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

of photometric long-term cycles in M dwarfs). Fig. 10 GLS periodograms of CARMENES activity indicators of G 264–012 in the VIS ( top panels ) and NIR ( bottom panels ) channels. Horizontal FAP lines and vertical signal indications are the same as in Fig. 3 . Since the CARMENES spectral range does not include the Ca II HK lines, we carried out an analysis of all available archival data covering this spectral region, namely the HARPS and HIRES spectra used in the RV analysis of this work (Sect. 4.1.2 ), 1 spectrum acquired by FEROS ( Kaufer et al. 1999 ), 3 spectra from UVES ( Dekker et al. 2000 ), and 10 from XSHOOTER ( Vernet et al. 2011 ). In order to extract from spectra acquired by all five instruments in a uniform manner, we rectified the spectra with a grid of PHOENIX spectra ( Husser et al. 2013 ) by extracting a set of narrow band passes around both the H and K lines. The approach will be described in detail in an upcoming publication (Perdelwitz et al., in prep.). The top panel of Fig. 12 shows the resulting time series. The smaller periodicity in the bottom panel of Fig. 12 , with a period of P = 36.6 ± 1.9 d, where the error was determined via a Monte Carlo approach, is in agreement with the stellar rotation period derived from Kepler K2 photometry (Sect. 3.1 ). Fig. 11 GLS periodograms of CARMENES VIS, HARPS, and HIRES activity indicators for Gl 393. Horizontal FAP lines and vertical signal indications are the same as in Fig. 4 . Fig. 12 Top panel : time series of Gl 393. The source of the archival data is denoted by the colour coding. Bottom panel : generalised Lomb-Scargle periodogram of the time series of Gl 393. The dashed vertical lines denote the periods of the activity cycle ( P ≈ 1226 d) and stellar rotation ( P ≈ 36.6 d) determined from the time series. FAP levels of 10 −3 and 10 −9 are displayed as horizontal lines. 4.5 Gaussian process analysis of RV data 4.5.1 G 264–012 All indicators gathered (Sect. 4.2 – 4.4 ) point towards activity as origin for the RV signal found at a period of P d = 92.70 d in G 264–012. There is, however, no indicator contradicting the planetary nature of the other two signals, P b = 2.30 d and P c = 8.05 d. In order to accurately determine their orbital parameters and masses, we account for stellar activity in our models. We choose a model with two circular Keplerian orbits (no evidence for eccentricity was found in these two very-short period planets, whose orbits should be almost certainly already circularised, see Sect. 4.1 ) for the planets combined with a Gaussian process (quasi-periodic kernel; see Eq. ( 1 )) for the activity signal. We present the fits in Fig. 13 , a summary of all parameters used in the model in Table 5 and the MCMC posterior distribution plot in Fig. A.1 . When subtracting the model from the VIS RV data we derived an rms of 2.22 m s −1 , which is consistent with the errorbars. Fig. 13 G 264–012. Panel a : GP fit (black solid line) and 1-sigma uncertainties (grey


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

of magnetic variability. We demonstrate here that using multiband photometric light curves contemporaneous to transit observations, we can overcome this problem, at least partly. The advantage of using several photometric bands is that, as discussed in the previous section (see also Mallonn et al. 2018 ), it is possible to independently estimate the zero point with respect to the unspotted photosphere ( z ), and thus the absolute spot filling factor (including any persistent level of spottedness during the rotation cycle), and the temperature contrast. This provides extremely valuable information to correct transit depth variations due to spots. We can actually use our StarSim simulations of WASP-52 to investigate at which level of accuracy we can correct the effects of starspots on transit data. The StarSim simulations of the WASP-52 planetary system discussed in the previous section reveal that, without additional information allowing us to infer the distribution and filling factor of spots during the transit, the planet-to-star radius cannot be determined with an accuracy better than a few percent. This is because in that case, we do not know at which rotation phase of the star the transit is observed, and therefore, we can only estimate a range of possible filling factors based on the amplitude of the photometric modulation. However, if we know the rotation phase at the time of transit, and we can estimate the spot properties and distribution from light curves, we can infer the correction that needs to be applied to the relative radius of the exoplanet (or transit depth) with an accuracy of order 10 −4 , as shown in Table 4 . Figure 10 illustrates the sequence of spot corrections on the transit depth as a function of time. We plot the planet radius variation due to unnoculted spots, Δ r SP ∕ R * , following our StarSim model describing the photometric light curve in Fig. 8 . Each panel corresponds to a different wavelength band of interest for the Ariel mission, as an example. To test all possible filling factors, we consider transits occurring at any time, but actual WASP-52 b transits are indicated as filled dots. The average value of all the simulations performed is displayed as a solid line. It is clear from this figure that the variation of the apparent planet radius depends on the projected spot filling factor and follows the stellar rotation period. The relative effect on the radius diminishes towards longer wavelengths due to the lower flux contrast of spots, from an average of ~4 × 10 −3 to ~8 × 10 −4 at 550 nm (VIS-Phot) and 6 μ m (AIRS-Ch1), respectively. In all cases it is possible to estimate the unocculted spot correction to the planetary relative radius with a precision close to 10%, estimated as the standard deviation ( σ ) of the sample of 21 optimal fits to the light curves. As explained above, this residual uncertainty is due to the variance of the stellar parameters and the randomness of the determination of the active


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ACKNOWLEDGEMENTS ZL thanks the support from the China Scholarship Council. This work is supported by the Strategic Pioneer Program on Space Science, Chinese Academy of Sciences, Grant No. XDA15052100. PTOB acknowledges support from STFC. JPO, PAE, and KLP acknowledge support from the UK Space Agency. This work made use of data supplied by the UK Swift Science Data Centre at the University of Leicester. This work is based on observation obtained with XMM–Newton , an ESA science mission with instruments and contributions directly fund by ESA Member States and NASA, based on observations made with the NASA/ESA Hubble Space Telescope , and obtained from the Hubble Legacy Archive, which is a collaboration between the Space Telescope Science Institute (STScI/NASA), the Space Telescope European Coordinating Facility (ST-ECF/ESA) and the Canadian Astronomy Data Centre (CADC/NRC/CSA) The Chandra data are obtained from the Chandra Data Archive . This research has made use of software provided by the Chandra X-ray Center (CXC) in the application packages ciao .


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

appropriate choice of the arbitrary function t 0 ( R ). Let us now find the general expressions for t sc ( R ) and t hor ( R ) for each type of solution ( 14 )–( 16 ). For further purposes, it is convenient to rewrite the solution in dimensionless unitsIn the rest of our paper, we shall use these dimensionless units and thus prefer to omit the tilde over the letters for brevity. The solution ( 14 )–( 16 ) in dimensionless units ( 26 ) reads: 1. The case of flat space: 2. The case of closed space: 3. The case of open space: By solving Equations ( 23 ) and ( 24 ) successively for each case ( 27 )–( 29 ), we obtain: 1. For k = 0 2. For k = 1 3. For k = −1 Here, in our calculations, we have used the fact that β is very large for realistic objects (estimative size of the universe is about m, and r g of a typical black hole in the galaxy center is about ∼10 10 m, i.e., one has β ∼ 10 16 ). Thus, we have employed the Taylor expansion for large β to get rid of transcendental equations for the open and closed spaces. Analyzing the expressions ( 30 )–( 32 ) and applying the condition ( 25 ), we come to the following form of the arbitrary functions t 0 ( R ): It is easy to see that these functions have the same limit at small R and the result ( 17 ) remains valid. To give an example of the behavior of the obtained model, we present in Figures 2 – 3 the profiles of the energy density and the Hubble parameter for the parabolic case of our solution ( 14 ). The energy density is given by ( 19 ). The Hubble parameter describing the rate of expansion of the inhomogeneous universe is generally defined according to Ellis ( 2009 ) as followswhere l is some characteristic length that corresponds to the scale factor a ( t ) in Friedmann models, and τ is proper time given in the standard way:Θ is the volume expansion given bywhere u α is usual 4-velocity of the dust particles. Figure 2. Energy density profile for the parabolic case of our solution given in dimensionless units. (a) At the initial (Big Bang) singularity there is an infinite value of the energy density that is decreasing with time, while in the center of symmetry there always holds ε ( R = 0, t ) = 0. (b) The structure of the profile at the initial singularity in the vicinity of the central point . There is a punctured infinite energy density cusp, which is rapidly decreasing with time. (c) The character of the energy density dependence on time for different values of R . The dashed curve corresponds to the parabolic Friedmann solution. The bigger is R , the closer is ε ( R , t ) to the Friedmannian line. (d) The picture at the epoch of the Milky Way birth. The number of each curve corresponds to the index of the value t (or R ) listed in the legend. Figure 3. The Hubble parameter profile for the parabolic case of our solution given in dimensionless units. (a) The line t = t +0( R ) marks the line of the initial singularity. (b) The structure of the profile at the initial singularity in the vicinity


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This research has made extensive use of data and/or software provided by the HEASARC, which is a service of the Astrophysics Science Division at NASA/GSFC and the High Energy Astrophysics Division of the Smithsonian Astrophysical Observatory. We also acknowledge the use of NASA's Astrophysics Data System (ADS) and of the SIMBAD data base, operated at CDS, Strasbourg, France. In this work, we analysed observations performed with INTEGRAL , an ESA project with instruments and science data centre funded by ESA member states (especially the PI countries: Denmark, France, Germany, Italy, Switzerland, Spain) and with the participation of Russia and the USA, and we also analysed observations obtained with XMM–Newton , an ESA science mission with instruments and contributions directly funded by ESA Member States and the USA (NASA). Moreover, this research has made use of data obtained from the Chandra Data Archive.


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

by stochastic kicks as seen by the torque fluctuating about zero, resulting in small inward and outward migration. The bottom panel is in a constant β disc, where it is clear that despite stochastic kicks being present, the torque, on average, is slightly negative, resulting in inward migration. Figure 14. Total z -component of the torque exerted by the disc material on a 1 M ⊕ planet is shown on the left-hand panels. The migration tracks are shown on the right-hand panels. The top five panels are simulations modelled with a variable β , whilst the bottom panel is one of the simulations modelled with a constant β . The migration is dominated by stochastic kicks, which results in random outward or inward motion, as seen in the torques oscillating about zero. 5.2 Comparison with previous work We compare our results to those of Baruteau et al. ( 2011 ) and Malik et al. ( 2015 ), who considered giant planet migration in self-gravitating discs. There are a couple of differences in this work. They perform their simulations using a 2D grid-based code, whereas the simulations in this work are performed using a 3D SPH code. The initial conditions of the disc are also slightly different. Their disc surface density and temperature decreased as R −2 and R −1 , respectively, whereas in this work, the disc surface density and temperature decreases as R −1 and R −0.5 , respectively. To ensure that the different initial conditions and codes used do not influence the results, we initially performed a set of simulations with a constant value of β = 15 to compare with Baruteau et al. ( 2011 ) and Malik et al. ( 2015 ). Despite the differences, we agree with their findings that using a constant β results in the giant planets rapidly migrating towards the inner edge of the disc. However, the crucial development in our work is how the disc is cooled: Whilst we also utilize the β -cooling approach, we allow β to vary with radius to mimic a more realistic disc such that our discs remain gravitationally stable in the inner regions, and the spiral structure only forms in the outer regions. The planets are able to slow down their rapid inward migration before reaching the inner disc edge. Using radiative transfer, Stamatellos Inutsuka ( 2018 ) migrate 1 M Jup planets in self-gravitating discs using a 3D SPH code. The initial conditions in their discs are more similar to ours: the same initial disc mass and surface density profile. Their usage of a realistic treatment of the disc thermodynamics meant that like in the simulations presented here, the planets are evolved in discs that are only gravitationally unstable in the outer regions. Using realistic thermodynamics to allow the gas to control its cooling/heating based on its properties, the thermodynamics of the circumstellar and circumplanetary disc is better captured. This enhances the mass growth of the planet causing the planets to grow to beyond the brown dwarf limit. This resulted in their planets opening up


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feature tracking. 3. Results The polar crown prominence in the northwest limb (N63W72) had been stable for over two weeks before our observation began on 2016 June 4. Figure 1 shows the prominence of interest in the H α , AIA 131, 171, and 304 Å images. The prominence is bright in H α , but dark in the AIA images. Note that the intensity of the H α image is reversed, and that darker features in the image are actually brighter. The dark structure in the AIA images is wrapped in the bright emissive gas, indicating that the cool material is wrapped inside the hot coronal gas. The main body of the polar crown prominence lies vertical to the plane of the sky, and an arch-shape prominence newly emerged in the northern region of the polar crown prominence after 04:00 UT. The arch-shape prominence was brighter than the main prominence in the AIA images, but it was fainter in the H α images. The observed arch-shape prominence looks similar to the hedgerow prominence overlying a cavity observed by Ryutova et al. ( 2010 ) except for the horizontal flow within it. We detected a horizontal flow, which is highlighted in the next paragraph. Figure 1. Polar crown prominence observed on 2016 June 4, with the newly emerged arch-shaped prominence in the region north of it. The SDO/ AIA 131, 171, 304 Å data, and H α images (upper-left, upper-right, lower-left, and lower-right panels, respectively) show the formation of a vortex on the arch-shaped prominence. Note that the intensity of the H α image is reversed. The dashed boxes in the H α image represent the FOV of the Figure 3 (small red), Figure 4 (big red), and Figure 8 (blue). A full sequence of the images, covering from 08:09:43 to 09:31:46 UT, is provided in the animation. In the animation, 08:09:43 UT is taken to be the epoch of time t . (An animation of this figure is available.) Figure 2 shows H α images of the observed region at the moment of the rising of the arch-shape prominence and the formation of a 5′′-size vortex. The arrows plotted on the H α images represent the velocity vectors of the stream on the plane of the sky. Near the left-hand side leg of the arch-shape prominence (690′′, 680′′), the dominant direction of the plasma flow is upward along the arch. It rose up slowly (20 km s −1 ) and reached about a height of 20′′. Then it streamed to the right along the arch toward the polar crown prominence with a speed of 20–60 km s −1 . The stream can be clearly identified in both panels of Figure 2 . The measured collimated flow speed of about 20–60 km s −1 is higher than the typical speed of the prominence plasma (10–20 km s −1 ; Zirker et al. 1998 ; Ahn et al. 2010 ; Vial Engvold 2015 ). The full dynamics of the evolution can be seen better in the animation of Figure 1 . Very luckily, the flow lies parallel to the plane of the sky considering the narrow bandpass width of the filter. Figure 2. Velocity vectors superposed on H α images at the moment of the rising of the arch-shape prominence (left)


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Acknowledgements Daniel Thun was funded by grant KL 650/26 of the German Research Foundation (DFG). Most of the numerical simulations were performed on the bwForCluster BinAC. 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. All plots in this paper were made with the Python library matplotlib ( Hunter 2007 ).


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M. Vitale is supported by the International Max-Planck Research School (IMPRS) for Astronomy and Astrophysics at the Universities of Bonn and Cologne. This work is partly based on observations with the 100-m telescope of the MPIfR (Max-Planck-Institut für Radioastronomie) at Effelsberg. The authors would like to thank Alex Kraus for support and help during observations and data reduction. The FIRST Survey is supported in part under the auspices of the Department of Energy by Lawrence Livermore National Laboratory under contract W-7405-ENG-48 and the Institute for Geophysics and Planetary Physics. The Sloan Digital Sky Survey is a joint project of the University of Chicago, Fermilab, the Institute for Advanced Study, the Japan Participation Group, Johns Hopkins University, the Max Planck Institute for Astronomy, the Max Planck Institute for Astrophysics, New Mexico State University, Princeton University, the United States Naval Observatory, and the University of Washington. Apache Point Observatory, site of the SDSS, is operated by the Astrophysical Research Consortium. Funding for the project has been provided by the Alfred P. Sloan Foundation, the SDSS member institutions, NASA, the NSF, the Department of Energy, the Japanese Monbukagakusho, and the Max Planck Society. The SDSS Web site is http://www.sdss.org . This work was supported in part by the Deutsche Forschungsgemeinschaft (DFG) via the Cologne Bonn Graduate School (BCGS), and via grant SFB 956. We had fruitful discussions with members of the European Union funded COST Action MP0905: Black Holes in a violent Universe and the COST Action MP1104: Polarization as a tool to study the Solar System and beyond. We received funding from the European Union Seventh Framework Program (FP7/2007-2013) under grant agreement No. 312789.


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

) and the 321 GHz H 2 O maser line (Hirota et al. 2014a ). The brightness temperature of the 658 GHz H 2 O line, >2 × 10 4 K, would indicate that the line is emitted via maser action. In contrast, the 336 GHz H 2 O line shows a different distribution (Figure 5 (c)) and a spectral profile with symmetric double-peaks (Figure 6 ). As discussed in Alcolea Menten ( 1993 ) and Hirota et al. ( 2014a ), the 336 GHz H 2 O line could be a thermal line and trace the edge-on disk, unlike emission from the other transitions at 232, 321, and 658 GHz. We note that via Menten et al. ( 2006 ) the 336 GHz H 2 O line toward the red supergiant VY CMa also favors a thermal nature for this line’s excitation. Recently, theoretical model calculations for the excitation of the 658 GHz H 2 O line have been presented (Richards et al. 2014 ; Nesterenok 2015 ) for the red supergiant VY CMa and for asymptotic giant branch (AGB) stars. The 658 GHz H 2 O maser is predicted to be pumped efficiently for H 2 densities of 10 9 –10 10 cm −3 , a fractional H 2 O abundance of 10 −4 –10 −5 , and gas temperatures of 1000–1500 K (Nesterenok 2015 ). Although the physical and chemical properties of the 658 GHz H 2 O maser emitting region associated with Source I would be expected to be different from that of a late-type star, these parameters are roughly consistent with those for the disk-outflow system of Source I associated with SiO masers and other submillimeter H 2 O lines (Reid et al. 2007 ; Goddi et al. 2009 , 2011 ; Hirota et al. 2014a , 2015 ). 4.4. Relationship between the 22 GHz H2O Supermaser We note that our ALMA observation was done during an active phase of the 22 GHz H 2 O maser burst in Orion KL, called a “supermaser” (Hirota et al. 2011 , 2014b ). The 22 GHz supermaser is located at the phase tracking center of the ALMA observation, corresponding to the Compact Ridge region (Hirota et al. 2011 , 2014b ). While the total flux density of the 22 GHz H 2 O supermaser was decreasing on 2012 August 25, it was still Jy (Hirota et al. 2014b ). In spite of the extremely high flux density of this maser, there was no significant emission of the 658 GHz H 2 O line in the Compact Ridge region in our ALMA data with the 1 σ noise level of 2 Jy beam −1 in the relevant region. Thus, the 658 GHz H 2 O line is not related to the 22 GHz supermaser phenomenon, as has also been reported for other transitions at 321 and 336 GHz (Hirota et al. 2014b ). 4.5. Continuum Emission at ALMA Band 9 To measure the position of the 658 GHz H 2 O emission, we also made a map of the ALMA band 9 continuum by applying the self-calibration solutions. Continuum emission is detected at the position of Source I, although it is outside the primary beam of the 12 m antenna. The flux density and the peak intensity are derived from Gaussian fitting, respectively 1.04 ± 0.04 Jy and 0.77 ± 0.03 Jy beam −1 (Figure 2 (b)). If we correct the primary beam attenuation by a factor of 8.6, the absolute flux


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


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

complex structure within the overall galaxy mass functions. For instance, in increasingly higher halo mass regimes, we observe high-mass drop-offs that occur at higher galaxy mass. These drop-offs mark the natural boundary in the largest central galaxy mass for a given halo mass, as seen in Figure 5 (b). Figure 13. Breakdown of RESOLVE-B (top row), L r ,tot -limited RESOLVE-B (middle row), and ECO (bottom row) SMF and BMF into different group halo mass regimes. The four group halo mass regimes are M halo 10 11.4 M ⊙ (low-mass groups, often solo centrals, blue), M halo between 10 11.4 and 10 12 M ⊙ (intermediate-mass groups, green), M halo between 10 12 and 10 13.5 M ⊙ (large group, orange), and M halo > 10 13.5 M ⊙ (cluster, red, ECO only). The dark gray marks the overall SMF or BMF for each data set, and the gray hash-marked region denotes the incomplete regions for each mass function. The conditional mass functions are much more complex than the overall mass function, with pronounced bumps and dips. In both the RESOLVE-B and ECO intermediate (green) and large (orange) group halo mass regimes, we observe a peak in the SMF and BMF. For intermediate group halos, this peak occurs at ∼10 10.2 M ⊙ , right between the gas-richness threshold and bimodality mass scales. For large group halos, the peak occurs at ∼10 10.8 M ⊙ , above the bimodality mass scale and also near the knee of the overall mass functions. The peak in the large group halo regime appears to be causing the pronounced “bump” seen in the overall RESOLVE-B SMF and BMF. While the ECO data set also has this characteristic peak in its large group halo regime mass functions, the more gradual decline in the overall SMF and BMF for large galaxy masses seems to be due to the cluster galaxy population in ECO. The cluster population makes up 17% of galaxies with M star > 10 10.5 M ⊙ and 31% of galaxies with M star > 10 11 M ⊙ for ECO. The RESOLVE-B “bump,” meanwhile, may be emphasized by the lack of a cluster galaxy population and the overabundance of halos in this large group halo regime, as shown in Figure 5 (a). It is clear that the low-mass slopes of the SMF and BMF in these different halo mass regimes are very different. In the cluster halo mass regime, the rise is quite smooth, although the slope appears to flatten out at lower masses, at least in part due to incompleteness in galaxy counts around the clusters caused by high fiber collision rates and by missing ultradiffuse galaxies like those recently found in the Coma and Virgo clusters (Mihos et al. 2015 ; van Dokkum et al. 2015 ). We note that Yamanoi et al. ( 2012 ) found a fairly flat slope for the Coma cluster LF down to galaxy magnitudes of M R = −14.0, with a steep upturn for galaxies below our luminosity limit. Also, not all previous studies of cluster LFs have found steep slopes, and many cluster LF studies have relied on using statistical counts to remove background galaxies (e.g., Dressler 1978 and Goto et al. 2002 ).


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effects of binary disruption and binary segregation resulting in a significant delay of complete mixing compared to the inert binary case. In order to further describe the overall effect of the different dynamical processes affecting the survival and evolution of binaries on the FG and SG binary mixing, we present in Fig. 3 the time evolution of the cumulative radial distribution of FG and SG binaries for four different models with different values of the initial hardness parameter of the binary population. Figure 3. Cumulative radial distribution of FG (red lines) and SG (blue lines) binaries. From left to right, each panel shows the evolution of radial distribution. Three different models with x g,0 = 3, 3-20, 50, and 800 are presented (top to bottom). The cumulative radial distribution of inert binaries are shown as dashed lines for comparison. The last panel for the model with x g,0 = 800, the distribution at t ∼ 40τ rh, 0 is shown to illustrate a complete spatial mixing of very compact binaries. The cumulative radial distribution of FG and SG inert binaries in the MPr5f1m2 are also shown in order to illustrate the spatial mixing of stars with the same mass of binaries but not affected by other processes such as ionization and ejection due to the recoil velocity acquired during interactions with other stars. At t = 0, FG and SG binaries in all models have the same cumulative distribution. For the MPr5f1x3 model, the rapid disruption of SG binaries significantly depletes the inner binary population as illustrated by the shift of the SG binary cumulative radial distribution towards the outer regions. For the MPr5f1x3-20 model, the evolution of the distribution of FG binaries is similar to that for the MPr5f1x3 model. However, the presence of harder binaries implies that the effects of ionization are not as strong as in the previous case and the evolution of the SG cumulative radial distribution is more similar to that of inert binaries. It is interesting to note the development of an approximately flat portion in the cumulative radial distribution of SG binaries in the MPr5f1x3 and x3-20 simulations. The formation of this shell depleted in SG binaries is due to the combined effect of binary segregation and disruption; binaries are rapidly disrupted in the inner regions but as the segregation time-scale increases with the distance from the cluster’s centre. Therefore, there is no replenishment of binaries beyond the gap because the segregation time-scale is longer than the cluster’s age (see also the fig. 7 in Paper I ). The development of this gap is also responsible for the rapid variation in R h,SG and R h,FG / R h,SG around t ∼ 23τ rh, 0 in Figs 1 and 2 for the MPr5f1x3-20. For the MPr5f1x800 model, the cumulative distributions for FG and SG binaries are, in general, very similar to those in the inert binary model. Eventually, the spatial mixing of FG and SG binaries is achieved at t ∼ 40τ rh, 0 . Fig. 4 shows the time evolution of the


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

we sample 10,000 parameter combinations and generate one random light curve with the time sampling of the observed Keck data (∼30 nights of historical data and four nights of 2019 data) for each parameter set. We present the complementary cumulative distribution function (CCDF, or 1—CDF) of these 10,000 light curves in Figure 5 in the form of the median CDF and 1, 2, and 3 σ credible intervals. Figure 5. Top panel: comparison of the complementary cumulative distribution function (CCDF) of the observed data (historical data and 2019 data; black) and the median CCDF (dashed blue line) and the 1, 2, and 3 σ contours calculated from 10,000 simulations. The simulations were drawn from the posterior in Witzel et al. ( 2018 ). The dashed section of the observed CCDF represents flux densities that occurred only during the brightest flux excursion on 2019 May 13. These simulations show that if we repeated the entire experiment with the time sampling of 30 historical nights of Keck observations 10,000 times, then the probability of observing a single night with flux levels as high as 6 mJy is less than 0.3%. Bottom panel: the same as for the top panel but contours determined from simulations based only on the time sampling of the four nights in 2019. Because three of the four nights have elevated Sgr A* flux levels, if an experiment with four nights of observations were repeated 10,000 times, the probability of observing Sgr A* flux levels similar to the nights in 2019 would be less than 0.05%. Based on simulations drawn from the Sgr A* flux model, we can compute the probability of observing the light curves presented here. The simulations show that if we repeated our experiment 10,000 times with the time sampling of the 30 nights of historical Keck observations (from 2005 to 2013) of Sgr A*, and including the 2019 nights, there is a 0.3% probability that we observe a single night with flux levels as high as seen on 2019 May 13. These simulations are the most consistent with our current observations because they have very similar noise properties, observed duration, and timing. If we also include VLT nights for a total of over 130 nights, then the probability is less than 1.5% to observe flux values higher than 6.18 mJy on a single night. We can also consider the probability of randomly drawing 4 light curves similar to the observations in 2019. This probability is considerably lower at less than 0.05%. We also note that we have only observed the decay of the light curve on 2019 May 13, which suggests the actual maximum was likely even higher (Figure 3 ). Here we examine two possibilities for explaining the very unusual brightness and variation of Sgr A* observed this year: (1) the statistical models need to be changed or updated, and (2) there is a physical change in the accretion activity of Sgr A*. Based on the statistical model of Witzel et al. ( 2018 ), with the four nights of observations, there is a probability of less than 0.05% to


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

1 . While ID11 and notably also A4.1 fall in the transition region between HII galaxies and giant HII regions, the rest of the lensed systems are among the most luminous star-forming systems and occupy a region of the diagram corresponding to the most luminous HII galaxies. The Lynx arc result from Fosbury et al. (2003) is shown with two different values of velocity dispersion; the lower value corresponds to the semi-forbidden lines and the higher value corresponds to the fit to the permitted CIV 1550 Å line. There is more scatter in the lensed systems, but on the whole the results tend to support the lensing model estimates of the magnification apart from the Lynx arc result when using semi-forbidden lines. There is a crucial general point that we need to make regarding this exercise: The luminosity in these young bursts of star formation evolves quickly in timescales of few megayears, which moves their position away from the correlation to lower luminosities and smaller equivalent widths (EW) of H or H ( Melnick et al. 2000 ; Bordalo Telles 2011 ). This fact may be related to the scatter of the Richard et al. (2011) points below the relation (Fig. 1 ). It is therefore highly recommended to obtain and publish the EW of H or H of star-forming systems to allow for an evolution correction to be performed. 4. Conclusions Regarding the intrinsic properties of ID11, Fig. 1 implies that it is either a luminous giant HII region, such as 30 Doradus in the Large Magellanic Cloud or those found in spiral galaxies like M101, or a low luminosity HII galaxy. Given that the deep HST images of ID11 show no indication of an underlying galaxy, we confirm the Vanzella et al. (2016) conclusion that ID11 is a low luminosity HII galaxy at . The relation spans more than three orders of magnitude in luminosity, which allows us to observe a significant number of objects at large redshifts without recourse to gravitational telescopes; hence this accounts for the power of the relation as a cosmological probe ( Plionis et al. 2011 ; Terlevich et al. 2015 ; Chávez et al. 2016 ). But, of course, only the most luminous HII galaxies can be observed at higher redshifts with the current generation of 810 m telescopes, so gravitationally amplified objects such as ID11 provide us with the unique chance to verify that at high redshift the correlation holds also for low luminosity and possibly low metallicity HII galaxies. We used an independent method, which is based on the standard candle provided by the relation that is valid for HII galaxies and giant HII regions, to measure the amplification affecting the star-forming system ID11 by the Abell cluster S0163 as ; this coincides with the value of 17 obtained by Vanzella et al. (2016) from the strong lensing model of Caminha et al. (2015) . Our result suggests that we can use giant HII regions in high- lensed galaxies as instruments for studying gravitational lenses and that the combination of the standard candle plus detailed


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We acknowledge the support of the staff of the Li-Jiang 2.4 m telescope (LJT), Xing-Long 2.16 m telescope, and Tsinghua-NAOC 0.8 m telescope (TNT). Funding for the LJT has been provided by the Chinese Academy of Sciences (CAS) and the People’s Government of Yunnan Province. The TNT is owned by Tsinghua University and jointly operated by the National Astronomical Observatory (NAOC) of the CAS. Financial support for this work has been provided by the National Science Foundation of China (NSFC, grants 11403096, 11178003, 11325313, 11133006, 11361140347, 11203034, 11473063, 11322327, 11103072, 11203078, 11303085, 11203070 and 11103078); the Major State Basic Research Development Program (2013CB834903); the Strategic Priority Research Program “The Emergence of Cosmological Structures” of the CAS (grant No. XDB09000000); the Key Research Program of the CAS (Grant NO. KJZD-EW-M06); the Western Light Youth Project; the Youth Innovation Promotion Association of the CAS; the Open Project Program of the Key Laboratory of Optical Astronomy, NAOC, CAS; and the key Laboratory for Research in Galaxies and Cosmology of the CAS.


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Acknowledgments DK and DP thank the Institut Lagrange de Paris, a LABEX funded by the ANR (under reference ANR-10-LABX-63) within the Investissements d'Avenir programme under reference ANR-11-IDEX-0004-02, where this work was started. AL acknowledges the NSF grant AST 1212096 and Center for Magnetic Self Organization (CMSO). He also acknowledges Institut D'Astrophysique de Paris and Institut Lagrange de Paris for hospitality during his visit.


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

also shows a U-shaped cavity wall encompassing the E1 and E2 lobes, extending to ∼17 from the source. The emission is from the swept-up material in the dense core, which has a radius of 1500 au (∼17). Unlike the observation, however, no emission is seen extending over ∼17 from the source to ∼4′′ to the northeast. In our model, that region is within the AGB halo. Thus, in order to produce detectable emission there, the AGB halo could be denser than currently assumed in our model. Figure 6 shows the comparison of HV CO J = 3–2 emissions between the observation and the model. Both panels are also plotted over the HST images of CRL 618. In the observation (Figure 6 (a)), HV (−183.2 to −114.1 km s −1 ) emission is observed at the tips of lobes E1, E3, E4, and E5, but not observed at the tip of lobe E2. In the model, HV emission, with a similar velocity range to that in the observation, is predicted to be seen at the tips of all the lobes. The velocity of the emission increases with distance, similar to that seen in the observation. Notice that the emission intensity in our model is about half of that observed, and thus more bullet mass may be needed to produce more HV CO emission. The reason for no HV emission being detected at the tip of and along the southern edge of lobe E2 is unknown. Figure 6. (a) Observed HV CO J = 3–2 emission (contours) plotted on the H α image of CRL 618 (from Figure 2 of Lee et al. 2013a ). The velocity ranges of the contours are −183.2 to −160.6 km s −1 (blue), −157.8 to −132.5 km s −1 (cyan), and −132.5 to −114.1 km s −1 (green). The contours start at 1 Jy beam −1 km s −1 , with a step of 2 Jy beam −1 km s −1 . The synthesized beam is ∼053 × 035 at a position angle of ∼83°. (b) Simulated HV CO J = 3–2 emission (contours) in Model 2, with similar velocity ranges to those in (a). HV CO emissions are seen at the tips of the E1 to E5 lobes. The contours start at 0.5 Jy beam −1 km s −1 , with a step of 1 Jy beam −1 km s −1 . The synthesized beam is the same as that in (a). Figures 7 (a) and (b) show the PV diagrams of the observed outflows in CO, cut along the axis of lobes E1 and E2. Figures 7 (c) and (d) show the same PV diagrams but calculated from the model. In these diagrams, lobes E4 and E5 can also been seen. In Figures 7 (c) and (d), the bullets are seen at the tips of the outflows and thus are traced by the fast molecular outflows. The diagrams also show that the velocities of the molecular outflows decrease linearly with the distance to the central source, as seen in the observations. The slope of the PV structure is roughly equal to , where p is the projected distance from the center and t is the time after ejection. Therefore, different slopes are seen for the outflows in different ejection episodes. Notice that since no molecular emission is detected for lobe E2, no PV structure is detected for lobe E2 in the observation, as shown in Figure 7 (b). Figure 7. (a and b) PV diagrams of CO J = 3–2 emission of CRL


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This work was supported by NSF grants AST-1211995, AST-1616598, and AST-1616710. A.W. gratefully acknowledges the assistance of Vincenzo Pota in prototyping the modeling code. A.J.R. is a Research Corporation for Science Advancement Cottrell Scholar. A.V. is supported by the NSF Graduate Research Fellowship Program. D.F. thanks the ARC for financial support. J.S. acknowledges support from a Packard Fellowship. S.B. acknowledges the support of the AAO PhD Topup Scholarship. We thank the anonymous referee for useful suggestions. Some of 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 research has made use of NASA’s Astrophysics Data System and 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. Funding for the DEEP2/DEIMOS pipelines (Cooper et al. 2012 ; Newman et al. 2013c ) has been provided by NSF grant AST-0071048. The DEIMOS spectrograph (Faber et al. 2003 ) was funded by grants from CARA (Keck Observatory) and UCO/Lick Observatory, a NSF Facilities and Infrastructure grant (ARI92-14621), the Center for Particle Astrophysics, and by gifts from Sun Microsystems and the Quantum Corporation. This research made use of Astropy, a community-developed core Python package for Astronomy (Astropy Collaboration et al. 2013 ). We acknowledge the use of other open source Python packages, including Numpy (Walt et al. 2011 ), Scipy (Jones et al. 2001 ), Matplotlib (Hunter 2007 ), Pandas (McKinney 2010 ), IPython (Pérez Granger 2007 ), and Scikit-learn (Pedregosa et al. 2011 ). The authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community. We are most fortunate to have the opportunity to conduct observations from this mountain. Facility: Keck:II (DEIMOS).


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

other hand, in more extended BCDs the observed CO emission corresponds only to the most luminous regions of the starbursts, so it may be considered as a lower limit to the total CO emission. For all the observed galaxies our CO observations cover at least 50% of the projected size of starburst (, see Sect. 3), while for three objects this fraction, , is 100%. Table 3 Multiwavelength properties of galaxies from the literature. Notes. Columns: (2) radius at the mag arcsec isophote; (3) radius of the star-forming region from 2D surface photometry ( Amorín et al. 2009 ); (4) inclination; (5) H i gas mass; (6) logarithm of H i gas surface density; (7) stellar mass; (8) logarithm of star-formation rate density; (9) gas-phase metallicity; (10) logarithm of H surface density assuming a Galactic CO-to-H conversion factor. Reference. For metallicity: References for : (a) Shi et al. (2006) ; (b) Cairós et al. (2007) ; (c) García-Lorenzo et al. (2008) ; (d) López-Sánchez Esteban (2010) ; (e) Zhao et al. (2010) ; (f) Pérez-Montero Díaz (2003) ; (g) Gil de Paz et al. (2003) . (h) Frayer et al. (1998) ; (i) Leroy et al. (2005) ; (j) Barone et al. (2000) ; (k) Albrecht et al. (2004) ; (l) Leon et al. (1998) ; (m) Verdes-Montenegro et al. (1998) ; (n) Taylor et al. (1998) ; (o) Tacconi Young (1987) ; (p) Wei et al. (2010) . A wobbler switching mode was used at a frequency of 0.5 Hz with a beam throw up to 240 in azimuth. Two independent SIS receivers at each frequency (A100, B100 at 3 mm and A230, B230 at 1.3mm) were used to observe both polarizations of the CO 10 and 21 transitions simultaneously. The receivers were tuned in single side band, using only the lower band (LSB), which improves the calibration procedure and avoids contamination by other spectral lines in the image band, especially in observations of the calibration sources. The SIS receivers were connected to two filter-banks with resolutions of 1 MHz and 4 MHz at 115 GHz and 230 GHz, respectively, which yields velocity resolutions, , of 2.6 km s and 5.2 km s, respectively. These backends have 512 channels each, covering a total instantaneous velocity range of 1300 and 2600 km s. Pointing and focus were checked before each integration on continuum sources and planets. The pointing accuracy was better than 2.5 on average. Calibrations were done using the chopper wheel method, by observing strong radio sources (Orion A, CW Leo, and W51d). System temperatures at both frequencies varied between 290–470 K, the higher values being due to bad weather conditions. Thus, we obtained a range for the antenna temperature sensitivity of –15 mK. The CO spectra was reduced using CLASS 2 ( Buisson et al. 1997 ). To increase the signal-to-noise ratio (S/N), averaged spectra for each pointing were obtained by weighting multiple scans by a factor of , where is the integration time and is the system temperature. In all cases, baselines of zeroth order were subtracted. The spectra was then scaled to a main-beam


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

; Pierrehumbert et al. 2011 ; Rose 2015 ). The substellar point in our simulations was located over the pacific ocean. We also tried aquaplanet simulations and achieved similar results. It would be interesting to test whether locating the substellar point over a large land mass such as Asia would significantly affects our results. 5. Conclusions The main conclusions of this paper are as follows. 1. Tidally locked planets will not necessarily exhibit a snowball bifurcation. This is in contrast with planets with Earth-like rotation and orbit, which are required to exhibit a snowball bifurcation except in unrealistic edge cases. 2. The reason tidally locked planets are not required to exhibit a snowball bifurcation is that the insolation increases strongly as the substellar point is approached. 3. Whether a tidally locked planet actually does exhibit a snowball bifurcation depends on the heat transport and ice/ocean albedo contrast. Higher heat transport and a higher ice/ocean albedo contrast both favor a snowball bifurcation. 4. We performed GCM simulations that suggest that realistic tidally locked planets will not exhibit a snowball bifurcation. Although they should tend to have stronger heat transport, the top-of-atmosphere ice/ocean albedo contrast should tend to be smaller. 5. This work suggests that we will not find habitable tidally locked exoplanets with an active carbon cycle in a snowball state. This should be verified with a model incorporating an active carbon cycle. Appendix Artifact of Sea Ice Scheme Interestingly, and in contrast with Lucarini et al. ( 2013 ), we find a small separation between equilibrated states that were initialized from Warm and Cold Starts. The most likely reason Lucarini et al. ( 2013 ) did not observe this separation is that they started their Cold Start simulations from a significantly warmer climate state that was not globally glaciated. To further investigate this issue, we show global mean sea ice concentration for a series of simulations with a stellar flux of 1200 W m −2 and initial conditions representing equilibrated climate states with different amounts of sea ice (Figure 9 ). We find that states started within the apparent separation band are stable, rather than approaching separate stable warm and cold states at the upper and lower ends of the separation band. This means that the separation band actually represents a smeared state, which we can consider a continuum of states rather than two distinct “fixed points” with separate basins of attraction. This is clearly unrelated to a snowball bifurcation and does not alter our interpretation that there is no snowball bifurcation in PlaSim in the tidally locked configuration. Figure 9. Time series of global mean sea ice concentration for a series of simulations at a stellar flux of 1200 W m −2 with different initial conditions. This plot demonstrates that the solution is a continuum of states rather than two distinct fixed points. By further


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We thank the referee for the very helpful suggestions that have significantly improved our manuscript. We acknowledge financial support from the 973 Program of China under grant 2013CB837000, the National Natural Science Foundation of China (No. 11573060), the Key Laboratory of Astroparticle Physics of Yunnan Province (No. 2016DG006), the Strategic Priority Research Program, and the Emergence of Cosmological Structure of the Chinese Academy of Sciences, Grant No. XDB09000000. The work of D.H.Y. is supported by CAS “Light of West China” Program.


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

for studies of the host galaxies of AGN ( Kauffmann et al. 2003 ). When using the Vanden Berk et al. (2001) template, it is important to know that it contains significant host galaxy contamination, in particular at redder wavelengths (e.g., Fynbo et al. 2013 , their Fig. 5). In this paper, we use data from the X-Shooter spectrograph to build a new QSO template that covers the full range from rest-frame Ly to 11 350 Å based on bright SDSS QSOs and observed with a single instrument over this spectral range. In Sect. 3 we describe the sample selection and provide details of the observations. Section 4 describes the construction of the composite spectrum, and Sect. 5 presents the resulting composite. In Sect. 6 we perform a comparison with existing composites and in Sect. 7 offer our conclusion. We use the cosmological parameters from Planck Collaboration XVI (2014) with km s Mpc, , and throughout the entire paper. 2. Issues with the currently most used template The currently most widely used QSO template is the SDSS template of Vanden Berk et al. (2001) . As described very clearly in that paper, the SDSS composite spectrum shows a significant change to a shallower spectral slope around 5000 Å. This is mainly attributed to contaminating light from the underlying host galaxies. Other effects also contribute, such as the emission from a hot dust component, but the dominating factor is probably the host contamination. Our main motivation for building a QSO template without this problem is the following: a search for dust-reddened QSOs using near-IR (NIR) selection of QSOs has been initiated ( Fynbo et al. 2013 ; Krogager et al. 2015 ). In the central southern SDSS footprint, Stripe 82 ( Annis et al. 2014 ), about 50 such candidate QSOs have been studied using the New Technology Telescope (NTT). Although the search was for QSOs reddened by foreground absorber galaxies, most systems turned out to be QSOs reddened by dust in their host galaxies. The optical spectra can be matched well by the SDSS template spectrum reddened by SMC-like extinction curve, but the NIR (rest frame optical) photometry from the UKIRT Infrared Deep Sky Survey (UKIDSS) cannot be simultaneously fitted. The problem is illustrated in Fig. 6 in Fynbo et al. (2013) , where they attempt to model the EFOSC2 spectrum of a dust-reddened QSO at observed in that survey. As seen, the SDSS spectrum and photometry can be matched well with the Vanden Berk template, but the photometry from UKIDSS is much too blue for even the unreddened template. This is not a unique case, but a problem that is seen for all the dust-reddened QSOs found in their search. In the SDSS composite, the QSOs which contribute to the spectrum at 5000 Å have to be at fairly low redshifts (). These QSOs have absolute magnitudes that are three to four magnitudes fainter than bright QSOs (e.g., Vanden Berk et al. 2001 , their Fig. 1). They are also likely to have the brightest rest frame optical host galaxies, so host


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-Instrument O B-Citation I-Citation I-Citation O O O O O O O B-Telescope I-Telescope I-Telescope O O O B-Observatory I-Observatory I-Observatory B-Organization I-Organization I-Organization I-Organization B-CelestialObject I-CelestialObject 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 B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O O B-Instrument I-Instrument O O O O O B-Formula I-Formula I-Formula I-Formula I-Formula I-Formula I-Formula O O O O O 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 B-Instrument I-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 O O O O O O O O B-Instrument 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-Wavelength B-Wavelength I-Wavelength O O O B-Archive I-Archive 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-CelestialObject I-CelestialObject O O O O B-CelestialObject I-CelestialObject O B-Survey B-Wavelength I-Wavelength 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 B-Telescope 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 B-Telescope B-Wavelength B-Wavelength 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 B-CelestialObject I-CelestialObject O O O B-CelestialObject I-CelestialObject O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject O O O O O O O O O

by Proshina et al. (2020) . 2. Observations and the data involved Our long-slit spectral observations were made with a multi-mode focal reducer SCORPIO-2 ( Afanasiev Moiseev 2011 ) at the prime focus of the Russian 6 m telescope of the Special Astrophysical Observatory, Russian Academy of Sciences. UGC 5936 was observed on February 28, 2014. In these observations, the 1′′ slit was aligned with the major axis, PA(slit) = 83deg, and the total exposure time was 75 min; also the target was exposed in the orientation through the neighbouring dwarf galaxy, PA(slit) = 135deg, with the exposure time of 45 min. The seeing during these observations was mediocre, F W H M = 3′′. We used the VPHG1200@540 grism providing an intermediate spectral resolution F W H M ≈ 5 Å in the wavelength region from 3700 Å to 7200 Å. This spectral range includes a set of strong absorption and emission lines making it suitable to analyse both stellar and gaseous kinematics of the galaxy as well as the gas excitation and chemistry and properties of the stellar populations. The slit is 6′ in length, which allowed us to use the edge spectra for subtraction of the sky background. The CCD E2V CCD42-90, with a format of 2048 × 4600, using in the 1 × 2 binning mode provided a spatial scale of 0.357 ′′/px and a spectral sampling of 0.86 Å/px. The data reduction as well as kinematics and stellar population characteristics derivation were standard for our SCORPIO-2 data; see for example Sil’chenko et al. (2019) or Proshina et al. (2020) . We used the g r -band images from the SDSS/DR9 archive ( Ahn et al. 2012 ) to study the large-scale structure of the galaxy. 3. Photometric and dynamical structure of UGC 5936 The smooth lenticular galaxy UGC 5936 (see 2MASS J -band image, Fig. 1 , lower middle) possesses a large pseudo-ring which looks unclosed in the optical continuum bands (Fig. 1 , upper left). However Fig. 1 (lower right), which shows the WISE/W3 data in the 12 μ m band delineates a perfect closed ring; in this figure hot dust heated by star formation reveals its organisation in a typical S0 outer ring structure. The colour map (Fig. 1 , upper right), demonstrates the blue colour of the ring, which is narrow to the east and splits into two arcs to the west, as well as the red colour of the arc connecting the core and the western part of the ring. The western blue-ring structures correspond to the ring in the unsharp-masked image (Fig. 1 , upper middle). Interestingly, the GALEX near-ultraviolet (NUV) map (Fig. 1 , lower left), which also refers to young star concentration, shows UV flux excesses beyond the borders of the continuum-band pseudo-ring; the offset is seen especially well at the major axis. To the south-east from UGC 5936, a small satellite, SDSS J105009.10+362009.9 , is observed at 15 kpc projected onto the sky plane; this satellite is observed at 19 kpc if we suggest that it lies in the plane of the UGC 5936 disc. This small satellite is very blue. According to


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ACKNOWLEDGEMENTS We are grateful to Peter Camps for code support, interesting discussion, and the development of skirt . We thank Ricardo Amorin, Ilse De Looze, Hiroyuki Hirashita, Roberto Maiolino for helpful discussion, and Nicolas Laporte for providing data. AF acknowledges support from the ERC Advanced Grant INTERSTELLAR H2020/740120. We acknowledge the use of the python programming language (Van Rossum de Boer 1991 ), and use of the numpy (van der Walt, Colbert Varoquaux 2011 ), ipython (Perez Granger 2007 ), jupyter (Kluyver et al. 2016 ), matplotlib (Hunter 2007 ), and pymses (Labadens et al. 2012 ) packages. This research made use of astropy , a community-developed core python package for Astronomy (Astropy Collaboration et al. 2013 ).


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map maximum. The variability period is illustrated by the red dashed line in the global spectra. A possible aliasing period is displayed in blue. Lower panel: Lomb–Scargle periodogram for the same LC. The main peak labelled ‘A’ corresponds to the period of 18.86 d, whereas ‘B’ indicates the 39 d period. 3.3 KIC 4829241, KIC 10420021 and KIC 11822535 The stars KIC 4829241, KIC 10420021 and KIC 11822535 were classified as non-periodic by Maoz et al. ( 2015 ). Instead, our study suggests a periodic or quasi-periodic modulation well describing a variability signature in the wavelet representations. We attribute to this variability periodicities ranging between 11 and 16 d, which is reasonable for WDs (e.g. Kawaler 2015 ; Valeev et al. 2017 ; Braker et al. 2017 ). In fact, these periodicities are more persistent in KIC 4829241 and KIC 10420021 than in KIC 11822535, even though also noticeable in the latter star. We do not expect to find magnetic cool-spot signatures for KIC 4829241 and KIC 11822535 considering their high T eff and their log g values (Tremblay et al. 2015 ). However, a transient signature for KIC 10420021 is conceivable in view of its T eff and log g being close to the limit value (Tremblay et al. 2015 ), so the presence of a convection zone at the WD surface can still be envisaged. Combining the surface gravities and the effective temperatures, derived spectroscopically by Hallakoun et al. ( 2018 ), with evolutionary grids (Benvenuto Althaus 1999 ), we are able to estimate the mass of these three WDs. KIC 4829241 is a WD of mass . The Lomb–Scargle periodogram for this WD shows a period of 16.68 d with an amplitude of 286 ppm, as indicated by the letter A in the upper-right panel of Fig. 5 . Following Maoz et al. ( 2015 ), given its mass, rotation combined with an accretion of gas from the interstellar medium (ISM) and the material carried on the magnetic poles generating hotspots, with an accretion luminosity of Δ L acc ∼ 1.4 × 10 29 erg s −1 , could explain the periodic modulation. Fig. 5 (upper-left panel) displays the wavelet analysis, from which we noticed a long-lasting periodicity of 16.04 d over the time observation. With a lower amplitude, we also found a 40-d period that could be considered as a possible aliasing period. The 16-d period characterizes the WD photometric modulation since it is the main period obtained in both Fourier- and wavelet-based techniques. Based on a rough FAP estimation (see Section 2 ), these periods likely lie above the confidence level. Figure 5. Left: light curves with additional treatment, wavelet local and global spectra of Kepler WDs (from top to bottom) KIC 4829241, KIC 10420021, and KIC 11822535. Contour levels are 90 , 80 , 70 , ..., 20 , and 10 of the map maximum. The variability periods are illustrated by the red dashed lines in the global spectra. Possible aliasing periods are displayed in blue. Right: respective Lomb–Scargle periodogram for the same WD. The main peak labelled 'A'