disc, where open clusters show a much flatter gradient than predicted by the models. However, to study the reliability of the current models in more detail and to further improve them, two questions apparently need to be resolved first: a somewhat larger open cluster sample that covers a broader Galactic distance range in each age group is needed. Most studied objects are still concentrated in the solar circle. Second, a more detailed investigation of the relative metallicity scale between open clusters and Cepheids is required. This should include an improvement of the knowledge of the metallicities of Cepheids that are cluster members and of the metallicities of the hosting open clusters, but the latter based on photometric methods or spectroscopic data of red giants or main-sequence stars. Finally, we showed the Galactic metallicity distribution of a 8 kpc large area in a contour plot and the metallicity as a function of the azimuthal angle. Both confirm that there is little difference between the Galactic quadrants. However, there are some slightly overabundant areas that might be related to the spiral structure of the Galaxy. The final results of spectroscopic surveys, such as the Gaia -ESO Public Spectroscopic Survey or the Gaia mission itself, will certainly provide the basis for more detailed studies. Furthermore, the availability of current and future homogeneous photometric datasets will allow complementing the spectroscopic metallicities with results obtained by photometric methods. We wish to mention in particular the DG method, which can be applied to any photometric data with available temperature calibrations. Appendix A: Additional table Table A.1 Metallicities of the open clusters. Notes. The HQS, LQS, and weighted photometric metallicities, respective errors, number of stars used for the calculation (N), and the total weight (W) of the photometric metallicities. The code for the photometric systems is 1 (DG), 2 (DDO), 3 (Friel), 4 (Ca


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We thank the anonymous referee for helpful comments that greatly improved the presentation of this paper. We also thank H.J. Mo and Frank C. van den Bosch for useful discussions. This work is supported by the 973 Program (No. 2015CB857002), NSFC (Nos. 11128306, 11121062, 11233005), the Strategic Priority Research Program “The Emergence of Cosmological Structures” of the Chinese Academy of Sciences, Grant No. XDB09000000, and a key laboratory grant from the Office of Science and Technology, Shanghai Municipal Government (No. 11DZ2260700).


O O O O B-Person I-Person I-Person O B-Person I-Person I-Person O O O O O O O O O O O O O O O O O O O O O B-Observatory O B-Identifier I-Identifier O B-Identifier I-Identifier B-Observatory O O O O B-Observatory I-Observatory I-Observatory I-Observatory O O O O B-Organization I-Organization I-Organization B-Location O B-Organization I-Organization I-Organization I-Organization I-Organization B-Location O O B-Organization I-Organization I-Organization B-Location B-Organization I-Organization I-Organization O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization B-Location O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization B-Location O O O O B-Location I-Location I-Location O B-Observatory I-Observatory I-Observatory O O O B-Observatory B-Organization I-Organization I-Organization B-Observatory I-Observatory I-Observatory B-Observatory O B-Observatory I-Observatory I-Observatory I-Observatory I-Observatory O O O O O O O B-ComputingFacility O B-Organization I-Organization I-Organization I-Organization I-Organization B-Organization O O O O B-Observatory I-Observatory I-Observatory I-Observatory B-Observatory O B-Organization I-Organization O B-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility B-Observatory I-Observatory I-Observatory I-Observatory I-Observatory B-Observatory O O O O O O O B-ComputingFacility I-ComputingFacility I-ComputingFacility I-ComputingFacility O O O B-ComputingFacility I-ComputingFacility I-ComputingFacility B-ComputingFacility B-Observatory I-Observatory I-Observatory I-Observatory I-Observatory B-Person O O O O B-Grant I-Grant I-Grant I-Grant I-Grant I-Grant I-Grant I-Grant I-Grant O O O O O B-Organization B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization B-Grant I-Grant I-Grant I-Grant B-Grant O O B-Grant B-Grant B-Grant B-Grant B-Grant B-Grant O B-Grant O B-Fellowship I-Fellowship I-Fellowship I-Fellowship O O B-Fellowship O O O B-Organization I-Organization I-Organization O B-Organization I-Organization I-Organization I-Organization I-Organization B-Organization O O B-Grant B-Person O O O O B-Organization I-Organization I-Organization I-Organization I-Organization B-Organization O O B-Grant O O O O O B-Organization I-Organization I-Organization O O B-Organization I-Organization I-Organization O O O O O O O B-Organization B-Grant B-Person O B-Observatory I-Observatory O O

We are grateful to Professor Inga Kamp and Dr. Satoshi Okuzumi for their useful comments. We thank the referee for many important suggestions and comments. This paper makes use of the following ALMA data: ADS/JAO.ALMA #2015.1.01259.S and ADS/JAO.ALMA #2013.1.00601.S. ALMA is a partnership of European Southern Observatory (ESO) (representing its member states), National Science Foundation (USA), and National Institutes of Natural Sciences (Japan), together with National Research Council (Canada), National Science Council and Academia Sinica Institute of Astronomy and Astrophysics (Taiwan), and Korea Astronomy and Space Science Institute (Korea), in cooperation with the Republic of Chile. The Joint ALMA Observatory is operated by ESO, Associated Universities, Inc/National Radio Astronomy Observatory (NRAO), and National Astronomical Observatory of Japan. Our numerical studies were carried out on SR16000 at Yukawa Institute for Theoretical Physics (YITP) and computer systems at Kwasan and Hida Observatory (KIPS) in Kyoto University, and PC cluster at Center for Computational Astrophysics, National Astronomical Observatory of Japan. ALMA Data analysis was carried out on the Multi-wavelength Data Analysis System operated by the Astronomy Data Center (ADC), National Astronomical Observatory of Japan. S.N. was supported by the ALMA Japan Research Grant of NAOJ Chile Observatory, NAOJ-ALMA-211. This work is supported by JSPS (Japan Society for the Promotion of Science) Grants-in-Aid for Scientific Research (KAKENHI) (grants No. 25108004, 25108005, 25400229, 15H03646, 15K17750, 17K05399, and 17K14244), by Grants-in-Aid for JSPS fellows (grant No. 16J06887), and by the Astrobiology Center Program of National Institutes of Natural Sciences (NINS; grant No. AB281013). C.W. acknowledges support from the Science and Technology Facilities Council (STFC; grant No. ST/R000549/1) and start-up funds from the University of Leeds. Astrophysics at Queen’s University Belfast is supported by a grant from the STFC (ST/P000312/1). T.J.M. thanks Leiden Observatory for hospitality.


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ACKNOWLEDGEMENTS We thank Volker Springel for access to arepo . We thank Joel Primack for insightful questions about the spin parameter that motivated this project. We also thank Ramesh Narayan, Re’em Sari, Rachel Somerville, James Bullock, and Michael Boylan-Kolchin for enlightening discussions. We also thank Cristina Popa for her help in the early stages of this project.


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

2011a ; Konopacky et al. 2013 ; Lee et al. 2013 ; Barman et al. 2015 ). The good fit in the 35 m region 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


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O.P. sincerely thanks Dr. A. Retinò and Prof. W. H. Matthaeus for the fruitful discussions. O.P. and S.S. are partly supported by the International Space Science Institute (ISSI) in the framework of the International Team 405 entitled “Current Sheets, Turbulence, Structures and Particle Acceleration in the Heliosphere.” D.P. was partially supported by STFC grant ST/N000692/1. Numerical simulations discussed here have been performed on the Marconi cluster at CINECA (Italy), within the projects IsC53_RoC-SWT and IsC63_RoC-SWTB, and on the Newton cluster at the University of Calabria (Italy). This project (F.V., S.S.) has received funding from the European Unions Horizon 2020 research and innovation program under grant agreement No. 776262 (AIDA, www.aida-space.eu ).


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

to the general population of XRT light curves of GRBs. The gray shaded region is a compilation of long GRB light curves ( Evans et al. 2007 , 2009 ), where the color represents density. The light blue lines are sGRB light curves from bursts with duration of T 90 ≲ 2 s and those that were classified as short in Kann et al. (2011) , Berger (2014) and D’Avanzo et al. (2014) . The thick green line is GRB 111117A. Despite the remarkably high redshift, the luminosity is comparable to the bulk of the short burst population, and subluminous compared to the lGRB population. The inset shows the X-ray luminosity distributions of sGRBs and lGRBs at 10 ks, indicated by the vertical dashed line in the main panel. The separation of GRBs in two distinct classes based on their high-energy observables points to an intrinsically different physical origin. Long GRBs are typically interpreted as collapsars ( MacFadyen Woosley 1999 ) in which a single, massive star undergoes gravitational collapse, where the currently preferred model for sGRBs is the merger of two NSs ( Eichler et al. 1989 ; Nakar 2007 ). Bromberg et al. (2013) investigated the degree to which high-energy observables of the long and short GRB populations overlap, and quantified the certainty in class membership. According to Bromberg et al. (2013) , GRB 111117A has a probability of being a sGRB. Compared to two other sGRB candidates at high redshift, GRB 060121 ( T 90 = 1.97 ± 0.06 s; de Ugarte Postigo et al. 2006 ; Levan et al. 2006 ) at 1.7 ≲ z ≲ 4.5 () and GRB 090426 ( T 90 = 1.28 ± 0.09 s; Antonelli et al. 2009 ; Levesque et al. 2010 ; Thöne et al. 2011 ) at z = 2.609 (), the certainty in class membership for GRB 111117A is much higher. Additionally, Horváth et al. (2010) classify both GRB 060121 and GRB 090426 as intermediate-duration bursts because both events have very soft spectra, compared to the hard ones typically seen in sGRBs. Intermediate bursts are very clearly related in their properties to lGRBs ( de Ugarte Postigo et al. 2011 ), so they are unlikely to come from compact object mergers. GRB 111117A is also securely classified as a sGRB according to the Horváth et al. (2010) classification scheme. A number of other GRBs are thought to be short and at relatively high redshift ( z ≳ 0.9). If we consider bursts with a probability of being short f NC > 50%, according to the Bromberg et al. (2013) classification scheme, five sGRB are found: GRB 051210 at z ~ 1.3 ( f NC = ; Leibler Berger 2010 ), GRB 060801 at z = 1.131 ( f NC = ; Berger et al. 2007 ), GRB 070714 at z = 0.923 (no f NC due to extended emission; Graham et al. 2009 ), GRB 090510 at z = 0.903 ( f NC = ; McBreen et al. 2010 ), and GRB 100117 at z = 0.915 ( f NC = ; Fong et al. 2011 ). Although in individual cases a secure host association (hence redshift determination) is uncertain, there does seem to be a number of sGRBs at z ~ 1. This certainly makes GRB 111117A, by far, the highest redshift sGRB detected to date. The


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We thank Nick Battaglia, Rychard Bouwens, Camila Correa, Jeremy Tinker, and Rogier Windhorst for helpful discussions. We thank Rick Costa, Roberto Gomez, and Bill Wichser for help with computing. H.T. acknowledges support from NASA grant ATP-NNX14AB57G and NSF grant AST-1312991. R.C. acknowledges support from NASA grant NNX12AF91G. Simulations were run at the NASA Advanced Supercomputing (NAS) Division, Pittsburgh Supercomputing Center (PSC), and the Princeton Institute for Computational Science and Engineering (PICSciE).


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This work was supported by “Programme National de Physique Stellaire” (PNPS) and “Programme National de Physique Chimie du Milieu Interstellaire” (PCMI) from INSU/CNRS. WV’s research is funded by an ERC starting grant (3DICE, grant agreement 336474). D.S. acknowledges support by the Deutsche Forschungsgemeinschaft through SPP 1385: “The first ten million years of the solar system a planetary materials approach” (SE 1962/1-3)” This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France.


O O O O O O O O O O O O B-Instrument O O O B-Software I-Software 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-Software O B-Citation I-Citation I-Citation O O O O O O O O O O O O B-Dataset I-Dataset I-Dataset I-Dataset I-Dataset O B-Citation I-Citation I-Citation I-Citation O O O O O O O O B-Formula I-Formula I-Formula I-Formula B-CelestialObject O O B-Citation I-Citation I-Citation I-Citation O B-Citation O O O O B-Model I-Model I-Model O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O B-Telescope O O O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O 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-Dataset I-Dataset I-Dataset I-Dataset I-Dataset O B-Citation I-Citation I-Citation 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 B-Citation I-Citation I-Citation O O O O B-Wavelength I-Wavelength O O O B-Citation I-Citation I-Citation I-Citation O O B-Telescope 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 B-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 O O B-Citation I-Citation I-Citation I-Citation O O B-Survey B-Citation I-Citation I-Citation O O B-Survey B-Citation I-Citation I-Citation I-Citation O O B-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O B-Software I-Software I-Software I-Software 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 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 B-Mission O 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-Mission O O O O O 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

as Giles et al. (2018) . A pipeline was built for the CORALIE spectra based on iSpec 1 ( Blanco-Cuaresma et al. 2014a ). All observations were aligned and co-added to increase the signal-to-noise ratio (S/N), were reduced and spectrally fitted using the code SPECTRUM ( Gray Corbally 1994 ) as the radiative transfer code. Atomic data were obtained from the Gaia -ESO Survey line list ( Heiter et al. 2015b ). We selected the line based on an R ~ 47 000 solar spectrum ( Blanco-Cuaresma et al. 2016 , 2017 ), and we used MARCS model atmospheres ( Gustafsson et al. 2008 ). The resulting errors were increased by quadratically adding the dispersions found when analysing the Gaia benchmark stars ( Heiter et al. 2015a ; Jofré et al. 2014 ; Blanco-Cuaresma et al. 2014b ) with the same pipeline. This resulted in an effective temperature of 4877 ± 68 K, a log g of 3.41 ± 0.07 dex, and [Fe/H] = −0.24 ± 0.04 dex. In the second data release of Gaia ( Gaia Collaboration 2018 ), EPIC248847494 has a measured parallax (see Table 2 ) based on which we can determine an independent stellar radius using bolometric absolute magnitudes and the spectroscopically determined effective temperature for EPIC248847494 following the method detailed in Fulton Petigura (2018) . We took the K -band apparent magnitude ( Skrutskie et al. 2006 ), the Gaia distance, and a bolometric correction (BC K , from Houdashelt et al. 2000 ) of 1.91 ± 0.05, which was interpolated from the range within the coarse grid. We chose not to include an extinction correction as this only introduces an uncertainty of 0.5% ( Fulton Petigura 2018 ). This gave a radius of 2.70 ± 0.12 R sol . Table 2. Properties of the EPIC248847494 system. Notes. (a) Huber et al. (2016) , (b) APASS: Henden Munari (2014) , (c) 2MASS: Skrutskie et al. (2006) , (d) Gaia Collaboration (2018) , (†) This Work. Taking the spectrally determined metallicity and effective temperature and the measured radius as observational constraints, we input them into the Geneva stellar evolution code ( Eggenberger et al. 2008 ). This resulted in a stellar mass of 0.9 ± 0.09 M sol . These values of mass and radius would therefore indicate a log g of 3.52 dex. When we fixed the iSpec analysis to this log g , the metallicity and effective temperature were very similar to the initial results (see Table 2 ). Log g is not well constrained spectroscopically, and changes have a very limited effect on other parameters. Therefore we adopt the parameters based on log g = 3.52. 3.2. Eliminating the photometric systematics of K2 The possibility for false positives is high in monotransits. We therefore endeavored to eliminate all causes for false positives. All objects listed as “stars” with K2 light curves within 25 arcmin were checked for similar artefacts. Of the 61 objects, none showed odd behaviour at the same epoch as the monotransit. Additionally, the location of EPIC248847494 was not near the edge of the CCD, which suggests that no near-edge effects


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This work was carried out using the computational facilities of the Advanced Computing Research Centre, University of Bristol. 1 PC thanks the University of Bristol for a Postgraduate Research Scholarship. We thank the referee, Erin Kara, for very useful comments which have improved the paper.


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The authors are thankful to A. A. Osipova (Pulkovo Observatory) for preparing sunspot data for Table 1 . AN and LK acknowledge the support by the Russian Foundation for Basic Research (project 17-52-80064_BRICS) and by budgetary funding of Basic Research programme II.16. DB acknowledges the support from DST, India for BRICS project grant.


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We would like to thank Dr Marco Milla, the director of Radio Observatorio de Jicamarca, for his help of initiating this experiment and collecting data. We thank JRO staff for performing the observations. We also thank Karim Kuyeng and Ramiro Yanque for clarifying many questions regarding these observations. JRO is a facility of the Instituto Geofisico del Peru operated with support from the National Science Foundation (NSF) grant AGS-0905448 through Cornell University. This effort was supported under NSF grants ATM 07-21613 and AGS 12-02019 to Pennsylvania State University.


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We are grateful to the anonymous referee and to Junichi Baba, Tomoki Morokuma, Shinya Komugi, Robert C. Kennicutt, Daisuke Iono, Tetsuo Hasegawa, Tsutomu Takeuchi, Kouichiro Nakanishi, Fumi Egusa, and Nario Kuno for enlightening comments and suggestions. We thank Eric Feigelson for his constructive comments on the statistical aspects. K.M. thanks Ross Burns for the English language editing.


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

The Liverpool Telescope is operated on the island of La Palma by Liverpool John Moores University in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias with financial support from the UK Science and Technology Facilities Council (STFC). C.G.M. acknowledges support from the Royal Society, the Wolfson Foundation, and STFC. D.K. acknowledges the financial support from the Slovenian Research Agency (research core funding No. P1-0188).


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The authors would like to thank Heather Knutson for helpful discussions throughout the preparation of this manuscript. The authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Maunakea has always had within the indigenous Hawaiian community. We are most fortunate to have the opportunity to conduct observations from this mountain. The data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California, and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation. This work was partially supported by funding from the NSF Astronomy Astrophysics and NASA Exoplanets Research Programs (grants AST-1109857 and NNX16AI14G, G.A. Blake P.I.), and the Center for Exoplanets and Habitable Worlds, which is supported by the Pennsylvania State Unviersity, the Eberly College of Science, and the Pennsylvania Space Grant Consortium. Basic research in infrared astrophysics at the Naval Research Laboratory is supported by 6.1 base funding. Finally, we thank an anonymous reviewer for helpful insights which improved the content of this paper.


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

represents an effect of the detailed distribution of the magnetic field over the visible hemisphere of the MWD as viewed during our second observation. It is very interesting to consider our data for WD 1703−267 in the light of our earlier effort (see Sect. 3 ) to calibrate the order-of-magnitude relationship of Eq. ( 1 ). For this star, we have accurate values of both ⟨| B |⟩ and ⟨ B z ⟩, and judging from the simple V ∕ I profiles through the Balmer lines, a relatively simple magnetic field structure. With ⟨ B z ⟩≈ 2.75 MG in our second observation, we predict from our calibration of the coefficient of Eq. ( 3 ) that there should be continuum circular polarisation of V ∕ I ≈ 0.2%. Rebinning our data over 16 pixels to improve the S/N, we find that a weak continuum circular polarisation is indeed present at the level of V ∕ I ≈ 0.25% (see Fig. 4 ). Furthermore, the relationship between the sign of the continuum polarisation and that of V ∕ I in the blue σ components of the Balmer lines is the same as in G99-47. This result reinforces confidence in the use of broad-band circular polarisation as a tool for measuring ⟨ B z ⟩, and in the general accuracy of our calibration. 5.3 DA star WD 0810−353 = UPM J0812−3529 This star was identified as a WD within the 20 pc volume by the USNO parallax programme ( Finch et al. 2018 ), and is included in the list of WDs within the 20 pc volume established by Hollands et al. (2018) on the basis of the Data Release 2 dataset ( Gaia Collaboration 2018 ). For the reasons discussed in connection with WD 1703−267, we provide an (unofficial) Villanova WD number for the star, which we use throughout. Physical parameters based on Gaia astrometry and photometry ( Hollands et al. 2018 ; Gentile Fusillo et al. 2019 ) are provided in Table 1 . The spectral class is unknown, asno previous spectra of the star appear to have been taken. As the effective temperature is high enough that we would expect H β to be observable if the atmosphere of this object is H-rich, while possible metal lines around 4000 Å would be of great interest, we obtained circular spectropolarimetry of this star with the 600B and 1200B grisms. WD 0810−353 is bright enough that we were able to obtain quite high S/N (450–500 per Å). Fig. 4 WD 1703−267 (observations of September 30, 2019): circular polarisation (red solid line) and null profile (blue solid line, offset by − 3%), superposed to the V ∕ I error bars (light blue). Spectra are rebinned at about 11 Å. The three Stokes V ∕ I spectra are in quite good agreement in that they all show a region of weak positive polarisation around 4720 Å, and two regions of stronger negative polarisation around 4830 and 4960 Å. The I spectra obtained with grism 1200B show a series of weak depressions (with an amplitude of the order of 1%) that are hard to disentangle from instrumental features, even after applying flat-fielding correction. After an approximate normalisation to the continuum (clearly difficult to


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 O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O B-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 B-Dataset I-Dataset O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 B-Formula I-Formula I-Formula I-Formula 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 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 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 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 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 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 B-Wavelength 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 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 B-CelestialObject I-CelestialObject O O O O O O O O O O O B-Wavelength O O O O O O O O O O O

Note the systemic velocity is relative to the template file used. We do note the systemic velocity we find, Γ = 321.18 ± 0.85 km s −1 , is higher than the 295.8 ± 2.0 km s −1 previously reported (Corbet et al. 2016 ). We undertook additional analysis to confirm that the wavelength calibration performed by the HRS pipeline was correct and that the comparison to the radial velocity standards was accurate to within the expected performance. We found no evidence of any discrepancy in the calibration, nor any long-term systematic shift in the wavelength calibration over the period of observations. We believe that there are two possible reasons for this difference; there may be a possible systematic offset arising from the fit of the high-resolution template fit to the medium resolution MILES libraries (however there is no significant offset to the radial velocity standards) or there may have been a systematic offset in the zero velocity in the field O-type star used in the previous analysis. However, this difference does not change the main results of determining the orbital parameters. If this difference in the systemic velocity is removed the radial velocities are in agreement with the previous results. The final radial velocity curve and the best-fitting model are shown in Fig. 2 , and the data are given in Table 3 . We find the binary is slightly eccentric, e = 0.40 ± 0.07, and place the time of periastron at HJD 2457412.13 ± 0.29. Figure 2. Radial velocity of the O5 iii star in the binary system, as determined from cross-correlation. The solid line shows the best fit to this data, with a fixed 10.301 d orbital period. For clarity the plot is reported over two orbital phases with phase φ = 0 at MJD = 57 410.25, which corresponds to the given phase in Corbet et al. ( 2016 ). The error bars show the statistical errors reported by the rvsao package. Table 3. Radial velocities obtained from cross-correlation. 3 DISCUSSION 3.1 Mass of the compact object The binary parameter solution gives a mass function of f = 0.0010 ± 0.0004 M and the constraints on the mass of the compact object are shown in Fig. 3 . The mass of an O5 iii (f) star is ∼40 M (Martins et al. 2005 ), but due to the binary evolution could be different. A mass range of 25–42 M was considered in Seward et al. ( 2012 ). For this mass range both a neutron star or black hole mass is compatible with the mass function, though a neutron star is favoured. For a 40 M star, the mass of the compact object will be >5 M for inclinations i ≤ 15 ± 2°, and while the inclination must be i ≤ 11 ± 1° for a 25 M optical companion. Assuming the compact object is a pulsar, with a mass of 1.4 M , the inclination will lie between i = 39 ± 6° and i = 59 ± 11° for a 25 M and 40 M star, respectively. Figure 3. The constraint on the mass of the compact object in LMC P3. The shaded area marks the range of assumed masses of the optical companion. 3.2 Binary orientation The orientation of the binary system is


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ACKNOWLEDGEMENTS This work is supported in part by the National Natural Science Foundation of China under Grant No. 11675032 (China).


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

luminosity wind that emanates from a stellar surface or a galaxy, the interaction between the wind supplying the energy and the ambient gas forms a ‘bubble’. This bubble consists of a forward shock that advances into the ambient gas, a reverse shock that travels back through the wind, and a contact discontinuity that separates the two shocked fluids (e.g. Castor, McCray Weaver 1975 ; Weaver et al. 1977 ; Fielding, Quataert Martizzi 2018 ); a self-similar solution describes the fluid properties between the forward shock and the contact discontinuity. When the shock generated from a supernova explosion or an expanding wind is young, the large temperatures and densities at the shock front imply that the effective adiabatic index γ of the multispecies fluid satisfies 4/3 ≲ γ ≲ 5/3. As the shock further expands and cools, however, the high-density regions near the shock front start to radiate efficiently, which reduces γ below the value of an ideal monatomic gas. On the other hand, the low-density material in the interior of the explosion remains hot, and the expanding blast wave is characterized by a thin, nearly isothermal shell that is in contact with a hot, pressurized interior (e.g. McKee Ostriker 1977 ; Shull McKee 1979 ; Blondin et al. 1998 ). Once the shock becomes radiative, it is no longer strictly correct to treat the fluid as adiabatic. However, it is reasonable to assume that the gas exhibits an increased degree of compressibility prior to radiating a significant amount of its total energy, during which time it is still approximately correct to use the self-similar scalings (derived under strict adiabaticity) for the relationship between the shock position and velocity, but with an appropriately reduced adiabatic index. At the very least, the self-similar solutions with a nearly isothermal equation of state describe the qualitative behaviour we expect at this stage – a nearly evacuated, extremely hot interior in contact with a thin shell that contains the majority of the mass (e.g. Fig. 1 ). Understanding the structure of these shocks therefore yields insight into the more general behaviour of radiative, thin shells formed from explosions and stellar and galactic winds. Figure 1. A comparison of the exact solution for the Sedov–Taylor blast wave (solid lines) and the analytical, leading-order solution in the quantity δ = (γ − 1)/(γ + 1) given by equation ( 13 ) (dashed lines) for the adiabatic indices γ and radial power-law indices n shown in the legend; recall that the ambient density follows the power-law decline ρ ∝ R − n , where R is the shock position. Solutions are plotted as functions of the standard Sedov–Taylor variable ξ = r / R ( t ), which is just the spherical radius at a given time normalized by the shock radius. The left, middle, and right columns show the radial velocity normalized by the shock velocity, the density normalized by the ambient density (and reduced by the factor δ), and the pressure normalized by the


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This work is based on observations carried out with the IGN 40 m radio telescope. We thank all the staff in Yebes Observatory for the generous help received. We acknowledge partial funding from MINECO grants AYA2012-32032, FIS2012-32096, and FIS2012-38160. This work has made use of the American Association of Variable Star Observers (AAVSO) database and of the SIMBAD Astronomical database (operated at CDS, Strasbourg, France).


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2009 , 2010 ; Lyra Klahr 2011 ; Lyra 2013 ). Yet it remains 2–3 orders of magnitude lower than the radial and azimuthal velocity components Figure 3 . Figure 1. Vertical vorticity of unstratified gas in a three-dimensional box at different times. The convective overstability amplifies seed noise into growing vortices, that intensify and merge viscously, saturating at radial length similar to the pressure scale height. Figure 2. The vertical component of the vorticity () in the midplane (left panels) and a vertical cut through the disk at the azimuthal location of maximum negative vorticity (indicated by the dashed line). Also shown is the vertical velocity component ( u z ). The upper plots are shown at 100 local orbits (top panels). At this time, there is no structure in the vertical direction: the flow is essentially two-dimensional. After 275 local orbits (bottom panels) there is now turbulent motion inside the vortex, which is also visible in the vertical velocity. This is an indication of the growth of elliptical instability in the vortex cores. Figure 3. Gas velocity components averaged over the entire box. As the elliptical instability sets in, the vertical velocity component (green line) increases steeply, while staying well below the azimuthal and radial components. 4.1. Pebble Evolution We first discuss the run without baroclinic driving. As expected, the pebbles settle to the midplane, soon triggering the streaming instability. Although we assumed an overall initial dust-to-gas ratio of , the initial dust-to-gas ratio in the midplane is already and increases even further as the sedimentation progresses. The maximum dust-to-gas ratio of the simulation with pebbles and without baroclinic driving can be seen in Figure 4 (red line). The other two lines in the plot show the maximum dust-to-gas ratios for the simulation with , (green line), and pebbles (black line). Figure 4. Maximum dust-to-gas ratio for simulations with (black line) and (green and red line) pebbles. indicates that no baroclinic feedback was included in the simulation. Although the high pebble concentration in the simulations with baroclinic effects disrupts the vortex in the midplane, where pebbles are located, the maximum dust-to-gas ratio is higher than without baroclinic effects. The maximum concentration for with baroclinic driving is only a factor of 2 higher than in the case without it. Inspecting the evolution of the dust density in the baroclinic simulation, we see that, as expected, the pebbles settle to the midplane and get trapped in the vortices. Yet, after only three orbits, the pebbles start to perturb the vortices in the midplane via the streaming instability. Like in the 2D simulations of Raettig et al. ( 2015 ), high pebble densities generate local perturbations of the velocity field visible as steep vorticity gradients. Figure 5 shows the vorticity at the top of the box (left panel) and in the midplane (middle panel), and the dust-to-gas ratio in the


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We thank Michele Bellazzini for a very careful reading and useful comments and suggestions that helped to improve the quality of the paper. We also thank the anonymous referee for comments that greatly improved this paper. Based on data products from observations made with ESO Telescopes at the La Silla Paranal Observatory under program ID 188.B-3002. These data products have been processed by the Cambridge Astronomy Survey Unit (CASU) at the Institute of Astronomy, University of Cambridge, and by the FLAMES/UVES reduction team at INAF/Osservatorio Astrofisico di Arcetri. These data have been obtained from the Gaia -ESO Survey Data Archive, prepared and hosted by the Wide Field Astronomy Unit, Institute for Astronomy, University of Edinburgh, which is funded by the UK Science and Technology Facilities Council. This work was partly supported by the European Union FP7 program through ERC grant number 320360 and by the Leverhulme Trust through grant RPG-2012-541. We acknowledge the support from INAF and Ministero dell’ Istruzione, dell’ Università’ e della Ricerca (MIUR) in the form of the grant “Premiale VLT 2012”. The results presented here benefit from discussions held during the Gaia -ESO workshops and conferences supported by the ESF (European Science Foundation) through the GREAT Research Network Program. I.S.R. gratefully acknowledges the support provided by the Gemini-CONICYT project 32110029. C.M. acknowledges the support from CONICYT-PCHA/Doctorado Nacional/2014-21141057. D.G. gratefully acknowledges support from the Chilean BASAL Centro de Excelencia en Astrofísica y Tecnologías Afines (CATA) gran PFB-06/2007. S.V. gratefully acknowledges the support provided by Fondecyt reg. 1130721. A.K. and N.K. acknowledge the Deutsche Forschungsgemeinschaft for funding from Emmy-Noether grant Ko 4161/1. U.H. acknowledges support from the Swedish National Space Board (SNSB). S.G.S acknowledges the support from the Fundação para a Ciência e Tecnologia, FCT (Portugal) and POPH/FSE (EC), in the form of the fellowships SFRH/BPD/47611/2008.


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

values similar to the ratios in local SFGs with high EW(H α ). This increase results in high EW([O III ] 5007) > 1000 Å for bright z > 7 galaxies ( Roberts-Borsani et al. 2016 ). The properties of the high- z SFGs discussed above contrast with those of the z ∼ 0 main-sequence SFGs. The latter exhibit, on average, a much lower level of star formation activity and, thus, they are characterised by much lower SFRs, sSFRs, and EWs of emission lines at a fixed stellar mass. However, there is a small fraction of low- z SFGs with compact structure and other properties that are similar to those of high- z galaxies. Subsets of these compact star-forming galaxies (CSFGs) are variously called blue compact dwarf (BCD) galaxies, ‘green pea’ (GP) galaxies, and luminous compact galaxies (LCGs) in the literature. BCDs are nearby dwarf galaxies at redshift z typically ≤0.01 − 0.02, with strong bursts of star formation and low oxygen abundances of 12 + logO/H ≲ 7.9, extending down to 12 + logO/H ≲ 7.0 (e.g., Thuan Martin 1981 ; Izotov et al. 1994 , 2018c ). Typically, their stellar masses are low, M ⋆ ≲ 10 8 M ⊙ . A subset of the lowest-mass BCDs with M ⋆ ∼ 10 6 –10 7 M ⊙ and very high O 32 ≳ 10 has been named ‘blueberry’ galaxies by Yang et al. (2017) because of their compact structure and intense blue colour on the Sloan Digital Sky Survey (SDSS) composite images. Cardamone et al. (2009) considered ‘green pea’ (GP) galaxies in the redshift range of z = 0.112–0.360, with unusually strong [O III ] λ 5007 emission lines, similar to high- z UV-luminous galaxies, such as Lyman-break galaxies and LAEs, with sSFRs up to ∼10 Gyr −1 . The green colour of these galaxies on composite SDSS images is due to strong [O III ] λ 5007 line emission (with EWs up to 1000 Å), redshifted into the SDSS r -band. Izotov et al. (2011) studied a much larger sample of 803 star-forming luminous compact galaxies (LCGs) in a larger redshift range z = 0.02–0.63, with global properties similar to those of the GPs. The sSFRs in LCGs are extremely high and vary in the range ∼1–100 Gyr −1 , comparable to those derived in high- z SFGs. Similar high sSFRs were also found for low- z LyC leaking galaxies ( Izotov et al. 2016a , b , 2018a , b ; Schaerer et al. 2016 ) and some other low- z SFGs (e.g., Senchyna et al. 2017 ; Yuma et al. 2019 ). Izotov et al. (2016c) found sSFRs of up to 1000 Gyr −1 in a sample of ∼13 000 CSFGs selected from the SDSS. Compared to BCDs, both GPs and LCGs are characterised by higher stellar masses ∼10 9 M ⊙ and higher oxygen abundances ∼7.8–8.3. Izotov et al. (2015) selected ∼5200 low-redshift (0 z 1) CSFGs and studied the relations between the global characteristics of this sample, including absolute optical magnitudes, SFRs, stellar masses, and oxygen abundances. They found that for all relations, low- z and high- z SFGs are closely related, indicating a very weak dependence of metallicity on stellar mass, redshift, and star formation rate. This finding favours the


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We thank the anonymous referee for helpful suggestions. This work is supported by the National Basic Research Program (973 Programme) of China (grant No. 2014CB845800), the National Natural Science Foundation of China (Grants 11673006, 11533003, 11303005, 11851304, 11773007, U1731239), the Innovation Team and Outstanding Scholar Program in Guangxi Colleges, the One-Hundred-Talents Program of Guangxi colleges, and the Guangxi Science Foundation (2016GXNSFFA380006, AD17129006, 2016GXNSFDA380027).


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). Tanvir et al. ( 2005 ) found a stronger correlation of BATSE SGRBs with early-type than late-type galaxies, which is not expected if nearby SGRBs arise from giant magnetar flares in nearby galaxies, but is consistent with a BNS origin. GRB 170817A is clearly associated with a BNS merger, but even without the connection to GW170817, the spectrum of GRB 170817A in the GBM data strongly disfavored the BB fit expected for a giant magnetar flare. Another possible signature of a giant magnetar flare is the ringing in its tail at the NS period of a few seconds, which could be detected by GBM or by SPI-ACS for flares outside our galaxy providing it was close enough. A search for periodic or quasi-periodic emission in the GBM data for GRB 170817A (Goldstein et al. 2017 ) found no periodic modulation, providing another discriminant between SGRBs and nearby extragalactic giant magnetar flares that might be masquerading as SGRBs. 6.4. Predicted Detection Rates The intrinsic specific volumetric SGRB rate is often quoted to be around 10 Gpc (see, e.g., Guetta Piran 2006 ; Coward et al. 2012 ; Fong et al. 2015 ; with the value originating from Nakar et al. 2006 , who noted that the true rate could be much higher). However, unlike GW signals, SGRBs do not have a clear relationship between the observed distance and brightness. As discussed in the previous section, this can be due to intrinsic variations in SGRB luminosities, as well as structure in the jet. In this Section, we investigate the former scenario presenting the implications of GW170817/GRB 170817A for future GW and SGRB observations in terms of a simple standard model for the SGRB luminosity distribution. Similar interpretations for other, perhaps more elaborate, models are straightforward. We model the SGRB luminosity function as a broken power law, with a logarithmic distribution 186 186 To get the linear distribution of luminosities, both and must be increased by 1. where is the peak isotropic luminosity (in the source frame) between 1 keV and 10 MeV, and and give the power law decay below and above the break at L ⋆ . 187 187 Other studies use a smaller energy band when defining the luminosity, and this has an impact on the value of L ⋆ , although not on the slopes of the power law components. Here, we follow Wanderman Piran ( 2015 ) in using erg s −1 , , and . The other important parameter is the minimum SGRB luminosity, which determines the lower cutoff of the luminosity distribution. This is poorly constrained as only nearby low luminosity SGRBs are observable. In Wanderman Piran ( 2015 ) the minimum luminosity is taken to be erg s −1 , while other studies use values ranging from erg s −1 to erg s −1 (Regimbau et al. 2015 ). We assume a threshold value for detectability in GBM of 2 photons cm −2 s −1 for the 64 ms peak photon flux in the 50–300 keV band, which is higher than the minimum detectability value to account for the sky-dependent sensitivity of GBM. Furthermore, we model the SGRB