O O O O O O O B-Grant O B-Grant B-Person B-Person O O B-Grant I-Grant I-Grant O O O O B-Location O B-Location B-Person O B-Person B-Person O O O O O B-Grant I-Grant O B-Organization I-Organization I-Organization O B-Location O B-Grant I-Grant I-Grant I-Grant I-Grant O B-Location O O O O B-Grant O O O O O O O O O O O O B-Organization I-Organization I-Organization O O B-Location I-Location B-Grant I-Grant I-Grant B-Grant O B-Organization O O B-Grant B-Person B-Person O O O O O B-Fellowship I-Fellowship I-Fellowship O O B-Organization I-Organization O O O O O O O O O O O O O O

This research was supported by the grant AYA2012-39303 and SGR2014-1073 (DV, DFT) and the Project Formosa TW2010005, for bilateral research between Taiwan and Spain (KH and DFT). KH is partly supported by the Formosa Programme between National Science Council in Taiwan and Consejo Superior de Investigaciones Cientificas in Spain administered through grant number NSC100-2923-M-007-001-MY3. The research leading to these results has also received funding from the European Research Council under the European Union's Seventh Framework Programme (FP/2007-2013) under ERC grant agreement 306614 (MEP). MEP also acknowledges support from the Young Investigator Programme of the Villum Foundation. We are grateful to the referee for a careful reading and the useful comments.


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ACKNOWLEDGEMENTS Based on observations with the NASA/ESA Hubble Space Telescope , obtained at the Space Telescope Science Institute, which is operated by AURA, Inc., under NASA contract NAS 5-26555. JA and AB acknowledge support from STScI grants GO-13297. This work has been supported by the European Research Council through the ERC-StG 2016, project 716082 ‘GALFOR’ and by the MIUR through the the FARE project R164RM93XW ‘SEMPLICE’. AFM acknowledges support by the Australian Research Council through Discovery Early Career Researcher Award DE160100851.


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

the acceleration of the star induced by the companion. We augment the internal RV uncertainties listed in Table 2 with a “jitter” term that is added in quadrature to accommodate the expected level of noise introduced from stellar variability. HD 169889 A has a visual color of B − V = 0.75 and median activity index of S HK = 0.18. We assume a B − V uncertainty of 0.01 mags. Using empirical relations from Isaacson Fischer ( 2010 ), we estimate a jitter value of σ = 2.2 ± 0.1 m s −1 . Table 2 Relative Stellar RV Measurements for HD 169889 Note. A horizontal line denotes the division between data sets when the HIRES detector was upgraded (2004 August) requiring a relative RV offset for analysis. Table 3 Summary of High-contrast Imaging Observations Showing the Number of Coadds ( n coadds ), Frames ( N fr ), Exposure Time per Frame (Δ t exp ), and Change in Parallactic Angle Δ π (°) Table 4 Summary of Astrometry Measurements Note. Position angle (PA) measurements are referenced with respect to the equator of the observing epoch. Table 5 Photometric Results and Companion Properties A straight line is fit using Markov Chain Monte Carlo (MCMC) methods to find a best-fitting RV slope of m s −1 yr −1 . Combined with the direct imaging projected separation, we use the relation from Liu et al. ( 2002 ) to derive a lower limit to the companion mass of m min = 0.369 ± 0.010 M ⊙ (Table 5 ). Clearly, HD 169889 B cannot be a brown dwarf but is instead a compact object provided the acceleration is attributable to this object. A more careful treatment would consider the slight curvature noticeable by eye in the RV time series as well as astrometry (Figure 1 ). However, only a small fraction of the companion orbit has been traced in either data set. Further, the astrometry measurements were obtained with two different instruments using three different filters. The systematics that result from combining these observations can be noticed in the zoomed-in portion of Figure 5 at the ≈1 σ level (10 mas). The WD companion appears to be moving to the north, i.e., closer to HD 169889 A, but further AO imaging is required to confirm. Figure 5. Proper motion and parallactic motion analysis taking into account the path that an infinitely distant background object would traverse is compared to the separation of HD 169889 B over time. Angular offsets are measured relative to the optical primary star. The inset shows a zoomed-in view of individual astrometric measurements. The companion shares the same space motion demonstrating that they are gravitationally associated. 5. WD Properties Empirical relations between WD masses and their progenitors have been established based on the study of open clusters, globular clusters, and common proper motion pairs (Weidemann 2000 ; Williams et al. 2009 ). Using the dynamical mass lower limit from RVs and imaging, we attempt to place a lower limit on the companion progenitor mass ( M i ) by inverting the empirical relationship from


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

time like Fig. 2 , we can conclude that the bending in period B must exist from the observational data about the periatron time T N in both periods A and C. It is possible to judge whether KL-oscillation had happened or not by using our analysis without the observational data in highly eccentric state. One may be worried about the spin evolution of the pulsar because the spin–orbit coupling appears in GR (Barker, O’Connell 1975 ), which may change the direction of the pulsar rotation axis. If the beaming direction of pulse signal is changed, the pulsar will disappear in stage B. However, following Liu Lai ( 2017 ), we find that the adiabaticity, which measures the ratio of the KL oscillation time-scale to the presession time-scale, is very small in the present system. It means that the spin of the inner binary evolves ‘non-adiabatically’ even when the KL-oscillation occurs. As a result, the spin will be parallelly transported just as Newtonian case, and then the beaming direction will not change so much even in the KL-oscillation stage and the whole period (A, B, and C) will be observed. Finally we shall discuss the possible parameter range that the bending of the cumulative-shift curve occurs within our lifetime. In addition to the model given in Table 1 , we have performed our calculation for 19 200 models by changing the outer semimajor axis () and the mass of the third body (10M ⊙ ≦ m 3 ≦90M ⊙ ) and analysed in which parameter ranges the above bending will occur within 100 yr. In order that such phenomenon occurs within 100 yr, the KL time-scale should be less than 100 yr. So we investigate how many years are required for the deviation of Δ P from Δ P 1 to become large enough. For 19 200 models with different m 3 and a out , we calculate Δ P and judge the deviation simply by the relative difference between Δ P and Δ P 1 , i.e. (Δ P − Δ P 1 )/Δ P 1 > 1. We define T d as the time when this criterion is satisfied. Fig. 3 shows the colour contour map of T d . The black region in the bottom-right corner of Fig. 3 corresponds to T d > 100 yr. The white dotted line shows the critical curve used by the Newtonian formula T KL = 100 yr given by equation ( 3 ). It is also found that the dependence of T d on a out and m 3 is consistent with the Newtonian formula T KL . This is because the semimajor axis in the present model is large and 1PN effect is quite small. The top-left white region in a out 15 au and m 3 > 30M ⊙ shows that the system becomes unstable if the initial parameters are in the region. The yellow solid line is the empirical criterion for the instability given by Blaes et al. ( 2002 ): From Fig. 3 , we find that if this binary has the third companion with m 3 = 10–100M ⊙ and , we may be able to see the bending of the cumulative-shift curve within 100 yr through the observation of radio pulse. Figure 3. The parameter range of third body with which the bending of cumulative-shift curve occurs within 100 yr. The horizontal and


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ACKNOWLEDGEMENTS The authors are grateful to the referee, Dylan Nelson, for detailed and constructive comments that enhanced this work. The authors would like to thank Joe Burchett, Ryan Horton, Ali Rahmati, John Stocke, Todd Thompson, Colin Norman, Ryan O’Leary, and Jason Prochaska for interesting conversations that contributed to this manuscript. We acknowledge discussions with Sean Johnson and Zachary Hafen about the Turnbull statistical method. Support for Oppenheimer was provided through the NASA ATP grant NNX16AB31G. This work was supported by the Netherlands Organisation for Scientific Research (NWO) through VICI grant 639.043.409. This work used the DiRAC Data Centric system at Durham University, operated by the Institute for Computational Cosmology on behalf of the STFC DiRAC HPC Facility ( www.dirac.ac.uk ). This equipment was funded by BIS National E-infrastructure capital grant ST/K00042X/1, STFC capital grants ST/H008519/1 and ST/K00087X/1, STFC DiRAC Operations grant ST/K003267/1, and Durham University. DiRAC is part of the National E-Infrastructure. RAC is a Royal Society University Research Fellow. AJR is supported by the Lindheimer Fellowship at Northwestern University.


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

of destroyed GCs to field populations in the bulges of MW-like galaxies, comparing the predictions with the measurements by Schiavon et al. ( 2017 ). They show that, for most of the MW-like galaxies in their simulated volume, the prediction for f Nr of the metal-poor stellar population is lower than the observations by factors of ∼2–30 (bottom panel of their fig. 4). However, for a few simulated galaxies, the predicted f Nr are in good agreement with the observations. Like the MW, the disc populations of those galaxies are characterized by a bimodal distribution in the α–Fe plane, which is a distinctive feature of the MW disc populations (e.g. Hayden et al. 2015 ; Mackereth et al. 2017 ). Mackereth et al. ( 2018 ) showed that this feature is associated with an atypical accretion history, characterized by intense merging in early times and relative calm since z ∼ 1–1.5. It is noteworthy, however, that Hughes et al. ( 2020 ) predictions for these few MW-like galaxies differ from our measurements with regard to the dependence of f Nr on position in the α–Fe plane. The high frequency of ex-GC stars in the field of simulated galaxies is predominantly due to high-α in situ GC formation and destruction, whereas our data show that the high f Nr in the MW bulge is due to the contribution by the dissolution of low-α accreted GCs. This discrepancy would be alleviated if some of the stars in the accreted region in Fig. 2 were in fact formed in situ (see fig. 2 of Hughes et al. 2020 ), but it is not clear that accounting for such a contamination would completely eliminate the disagreement. 4 SUMMARY The results presented in this paper make use of elemental abundances from APOGEE DR16 along with data from Gaia DR2 to study the chemical and kinematic properties of 146 N-rich stars located within the inner 4 kpc of the Galaxy. Our conclusions can be summarized as follows: We find that there are likely accreted and in situ components to the N-rich population within 4 kpc of the Galactic Centre, identified via chemistry by making a cut in [α/Fe]–[Fe/H] space towards low metallicities (as shown in Fig. 2 ; e.g. Hayes et al. 2018 ; Mackereth et al. 2019 ; Das et al. 2020 ). By making this cut and removing stars without proper motions in Gaia , we select 428 and 13 982 bulge stars that lie in the accreted and in situ positions, respectively, with 25 N-rich stars being located in the accreted and 58 located in the in situ locus. We show that our sample of N-rich stars occupies the same locus as so-called SG GC stars, supporting the idea that they are the by-products of GC destruction/evaporation. We find that there is a significant difference in the kinematic properties of chemically defined accreted and in situ populations. This shows that our chemistry-based distinction of these populations maps into differences in kinematic space. We also find that the accreted bulge field population includes stars that share orbital properties with the GE/S system, although no


O O O O O 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Wavelength O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-CelestialObject I-CelestialObject 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 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-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 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 B-Citation I-Citation I-Citation I-Citation O O O O O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O 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 the PA and inclination angle. Since the residual clumps proved to be not sensitive to the choices of the assumptions, we preferred the best fit model with PA major = 330° and i = 30°. The results from the kinematical analysis are presented in Fig. 6 , where the comparison between the model and the data in terms of intensity, velocity, and velocity dispersion maps (zeroth, first, and second moments, respectively) is shown. The model and data (smoothed by a factor of (2) are in excellent agreement. The kinematics of the molecular gas, traced by the CO(2–1) emission line, are clearly dominated by purely circular rotational motions around the nucleus. This is the first successful attempt to model the kinematics of the molecular gas in NGC 7213, thanks to the software 3D BAROLO, which is able to handle the available data, despite the sparse information that limited previous attempts to model the kinematics (e.g. Ramakrishnan et al. 2019 ). As is clearly visible in Fig. 7 , we found that the data cube central velocity was offset by an additional δV sys = 36 ± 10 km s −1 . This means that the systemic velocity of the source with respect to our rest frame has to be V sys = 1716 ± 10 km s −1 , which is a mean value between the systemic velocity obtained from the study of the stellar kinematics by Schnorr-Müller et al. (2014) , and the results by Ramakrishnan et al. (2019) . Fig. 7. PV diagrams on the major ( top panel ) and minor ( bottom panel ) axis, respectively. On the x -axis, the angular distance from the centre of the source along the direction fixed by the angle φ is reported. On the right y -axis, the velocity along the line of sight (in km s −1 ) is shown as it appears in the input data cube. On the left, the velocity along the line of sight (in km s −1 ) is shown, once the systemic velocity ( v sys = 36 ± 10 km s −1 ) has been subtracted. The data are represented in grey scales with blue contours, while the best fit model is identified by red contours. Upper panel : the blue circles represent the projected best fit rotational velocity at different radial distances from the centre, associated with each ring of the disc. Bottom panel : the signature of the potential outflow associated with Region C (bordered in green; see Sect. 4.3 ) is clearly visible as an excess of CO emission over 100 km s −1 with an offset of ∼1′′ from the centre of the rotation. Looking at the lower panels, the velocity dispersion in the best-fit model is ∼10−15 km s −1 : values that are expected from a rotation-dominated disc. The highest values of the velocity dispersion ( v disp ∼ 70 km s −1 ) are not reproduced by the best fit model. The high-dispersion residual clumps proved to be insensitive to our assumptions (e.g. PA, i , v rad ), therefore we suggest that the residuals are associated with non-rotational motion, as we discuss in Sect. 4.3 . 4.3. Complex structures in the CO emission The results of the modelling of the gas kinematics presented in Sect. 4.2


O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation I-Citation O B-Citation O B-Citation I-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O B-Telescope O O O O O O O B-Survey B-Wavelength O O O O O O B-Wavelength O B-Wavelength O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Dataset I-Dataset I-Dataset B-Dataset O O B-Observatory O O O O O B-Model I-Model I-Model O O O O O O O B-Model I-Model I-Model O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Formula I-Formula I-Formula I-Formula O O O O O O O O O B-CelestialRegion I-CelestialRegion O O O O O O O O O O O B-CelestialRegion I-CelestialRegion 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 B-CelestialRegion I-CelestialRegion O O O O O O O O O O O O O O O O B-CelestialRegion I-CelestialRegion O O O O O B-Wavelength O O B-Survey O O O O O O O B-Survey O B-Survey 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 B-Wavelength O O O O O B-Telescope O O O O O O O B-Dataset O O O B-Telescope O B-Mission O O B-Citation I-Citation O B-Citation I-Citation I-Citation I-Citation O O O O O O B-Survey O B-Dataset 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 B-Survey B-Citation I-Citation I-Citation I-Citation I-Citation O B-Survey B-Citation I-Citation I-Citation I-Citation O O O O O O O O O O O O O O B-Survey O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Dataset 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 B-Dataset I-Dataset O O O O O O B-Formula I-Formula I-Formula I-Formula O

a clear separation can be made between most stars and extragalactic objects, the separation between galaxies and quasars is less abrupt, reflecting the range of central-engine-dominated and host-dominated emission in these objects. As simple colour-magnitude diagnostics show, our classification generalises simple colour-magnitude cuts that have been proposed in the literature ( Stern et al. 2005 , 2012 ; Assef et al. 2013 ; Chehade et al. 2018 ), but it also deals with the overlap of different classes, especially towards the faint end in WISE magnitudes and towards higher-redshift quasars. The three VEXAS optical+IR classified tables, with object IDs, coordinates, optical and infrared magnitudes, including the imputed ones (which are flagged), and the probability to belong to each class, are publicly released and available for the scientific community as part of the VEXAS Data Collection (DR2) via the ESO Phase 3 13 . The machine-learning classification pipeline and the code are available on the VEXAS Github repository 14 . The density of extragalactic objects varies with the survey depth, quite expectedly. Considering only high-confidence objects for which the probability of belonging to the class is p class ≥ 0.7, our classifiers yielded 111 qso deg −2 for the VEXAS-DESW footprint (≈4900 deg 2 ) and 103 qso deg −2 for the VEXAS-PSW footprint (≈3800 deg 2 ). These numbers roughly meet the requirements of the 4MOST selection for baryon acoustic oscillations measurements. The density drops to ≈10 qso deg −2 for the VEXAS-SMW footprint (≈9300 deg 2 ). All in all, the combined survey depth is the limiting factor. The VEXAS-SMW footprint is limited by the shallow optical coverage of SkyMapper, which also has less uniform coverage than DES and PanSTARRS. A solution would be to consider only NIR and mid-IR magnitudes in the classification, which in turn requires more uniform coverage in the NIR and smaller uncertainties in the WISE magnitudes. This may be obtained thanks to unWISE, a re-processing of WISE, and NeoWISE imaging ( Lang 2014 ; Schlafly et al. 2019 ), or with forced photometry (on SkyMapper and unWISE cutouts) based on VISTA detections. We caution that these samples are only the very first step for cosmological measurements, which also require spectroscopic redshifts and well characterised coverage maps. The spectroscopic follow-up is already the aim of southern surveys (4MOST, de Jong et al. 2012 ; SDSS-V, Kollmeier et al. 2017 ). Since this paper is focused on the classification techniques and their application to VEXAS, the coverage maps are outside the scope of this work and will be part of a subsequent release. Appendix A: Testing the classification at the faint end of VEXAS-DESW As described in detail in Sect. 6.1 , the r -band magnitude ranges covered by the input tables are larger than those covered by the corresponding tables. This is particularly true for the VEXAS-DESW table, which reaches magnitudes as faint as r ∼ 25 m (see


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The authors acknowledge funding support for CMOR from NASA cooperative agreement NNX11AB76A, the Natural Sciences and Engineering Research Council of Canada and the Canada Research Chairs Program. Visual IMO Eta Aquariid data provided by J. Rebdtel is gratefully acknowledged.


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We thank the anonymous referee for helpful comments that greatly improved the presentation of this paper. This work is supported by the 973 Program (No. 2015CB857002), national science foundation of China (grant Nos. 11128306, 11121062, 11233005, 11073017), NCET-11-0879, the Strategic Priority Research Program “The Emergence of Cosmological Structures” of the Chinese Academy of Sciences, Grant No. XDB09000000 and the Shanghai Committee of Science and Technology, China (grant No. 12ZR1452800). We also thank the support of a key laboratory grant from the Office of Science and Technology, Shanghai Municipal Government (No. 11DZ2260700). H.J.M. would like to acknowledge the support of NSF AST-1517528, and F.v.d.B. is supported by the US National Science Foundation through grant AST 1516962. A computing facility award on the PI cluster at Shanghai Jiao Tong University is acknowledged. This work is also supported by the High Performance Computing Resource in the Core Facility for Advanced Research Computing at Shanghai Astronomical Observatory.


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I thank A. Partl, D. Streich, C. Vocks, H. Enke, and L. Strandberg for comments on the manuscript. C.S. was supported by funds of PTDESY-05A12BA1 and 05A14BA1, and the BMBF VIP program 03V0843. This research has made use of the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration.


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T. Rauch was supported by the German Aerospace Center (DLR) under grant 05 OR 1402. This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France, and of NASA’s Astrophysics Data System Bibliographic Services. Some of the data presented in this paper were obtained from the Mikulski Archive for Space Telescopes (MAST).


O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Tag I-Tag I-Tag I-Tag I-Tag I-Tag I-Tag O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Model I-Model O O O O O O O O O O O B-Model I-Model I-Model O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 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 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 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 O O O O O O O O O O O O O B-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O O O O O

gas density was measured from either the neutral or ionized hydrogen and includes a correction factor of 1.36 in order to account for the contribution from helium. Since we are interested in a comparison with our 1D cosmic ray transport models, we neglect the error contribution from the variance of the density; thus, the error is solely determined by uncertainty of the velocity dispersion and is δ U kin = 2δσ v /σ v , equivalent to 50 per cent. We find that the kinetic energy densities are within a factor of a few in energy equipartition with the magnetic field. This is the expected result since the magnetic field in dwarf galaxies is usually assumed to be amplified by the turbulent gas motions via the ‘fluctuating dynamo’, which saturates when the energy densities become comparable (Schleicher Beck 2013 ). Interestingly in the central part, the magnetic energy density shows a peak that corresponds best to kinetic energy density of the ionized gas, whereas that of the atomic gas is practically constant. The cosmic ray energy density is important everywhere. Such a behaviour has been suggested by theoretical works of cosmic ray-driven galactic winds (e.g. Salem Bryan 2014 ), where the cosmic rays can diffuse outwards from star formation sites and thus create an overpressure in the surrounding area. A galactic wind can then be launched in the disc–halo interface, where the cosmic rays dominate the energy density and are able to transfer some of their energy and momentum to the ionized gas via the streaming instability. Figure 5. Minor axis profiles with distances from the major axis as seen on the sky. Shown are the energy densities of the magnetic field, warm neutral medium (H i ), warm ionized medium (Hα), and cosmic rays (CRs). The blue line shows the magnetic energy density for the accelerating wind model. The blue-shaded area indicates the range of magnetic field strengths for which we tested our models, where the upper limit is for energy equipartition and the lower limit is for almost freely escaping CREs (only radiation losses). Error bars show ±1σ uncertainties, which are assumed to be 50 per cent stemming solely from the uncertainty of the velocity dispersion (see text for details). While the diffusion model (Model I) fits our data best, the resulting diffusion coefficients are at least one order of magnitude lower than both estimates for the expected diffusion coefficient. On the other hand, the projected extent of CRE injection sources due to the mild inclination can account for the parabolic shape of the spectral index profile, where the spectral index remains flat out to distances of up to 0.5 kpc away from the major axis; so the abrupt steepening of the spectral index at | z | > 0.5 kpc that distinguishes diffusion from advection may be the result of geometry. Hence, we explored the possibility of dominating advection and found that constant wind speeds are only a factor of 2 below the escape velocity – the expected advection


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Acknowledgments CJ acknowledges a Durham COFUND Junior Research Fellowship (4958) supported by the DIFeREns2 project that has received funding from the European Union's Seventh Framework Programme for research, technological development and demonstration under grant agreement no 609412. This work was supported by the Science and Technology Facilities Council [ST/L00075X/1]. This work used the DiRAC Data Centric System at Durham University, operated by the ICC on behalf of the STFC DiRAC HPC Facility ( www.dirac.ac.uk ). This equipment was funded by BIS National E-infrastructure capital grant ST/K00042X/1, STFC capital grant ST/H008519/1, and STFC DiRAC Operations grant ST/K003267/1 and Durham University. DiRAC is part of the UK's National E-Infrastructure.


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

shape of the -band excess or to place meaningful constraints on the frequency of such events, as we are insensitive to many periods. Nonetheless, this opens up the possibility that some of the variations in the CoRoT and Spitzer light curves could be due to short-lived, and potentially recurring, hot spot emission or to stellar flares. We recall that the system’s average colours do not show a significant -band excess (Paper I, Sect. 4.5), implying a low-to-negligible average accretion rate consistent with the majority of the CFHT light curve. It is important to note that the raw CFHT light curves consist of short runs of 5 consecutive observations and these have been binned using a weighted average to give each data point seen in the figure. Hence, the three points indicating a -band excess are derived from 15 independent observations conducted in photometric conditions. 3.2. Spot model Young, low-mass stars tend to be heavily spotted (e.g. Donati et al. 2000 ) and can display photometric modulations not unlike the LAV-dominated parts of the CoRoT light curves. However, changes in such modulations from starspot evolution tend to occur smoothly on timescales of many days to weeks (e.g. Roettenbacher et al. 2013 ), rather than the sudden morphological changes we see at rHJD 4544 in the 2008 CoRoT light curve (when the LAVs appear) and at rHJDs 5910 and 5930 in the 2011/2012 CoRoT light curve (when the LAVs appear and disappear, respectively). Given that there are two stars in this system, both of which are likely spotted, an appealing solution could be constructive and destructive interference of starspot signals where different active regions have slightly different periods. This could arise due to either different stellar rotation periods (the system is young and not necessarily synchronised) or, even if the stars are synchronised, different latitudes (differential rotation) 4 . We therefore sought to understand how much of the large scale structure in the CoRoT and Spitzer OOE light curves could be attributed to the constructive and destructive interference of starspot signals. We set up a simple two-spot model following the formalism of Dorren (1987) . Due to the well known degeneracy of spot models we allow only two starspots and opt for a single spot on each star rather than two spots on one star 5 . Before modelling the data we need to: a) determine the flux contributions of each star in each band; b) determine the flux ratios between the spotted and unspotted stellar photospheres and c) mask variations that are clearly not due to spot modulation: a) the observed flux in each band is the sum of the stellar and dustcomponents. Our SED modelling indicated that the dustcontributes negligible flux in the CoRoT band but makes up asignificant fraction at 3.6 and 4.5 m. The stellar flux fractions are 0.62, 0.41 and 0.34 for the primary star in the CoRoT, Spitzer 3.6 and 4.5 m bands respectively, and 0.38, 0.25 and 0.21 in turn for the secondary.


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

of alignment of the outer disc with the equatorial plane of the BH. Caproni et al. (2004) and Liu Melia (2002) show that the misalignment between the rotation axis of the accretion disk and of the Kerr BH can explain jet precession in AGN as well. The authors successfully apply this scenario to quasars, Seyfert galaxies, and Sgr A*. However, Nixon King (2013) show that it is unlikely for any type of BH system to explain jet precession due to the result from a combination of the Lense-Thirring effect and accretion-disc viscosity. In general, the angular momentum of the disc is not sufficiently large, in comparison with that of a spinning hole, to cause jet precession on short timescales. AGN accretion events are assumed to be uncorrelated and this would cause changes in the jet directions on comparably large timescales of yr. Chaotic accretion is discussed in more detail by King Nixon (2015) – the effect of self-gravity limits the mass of any accretion-disc feeding event. They suggest that additional disc physics is needed to explain any jet precession on timescales short compared with the accretion time. Lubow et al. (2002) investigate the shape of a warped disc around a Kerr BH. In particular, they investigate this system under conditions such that the warp propagates in a wave-like manner. One of their results is, that in low-viscosity discs around a Kerr BH, the inner parts of the disc are not necessarily aligned with the BH. Because the shape of the inner disc depends sensitively on the radial dependence of the disc properties, a changing accretion rate can cause a change in the inner disc warp even without effecting the tilt of the outer disc. This finding contrasts with that of, for example, Nelson Papaloizou (2000) . Lai (2003) discusses the warping of an accretion disc due to magnetically driven outflows as a possible origin for jet precession. In particular, the disk might be subject to warping instabilities and retrograde precession driven by magnetic torques associated with the outflow. The growth timescale for the disk warp and the precession period would be of order the radial infall time of the disk. Caproni et al. (2006) study four mechanisms (tidal, irradiative, magnetic, and Bardeen-Petterson) that can cause warping and precession in galactic and extragalactic accretion disks. Eight X-ray binaries and four AGN are examined. The paper demonstrates how difficult it is to constrain the precession mechanism. The Bardeen-Petterson effect can produce the precession timescales compatible with the observed ones in all sources of their sample. However, the work by Caproni et al. (2006) shows that due to the many observational uncertainties other precession mechanisms can not easily be ruled out. In addition, more than one mechanism might be at work in a given source. Another alternative explanation of the observed effects could result from a magnetically-arrested disc (e.g., McKinney et al. 2012 ; Tchekhovskoy et al. 2014 ). Using


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the albedo of ice on M-dwarf hosted Earth-like planets () and the ice-albedo feedback becomes a negative, or stabilizing, feedback. We tested particular land distribution scenarios (i.e., configurations in which a nonuniform (by latitude) land fraction is used as input), including present-day Earth, and equatorial ocean and land belt configurations, that revealed some variation in surface temperature (∼5 K) and albedo (∼0.02), depending on the distribution of land in the model (see labeled markers in Figure 2 ). In particular, ocean-dominated planets and land planets with oceanic equatorial belts (with otherwise identical orbital parameters) have higher global mean surface temperatures when compared to a uniform land fraction distribution by latitude. Distinct zones of high and low albedo surfaces are shown in the simulations pertaining to planets orbiting F-, G-, and K-dwarfs (Figure 3 ), which represent frozen (yellow) and unfrozen surfaces, respectively. Note that these frozen surfaces exhibit different Bond albedos depending on SED, despite being comprised of the same 50% snow/ice mixture. Transitional boundaries between these zones are also evident, representing the maximum equator-ward extent of glacial conditions, i.e., a nominal “ice line” in latitude and F L parameter space, from which we can approximate the minimum land fraction required to force mid-latitude glaciations, assuming default model conditions: 0.46, 0.55, and 0.68 F L for planets around F-, G-, K-dwarfs, respectively. Figure 2. Mean global surface temperature (; left) and Bond albedo (; right) by average planetary land fraction ( F L ) for terrestrial planets orbiting FGK- and M-dwarf stars. The dashed lines labeled denote the albedos of the 50% ice-snow mixture on planets orbiting X -type star; the gray shaded area depicts a range of F L over which exceeds the blue shaded region illustrates temperatures below 271 K. Markers indicate some individual land fraction and latitudinal distribution configurations tested in the model: present-day “Earth” (E), a scenario in which a narrow band of equatorial ocean exists on an otherwise land-covered planet (Eq. O), an “oceanplanet” (O) configuration with some land at the poles, an equatorial land belt on an otherwise oceanic planet (Eq. L.), and a “landplanet” (L) distribution with some open ocean at the poles. See text for discussion. Figure 3. Contours of albedo (dimensionless) as a function of land fractional coverage ( F L ) and latitude ( φ ) for a planet orbiting an M-dwarf (top left), K-dwarf (top right), G-dwarf (bottom left), and F-dwarf (bottom right) stars. M-dwarf planets do not exhibit mid-latitude glaciations in any of these simulations, although albedos for high land fraction planets in orbit around these stars are proportionally higher in the mid-latitudes and equatorial regions than at the frozen polar regions. A general trend exists in that planetary broadband albedo increases with increasing F L regardless of


O O O O O O O O O O O B-EntityOfFutureInterest O O O O O O O O O O O O O O O O O O O B-Wavelength O O O O B-Citation I-Citation I-Citation 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 B-Dataset I-Dataset O O O O O B-EntityOfFutureInterest O O O O B-Survey I-Survey I-Survey I-Survey I-Survey I-Survey B-Survey O O B-Collaboration I-Collaboration I-Collaboration I-Collaboration I-Collaboration I-Collaboration I-Collaboration O O O O O O B-Citation I-Citation I-Citation O B-Citation I-Citation I-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-Wavelength O O O O O O O O O O O O O B-Wavelength O O O O O O O O O O O O O O O O O O O O O O B-EntityOfFutureInterest O B-Citation I-Citation I-Citation O B-Citation O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O B-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 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 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 B-EntityOfFutureInterest O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O

BBHs in the sense that it can be detectable before the GW signal and that the luminosity increases with time. However, the luminosity and total energy are too low to explain GRBs or the GBM event (Connaughton etal. 2016 ). Figure 3. The photon fluxes from the internal shocks for models A, B and C. The 1-scan sensitivity of MAXI GSC is also plotted. It is found that all the models are detectable by MAXI GSC. The optical follow-up surveys for GW counterparts, such as the Panoramic Survey Telescope Rapid Response System (Pan-STARRS) and the Japanese collaboration for Gravitational wave ElectroMagnetic follow-up, have a sensitivity of 1921 mag (Morokuma etal. 2016 ; Smartt etal. 2016 ). Assuming that 10percent of L AG is in the optical range, we calculate the distance of detection limit for the afterglow, d L, lim , whose values are tabulated in Table 1 . While the afterglow is fainter than the internal shock emission, the optical follow-up surveys of afterglows can detect comparable or more distant events than the X-ray monitoring systems owing to the good sensitivity. However, d L, lim for model A is shorter than the distance of the observed GW events (Abbott etal. 2016a , b ). Although d L, lim is larger for IMBHs and MBHs, we do not discuss the detection probability because the merger rates of IMBH and MBH binaries are very uncertain. 4 SUMMARY and DISCUSSION We study evolution of an accretion disc in BBH systems and propose an evolutionary track of the disc, which leads to different conclusion from the previous work (Perna etal. 2016 ). At first, the disc viscously expands outward but the companion prevents the disc from expanding beyond r out due to the tidal torque. The evolution of viscous disc results in the decrease of the disc mass and the temperature. When the disc sufficiently cools down (typically less than 3000K), the dead disc forms because MRI becomes inactive. Since the thermal instability causes the rapid drop of the disc temperature, the disc becomes dead when the temperature becomes less than a few tens of thousands K. This dead disc remains until the binary separation sufficiently decreases. As the binary separation decreases, the position at which the tidal torque is effective moves inward, and the mass of the outer rim increases. Then, the angular momentum is transported by the tidal torque, which induces the mass inflow from the outer rim to the dead disc. When the mass inflow by the tidal torque becomes higher than , the accretion heating activates MRI, restarting the mass accretion from the disc to the central BH (the disc revival). This disc revival typically happens tens of thousands years before the merger event. The evolution of the revival disc is determined by the tidal torque, keeping t vis t GW . The mass accretion rate of the revival disc increases with time. In the late phase of the revival disc evolution, the mass accretion rate can exceed Eddington rate, and a relativistic jet is expected to be


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Wyatt ( 2014 ) showed that the disc should be truncated at a semimajor axis (at apocentre) where R H,Q is the Hill radius of an eccentric planet at apocentre. Under the influence of two planets, that inner edge could be further out if the strongest secular resonance is located beyond a in . Using equation (24) in Yelverton Kennedy ( 2018 ), we find the approximate location of the resonance to be (dashed blue line in Fig. 9 ) where a 1 and a 2 are the semimajor axes of the companions, while m 1 and m 2 are their masses, with 1 designating the innermost companion and 2 the BD. In order to effectively truncate the disc, the width of the secular resonance would need to be large enough to excite solids from a in to . This width is controlled primarily by the eccentricity of HD 206893 B. Given the uncertainties on its eccentricity () and its mass (12–50 M Jup ; Delorme et al. 2017 ), it is hard to assess whether secular resonances will deplete a region that otherwise would be stable. Assuming the best-fitting values of our fit and using equation (36) in Yelverton Kennedy ( 2018 ), we expect a secular resonance width of ∼20 au. Therefore, it is possible that the disc inner edge is not truncated by the BD, but by secular resonances. Note that HD 206893 B alone could truncate the disc out to ∼30 au if it has an eccentricity of 0.2 and a mass of 50 M Jup . A more precise characterization of HD 206893 B’s orbit and the disc inner edge could provide better estimates of the masses of the two inner companions. 5.2.3 Secular resonances between BD and planet at 30 au Secular resonances might still have created the gap if instead the resonance at 74 au was with the BD at 11 au and an outer planet sitting just interior to the disc inner edge at around 30 au. We find that this putative planet and the BD at 11 au would need to have a mass ratio close to unity (see solid orange line in Fig. 9 ). This is ruled out by direct imaging, which did not detect any companion more massive than 3 M Jup at 30 au (assuming an age of 50 Myr; Delorme et al. 2017 ). 5.2.4 Secular resonances between BD and planet at 74 au If the gap at 74 au was indeed carved by a planet at that distance, it is worth discussing what other observables could confirm this scenario. Given the range of possible masses of the BD and the putative planet c at 74 au, a secular resonance could be located in between the orbit of b and the disc outer edge. We find that the resonance would be located between 85 and 170 au for mass ratios of b and B of 10 −3 to 0.3 (solid green line in Fig. 9 ). This resonance would create a gap that is expected to be very wide (≳50 au if b has an eccentricity of ≳0.05 according to equation 36 in Yelverton Kennedy 2018 ), and thus noticeable by our observations. This means that we can already rule out that this gap is present in between 74 and 110 au, otherwise the gap would be seen to be much wider and the disc peak further out (or at 40 au instead). Therefore, imposing


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

ACKNOWLEDGEMENTS ER, IFS and VL acknowledge the Programme National ‘Physique et Chimie du Milieu Interstellaire’ (PCMI) of CNRS/INSU with INC/INP co-funded by CEA and CNES. They also thank for generous financial support from La Région Haute-Normandie via the GRR Electronique, Energie et Matériaux, from the ‘Fédération de Recherche Energie, Propulsion, Environnement”, and from the LabEx EMC 3 and FEDER via the projects PicoLIBS (ANR-10-LABEX-09-01), EMoPlaF and CO 2 -VIRIDIS. IFS and VL thank PHC GALILEE 2018 PROJET (39379SF) and the GdR THEMS. IFS and JZM acknowledge support from the IAEA via the Coordinated Research Project ‘Light Element Atom, Molecule and Radical Behaviour in the Divertor and Edge Plasma Regions”. JZM acknowledges support from USPC via ENUMPP and Labex SEAM. This work is supported by BATTUTA Project (Building Academic Ties Towards Universities through Training Activities) in the frame of the Erasmus Mundus program, at LOMC UMR-CNRS-6294 of Le Havre University. YM thanks the SRI department, especially Mrs. Martine Currie, for outstanding hospitality.


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We are indebted to an anonymous referee, who provided very useful comments and suggestions from which we have indeed learned a lot, and which led to significant improvement of this paper. This work has been supported by Hungarian OTKA Grant NN 107637 (PI Vinko).


O O B-Event I-Event I-Event O O O O O O O O O O B-Organization I-Organization O B-Event I-Event I-Event I-Event O O O O O O O O O O B-Organization B-Event I-Event I-Event I-Event O O O O B-Person B-Person B-Person B-Person B-Person O O O O O O O O O O O O O O O B-Person I-Person O O O O O B-Telescope O O O O O B-Person I-Person O B-Person I-Person 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 B-Person O O O O B-Organization I-Organization O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization O B-Organization I-Organization O O O O O O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization O O B-Organization I-Organization I-Organization B-Person O O O O O O B-Organization I-Organization I-Organization O O B-Organization B-Grant O B-Organization O B-Grant B-Person O O O O B-Organization B-Fellowship I-Fellowship I-Fellowship I-Fellowship I-Fellowship O O B-Fellowship B-Person O O O B-Organization I-Organization I-Organization I-Organization I-Organization I-Organization I-Organization B-Organization

We thank Astro Hack Week for providing the venue that started this project and the Lorentz Center workshop “The X-ray Spectral-Timing Revolution” (2016 February) that started the collaboration, as well as the Google Summer of Code Program for funding five coauthors (O.M., U.M.K., H.M., H.R., S.S.) who implemented a large fraction of the various library components over three summers. We thank Jim Davenport for providing help with the Kepler data. We are grateful to Matteo Guanazzi and Erik Kuulkers for their advice and suggestions for implementing DAVE, as well as the anonymous referee and the AAS statistics editor for their helpful comments on the manuscript. D.H. acknowledges support from the DIRAC Institute in the Department of Astronomy at the University of Washington. The DIRAC Institute is supported through generous gifts from the Charles and Lisa Simonyi Fund for Arts and Sciences and the Washington Research Foundation. M.B. is supported in part by the Italian Space Agency through agreement ASI-INAF n.2017-12-H.0 and ASI-INFN agreement n.2017-13-H.0. A.L.S. is supported by an NSF Astronomy and Astrophysics Postdoctoral Fellowship under award AST-1801792. E.M.R. acknowledges support from Conselho Nacional de Desenvolvimento Científico e Tecnológico (CNPq—Brazil).


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

some insight into the spatial structure of the segregated groups. The upper right panel of Fig. 2 shows the map of the median positions of the segregated bins over the grey-scale density contours of the group A sample, using the same colour code as in the spectrum and histogram. As could also be seen in the contour plot from Fig. 1 , group A has a bilobed northsouth double peak structure. The two bins associated with segregated group (in yellow) are located towards the western medium density structure and the lower density peak. When we consider the RV, these bins are consistent with the unconnected components A2 and A3 previously described. The bins associated with group (red) are spread over the bilobed structure: the two that comprise 1 are associated with the southern and the one constituting 2 with the northern lobe. The bins associated with segregated groups (green) and (cyan) are both clearly associated with the upper density peak. In the lower right panel of Fig. 2 , the median RV of each bin is shown against its median (with segregated bins coloured with the same colour code as in previously described panels), to highlight the trend of the whole of group A: bins with higher RV tend to be located towards the north. The trend is noisy but clear, particularly when we consider a linear fit of the data, shown in Fig. 2 as a purple line. Fig. 3 shows in four panels the coloured contours of the density distribution of the segregated groups plotted over the grey-scale density of the complete sample, following the same colour code as in Fig. 2 : from left to right and top to bottom the density map of (yellow), (red), (green), and (cyan). Once again, the relationship of the RV of each group with the declination of its density peak is clear. Group is mostly concentrated around the lower peak and the east/central structure, with the highest density in the western area. Group , which is also the one with the most members, is distributed along the whole bilobed structure, reaching the highest density in the lower peak. Group is dense in both peaks, but mostly in the northern one, with a trend towards the east. Finally, group is very clearly associated with the northern peak of the bilobed structure. When we consider the unconnected components of group 1 and 2 the general RV declination trend still applies: group 1, with a lower RV, peaks in the lower lobe, while 2 peaks around the centre of the spatial domain, between both lobes. Figure 3. Spatial distribution of the four segregated groups from the analysis of group A. In each panel, the coloured contours show the density star distribution of each segregated group, with the same colour code as in Fig. 2 , over the grey-scale density of the group A. From left to right and top to bottom: group , in yellow, with kms 1 , group , in red, with kms 1 , group , in green, with kms 1 , and group , in cyan, with kms 1 . We will now compare these results with those obtained by Frsz etal. ( 2006 ) and Tobin


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

f NS , One can see from these equations that the luminosity of seed photons produced by the NS surface anticorrelates with the NS spin. For higher spin sources, supply of the seed photons will be reduced resulting in stronger Comptonization, in particular in large Compton amplification factor A and Comptonization parameter y , in the same way as the dichotomy between BH and NS sources had been explained in Paper I . Furthermore, assuming that the emission from the boundary layer is well approximated by a blackbody spectrum, then L bl ∝ T 4 . This means that if all other factors (especially the emitting surface area) are similar across the population of NSs, the temperature of the seed photons . The maximal Keplerian frequency in the accretion disc is of the order of f K ∼ 1.6–1.8 kHz (Sibgatullin Sunyaev 2000 ), while NS spins vary in the f NS ∼ 0.3–0.6 kHz range. At L X / L Edd ∼ 0.05, the seed photon temperature varies in the kT BB ∼ 0.8–1.0 keV range, i.e. by a factor of about ∼1.2–1.3, between low- and high-spin sources (Fig. 3 ). In order to explain such a difference between f NS ∼ 0.3 kHz and f NS ∼ 0.6 kHz sources, α ∼ 0.7 is required, which is broadly consistent with the conclusions of Paper I that the accreting material reaches the surface of the NS still possessing about ∼1/3–1/2 of its initial energy. 5.1 Can fast spinning NSs mimic BHs? From equation ( 2 ), at sufficiently high NS spin frequencies the luminosity of the boundary layer can become small. Theoretically, this can make accreting NSs similar to BHs from the point of view of the energy balance in the Comptonization region. Indeed, such a behaviour is evident in Fig. 4 where the high-spin sources at high Eddington ratios progressively overlap with the region of Comptonization parameters, identified with BHs in Paper I . This dilutes the boundary between BHs and NSs. From Fig. 2 , one can see that three sources expand into y ≳1.0 regions at higher L / L Edd : 4U 1636–536, KS 1731–260 and, most significantly, GS 1826–238. This behaviour is strongest at the highest Eddington ratios (but still compatible with the source being in the hard spectral state). All these sources have high spin (>500 Hz), and their behaviour markedly contrasts with that of the slower sources at similar L / L Edd , such as 4U 1728–34, the Comptonizing strength of which remains low. For GS 1826–238, which has one of the highest spins and mass accretion rates in our sample, the Comptonization parameter can reach values y ∼ 1.25, which is inside the region occupied by BH LMXBs. 2 However, such high values of y are achieved only at high Eddington ratios, L X / L Edd ∼ 0.1, at which NSs are characterized by extremely low electron temperature, kT e ∼ 10–20 keV (Fig. 2 ), well below values typical for BHs (fig. 5 in Paper I ). This behaviour in kT e − y distinguishes NSs from BHs and is illustrated in Fig. 5 . A detailed comparison of high-spin NSs with BHs will be a subject of a separate study. As a final


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

is shown in Fig. 4(b) . Figure 4. (a) Reconstruction of B on a 1000-yr time-scale. Grey circles – 400-yr reconstruction. (b) Reconstruction of B on a 3200-yr time-scale. The grey areas are error intervals. 6 RESULTS AND DISCUSSION Here, using the MSR method, implemented by using equations ( 1 )–( 4 ), we have built reconstructions of the variations of the IMF strength B on different time-scales: 170, 400, 1000, and 3200 yr, based on different proxies ( IDV , Wolf numbers 2.0, and radiocarbon reconstruction). Typical correlation coefficients that we obtained are within the range 0.92–0.95, which on the one hand indicates the success of B reconstructions, and on the other hand, shows that we do not take into account all the factors of B variations in proxies, but only ∼85–90 per cent of the dispersion. Changes in IMF strength B vary significantly: for annual values , for a decade DB ∼ 2.5. Note that the cyclicity that makes the largest contribution to the dispersion in Fig. 4(b) is made by variations with the period slightly longer than 200 yr (Suess cycle). In order to see the most long-term trends in IMF variations, we compose the average values of B over 210 yr (see Fig. 5 ). Figure 5. The variations in the last 3200 yr of the B mean values for 210 yr (black circles) and their smoothed values over three points (white circles). Dashed line – linear trend, dotted line – mean, and dash–dotted line – forecast. Noteworthy is the long-term decreasing trend, which is −(0.39 ± 0.17) · 10 −3 nT· yr −1 over the past 3200 yr. Although its significance is ensured only by 2.3σ, the values smoothed over three points and the behaviour in time relative to the average (Fig. 5 ) speak of its reality. We also note that geomagnetic changes are unlikely to have contributed to the nature of this trend since we believe that the reconstruction used by Usoskin et al. ( 2014 ) is based on modern geomagnetic dipole moment data from the 1000 archeomagnetic field models (Licht et al. 2013 ). We also believe that the relationship between IDV and sunspot numbers depends only on the time-scale τ, but not on time t : any reconstruction is generally based on this assumption of time invariance. At the same time, an episode of increasing B has been observed over the past ∼400 yr. This means that the cosmic ray flux near the Earth’s orbit during these years decreases with time. As speculation, according to the Tinsley–Heelis mechanism (Tinsley Heelis 1993 ), this should lead to a decrease in the reflective properties (albedo) of the Earth’s atmosphere and an increase in the solar energy flux that reaches the Earth’s surface. All of this can make a partial contribution to the global warming effect. On the other hand, judging by the general nature of the variations in Fig. 5 , in the next 200–250 yr we will face an episode of a decrease in the IMF strength so that the increase in the Earth’s global temperature will decrease or decrease the increase rate if it is really related


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

M.C. is supported by the Scientific and Technological Research Council of Turkey (TÜBİTAK) through the Research Fellowship Programme for International Researchers (2216). We thank Amy Y. Lien for valuable suggestions regarding the analysis of Swift BAT survey data.


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

of horizontal field changes in shear region P is longer and followed by a flatter gradual increase. 4. Summary and Discussions We have presented a detailed study of the structural evolution of photospheric flow and magnetic fields associated with the 2015 June 22 M6.5 flare, concentrating on the central penumbrae of the δ region around the flaring PIL. We tracked the penumbral flow field with the DAVE method using high-resolution TiO images from BBSO/GST, and also analyzed the associated magnetic field using SDO /HMI vector magnetograms. The main results are summarized as follows. 1. Beginning from the first HXR peak, one segment of penumbra lying in the negative field experiences a distinct grow, extending northward at 1.5(±0.05) km s −1 for a distance of 26(±3); meanwhile, the flow velocity within the entire penumbral region becomes more pronounced, gaining a 115(±10)% increase reaching up to 2.2(±0.11) km s −1 at ∼90 minutes after first flare peak. These structure and flow field evolutions are accompanied by a step-like increase of horizontal magnetic field by 150(±15) G and decrease of inclination angle by 5(±0.5)° in 18 minutes from 17:50 UT to 18:08 UT. 2. A region of shear flow at the flaring PIL is found next to the location of a flare precursor brightening, and expands significantly after the flare. From the first HXR peak, the shear flow velocity increases gradually by 0.4(±0.1) km s −1 from 0.5(±0.05) km s −1 at 17:30 UT to 0.9(±0.05) km s −1 at 19:30 UT. As for magnetic field properties in the shear flow region, the horizontal field strength increases 220(±20) G while the inclination angle decreases 5(±0.5)° from 17:50 UT to 18:20 UT, both following a stepwise fashion. In comparison with the extending penumbral segment, the transition time of step-like magnetic field changes in the shear flow region is ∼10 minutes longer. Our results of the strengthening of penumbral structure at the center of the δ region, the associated increase (decrease) of horizontal field (inclination angle), and their close timing relationship to the flare energy release strongly favor the back-reaction of coronal fields as the underlying cause, which results in a more horizontal configuration of the photospheric field. While most other work only deals with the flare-related overall intensity change of penumbral structure, we are able to study the evolution of penumbral flow field owing to the high-resolution BBSO/GST data. We show that, although noticeable intensity darkening is only observed in the newly formed penumbral portion, the strengthening of penumbra is predominantly manifested as the enhancement of penumbral flow velocity and enlargement of penumbral area. The inconsistency of our results with the claim of darkening of the central penumbral feature in some previous studies may be due to the limitation of the low-resolution images used (e.g., Liu et al. 2005 ) and also the complexity of the central penumbral area near the PIL (Wang et al. 2017


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

part of the H α line is not due to chromospheric activity. Figure 11. Variations in the Mg ii λ2800 emission line in several IUE spectra. [A colour version of this figure is available in the online version.] Following the procedure explained in Acker Jasniewicz ( 1990 ) for Abell 35, we have tried to find a modulation of the H α profile with the rotation phase, by measuring the flux variation between the nearby continuum and the deepest point of the H α absorption core. In contrast to the results found by Acker Jasniewicz ( 1990 ) in the case of Abell 35, we do not find any correlation with the rotational period during the long time span of the observations in LoTr 5. We have inspected the H α profile in all the Hermes spectra and we found no correlation with the rotation period, see Fig. A1 , where we show the phased-folded trailed spectra of the H α line with the rotation period. The Ca ii HK λλ3933, 3967 lines appear in emission in the HERMES data set (as already reported by Strassmeier et al. 1997 ). We do not find variations in these lines at the level of the H α variations that we find in HERMES data. Hence, Ca ii HK lines are not useful for studying the activity of the star. 3.5.2 Accretion disc around the possible close companion The presence of an accretion disc in LoTr 5 has been also discussed in the literature. If indeed the H α emission was due to an accretion disc, then the double-peaked profile would come from two different emission regions instead of a broad emission feature. Montez et al. ( 2010 ) already discussed the system within this scenario: since the presence of an accretion disc would imply the RLOF, they calculated the Roche lobe radius ( R L ) assuming the synchronization of the rotation and orbital period (i.e. they assumed a close 0.6 M ⊙ companion at ∼5.9 d). Under this configuration, the R L was below the radius of the star, filling its Roche Lobe and allowing the existence of an accretion disc. However, from our radial velocity analysis, we know that it is not possible to have such a massive companion so close to the G-type star. On the contrary, in the case we would have a low-mass object at ∼5.9 d, the R L would be ∼16 R⊙, and the star could not have filled its Roche lobe. Therefore, with the current data, we conclude that it is not possible to have an accretion disc via RLOF in LoTr 5. Finally, and as mentioned in Section 3.2 , no significant infrared excess is identified in the SED, which would point out that an accretion disc is highly unlikely in the system. 3.5.3 The possible resemblance with symbiotic stars Finally, it is worth mentioning that this double-peaked H α emission resembles the observed in some symbiotic stars (SS), the widest interacting binary systems known with a white dwarf accreting material from a giant star via stellar winds (see e.g. Van Winckel, Duerbeck Schwarz 1993 ; Burmeister Leedjärv 2009 ). These systems are sometimes embedded into a bipolar nebula similar to some PNe (like


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We gratefully acknowledge funding for Swift at the University of Leicester by the UK Space Agency. AAR's studentship funding is provided by the STFC. We thank the referee for their useful comments and suggestions.


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ACKNOWLEDGEMENTS We thank Scott Tremaine for helpful discussion, and comments on the manuscript. We also thank Stephen McMillan and Piet Hut for stimulating discussions, Douglas Heggie, Fred Rasio, Evgeni Grishin, and Chris Hamilton for comments on the manuscript, and the anonymous referee for a helpful report. ASH gratefully acknowledges support from the Institute for Advanced Study, and the Martin A. and Helen Chooljian Membership. JS acknowledges support from the Lyman Spitzer Fellowship.


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 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-Wavelength B-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 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 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 B-Wavelength B-Wavelength O O O O O O O O O O O O O O O O O O O O O B-Wavelength B-Wavelength O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 O O O O O O O O O O O O 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 O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O O 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 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 B-Wavelength I-Wavelength O O O O O O B-Wavelength I-Wavelength O O O O O O O O B-Database I-Database I-Database I-Database O O O O O O O O O O O O O B-Wavelength B-Wavelength 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 O O O O O O O O O O O O O O

from the Coronene, Ovalene, Anthracene, Tetracene families ( Ricca et al. 2012 ; Peeters et al., in prep.), and the Pyrene family. While most of these PAHs are very compact, we have also included a few non-compact ones, to investigate the effect of the molecular structure on the 3.3/11.2 m ratio. All the PAHs in the selections are shown in Fig. A.2 ). As the UV absorption cross-sections for all PAHs are very similar, the uncertainty produces a vertical shift in Fig. 5 . Figure 5 shows that the intensity ratio is sensitive to the size of the PAHs ( Allamandola et al. 1985 ) and forms a well-defined relationship ( Ricca et al. 2012 ). The slope and shape of the relation is not sensitive to the stellar input spectrum or the extinction in the PDR. However, it can be shifted vertically for different average photon energies or adopted intrinsic strength. 4.2. PAH size With the model (Fig. 5 ), the 11.2/3.3 m ratio can be converted to PAH sizes. The horizontal colored lines in the figure have the same color coding as the 11.2/3.3 m intensity ratio in the maps. In the NW PDR (orange in Fig. 2 ), PAHs have a typical size of 50 carbon atoms. Moving towards the star, the ratio drops (green) and PAHs have average sizes of . In the south PDR the average sizes of the PAHs are slightly larger with respect to the NW PDR, (green) in the PDR surface, and increase towards the left up to (light blue) in Fig. 4 . The results show that the 11.2/3.3 m ratio is mainly sensitive to size and has little dependence on molecular structure. Matrix isolation and gas phase experiments, as well as theoretical density functional theory (DFT) calculations show that the intrinsic band-strength ratio of the C-H stretch to the out-of-plane-bending mode (3.3/11.2 m) is very constant between different PAHs ( Joblin et al. 1994 ; Langhoff 1996 ). However, the ratios derived by these different techniques are different. Further experiments are needed to settle this issue. However, we emphasize that a different adopted band-strength ratio will only produce a vertical shift in Fig. 5 . An uncertainty in the adopted intrinsic strength translates to a systematic uncertainty in quantifying PAH size. If the adopted intrinsic strength is increased by a factor of 2, then a range of –105 would change to a range of –80. The trends, however, remain. Fig. 5 11.2/3.3 m intensity ratio of 27 different PAHs with blue coronene family, green compact, and red non-compact PAHs. The average photon energy absorbed by the PAHs used is 6.5 eV (see Eq. ( 1 )). The ratio of the C-H stretching mode (3.3 m) over the solo out-of-plane bending mode (11.2 m) is extracted from the theoretical spectra from the NASA Ames PAH database. The -axis is linearly proportional to the adopted intrinsic band strength of the 11.2/3.3 m ratio. The colored horizontal lines have the same color coding as used for the 11.2/3.3 m ratio maps in Figs. 2 and 4 . 4.3. PAH size variation in the PDR Figure 6 shows a crosscut across


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Stokes V spectrum profile characteristic of circular polarization generated by the Zeeman effect, but it can arise from scattering of a cloud outside the velocity range of the maser. Non-anti-symmetric Stokes V spectra were observed by Herpin et al. (2006) , but they can also be explained by a velocity gradient along the propagation path of the maser, or the lack of spatial resolution from the single-dish observations. 5.2. H2O masers 5.2.1. Simulations The relevant characteristic maser luminosities are given in Table 4 . We list the relevant luminosities for individual hyperfine transitions as well as the blended line. Compared to the SiO maser, radiative interactions remain relatively weak with respect to magnetic interactions up to high maser luminosities. This is due to the much smaller line strength of this maser transition, which means that the Zeeman effect will be the dominating polarizing mechanism up to high maser luminosities, and will thus follow Eq. ( 38 ) up to high maser brightness. Linear polarization will also remain rather low because the isotropic decay will be a dominant de-polarizing entity up to ( T b ΔΩ) sat (Ksr) at about ∼10 10 Ksr. Strong linear polarization is thus only seen for the strongest masers. Table 4. Characteristic maser luminosities for the 22 GHz water maser. For the regime T b ΔΩ ≪ ( T b ΔΩ) mag. sat , the maser circular polarization can be described by Eq. ( 38 ). A LTE analysis of the constant A F F ′ , gives for the individual hyperfine transitions, A 76 = 13.3, A 65 = 8.3, and A 54 = 1.0. A LTE analysis of a completely blended water maser line gives A blend = 8.2. Figure 15 shows the results of our full radiative transfer analysis of the circular polarization constants. Apart from the standard maser line profile, water masers are further broadened by their hyperfine structure. This leads to an overestimation of Δ v L . The dominant Zeeman effect though, comes from a single hyperfine transition. This produces higher Zeeman A F F ′ coefficients with respect to a LTE analysis. We observe that this effect is most pronounced for masers pumped at v th = 0.6 km s −1 , where the hyperfine transitions are minimally mixed. At v th = 2.0 km s −1 , the hyperfine broadening is negligible and the LTE value for the A F F ′ coefficient of the F = 7−6 hyperfine transition is returned for the weakest masers. Fig. 15. The A F F ′ coefficients of an isotropically pumped water maser at B = 20 G as a function of the magnetic field-propagation direction angle cos θ . The panels show different thermal widths v th = (a) 0.6 km s −1 , (b) 1.0 km s −1 , and (c) 2.0 km s −1 . Plots are given for different log( R / g Ω). The LTE solutions (constant over cos θ ) of the different hyperfine sub-transitions are indicated by a dotted line. Paradoxically, the preferred pumping of the hyperfine component with the strongest Zeeman effect has a consequence that the linear polarization deteriorates in the relevant intensity