Wind Properties of Variable B Supergiants Evidence of Pulsations Connected with Mass-Loss Episodes? M

A&A 614, A91 (2018) Astronomy https://doi.org/10.1051/0004-6361/201731678 & © ESO 2018 Astrophysics

Wind properties of variable B supergiants Evidence of pulsations connected with mass-loss episodes? M. Haucke1,2, L. S. Cidale1,2,3, R. O. J. Venero1,2, M. Curé3, M. Kraus4,5, S. Kanaan3, and C. Arcos3

1 Departamento de Espectroscopía, Facultad de Ciencias Astronómicas y Geofísicas, Universidad Nacional de La Plata, Paseo del Bosque s/n, La Plata, Argentina e-mail: [email protected] 2 Instituto de Astrofísica de La Plata, CONICET-UNLP, Paseo del Bosque s/n, La Plata, Argentina 3 Instituto de Física y Astronomía, Facultad de Ciencias, Universidad de Valparaíso, Av. Gran Bretaña 1111, Casilla 5030, Valparaíso, Chile 4 Astronomický ústav, Akademie vedˇ Ceskéˇ republiky v.v.i., Fricovaˇ 298, 251 65 Ondrejov,ˇ Czech Republic 5 Tartu Observatory, Tõravere, 61602 Tartumaa, Estonia

Received 31 July 2017 / Accepted 25 December 2017

ABSTRACT

Context. Variable B supergiants (BSGs) constitute a heterogeneous group of stars with complex photometric and spectroscopic behaviours. They exhibit mass-loss variations and experience different types of oscillation modes, and there is growing evidence that variable stellar winds and photospheric pulsations are closely related. Aims. To discuss the wind properties and variability of evolved B-type stars, we derive new stellar and wind parameters for a sample of 19 Galactic BSGs by fitting theoretical line profiles of H, He, and Si to the observed ones and compare them with previous determinations. Methods. The synthetic line profiles are computed with the non-local thermodynamic equilibrium (NLTE) atmosphere code FAST- WIND, with a β-law for hydrodynamics. Results. The mass-loss rate of three stars has been obtained for the first time. The global properties of stellar winds of mid/late B supergiants are well represented by a β-law with β > 2. All stars follow the known empirical wind momentum–luminosity relationships, and the late BSGs show the trend of the mid BSGs. HD 75149 and HD 99953 display significant changes in the shape and intensity of the Hα line (from a pure absorption to a P Cygni profile, and vice versa). These stars have mass-loss variations of almost a factor of 2.8. A comparison among mass-loss rates from the literature reveals discrepancies of a factor of 1 to 7. This large variation is a consequence of the uncertainties in the determination of the stellar radius. Therefore, for a reliable comparison of these values we used the invariant parameter Qr. Based on this parameter, we find an empirical relationship that associates the amplitude of mass-loss variations with photometric/spectroscopic variability on timescales of tens of days. We find that stars located on the cool side of the bi-stability jump show a decrease in the ratio V /Vesc, while their corresponding mass-loss rates are similar to or lower than the values ∞ found for stars on the hot side. Particularly, for those variable stars a decrease in V /Vesc is accompanied by a decrease in M˙ . Conclusions. Our results also suggest that radial pulsation modes with periods longer∞ than 6 days might be responsible for the wind variability in the mid/late-type. These radial modes might be identified with strange modes, which are known to facilitate (enhanced) mass loss. On the other hand, we propose that the wind behaviour of stars on the cool side of the bi-stability jump could fit with predictions of the δ slow hydrodynamics solution for radiation-driven winds with highly variable ionization. − Key words. stars: early-type – supergiants – stars: mass-loss – stars: winds, outflows

1. Introduction are mainly driven by line scattering of UV photons coming from the stars’ continuum radiation. The standard stationary radiation- Massive stars have a significant impact on the ionization, struc- driven wind theory (Castor et al. 1975; Pauldrach et al. 1986; ture, and evolution of the interstellar medium (ISM). Via their Friend & Abbott 1986) predicts mass-loss rates and terminal stellar winds they release energy and momentum to their sur- velocities as a function of stellar parameters (Abbott 1982; Vink roundings leading to the formation of stellar wind-blown bub- 2000; Vink et al. 2001), and also predicts a wind momentum– bles, bow-shocks, and circumstellar shells. On the other hand, luminosity relationship (WLR; Kudritzki et al. 1995). This the huge amount of mass lost from the stars modifies the later theory has long described the global properties of stellar winds stages of their lives (evolutionary timescales and final core of OBA stars. However, advances in the observational techniques masses) as well as the type of SN remnants left (Meynet et al. (high spatial and spectral resolution) together with progress in 1994; Woosley et al. 2002). Stellar winds of hot massive stars more realistic atmosphere models have revealed discrepancies in the mass-loss rates not only between theory and observations, ? Based on observations taken with the J. Sahade Telescope at Com- plejo Astronómico El Leoncito (CASLEO), operated under an agree- but also amongst different observational tracers themselves (cf. ment between the Consejo Nacional de Investigaciones Científicas Puls et al. 2008). y Técnicas de la República Argentina, the Secretaría de Ciencia y As Owocki & Rybicki(1984), Owocki et al.(1988), and Tecnología de la Nación, and the National Universities of La Plata, later Feldmeier(1998) demonstrated, strong hydrodynamic Córdoba, and San Juan. instabilities occur in line-driven winds of massive O-type stars,

Article published by EDP Sciences A91, page 1 of 28 A&A 614, A91 (2018) leading to the formation of wind-clumping (small-scale density as a consequence of the instabilities, the non-linear simulations inhomogeneities distributed across the wind). Wind-clumping revealed finite amplitude pulsations consistent with the observa- has traditionally been used to solve the issue of the mass-loss tions. discrepancy found in O supergiants (OSGs) when different To obtain further insights on the wind structure and wind diagnostic methods are employed (Bouret et al. 2005; Fullerton variability of evolved stars of intermediate mass, we studied a et al. 2006). Mass-loss rates derived from unsaturated UV sample of Galactic BSGs (classified either as irregular or α Cyg- resonance lines give lower values than those obtained from the type variables) in both the blue and the Hα spectral region. The Hα line (cf. Puls et al. 2008). The macroclumping approach Hα emission is quite sensitive to the wind properties (Puls et al. (with optically thick clumps at certain frequencies) has provided 2008), and we used the Fast Analysis of STellar atmospheres a fully self-consistent and simultaneous fit to both UV and with WINDs (FASTWIND; Santolaya-Rey et al. 1997; Puls et al. optical lines (Oskinova et al. 2007; Sundqvist et al. 2010, 2011; 2005; Rivero González et al. 2012) to obtain terminal velocities Šurlan et al. 2013). In previous works an artificial reduction of and mass-loss rates by fitting synthetic to observed line profiles. the stellar mass-loss rate, an extremely high clumping factor, or In addition, to derive proper stellar parameters (effective tem- an anomalous chemical abundance for specific elements were perature and surface gravity) we consistently modelled the Si invoked; however, the macroclumping approach means that ionization balance together with the photospheric lines of H and none of these is necessary. He. Determining new, more accurate values of mass-loss rates and stellar parameters of massive stars is crucial for our under- The situation is different for B supergiants (BSGs). Stellar standing of stellar wind properties, wind interactions with its winds of BSGs are fairly well represented by the analytical veloc- surroundings, and the plausible mechanisms related with wind ity β-law, with β typically in the range 1 3 (Crowther et al. 2006; variability and its evolution. Markova & Puls 2008). Moreover, these− stars show a drop in the The paper is structured as follows: our observations and mod- ratio V /Vesc of a factor of two at around 21 000 K, from 2.6 for ∞ elling are described in Sects.2 and3, respectively. The results stars on the hot side to 1.3 for stars on the cool side (Pauldrach & of our line profile fittings are shown in Sect.4 and the new Puls 1990), known as the bi-stability jump (Lamers et al. 1995). wind parameters are compared with previous determinations. This jump was interpreted as a consequence of the recombina- For three objects (HD 74371, HD 99953, and HD 111973) the tion of Fe IV to Fe III in the wind (Vink et al. 1999). These wind parameters are reported for the first time. As most of the authors predicted that this drop should be accompanied by a stars in our sample exhibit photometric and spectroscopic vari- steep increase in the mass-loss rate of about a factor of 3 to 5, ability, on timescales from one to tens of days, and have been whilst some observations indicate a decrease (Crowther et al. modelled previously by several authors, in Sect.5 we discuss 2006; Markova & Puls 2008). A recent study of the effects of changes in the wind parameters. We present for mid and late micro- and macroclumping in the opacity behaviour of the Hα BSGs an empirical relationship between the amplitude of mass- line at the bi-stability jump suggests that the macroclumping loss variations and light/spectroscopic periods, indicating that should also play an important role at temperatures on the cool a significant percentage of the mass-loss could be triggered by side of the jump (Petrov et al. 2014). pulsation modes. Our conclusions are given in Sect.6. Finally, It is also well known that stellar winds of BSGs are highly AppendixA shows model fittings to the photospheric lines of variable. To date, most of the studies related to α Cyg vari- each star and summarizes the stellar and wind parameters found ables and hypergiants have focused on their photometric and/or in the literature and in this work. optical spectroscopic (mainly Hα) variability, and many sur- veys have been carried out to search for variability periods 2. Observations (e.g. Kaufer et al. 2006; Lefever et al. 2007; Lefèvre et al. 2009). Long-term space-based photometry and spectroscopy We took high-quality optical spectra for 19 Galactic BSGs of have linked this variability to opacity-driven radial and non- spectral types between B0 and B9. These stars were selected radial oscillations (Glatzel et al. 1999; Lefever et al. 2007; Kraus from the Bright Star Catalog (Hoffleit & Jaschek 1991). The et al. 2015), and new instability domains have been established observations were performed in January 2006, February 2013, in the Hertzsprung–Russell (HR) diagram covering the region of April 2014, and February and March 2015. We used the REOSC BSGs (Saio et al. 2006; Godart et al. 2017). As massive stars can spectrograph (in crossed dispersion mode) attached to the cross the BSGs’ domain more than once, the pulsation activity Jorge Sahade 2.15 m telescope at the Complejo Astronómico El of these stars can drastically change between their red- and blue- Leoncito (CASLEO), San Juan, Argentina. The adopted instru- ward evolution (Saio et al. 2013). BSGs on a blue-ward evolution mental configuration was a 400 l/mm grating (#580), a single (i.e. a post-red supergiant) tend to undergo a significantly larger slit of width 250 µ, and a 1024 1024 TEK CCD detector with × number of pulsations, even including radial strange-mode pul- a gain of 1.98 e−/ adu. This configuration produces spectral res- sations with periods between 10 and 100 days (or more). It has olutions of R 12 600 at 4500 Å and R 13 900 at 6500 Å. ∼ ∼ been suggested that strange-mode pulsations cause time-variable Spectra were reduced and wavelength calibrated following the mass-loss rates in very luminous evolved massive stars (Glatzel standard procedures using the corresponding IRAF1 routines. et al. 1999). The first observational evidence of the presence of The resulting spectra have an average signal-to-noise ratio (S/N) strange modes have been found in two BSGs: HD 50064 (Aerts of 300. ∼ et al. 2010) and 55 Cyg (Kraus et al. 2015). In the latter, the Table1 lists the stars in our programme, indicating name authors found line variations with periods in the range of 2.7 h to and HD number, spectral type, and variable type designation 22.5 days. They interpreted these variations in terms of oscilla- according to (Lefèvre et al. 2009); from a study of HIPPARCOS tions in p-, g-, and strange modes. The last could lead to phases light curves or the VSX database. In addition, binary stars are of enhanced mass loss. Finally, the connection between pulsa- 1 IRAF is distributed by the National Optical Astronomy Observatory, tion and mass loss in 55 Cyg was confirmed by Yadav & Glatzel which is operated by the Association of Universities for Research in (2016), based on a linear non-adiabatic stability analysis with Astronomy (AURA) under cooperative agreement with the National respect to radial perturbations. These authors demonstrated that, Science Foundation.

A91, page 2 of 28 M. Haucke et al.: Wind properties of variable B supergiants

Table 1. Log of observations

Star name HD number Sp. type a Var. type b Period Observation date Wavelength interval days (YY/MM/DD) [Å:Å] β Ori 34 085 B8Iae ACYG 2.075 c, 1.22–74.74 d, 28 n 2006/01/15 [3500:7850] κ Ori 38 771 B0.5Ia IA 1.047 e, 1.9 f , 4.76 e, 6.5 f , 9.5 f 2006/01/15 [3500:7850] χ2 Ori 41 117 B2Ia ACYG 0.92 e, 0.95 e, 2.869 e, 20 e, 40 e, 200 e 2 006/01/15 [3500:7850] PU Gem 42 087 B4Ia g ACYG 6.807 o, 25 e 2006/01/15 [6500:7850] V731 Mon 47 240 B1Ib ACYG, SB q 1.73 j, 2.742 c, e, 133 e 2006/01/15 [3500:7850] CMa 52 089 B1.5 II IA 2013/02/05 [4300:6850] V820 Cas 52 382 B1Ia IA 2006/01/15 [3500:7850] σ2 CMa 53 138 B3Ia ACYG 3.69 j, 24.39 j, 24.44 c 2006/01/15 [3500:7850] 2013/02/05 [4300:6850] η CMa 58 350 B5Ia L, ACYG 4.7 i, 6.631 j 2006/01/15 [3500:7850] 2013/02/05 [4300:6850] J Pup 64 760 B0.5Ib IA 1.2 m, 1.8 j, 2.4 m, 2.8 j, 6.8 m 2013/02/05 [4300:6850] LN Vel 74 371 B6Iab/b IA, ACYG 8.29 c, i, 1, 15–20 h 2006/01/15 [3500:7850] OP Vel 75 149 B3Ia SPB?, ACYG 1.086 c, n, 1.215 j, 2.214 j 2006/01/15 [3500:7850] 2013/02/05 [4300:6850] 2013/02/07 [4300:6850] 2014/04/14 [4200:6650] GX Vel 79 186 B5Ia IA 2006/01/15 [3500:7850] PV Vel 80 077 B2Ia+e GCAS?, SDOR 3.115 c, 21.2 l, 41.5 h, 55.5 h, 66.5 h, 76.0 h 2006/01/15 [3500:7850] 2014/04/12 [4200:6650] V519 Car 92 964 B2.5Ia ACYG 2.119 j, 4.71 c, 14.706 j 2013/02/05 [4300:6850] V808 Cen 99 953 B1/2Iab/b IA 17.7 k 2014/04/14 [4200:6650] 2015/02/13 [4200:6650] 2015/03/13 [4200:6650] κ Cru 111 973 B2/3Ia ACYG?, IA, SB p 9.536 i, 57.11 i 2014/04/11 [4200:6650] 2014/04/12 [4200:6650] ALS 3038 115 842 B0.5Ia/ab ACYG?, IA 10.309 c, 13.38 i 2014/04/11 [4200:6650] V1058 Sco 148 688 B1Iaeqp ACYG 1.845 c, 6.329 c 2014/04/11 [4200:6650]

Notes. (a)From the SIMBAD astronomical database (Wenger et al. 2000). (b)IA (poorly studied variables of early spectral type), L (slow irregular variables), GCAS (γ Cassiopeiae type), ACYG (α Cyg type), SPB (slowly pulsating B stars), and S Dor (S Doradus type) (classification taken from Lefèvre et al. 2009). The question mark means that the classification was based on a new period determination or that the type could not be clearly identified from the HIPPARCOS light curve. (c)Lefèvre et al.(2009), (d)Moravveji et al.(2012), (e)Morel et al.(2004), ( f )Prinja et al.(2004), (g)this work, (h)van Genderen et al.(1989), (i)Koen & Eyer(2002), ( j)Lefever et al.(2007), (k)Sterken(1977), (l)Knoechel & Moffat(1982), (m)Kaufer et al. (2006), (n)Aerts et al.(2013), (o)International Variable Star Index (VSX) database, and (p,q)SB (spectroscopic binary Chini et al. 2012; Prinja et al. 2002). also listed. In the following columns of the same table, we list We opted for a wind model without clumping because we the reported light curve or spectroscopic periods, the observa- want to compare the derived mass-loss rates with previous tion dates of our data, and the achieved wavelength coverage. determinations and to study the wind variability. In general, Here, we present six stars with more than one observation mass-loss rates found in the literature are mainly estimated using (HD 53138, HD 58350, HD 75149, HD 80077, HD 99953, and unclumped models. A second reason is that we do not have con- HD 111973). One of these stars (HD 111973) was observed on temporaneous data in the UV and IR spectral regions to evaluate two consecutive nights. the importance of the clumping factor in the Hα line modelling. As a consequence, our results will provide upper-limit values for 3. Stellar and wind parameters the mass-loss rates. In addition, to obtain an optimal fit to the To derive stellar and wind parameters we made use of the code observed line profiles, it was necessary to consider not only a FASTWIND (v10.1.7). The code computes a spherically expand- microturbulence velocity, Vmicro, and broadening due to the star’s ing line-blanketed atmosphere. All background elements are rotation, V sin i, but also the effect of a macroturbulence veloc- considered with their solar abundances (taken from Grevesse & ity, Vmacro. The latter contributes mainly to line broadening and Sauval 1998). H , He, and Si atoms are treated explicitly with shaping as noted by Simón-Díaz & Herrero(2007). high precision by means of a complete NLTE approach and an To obtain the effective temperature (Teff) we evaluated accelerated lambda iteration (ALI) scheme is applied to solve the the ionization balance (e.g. Si II Si III, Si III Si IV, if the − − comoving-frame equations of radiative transfer (Puls 1991). The Si IV λ 4089, 4116 lines are observed, and He I He II). We used atmospheric stratification is modelled considering a smooth tran- solar abundances for He and Si (log N /N− = 1.07 and He H − sition between a quasi-hydrostatic photosphere and an analytical log NSi/NH = 4.45) and found good fits for each modelled wind structure described by a velocity β-law. The temperature spectrum (see details− in Sects. 4.1 and5). To derive an accu- structure is calculated from the electron thermal balance and is rate determination of the surface gravity (log g), we modelled consistent with the radiative equilibrium condition. the Hγ and Hδ lines. We used a “by eye” fitting procedure

A91, page 3 of 28 A&A 614, A91 (2018) to find the best-fitting synthetic line spectrum to the observed modelled and ∆Teff = 500 K if the Si II, Si III, and He lines were one. used to derive the temperature. From measurements of the wings The Hα line was modelled to derive the mass-loss rate (M˙ ) of Hγ and Hδ we estimated an error of 0.1 dex in log g. The error 1 1 and the parameters of the velocity β-law: the power index β and bars in Vmicro are 2 km s− for Si and He lines, 5 km s− for the 1 the terminal velocity V . It is important to stress that the major H lines, and 10 km s− for Hα. We model all the photospheric ∞ source of uncertainty in the determination of mass-loss rates is lines using the same Vmicro. due to ambiguous determinations of stellar radii and distances Both V sin i and Vmacro affect the line broadening. To model 1 to the stars (see Markova et al. 2004). On the other hand, Hα is the line profiles we adopted the value of V sin i ( 10 km s− ) rather sensitive to V (Garcia et al. 2017). Therefore, the stellar found in the literature (Kudritzki et al. 1999; Lefever± et al. 2007; radius and terminal∞ velocities should be derived independently, Fraser et al. 2010, given in Table A.1) or average values when as explained below. large scatters are present. We derived the macroturbulent veloc- Prior to the modelling process, the stellar radius, R?, was ity that reproduces the extra broadening seen in the line profile. 1 1 derived using the measured angular diameter, the distance to the We found uncertainties in Vmacro from 10 km s− to 20 km s− . star, or a fit to the observed spectral energy distribution (SED). Regarding the wind parameters, we searched for V mea- We model the SED with the interactive user interface BeSOS2 surements from UV observations (Prinja & Massa 2010∞; Prinja using the best-fitting atmospheric (TLUSTY or Kurucz) model et al. 1998; Howarth et al. 1997) and used them as initial values to (Hubeny & Lanz 1995; Kurucz 1979). BeSOS reads stellar pho- reproduce the Hα line. Good fits were obtained using slight vari- tometry data of any star from photometry-enabled catalogues ations of V . These values are listed in Table3 with deviations ∞ in VizieR and the parallax distance from the HIPPARCOS cat- of 10% with respect to the UV measurements. In a few cases, alogue. Sometimes, it was not possible to obtain a good fit to the UV terminal velocities were not able to fit the observed Hα the SED (see details below on the procedures applied to each line profile and our own V determinations are provided. In this star). In these cases, a slightly modified distance was required case, the discrepancies obtained∞ between the derived V and the to improve the fit. To find the best model, BeSOS uses the IDL UV data can increase up to 30%. Therefore, we consider∞ errors program mpfit that searches for the best fit using the Levenberg– for V between 10% and 30%. Marquardt (LM) method, which is a standard technique for Adopting∞ the values for V from the UV and errors of 10% ∞ solving non-linear least squares problems. for R?, we observe that late and early B supergiants may show BeSOS provides, for a given colour excess, a new set of stel- a change of up to 20% in the equivalent width (EW) of the lar parameters: Teff, log g, and R? (in units of the solar radius). emission component of a P Cygni line profile if M˙ varies by an As initial entries we use the Teff and log g values obtained in this amount of 10% and 20%, respectively. Larger errors in V might work from the Si and He ionization balances and from Hγ and give uncertainties on M˙ of 30%. However, when the Hα∞ line is Hδ line widths, respectively, and the code searches then for the seen in absorption, the error on M˙ can be larger than a factor of most optimal values to fit the SED. The colour excess, E(B V), 2 (Markova et al. 2004). was calculated using the observed B V and the interpolated− − (B V)0 values calculated from the Teff (B V) scale for super- 4. Results giants− (Flower 1996). In a few cases,− the E−(B V) had to be − modified to fit the 2200 Å bump properly. The stellar radius For each star in our programme we modelled the line profiles obtained using the SED is averaged with the values derived of H, He, and Si. Figure1 shows the observed H α line and the via the angular diameter and Mbol (the bolometric magnitude). best-fitting synthetic model. The rest of the lines and their cor- Table2 lists the intrinsic properties for each star: the spec- responding fits are shown in the Appendix (Figs. A.1–A.7). In tral type found in the literature, the visual apparent magnitude, Table3 we list the obtained stellar and wind parameters with the observed and intrinsic colour indexes, the calculated colour their errors: Teff, log g, β, M˙ , V , Vmicro, Vmacro, V sin i, R? ∞ excess, the stellar parameters (Teff and log g, obtained from fit- (stellar radius in solar radius units), log L/L (stellar luminos- tings to the SED), the distance to the star (d), the bolometric ity referenced to the solar value), log Dmom (the modified wind correction (BC), the calculated bolometric magnitude, and the momentum rate, see Sect.5), log L/M (both L and M in solar stellar radius. The computed parameters are listed with their units, where M is the stellar mass), and Q (optical-depth invari- corresponding errors. Details on the best-fitting model obtained ant, discussed in Sect.5). The parameters Vmicro and Vmacro are with BeSOS are given in Sect.4. related to the photospheric lines. Initially Vmicro was fixed at 1 1 Typical error bars derived using the BESOS code for T , 10 km s− and then varied by 5 km s− to achieve the best fit eff ± log g, and R? are of about 1%, 2%, and 5%, respectively. Nev- to the intensities of He and Si lines. ertheless, considering uncertainties in the distance estimates The photospheric lines of each star (at a given epoch) were between 5% and 30% (from parallax measurements) and differ- modelled with the same set of Vmicro and Vmacro values. However, ences less than 1500 K between the Teff values derived from to properly match the Hα and He I λ 6678 line widths and shapes BESOS and FASTWIND, the propagation of errors yields an we often needed a different turbulent velocity since these lines uncertainty in Mbol lower than 10% and a margin of error are affected by the velocity dispersion of the flow. between 5% and 20% for R?. Regarding the determination of the wind parameters, we In relation to the stellar parameters derived from FAST- were able to reproduce the Hα line of many BSGs. However, in a WIND, we can adopt the parameter errors estimated by Kraus few cases, even when the emission component of the P Cygni et al.(2015), who found uncertainties in Teff of 300–500 K from profile looked well fitted, it was impossible to reproduce the the Si ionization balance and 1000 K for the He lines. We adopt intensity of the absorption component. Similar problems were ∆Teff= 1000 K for those stars for which only the He lines were found when reproducing the Hα absorption profile of HD 38771, HD 75149, and HD 111973 because of the presence of an incipi- 2 The Be Stars Observation Survey (BeSOS; http://besos.ifa. ent emission in the line core. Contrary to the Hα line, synthetic uv.cl) is a database containing reduced spectra acquired using the line profiles of the photospheric lines match the observations echelle spectrograph PUCHEROS. very well.

A91, page 4 of 28 M. Haucke et al.: Wind properties of variable B supergiants

Table 2. Stellar parameters calculated in the present study or adopted from the literature.

a a HD ST mV B V (B V)0 E(B V) Teﬀ log g d BC Mbol R? number− − − K dex pc R

34 085 B8 Iae 0.13 0.03 0.07 0.04 11 760 120 2.00 0.10 264 23 0.83 7.9 0.2 72 10 38 771 B0.5 Ia 2.06 −0.18 −0.23 0.05 25 700 ± 260 2.70 ± 0.05 198± 8 −2.41 −7.0 ± 0.2 13 ± 1 41 117 B2 Ia 4.63− 0.28 −0.18 0.42 17 940 ± 180 2.20 ± 0.05 552 ± 85 −1.84 −7.3 ± 0.4 23 ± 3 42 087 B4 Ia 5.78 0.20 −0.14 0.34 15 000 ± 150 2.39 ± 0.05 2075 ±1 000s −1.46 −8.3 ± 1.1 55 ± 11 47 240 B1 Ib 6.15 0.15 −0.18 0.33 17 500 ± 180 2.40 ± 0.05 1598± 479 −1.78 −7.7 ± 0.7 30 ± 4 52 089 B1.5 II 1.50 0.21 −0.20 0.00 21 000 ± 210 3.00 ± 0.05 124± 2 −2.20 −6.2 ± 0.1 11 ± 1 52 382 B1 Ia 6.50− 0.19 −0.20 0.44 23 140 ± 230 2.47 ± 0.05 1301 ±430s −2.05 −7.5 ± 0.8 21 ± 2 53 138 B3 Ia 3.02 0.08 −0.21 0.10 18 000 ± 180 2.20 ± 0.05 847 ± 287 −1.66 −8.5 ± 0.8 46 ± 4 − − ± ± ± − − ± ± 58 350 B5 Ia 2.45 0.09 0.17 0.08 15 000 500 f 2.00 0.10 f 609 148 1.24 8.0 0.6 54 6 64 760 B0.5 Ib 4.24 −0.14 −0.21 0.07 22 370± 220 2.50± 0.05 507± 25 −2.21 −6.7 ± 0.2 12 ± 1 74 371 B6 Iab/b 5.23− 0.20 −0.15 0.35 13 800 ± 140 2.00 ± 0.05 1800 ±360s −1.02 −8.7 ± 0.5 73 ± 10 75 149 B3 Ia 5.46 0.27 −0.19 0.46 15 000 ± 150 2.12 ± 0.05 1642 ± 330s −1.39 −8.4 ± 0.7 61 ± 13 79 186 B5 Ia 5.00 0.22 −0.13 0.35 15 000 ± 150 2.12 ± 0.05 1449± 609 −1.36 −8.3 ± 0.9 61 ± 7 80 077 B2Ia+e 7.56 1.34 −0.16 1.50 18 000 ± 180 2.17 ± 0.05 3600 ±600s −1.62 −11.6 ± 0.4 195 ± 47 92 964 B2.5Ia 5.38 0.27 −0.16 0.48 18 000 ± 180 2.19 ± 0.05 1851± 994 −1.66 − 9.1 ± 1.2 70 ± 14 99 953 B1/2 Iab/b 6.5 0.38 −0.18 0.56 18 830 ± 190 2.30 ± 0.05 1075 ± 427 −1.78 −7.2 ± 0.9 25 ± 3 111 973 B2/3 Ia 5.94 0.24 −0.14 0.38 17 180 ± 170 2.18 ± 0.05 1660 ±350s −1.46 −7.8 ± 0.5 46 ± 9 115 842 B0.5 Ia/ab 6.03 0.30 −0.23 0.60 25 830 ± 260 2.75 ± 0.05 1538± 750 −2.46 −9.2 ± 1.2 35 ± 4 148 688 B1 Iaeqp 5.31 0.35 −0.20 0.54 20 650 ± 210 2.20 ± 0.05 833 ± 229 −2.00 −8.0 ± 0.7 31 ± 4 − ± ± ± − − ± ±

Notes. Values of Teff and log g were derived by fitting the SED, with the exception of those flagged with f . R? is an average of the values obtained from three different approaches (see details in Sect.3). d is the distance to the star in parsecs derived from HIPPARCOS, with the exception of those data flagged with s. (a) BC: bolometric correction taken from Flower(1996), ( f )parameters derived with the code FASTWIND (see details in Sect. 4.1), and (s)distances derived from other methods, see details in the text.

4.1. Comments on individual objects The Hα line is highly variable (with profiles in absorption, filled in by emission, double-peaked, or inversed P Cygni; In the following we summarize the previous and current data Chesneau et al. 2014). Our observation resembles the one of the stellar and wind parameters for each star. In addition, reported by Przybilla et al.(2006). Our estimate for the mass-loss we complement our data with images taken with the Wide-field rate is 1.5 times lower than the value obtained by Markova et al. Infrared Survey Explorer (WISE; Wright et al. 2010), covering (2008),∼ and at least one-third of the value reported by Chesneau the range 3.5 µ–22 µ (bands W1 and W4). These images are et al.(2014). We also derive a low terminal velocity. This value used to identify former phases of strong stellar winds. is quite uncertain because we were not able to fit the absorption HD 34085 is a B8 Iae star that has been studied by many component of the Hα P Cygni profile. Moreover, a high Vmacro 1 authors. Using FASTWIND we estimated the stellar fundamen- (85 km s− ) was needed to reproduce the emission component of tal parameters: Teff = 12 700 K and log g = 1.7. The best-fitting the line profile. The photospheric lines were modelled very well, model to the SED was obtained with the BeSOS interactive inter- with the exception of the He I λ 4471 line, which does not exhibit face using a Kurucz model with Teff = 11 760 K, log g = 2.0, the forbidden component. R? = 71 R , E(B V) = 0.044 (calculated as explained in HD 38771 (B0.5Ia): the stellar and wind parameters of − Sect.3), and d = 259 pc (which is close to the distance of 264 pc this star have been derived by many authors: Nerney(1980), given by HIPPARCOS). Using the HIPPARCOS distance and the Garmany et al.(1981), Lamers et al.(1982), Kudritzki et al. mentioned value of E(B V), we estimated Mbol = 7.95 mag (1999), Crowther et al.(2006), Searle et al.(2008), Zorec − − and R? = 70 R ; instead, from the angular diameter (2.713 mas, et al.(2009). Their values range from 26 000 K to 27 500 K

Zorec et al. 2009) we calculated R? = 76 R . We assumed a mean in Teff, from 2.9 dex to 3.07 dex in log g, from 13.0 R to 6 1 value for this star of R? = 72 R . 28 R in stellar radius, from 0.27 10− M yr− to 1.20 6 1 × 1 1 × Our estimate of Teff (12 700 K) agrees with previous deter- 10− M yr− in M˙ , and from 1350 km s− to 1870 km s− in minations (see Table A.1). This star shows a photometric vari- V (see Table A.1). The star presents a variable magnetic field ation of 2.075 days (Lefèvre et al. 2009). Moreover, Moravveji (Nerney∞ 1980) and exhibits spectral variations. The observed et al.(2012) found 19 significant pulsation modes from radial Hα line is quite variable, showing a pure absorption profile velocities with variability timescales ranging from 1.22 days to (Kudritzki et al. 1999), a double-peaked absorption profile with 74.74 days. It also presents a variable stellar wind with changes a central emission (Rusconi et al. 1980), and an absorption pro- in the mass-loss rate of at least 20% on a timescale of one file with a strong central emission (Crowther et al. 2006). Morel year (Chesneau et al. 2014). Using spectro-interferometric mon- et al.(2004) also described changes in morphology and profile itoring, these authors found time variations in the differential amplitude of 32.6%, which clearly suggest variations in the wind visibilities and phases. For some epochs, the temporal evolution conditions. These authors also found two photometric periods of of the signal suggests the rotation of circumstellar structures. 4.76 and 1.047 days and Prinja et al.(2004) reported additional However, at some periods, no phase signal was observed at spectroscopic periods of 1.9, 6.5, and 9.5 days. all. This result was interpreted in the context of second-order At the time of our observation the Hα line displayed an perturbations of an underlying spherical wind. asymmetric absorption profile with a weak emission in the core

A91, page 5 of 28 A&A 614, A91 (2018)

Fig. 1. Hα line and the best-ﬁtting synthetic model. For HD 53138, HD 58350, HD 75149, HD 80077, HD 99953, and HD 111973 more than one plot are shown due to multi-epoch observations.

(Fig.1). We also note that the core of the H β line might be filled found Hα variations, both in shape and intensity, and reported in by an incipient emission (Fig. A.1). Extra broadening effects light variability on a period of 2.869 days. Additional sets of are present in all the photospheric lines leading to large Vmacro spectroscopic periods were obtained from Hα equivalent width estimates (of the order of the projected rotation velocity). time series (Morel et al. 2004, Table 1). We derived Teff = 25 000 K using Si II and Si III, and also From Table A.1 we see that the fundamental stellar parame- using the He I and He II ionization balance. The best-fitting SED ters and luminosity are consistent with those already published, model was obtained with a TLUSTY model for Teff = 25 700 K, in contrast with the wind parameters (mainly M˙ ) that are quite log g = 2.70, R? = 14 R , and E(B V) = 0.05 mag for a dis- different (by a factor of up to 3.6). − tance of 191 pc (close to the HIPPARCOS distance of dH= 198 pc; The best-fitting SED model corresponds to Teff = 17 940 K, van Leeuwen 2007). Using dH in Pogson’s formula we derived a log g = 2.2, R? = 26 R , E(B V) = 0.42 mag (instead we − Mbol = –6.99 mag and R? 12 R . The adopted stellar radius, derived E(B V) = 0.46 mag from Flower 1996) and d = ∼ − R? 13 R , which is consistent with the value derived from the 477 pc (while the distance measured by HIPPARCOS is angular∼ diameter calculated by Zorec et al.(2009, 0.62 mas). 552 85 pc). On the other hand, using d = 447 pc and E(B V), ± − Our Teff value is 1000 K lower than the values obtained by the angular diameter (0.371 mas, Zorec et al. 2009), or the com- Searle et al.(2008) and Gathier et al.(1981), while the stellar puted bolometric magnitude ( 7.35 mag) the stellar radius is − radius is similar to that obtained by Kudritzki et al.(1999). Com- R? 23 R . We therefore adopted R? = 23 R , which is also ∼ pared with previous determinations, our mass-loss rate has the consistent with the stellar radius expected from the HIPPARCOS lowest value. As Hα is in pure absorption M˙ might have a large distance. The stellar radius is lower than the value derived by uncertainty. other authors (see Table A.1); however, if we use the distance Considering the diversity of the stellar and wind parameters derived by Megier et al.(2009, 1.62 kpc) we obtain R? = 62 R . found in the literature, and that the Hα line profile shows impor- From the line modelling we obtained higher Teff = 19 500 K tant variations, this star is a good candidate to search for pulsat- and log g = 2.3 values. Our spectrum reveals a Hα emission line ing or magnetic activity connected to cyclic wind variability. with a complex blue-ward absorption component (see Fig.1). HD 41117 (B2Ia): this is another deeply studied star (Zorec The mass-loss rate derived here is lower (by a factor of about et al. 2009; Crowther et al. 2006; Morel et al. 2004; Kudritzki 5) than the value reported by Kudritzki et al.(1999). We also et al. 1999; Scuderi et al. 1998; Nerney 1980). Morel et al.(2004) obtained a lower terminal velocity. The rest of the lines are in

A91, page 6 of 28 M. Haucke et al.: Wind properties of variable B supergiants

Table 3. Stellar and wind parameters derived from line ﬁtting procedures.

STAR Teﬀ log g β MV˙ Vmacro Vmicro V sin i R? log (L/L ) log Dmom log (L/M) log Q 6 1 ∞ 1 1 1 1 K dex 10− M yr− km s− km s− km s− km s− R dex dex dex dex

HD 34085 12 700 500 1.70 0.1 2.6 0.23 0.02 155 46 52 3 10 2 30 72 5.09 0.09 27.22 0.19 4.13 12.71 0.32 HD 38771 25 000 ± 1000 2.70 ± 0.1 1.5 0.14±0.04 1500±150 60±10 13 ± 2 80 13 4.78 ± 0.07 27.69 ± 0.19 4.27 −13.29 ± 0.24 HD 41117 19 000 ± 1000 2.30 ± 0.1 2.0 0.17±0.03 510±51 65±20 10 ± 5 40 23 4.84 ± 0.16 27.38 ± 0.15 4.27 −12.78 ± 0.23 HD 42087 16 500 ± 1000 2.45 ± 0.1 2.0 0.57±0.05 700±70 80±15 15 ± 5 80 55 5.31 ± 0.43 28.27 ± 0.13 3.83 −13.12 ± 0.23 HD 47240 19 000 ± 1000 2.40 ± 0.1 1.0 0.24±0.02 450±90 60± 10 10 ± 3 95 30 5.03 ± 0.30 27.57 ± 0.15 4.08 −12.82 ± 0.26 HD 52089 23 000 ± 1000 3.00 ± 0.1 1.0 0.02 ±0.006 900 ±270 65±5 8 ± 2 10 11 4.49 ± 0.05 26.58 ± 0.28 3.88 −13.69 ± 0.38 HD 52382 21 500 ± 1000 2.45 ± 0.1 2.2 0.24± 0.04 1000± 100 65±5 10 ± 2 55 21 4.94 ± 0.32 27.84 ± 0.14 4.29 −13.10 ± 0.21 HD 53138 18 000 ± 1000 2.25 ± 0.2 2.0 0.24±0.02 600 ±120 50±10 10 ± 4 40 46 5.31 ± 0.32 27.79 ± 0.14 4.19 −13.09 ± 0.22 18 000 ± 1000 2.25 ± 0.2 2.0 0.20±0.01 450±135 60±5 9 ± 3 40 46 5.31 ± 0.32 27.59 ± 0.14 4.19 −13.36 ± 0.22 HD 58350 15 000 ± 500 2.00 ± 0.1 3.0 0.15±0.02 233± 23 70±10 12 ± 2 40 54 5.13 ± 0.23 27.21 ± 0.13 4.12 −12.97 ± 0.20 16 000 ± 500 2.00 ± 0.1 3.0 0.12±0.01 175 ±18 50±10 10 ± 5 40 54 5.24 ± 0.23 26.98 ± 0.13 4.23 −12.90 ± 0.20 HD 64760 23 000 ± 1000 2.90 ± 0.1 0.5 0.42±0.06 1500 ± 150 100 ± 50 15 ± 5 230 12 4.57 ± 0.08 28.14 ± 0.12 3.73 −12.76 ± 0.18 HD 74371 13 700 ± 500 1.80 ± 0.1 2.0 0.28±0.03 200±60 60±10 10 ± 3 30 73 5.23 ± 0.19 27.48 ± 0.21 4.17 −12.80 ± 0.33 HD 75149 16 000 ± 1000 2.10 ± 0.1 2.5 0.09±0.03 400±40 62±12 9 ± 3 40 61 5.35 ± 0.29 27.25 ± 0.11 4.14 −13.63 ± 0.22 16 000 ± 1 000 2.10 ± 0.1 2.5 0.20±0.01 350±35 52±12 11 ± 1 40 61 5.35 ± 0.29 27.54 ± 0.11 4.14 −13.19 ± 0.22 16 000 ± 1000 2.10 ± 0.1 2.5 0.16±0.05 400±40 57±7 17 ± 1 40 61 5.35 ± 0.29 27.50 ± 0.11 4.14 −13.37 ± 0.22 16 000 ± 1000 2.10 ± 0.1 2.5 0.25±0.01 350±35 55±10 15 ± 1 40 61 5.35 ± 0.29 27.64 ± 0.11 4.14 −13.10 ± 0.22 HD 79186 15 800 ± 500 2.00 ± 0.1 3.3 0.40±0.02 400±40 53±7 11 ± 1 40 61 5.33 ± 0.38 27.90 ± 0.11 4.21 −12.98 ± 0.20 HD 80077 17 700 ± 1000 2.20 ± 0.1 3.2 5.4 ±0.50 200±20 ± ± 10 195 6.53 ± 0.17 28.86 ± 0.14 4.23 −12.15 ± 0.26 17 700 ± 1000 2.20 ± 0.1 3.0 5.4±0.50 150±15 60······5 10 5 10 195 6.53 ± 0.17 28.98 ± 0.14 4.23 −11.97 ± 0.26 HD 92964 18 000 ± 1000 2.20 ± 0.1 2.0 0.49± 0.03 370±55 40±5 11 ± 1 45 70 5.67 ± 0.49 27.98 ± 0.13 4.25 −12.93 ± 0.25 HD 99953 19 000 ± 1000 2.30 ± 0.1 2.0 0.08±0.01 250±50 50± 5 18 ± 2 50 25 4.87 ± 0.37 26.80 ± 0.14 4.22 −12.79 ± 0.24 19 000 ± 1000 2.30 ± 0.1 2.0 0.13 ± 0.01 500 ±100 50± 5 18 ± 2 50 25 4.87 ± 0.37 27.33 ± 0.14 4.22 −13.03 ± 0.24 19 000 ± 1000 2.30 ± 0.1 2.0 0.22 ± 0.01 700±140 50 ± 5 18 ± 2 50 25 4.87 ± 0.37 27.69 ± 0.14 4.22 −13.02 ± 0.24 HD 111973 16 500 ± 1000 2.10 ± 0.1 2.0 0.21 ± 0.01 350± 70 57±8 12 ± 2 35 46 5.16 ± 0.20 27.51 ± 0.15 4.19 −13.02 ± 0.27 16 500 ± 1000 2.10 ± 0.1 2.0 0.14 ±0.004 350 ±105 63± 3 12 ± 2 35 46 5.16 ± 0.20 27.33 ± 0.15 4.19 −13.19 ± 0.27 HD 115842 25 500 ± 1000 2.75 ± 0.1 2.5 1.80± 0.30 1700± 340 63±3 23 ± 7 50 35 5.67 ± 0.62 29.06 ± 0.18 4.30 −12.91 ± 0.28 HD 148688 21 000 ± 1000 2.45 ± 0.1 2.5 1.15±0.20 1200±360 65±5 11 ± 1 50 31 5.23 ± 0.28 28.69 ± 0.23 4.26 −12.80 ± 0.36 ± ± ± ± ± ± ± ± − ±

1 pure absorption and are excellently modelled (see Fig. A.1). Val- HD 47240 is a fast rotating B1Ib star (V sin i = 114 km s− , ues of Vmicro and Vmacro as high as the rotational velocity of Simón-Díaz & Herrero 2014) that lies behind the Mono- 1 the star (V sin i = 40 km s− ) are needed to reproduce the extra ceros Loop supernova remnant (SNR). It presents very broad broadening of the photospheric lines. A larger macroturbulent photospheric absorption lines, periodic light variations of 1 velocity (65 km s− ) was also obtained by Simón-Díaz & Herrero 2.742 days (Lefèvre et al. 2009; Morel et al. 2004), periodic (2014). motions due to binarity, and the presence of discrete absorption HD 42087 (B4Ia): the SED was fitted with a TLUSTY components (DACs; Prinja et al. 2002) that appear double. We model using Teff = 15 000 K, log g = 2.39, R? = 60 R , were able to get a good fit to the SED for the following param- E(B V) = 0.34 mag, and d = 2 021 pc (which is similar to the eters: T = 17 500 K, log g = 2.4, E(B V) = 0.33 mag, d = − eff − distance of 2 075 pc derived by Megier et al. 2009, from the 1515 pc, and R? = 35 R . The derived distance is consistent with Ca II H+K lines). On the other hand, a distance d = 2 075 pc the distance estimate of 1598 pc by Megier et al.(2009) and the gives Mbol = 8.33 mag and R? = 50 R , while the angular dis- − colour excess also matches the 2200 Å bump. We also calcu- tance calculated by Zorec et al.(2009) leads to R? = 57 R . We lated a Mbol = 7.66 mag and R? = 28 R (for d = 1598 pc). This adopted a mean value of R? = 55 R . − stellar radius agrees with the computed value from the angular Our optical spectrum only covers the Hα region; therefore, diameter (0.157 mas) by Zorec et al.(2009). We adopted a mean we were not able to determine the effective temperature using value of R? = 30 R . the silicon lines. Thus, we adopted the value calculated by Zorec The stellar parameters derived with FASTWIND agree with et al.(2009)( Teff = 16 500 K) and derived from the H lines the values given in one of the models calculated by (Lefever g = a log 2.45. This Teff is lower than the value reported by et al. 2007); who found two models with very different Teff Benaglia et al.(2007) and Searle et al.(2008), and greater than and log g and both fit very well the Si III line profiles. The Teff the value we derived from the fitting of the SED. obtained here is lower than the 21 670 K derived with the BCD This star presents a significant Hα variability of 91.2% and spectrophotometric method (Zorec et al. 2009). both the Hα and He I λ6678 lines show a cyclic behaviour on a To model the photospheric lines we needed a high Vmacro 1 periodicity of 25 days (Morel et al. 2004). This period is larger value (60 km s− ) that is much higher than those used to model ∼ 1 than the 6.807 day period found from the HIPPARCOS light curve the H lines formed in the wind (3 km s− ). As this object is (Morel et al. 2004; Lefèvre et al. 2009). a pulsating variable star of α Cyg type, a high Vmacro value is Our Hα line shows a P Cygni feature with a weak emission expected. and a broad absorption component with a superimposed incipi- Our spectrum shows Hα as a double-peaked emission line ent emission structure. We fitted the Hα line profile using a wind with an intense central absorption. This profile was considered 6 1 1 model with M˙ = 0.57 10− M yr− and V = 700 km s− . The ∞ as an indicator of the presence of a disk-like structure (Lefever obtained V is similar to the one measured in the UV spectral et al. 2007). We achieved a good fit to the observed Hα line ∞ 1 range (735 km s− , Howarth et al. 1997). for the intensity of the absorption and for the two emission

A91, page 7 of 28 A&A 614, A91 (2018)

7 1 components. We derived a mass-loss rate of 2.4 10− M yr− . This value agrees with the upper limit of the range× reported by 7 1 7 1 Lefever et al.(2007, 1.7 10− M yr− – 2.4 10− M yr− ) and is lower than the value× calculated by Morel× et al.( 2004, 7 1 0.31 10− M yr− ) from the theoretical WLR for OBA supergiants.× For the terminal velocity, we found almost half of the value obtained by Lefever et al.(2007), Prinja & Massa (2013). HD 52089 (B1.5 II) is the brightest EUV source in the sky. It shows a UV emission line spectrum and X-ray emission consistent with a wind-shock model. Using the observed SED and TLUSTY models we derived Teff= 21 000 K, log g = 3.0, and R? = 11 R for a distance of 124 pc (van Leeuwen 2007) and E(B V) = 0.00 mag. This stellar radius is also consis- − tent with the bolometric magnitude (Mbol= –6.18 mag) and the measurement of the angular diameter (0.801 mas, Zorec et al. 2009) that lead to 10 R and 10.6 R , respectively. From FAST- µ µ Fig. 2. WISE images (W2 = 4.6 m in blue, W3 = 12 m in green, and WIND we derived a higher Teff (23 000 K) that agrees with the W4 = 22 µm in red) showing a bow-shock structure near HD 52382. The values given in Lefever et al.(2010), Zorec et al.(2009), and white arrow indicates the direction of the proper motion of the star. Morel et al.(2008), but it is higher than the value of 20 100 K obtained by Fraser et al.(2010). Our surface gravity value (3 dex) agrees more closely with the Fraser et al.(2010) esti- 2010), which could be an indication of pulsation activity. The mate and it is 0.2 dex lower than those reported by Lefever He I λ 6678 line shows a little emission in the core. et al.(2010) and Morel et al.(2008). The star has a longitudinal HD 53138 (B3Ia) has been studied by several authors (see magnetic field of 149 G (Morel et al. 2008). Our spectrum dis- summary in Table A.1). The derived Teff ranges from 15 400 K plays all the H lines in absorption. Using optical data we derive to 18 500 K, and log g from 2.05 dex to 2.35 dex. The star 8 1 a mass-loss rate of 2 10− M yr− , which is similar to the exhibits irregular variations in the Hα line (Morel et al. 2004). value obtained phenomenologically× by Cohen et al.(1996, 3 Our spectra also show variations: in 2006 the star displayed a 8 1 8 1 × 10− M yr− –8 10− M yr− ). typical P Cygni profile, while in 2013 a double-peaked emission HD 52382 (B1Ia)× is an O-B2 runaway candidate (Peri was observed. We also note differences in the intensities and line et al. 2012). Using FASTWIND to fit the ratios of Si III/Si II widths of the Hγ and the Si III lines for the two epochs that lead g and He II/He I we obtain Teff = 21 500 K, and from the Hγ to slightly different values of Teff and log . We obtained for both and Hδ lines a log g = 2.45. From Flower(1996) we have epochs the following fundamental parameters: 18 000 K for Teff g . (B V)0 = –0.20 mag and derive a colour excess of 0.39 mag. and 2.25 dex for log with uncertainties of 1000 K and 0 2 dex. The− best fit to the observed SED corresponds to a TLUSTY The best-fitting SED model suggests Teff = 18 000 K, log g = . IPPARCOS model with Teff = 23 140 K, log g = 2.47, R? = 21.6 R , and d = 2 2, R? = 46 R , d = 822 pc (close to 847 pc the H estimate), and E (B V) = 0.10 (lower than the colour excess of 1.23 kpc, but with a larger E(B V) = 0.44 mag. This dis- − tance agrees with the 1.1 kpc found− by Megier et al.(2009) 0.131 mag we obtained from Flower 1996). The calculation of R? by means of the angular diameter gives 51 R , while the and with the parallactic distance of 1.3 kpc measured by Gaia computed Mbol (–8.53 mag) suggests R? = 46 R . We adopted (0.768 mas; Astraatmadja & Bailer-Jones 2016). Using the dis- R?= 46 R . tance of 1.3 kpc in Pogson’s formula and the expression for the angular radius (with 0.17 mas; Pasinetti Fracassini et al. 2001), The photospheric lines are broadened by macroturbulent 1 we obtain a stellar radius of 18.6 R and 23.5 R , respectively. motion (Vmacro = 60 km s− ). This high value of Vmacro, which is Therefore, we adopted R = 21 R as the mean value. We rejected almost twice the measured value for the projected rotation veloc- the HIPPARCOS distance of 471 pc because it leads to a value of ity, is consistent with the suggestion of a pulsating variable star α Mbol that is too low for a B supergiant. of Cyg type (Lefèvre et al. 2009). The derived Teff = 21 500 K is between the values calcu- Our model fits the observed He I λ 4471 Å line and the emis- lated by Lefever et al.(2010) and Krti ckaˇ & Kubát(2001). The sion component of Hα fairly well, although we were not able to star presents a variable Hα line profile displaying a P Cygni or a reproduce the absorption component of the profile in 2006 or pure emission feature (Morel et al. 2004). Our spectrum shows the emission component observed in 2013. Our mass-loss esti- a pure emission line. We were able to reproduce all the spectral mates are in the range of previously determined values of other 1 lines very well. We derived a terminal velocity of 1000 km s− , authors, being higher in 2006. The terminal velocity measured which is of the same order as that obtained from the UV region in the spectrum of 2013 is lower than in 2006 (see Table A.1). 1 1 (900 km s− , Prinja et al. 1990), but 200 km s− lower than the Barlow & Cohen(1977) measured a significant infrared excess value found by Krtickaˇ & Kubát(2001) for a similar mass-loss at 10 µm requiring a mass-loss rate a factor of 20 times higher rate. The combined WISE W2, W3, and W4 images (Fig.2, at than ours to account for it. 4.6 µm, 12 µm, and 22 µm, respectively) reveal a complex struc- HD 58350 (B5Ia) is an MK standard star. The set of stel- ture close to the star and in the direction of its proper motion lar parameters derived with FASTWIND (log g and Teff) agrees (indicated in Fig.2 with a white arrow) that could be related to very well with those derived by McErlean et al.(1999), Searle wind interactions with the ISM, due to high mass-loss episodes et al.(2008), Fraser et al.(2010). We were not been able to in the past. obtain a good fit to the SED using simultaneously the UV, visual, We want to stress that for an optimal fitting of the pho- and IR photometry. However, based on the distance given by tospheric lines we have to introduce high values of Vmacro, of HIPPARCOS (609 pc) and the calculated angular diameter 1 around 65 km s− (a similar result was reported by Lefever et al. (0.882 mas, Zorec et al. 2009), we obtained R? = 57 R . Based on

A91, page 8 of 28 M. Haucke et al.: Wind properties of variable B supergiants

HIPPARCOS parallax measurement and the derived E(B V) = 0.083 mag from Flower(1996), we obtained M = 7.97− mag bol − and R? = 51 R . We adopted R? = 54 R as the mean value. The Hα line shows a variable P Cygni profile with a tiny emission component. We were able to match the observations taken in 2006 and to get a fairly good fit to the one acquired in 2013. In the former, we failed to fit the absorption component. Our mass-loss rate is similar to that given by Lefever et al.(2007) and lower than the value given by Searle et al.(2008). As these last authors do not show the stellar spectrum in the Hα region, we cannot discuss the origin of the discrepancies. Morel et al. (2004) show time-series of the Hα line where the profile is seen in absorption, while the observations of Ebbets(1982) display a P Cygni profile. The light curve presents one period of 4.70 days (Koen & Eyer 2002) and another one of 6.631 days (Lefever et al. 2007). HD 64760 (B0.5Ib) is a rapid rotator. The star was exten- sively studied as part of the IUE “MEGA Campaign” by Massa et al.(1995) and Prinja et al.(1995). The line variability observed in the UV argues in favour of rotationally modulated wind variations. It displays a double-peaked emission line profile in Hα and very broad absorption lines of H and He. Direct obser- Fig. 3. WISE W1 and W4 images (3.4 µm and 22 µm, respectively) showing a double component or a wind-lobe structure. Yellow curves vations reveal a connection between multi-periodic non-radial represent regions with the same brightness. pulsations (NRPs) in the photosphere and spatially structured winds (Kaufer et al. 2006). These observations also seem to be compatible with the presence of co-rotating interaction regions. by Lefever et al.(2007), Fraser et al.(2010), and our SED The fitting of the SED provides the following parameters: model with Teff = 15 000 K, log g = 2.12, R? = 71 R ,

Teff = 22 370 K, log g = 2.50, R? = 15 R , E(B V) = 0.07 mag, E(B V) = 0.46 mag, and a distance of 1642 pc. We were not − − and d = 486 pc (while HIPPARCOS parallax gives d = 507 pc). able to fit the SED with the parallax measurement of 0.37 mas The derived Mbol = –6.72 mag gives R? = 12 R , while the angu- (2.7 kpc) given by HIPPARCOS, but our value of d is within lar diameter leads to R? = 10 R . We adopted R? = 12 R as the the corresponding error bars (1.4 kpc). The derived distance of mean value. 1.64 kpc leads to MV = 7.05, which is in agreement with the Our stellar parameters derived with FASTWIND value calculated by Humphreys− (1978, –7.0), M = 8.45 and bol − (Teff = 23 000 K, log g = 2.90) agree fairly well with those of R? = 56 R . The angular diameter also gives R? = 58 R . We

Lefever et al.(2007), but show large discrepancies with the adopted a mean value of R? = 61 R . The star displays small values reported by Searle et al.(2008). Although Searle et al. amplitude light variations with a period of 1.086 days (Koen & (2008) do not show the spectrum, they mentioned some difficul- Eyer 2002; Lefèvre et al. 2009), in addition to the variability of ties in deriving accurate values from the models due to the large 1.2151 days and 2.2143 days reported by Lefever et al.(2007). width of the spectral lines. From our spectra, we found that the star shows important In general, we obtained very good fits for all the lines, even variations in the Hα line: an absorption line profile with a weak for the double-peaked Hα profile. Our model gives the same set emission at the core is seen in 2006. A P Cygni feature is seen on of wind parameters as Lefever et al.(2007), but for the stel- 2013 February 5; it turns into an absorption profile two days later. lar radius a half of the value published by these authors was In 2014 a P Cygni profile with a complex absorption component obtained. The combined WISE W1 and W4 images (3.4 µm and is seen again (see Fig.1); this profile also presents two weak 22 µm, respectively) show the presence of either a double visual emission components that resemble the one published by Lefever component or a wind-lobe structure (see Fig.3). et al.(2007). Using the spectrum taken in 2006 the derived M˙ is HD 74371 (B6Iab/b) displays light variations with periods of similar to the value reported by these authors, although our value 5 20 days (van Genderen et al. 1989) and 8.291 days (Koen of V is a bit lower. To explain the line profile changes we have to − ∞ & Eyer 2002; Lefèvre et al. 2009). The stellar parameters Teff assume a variable wind structure with a mass-loss enhancement and log g were determined by Fraser et al.(2010) and agree of a factor of about 1.8 and 2.2, between our observation in 2013 very well with the values that we derived with FASTWIND and in 2014, and of a factor of about 2.8, between data taken in (Teff = 13 700 K, log g = 1.8). From the SED we were not able 2014 and 2006. We want to stress that in order to model the Hα to get a good fit using the HIPPARCOS parallax (0.43 mas, d = line width at various epochs, it was necessary to consider differ- 2.3 kpc van Leeuwen 2007). The best fit was achieved using a ent and also large values for Vmacro. The terminal velocity is very distance of 1.8 kpc, E(B V) = 0.35, T = 13 800 K, log g = similar in all the models. Nevertheless, we have to keep in mind − eff 2.0, and R? = 73 R . The found distance of 1.8 kpc agrees with that the values derived from a pure absorption profile are more the value reported by Humphreys(1978, 1.9 kpc). uncertain since the line is less sensitive to the wind conditions. Our model fits all the lines very well, with the exception of We were able to model all the photospheric lines quite well the absorption component of the Hα P Cygni profile and the core with the exception of the Hβ line core observed in 2006, which of the Hβ line, which seems to be filled in by an incipient emis- is weaker. sion (Fig. A.4). This is the first time that the wind parameters of HD 79186 (B5Ia): From the SED we were able to match HD 74371 have been determined (see Table3). a TLUSTY model using the following parameters: Teff = HD 75149 (B3Ia): the stellar parameters obtained in this 15 000 K, log g = 2.12, R? = 67 R , and d = 1.42 kpc, adopt- work with FASTWIND agree very well with those given ing E(B V) = 0.35 mag. The distance is very similar to that − A91, page 9 of 28 A&A 614, A91 (2018) obtained from the HIPPARCOS parallax (1.45 kpc), which yields 62 R , and from the angular diameter (0.37 mas Pasinetti

Mbol = 8.26 mag and R? = 53 R , while the angular diameter Fracassini et al. 2001) the derived stellar radius is 73 R . As a − (0.4 mas) gives 61 R . We used a mean stellar radius of 61 R . mean value we used R? = 70 R .

Our estimation of Teff is slightly higher than the values From our modelling of the photospheric lines, done with given by Fraser et al.(2010) and Prinja & Massa(2010), respec- FASTWIND, we obtained Teff = 18 000 K and log g = 2.2. The tively ∆ T 700 K and ∆ T 800 K, and even higher than those derived Teff agrees with the values adopted by Lefever et al. obtained by∼ Krtickaˇ & Kubát∼(2001) and Underhill(1984). (2007) and Krtickaˇ & Kubát(2001), but present large departures The logarithm of the surface gravity agrees with the unique from the atmospheric parameters derived by Fraser et al.(2010). value (2.0 dex) available in the literature (Fraser et al. 2010). To model the spectrum we used a projected rotational velocity 1 The stellar radius of our model is similar to that estimated by (40 km s− ) that is slightly higher than that given in the literature 1 1 Underhill(1984). (28 km s− and 31 km s− ). As did Lefever et al.(2007), we found Although we obtained a very good fit for all the photospheric a very high value for Vmacro. lines, we were not able to reproduce the shallow absorption Previous observations reveal important intensity variations component of the Hα P Cygni profile. Nevertheless, the termi- in both absorption and emission components of the P Cygni nal velocity agrees with the model parameters given by Krtickaˇ profile of the Hα line (see Fig. A.1 given in Lefever et al. & Kubát(2001) and with measurements from UV observations 2007). Comparing our spectra with that observation, the Hα line 1 (435 km s− , Prinja & Massa 2010). Our M˙ is slightly lower than observed in 2013 shows a wider absorption and a higher emission the value given by Krtickaˇ & Kubát(2001). component. We were able to reproduce only the emission com- HD 80077 (B2Ia+e) might be a member of the open cluster ponent of the Hα line profile. Nevertheless, we derived a value 1 Pismis 11, located at a distance of 3.6 kpc. With an absolute for V = 370 km s− that is close to that measured using the ∞ 1 bolometric magnitude of –10.4, this star is among the brightest of UV lines (435 km s− , Prinja et al. 1990) and lower than the val- the known B-type supergiants in our Galaxy (Knoechel & Moffat ues quoted in Table A.1. However, the mass-loss rate obtained in 1982; Marco & Negueruela 2009). this work is consistent with the value measured by Lefever et al. Carpay et al.(1989, 1991) detected light variations with an (2007). amplitude 0.2 mag and suggested that the star could be a lumi- HD 99953 (B1/2 Iab/b) is a poorly studied star. Fraser et al. ∼ nous blue variable (LBV). Using HIPPARCOS and V photometric (2010) derived the following stellar parameters: Teff = 16 800 K data, van Leeuwen et al.(1998) obtained a periodogram that and log g = 2.15. Based on our modelling of the Hγ, Si III, and reveals that the most significant peaks are those near 66.5 days He I lines, we obtained a new set of values: Teff = 19 000 K and 55.5 days, and peaks with much lower significance near and log g = 2.30. The latter are compatible with the assigned 76.0 days and 41.4 days. spectral classification and the observed SED (Teff = 18 830 K, Light variations of 0.151 mag over a period of 3.115 days log g = 2.30, E(B V) = 0.56, R? = 28 R ) for d = 1077 pc − were found by Lefèvre et al.(2009) and another period of (which is similar to the HIPPARCOS distance, 1075 pc). This dis- 21.2 days was determined from the polarimetry (Knoechel & tance and E(B V) predict Mbol = 7.16 mag and R? = 22 R . − − Moffat 1982). From the SED we obtained Teff = 18 000 K, We adopted R? = 25 R as the mean value. log g = 2.17, E(B V) = 1.5 mag, R? = 200 R , and d = The Hα line shows a P Cygni profile and its intensity varies − 3 600 pc (while the HIPPARCOS distance is 877 pc). Using the with time. The mass-loss and terminal velocity change by a fac- distance to the Pismis 11 cluster and the obtained E(B V), we tor of 2.8. Different values of V were needed to model the − ∼ macro calculated Mbol = 11.49 mag and R? = 187 R . We adopted a line at different epochs. stellar radius of 195− R . This work reports the first determinations of the mass-loss 6 1 6 1 Carpay et al.(1989) derived Teff = 17 700 K, log g = 2, rate of the star (M˙ = 0.08 10− M yr− , 0.13 10− M yr− , 6 1 6 1 and a mass-loss rate of 5.11 10− M yr− . Based on radio and 0.22 10− M yr− ). The terminal velocity derived from × 1 1 data, Benaglia et al.(2007) estimated a significantly lower mass- the Hα line ranges from 250 km s− to 700 km s− . This wide 6 1 loss rate (1.7 10− M yr− ) that agrees with our result (see range includes the value measured from the UV spectral region × 1 Table A.1). We observed a P Cygni Hα profile in the spectra (510 km s− , Prinja et al. 1990). The combined WISE W3 and taken in 2006 and 2014. Small changes in the emission and W4 images show warm dust heated by radiation coming from absorption components can be seen in Fig.1. the star (see Fig.4). Finally, we want to stress that HD 80077 is a very massive HD 111973 (B2/3Ia): this spectroscopic binary was classified post-main sequence star and, hence, an enhancement of the He as a B3 I star (Chini et al. 2012). It presents light variations with abundance is expected. However, we do not observe any notice- periods of 57.11 days and 9.536 days (Koen & Eyer 2002). able contribution of forbidden components and the He I λ 4471 The SED was fitted with a TLUSTY model with Teff = line was well-matched with a solar He abundance model. 17 180 K, log g = 2.18, R? = 48.6 R , E(B V) = 0.38 mag, and HD 92964 (B2.5Ia) displays photometric variations with d = 1660 pc. The colour excess was obtained− with the calibra- periods of 2.119 days and 14.706 days (Lefèvre et al. 2009). The tion scale of Flower(1996). Using this distance and the derived Hβ and Hγ line profiles show asymmetries. Lefever et al.(2007) colour excess, we obtained Mbol = 7.81 mag and R? = 40 R . ascribed this behaviour to the strong stellar wind which affects On the other hand, when using the− angular distance (0.26 mas) the photospheric lines. The same authors also noted that the He I we have R? = 46 R , a value that is close to that derived by

λ 6678 line requires a Vmacro that is two times higher than the Pasinetti Fracassini et al.(2001, 47 R ). We adopted R? = 46 R . value derived for the Si lines. We were able to match all the photospheric lines. The derived The stellar parameters derived from the SED are Teff and log g (16 500 K and 2.1 dex) are similar to those reported Teff = 18 000 K, log g = 2.19, E(B V) = 0.481, R? = 76 R , and by Fraser et al.(2010, 16 000 K and 2.3 dex) and Prinja & Massa − d = 1806 pc (close to the HIPPARCOS distance, d = 1851 pc). We (2010, 16 500 K). used the extinction curve from van Breda et al.(1974), who found The Hα line shows short-term variations: we observed a R = 3.5 for this star. From the HIPPARCOS distance and the cal- P Cygni profile with a weak emission component that turned culated A = 168˙ mag we obtained M = 9.14 mag and R = into an absorption profile the following night (see Fig.1). This V bol − ? A91, page 10 of 28 M.HauckeHaucke Haucke et et et al.: al.: al.: Wind Wind Wind Properties Properties properties of of of Variable Variable variable B B B Supergiants Supergiants supergiants

Fig.Fig.Fig. 4. 4. 4.WISEWISEWISE images images images showing showing showing wind-blown wind-blown wind-blown and and and warm warm warm dust dust dust structures structures structures Fig.Fig. 5. 5.Bow-shockBow-shock structure structure and and asymmetric asymmetric density density structure structure associ- associ- aroundaroundaround HD HD HD 99953. 99953. 99953. atedatedated with with HD HD 115842. 115842. The The The yellow yellow yellow curves curves curves represent represent represent regions regions regions with with with the the the samesamesame brightness, brightness, brightness, each each one one one associated associated associated with with with a a a different di difffferenterent lobe. lobe. variation is similar to the expected dynamical timescale of a distancedistance and andEE((BB VV)) predict predictMMbolbol== 77..1616 and andRR??==2222 R R .. We We prototypical BSG−− wind ( 1.3 days). We−− modelled the Hα line interaction phase. In addition, an asymmetric density structure adoptedadoptedRR?? ==2525 R R asas the the∼ mean mean value. value. and derived a lower terminal velocity than the value measured seemsThisThis to workbe work present reports reports in the the WISE mass-loss mass-loss W1 rateimage rate of of with the the star staran angular for for the the sizefirst first 1 66 11 66 11 fromTheThe the H HUVααlineline lines shows shows (520 a km a P P Cygni Cygnis− Prinja profile profile &Massa and and its its 2010 intensity intensity). Further- varies varies thattime.time. is It It greater ranges ranges than from from the 0 0..1414 WISE 10 10−− resolution.MM yryr−− totoMM˙˙==00..2121 10 10−− MM yryr−− .. 1 more,withwith time. time. a high The TheVmacro mass-loss mass-loss= 160–190 and and terminal terminalkm s− was velocity velocity needed change change to model by by a a this fac- fac- HD 148688 (B1Iaeqp) is a southern oxygen-rich supergiant tortor of of 2.8.2.8. Di Difffferenterent values values of ofVVmacromacro werewere needed needed to to model model the the that displays line variations// (Jaschek & Brandi 1973). Data from line. ∼∼ HDHD 115842 115842(B0.5Ia(B0.5Iaab)ab) presents presents photometric photometric variations variations on on a a linelineThis at at di diffff workerenterent epochs. epochs. reports the mass-loss rate of the star thetimescaletimescale HIPPARCOS of of 13.38 13.38satellite days days correlate (Koen (Koen & & with Eyer Eyer periods 2002) 2002) of and and 1.845 variations variations days and in in 6 1 for theThisThis first work work time. reports reports It ranges the the first first from determinations determinations0.14 10− of ofM the theyr mass-loss mass-loss− to M˙ 6.329thethe H Hα daysα lineline (Lefever profile. profile. et Our Our al. spectrum2007 spectrum). displays displays the the H Hαα lineline in in pure pure 6 1 66 11 × 66 11 =raterate0.21 of of the the10 star− starM ( (MM˙˙yr==−00..0808 10 10−− MM yryr−− ,, 0 0..1313 10 10−− MM yryr−− ,, and and emissionemissionThe best-fitting (see (see Fig. Fig. 1), 1), model while while to the the the one one observed given given by by CrowtherSED Crowther leads et et al. toal. × 66 11 00..2222HD 10 10 115842−− MM yryr(B0.5Ia/ab)−− ).). The The terminal terminal presents velocity velocity photometric derived derived variations from from the the on H H aαα T(2006)(2006)eff = 20 displayed displayed 650 K, log a a P Pg Cygni= Cygni 2.2, R feature. feature.? = 34 TheR The, and H Hββ linedline= 838 is is seen seen pc (while in in ab- ab- 11 11 timescalelineline ranges ranges of 13.38 from from days250 250 km ( kmKoen s s &toto Eyer 700 700 2002 km km) s s and.. Thisvariations This wide wide in range range the thesorptionsorption HIPPARCOS and and the thedistance profile profile is is is slightly slightly833 pc). asymmetric asymmetric This gives M (see (seebol = Fig. Fig.7 A.7).. A.7).97 mag −− −− − Hincludesincludesα line profile. the the value value Our measured measured spectrum fromdisplays from the the the UV UV H spectralα spectralline in region regionpure emis- (510 (510 and R? = 26 R , for E(B V) = 0.54 mag. This E(B V) was== 11 FromFrom the the SED SED we we derived derived− the the following following parameters: parameters:− TTeeffff sionkmkms (sees−− ,, Prinja Fig.Prinja1), et et while al. al. 1990). 1990). the one The The given combined combined by Crowther WISE WISE et W3 W3 al.( and2006 and W4 W4) obtained from Flower(1996) and also fits the deep bump at 2525 830 830 K, K, log log gg==2.73,2.73,RR?? ==3838 R R ,,EE((BB VV))==0.60.6 mag, mag, and anddd displayed a P Cygni feature. The Hβ line is seen in absorption 2200 Å present in the SED. We adopted R−−= 31 R as the mean imagesimages show show warm warm dust dust heated heated by by radiation radiation coming coming from from the the ==11 543 543 pc. pc. This This colour colour excess excess fits fits the the observed observed? 2 2 200 200 Å Å bump bump and the profile is slightly asymmetric (see Fig. A.7). value. starstar (see (see Fig. Fig. 4). 4). featurefeature and and is is higher higher than than the the EE((BB VV)) == 0.530.53 mag mag derived derived From the SED we derived the following parameters: The stellar parameters obtained from FASTWIND (T = fromfrom Flower Flower (1996). (1996). The The obtained obtained distance distance−− is is consistent consistent with witheff the the T = 25 830 K, log /g/ = 2.73, R = 38 R , E(B V) = 0.6 mag, 21 000 K and log g = 2.45) are in better agreement with the eff HDHD 111973 111973(B2(B23Ia):3Ia): This This? spectroscopic spectroscopic binary binary was was classi- classi- HIPPARCOSHIPPARCOS and and Gaia Gaia parallaxes parallaxes (0.65 (0.65 mas mas and and 0.6105 0.6105 mas, mas, − works of Fraser et al.(2010) and Lefever et al.(2007). The H α andfiedfiedd as as= a a 1543 B3 B3 I I pc. star star This (Chini (Chini colour et et al. al. excess 2012). 2012). fits It It presents presents the observed light light variations variations 2200 Å respectively)respectively) and and the the distance distance of of 1 1 583 583 pc pc estimated estimated from from the the line from Fig.1 shows a P Cygni feature that looks like the one bumpwithwith periods periodsfeature andof of 57.11 57.11 is higher days days than and and the 9.536 9.536E(B days daysV (Koen) (Koen= 0.53 & & mag Eyer Eyer derived 2002). 2002). CaCa II II H H and and K K lines lines (Megier (Megier et et al. al. 2009). 2009). Using Using the the HIPPAR- HIPPAR- − shown by Lefever et al.(2007), while the spectrum of Crowther fromTheThe Flower SED SED(1996 was was). fittedfitted The obtained with with a a TLUSTY TLUSTY distance is model model consistent with with T withTeeffff == COSCOS distance distance (1 (1 538 538 pc) pc) and andEE((BB VV))==0.60.6 mag, mag, we we calculated calculated the HIPPARCOSggand== Gaia parallaxes== (0.65 mas and= 0.6105= mas, et al.(2006) exhibits an emission line.−− 1717 180 180 K, K, log log 2.18,2.18, RR?? 48.648.6 R R ,, EE((BB VV)) 0.380.38 mag, mag, MMbolbol == 99..2222 mag mag and andRR?? ==3232 R R ,, while while from from the the angular angular di- di- respectively) and the distance of 1583 pc estimated−− from the We were−− only able to model the emission component of the andanddd==11 660 660 pc. pc. The The colour colour excess excess was was obtained obtained with with the the cal- cal- ameterameter (0.22 (0.22 mas, mas, Pasinetti Pasinetti Fracassini Fracassini et et al. al. 2001) 2001) we we derived derivedRR?? Ca II H and K lines (Megier et al. 2009). Using the HIPPARCOS P Cygni profile. We obtained a mass-loss rate that is a factor of ibrationibration scale scale of of Flower Flower (1996). (1996). Using Using this this distance distance and and the the de- de- ==3333 R R .. We We adopted adoptedRR??==3535 R R asas the the mean mean value. value. distancerivedrived colour colour (1 538 excess, excess, pc) weand we obtained obtainedE(B VMM) bolbol=== 0.677. mag,.8181 mag mag we and and calculatedRR?? ==4040 1.1 and 1.4 lower than the values found by Lefever et al.(2007) − −− WeWe achieved achieved very very good good fits fits to to the the photospheric photospheric lines lines for forTT ffff MRRbol.. On On= the the9.22 other othermag hand, hand, and whenR when? = using32 usingR the, the while angular angular from distance distance the angular (0.26 (0.26 and Crowther et al.(2006), respectively. The terminal velocity isee − == 2525 500 500 K K and and log log gg == 2.75.2.75. The The stellar stellar parameters parameters derived derived diametermas)mas) we we (0.22 have haveR mas,R?? == Pasinetti4646 R R ,, a a Fracassini value value that that et is is al. close close 2001 to to) thatwe that derived derived derived similar to the value found by Lefever et al.(2007, see Table A.1), inin this this work work agree agree with with those those found found by by Crowther Crowther et et al. al. (2006) (2006) Rbyby? = Pasinetti Pasinetti 33 R . WeFracassini Fracassini adopted et etR al. al.? = (2001, (2001, 35 R 47 47as R Rthe).). mean We We adopted adopted value. RR?? ==4646 although it is greater than measurements derived from the UV 1 linesandand FraserFraser (725 km et et s al. al.− , (2010) (2010)Prinjaet (see (see al. Table1990 Table). A.1). A.1). We We obtained obtained a a lower lower RR .We. achieved very good fits to the photospheric lines for 66 11 g mass-lossmass-loss rate rate (1.8 (1.8 1010−− MM yryr−− )) and and a a higher higher terminal terminal ve- ve- Teff =WeWe 25 were were500 K able able and to tolog match match= 2.75. all all the the The photospheric photospheric stellar parameters lines. lines. derivedThe The de- de- 1×1× in this work agree with those found by Crowther et al.(2006) and locitylocity (1 (1 700 700 km km s s−− )) than than the the Crowther Crowther et et al. al. (2006) (2006) estimates estimates rivedrivedTT ffff andand log log gg(16(16 500 500 K K and and 2.1 2.1 dex) dex) are are similar similar to to those those 4.2. Global66 properties11 of stellar1 and1 wind parameters ee (2.0(2.0 1010−− MM yryr−− ,, 1 1 180 180 km km s s−− ).). Fraser et al.(2010; see Table A.1). We obtained a lower mass- reportedreported by by Fraser Fraser et et6 al. al. (2010, (2010,1 16 16 000 000 K K and and 2.3 2.3 dex) dex) and and Prinja Prinja ×× loss rate (1.8 10− M yr− ) and a higher terminal velocity AlthoughOurOur value value the stars of of VV presentisis also also spectral higher higher variability, than than the the valueit value is still determined determined possible && Massa Massa (2010, (2010,1× 16 16 500 500 K). K). ∞∞ 11 (1700 km s− ) than the Crowther et al.(2006) estimates tofromfrom describe the the IUE IUE their data data global (1 (1 125 125 properties. km km s s−− ,, Evans Figure Evans6 et et shows al. al. 2004). 2004). a comparison These These au- au- TheThe H H6 αα lineline1 shows shows short-term short-term1 variations: variations: we we observed observed a a 11 (2.0 10− M yr− , 1180 km s− ). ofthorsthorsTeff reported reportedof BSGs a a with very very the high high spectral macroturbulence, macroturbulence, subtype. TheVVmacro samplemacro==225225 is a km collec-km s s−− ,, PP Cygni Cygni profile profile with with a a weak weak emission emission component component that that turned turned into into Our× value of V is also higher than the value deter- tionwhichwhich ofis data is twice twice from our our the estimate. estimate. literature and this work. With the exception anan absorption absorption profile profile∞ the the following following night night (see (see Fig. Fig. 1). 1). This This vari- vari- mined from the IUE data (1125 km s 1, Evans et al. of those stars that have T based on the BCD spectrophoto- ationation is is similar similar to to the the expected expected dynamical dynamical timescale timescale− of of a a proto- proto- WISEWISE W3 W3 and and W4 W4 images imageseff show show a a large large well-defined well-defined bow- bow- 2004). These authors reported a very high macroturbulence, metric calibration (Zorec et al. 2009), the effective temperature typicaltypical BSG BSG wind wind ( ( 11..33 days). days). We We modelled modelled the the H Hαα lineline and and shockshock structure structure (see (see Fig. Fig. 5) 5) related related to to a a possible possible strong strong wind-ISM wind-ISM V = 225 km s 1, which is twice our estimate. was obtained via adjustments of the line profiles of He I, He II, derivedderivedmacro a a lower lower terminal terminal− ∼∼ velocity velocity than than the the value value measured measured from from interactioninteraction phase. phase. In In addition, addition, an an asymmetric asymmetric density density structure structure WISE W3 and W4 images11 show a large well-defined bow- and line intensity ratio of Si IV/Si III or Si III/Si II, as in this thethe UV UV lines lines (520 (520 km km s s−− PrinjaPrinja & & Massa Massa 2010). 2010). Furthermore, Furthermore, a a seemsseems to to be be present present in in the the WISE WISE W1 W1 image image with with an an angular angular size size shock structure (see Fig.5) related11 to a possible strong wind-ISM work. highhighVVmacromacro ==160-190160-190 km km//ss−− waswas needed needed to to model model this this line. line. thatthat is is greater greater than than the the WISE WISE resolution. resolution. A91, page 11 of 28 ArticleArticle number, number, page page 11 11 of of 28 28 A&A proofs: manuscript no. 31678corr_last A&A proofs: manuscript no. 31678corr_last A&A 614, A91 (2018) HD 148688 (B1Iaeqp) is a southern oxygen-rich supergiant 30000 Searle et al. (2008) that displaysHD 148688 line(B1Iaeqp) variations (Jaschek is a southern & Brandi oxygen-rich 1973). Data supergiant from Crowther et al. (2006) We performed a fit to the observed Teff versus spec- 30000 Searle et al. BCD(2008) thethat HIPPARCOS displays line satellite variations correlate (Jaschek with & Brandi periods 1973). of 1.845 Data days from CrowtherOur et Sampleal. (2006) tral subtype relation using a third-order polynomial regression Markova Puls (2008)BCD the HIPPARCOS satellite correlate with periods of 1.845 days Our Sample and(a = 6.32954. days6 7 (Lefever.7, b = 973 et al..4 2007).107.5, c = 6054.3 387.9, and Markova Puls (2008) and 6.329− days± (Lefever et al.± 2007). − ± 25000 d =The27 316best-fitting331). model The derived to the coefficients observed SED agree leads very to wellTeff with= ± 25000 20the 650The fit K, obtained best-fitting log g = by2.2, modelMarkovaR? = to34the & R Puls observed, and(2008d =). SED838 We leadspc illustrate (while to Te ourtheff =

HIPPARCOSfit20 650in Fig. K,6 log with distanceg = a2.2, grey isR 833? wide= pc).34 band R This, and that givesd has=M838bol a dispersion= pc (while7.97 and ofthe (K) − 20000 RHIPPARCOS? = 26 R , for distanceE(B V is) 833= 0.54 pc). mag. This This givesEM(Bbol =V) was7.97 ob- and eff 1290 K. The supergiants of our sample, with the exception (K) − − − T 20000 R? = 26 R , for E(B V) = 0.54 mag. This E(B V) was ob- eff tainedof HD from 52089, Flower HD (1996) 53138, and and also HD fits 47240 the deep (indicated bump at in 2 the 200 fig- Å T − − presenturetained with fromin the large Flower SED. open We (1996) circles),adopted and alsoR fall? = fits inside31 the R deepas the the bump traced mean at relation.value. 2 200 Å

Moreover,present in the we SED. can Weobserve adopted thatR? the= 31 departures R as the of mean the value. effec- 15000 The stellar parameters obtained from FASTWIND (Teff= 15000 21tive 000The temperatures K stellar and log parametersg for= the2.45) early obtained are BSGs in better from range agreement FASTWIND roughly fromwith (T the0eff K= worksto21 5000000 of K Fraser K and and log et tend al.g (2010)= to2.45) be lower and are Lefeverin the better later et al.agreement the (2007). B spectral The with Hsub- theα type.works This of Fraser same et tendency al. (2010) is and also Lefever observed et al. in (2007).Markova The et H al.α 10000 line from Fig. 1 shows a P Cygni feature that looks like the one −2 0 2 4 6 8 10 12 shown(line2008 from). by Lefever Fig. 1 shows et al. (2007), a P Cygni while feature the spectrum that looks of like Crowther the one 10000 −2 0 2 Spectral 4 Type 6 8 10 12 etshown al.Figure (2006) by7 Lefever exhibitsdisplays et an al. a emission linear (2007), correlation while line. the between spectrumlog of CrowtherTeff and log g (the surface gravity corrected by the centrifugal accel- Spectral Type et al.Wecor (2006) were only exhibits able an to emission model the line. emission component of the eration, see Table A.1). This kind of correlation was previously Fig. 6. Effective temperature as function of the spectral subtype. The re- P CygniWe profile.were only We able obtained to model a mass-loss the emission rate that component is a factor of ofthe Fig. 6. Effective temperature as function of the spectral subtype. The reported by Searle et al.(2008). We performed a linear fit and lationFig. 6. wasEff fittedective using temperature a third-order as function polynomial of the spectralregression subtype. (dotted The line) re- 1.1P Cygni and 1.4 profile. lower We than obtained the values a mass-loss found by rateLefever that et is al. a factor (2007) of relation was fitted using a third-order polynomial regression (dotted obtained a slope of 3.98 0.31. This indicates that most of withlation a dispersion was fitted ofusing 1 290 a third-order K (grey band, polynomial see text regressionfor details). (dotted Compari- line) and1.1 Crowther and 1.4 lower et al. than (2006), the values respectively. found by The Lefever terminal et al. velocity (2007) sonwithline) of aour with dispersion sample a dispersion of of Galactic 1 290 of 1290 K B (grey supergiants K (grey band, band, see with text see a setfor text of details). for BSGs details). collected Compari- Com- the stars in our sample have± a similar luminosity-to-mass ratio isand similar Crowther to the et value al. (2006), found by respectively. Lefever et The al. (2007, terminal see velocity Table fromsonparisonof the our literature. of sample our sample Dataof Galactic points of Galactic B of supergiants our sampleB supergiants with with a large set with of open BSGs a set circles collected of BSGs are (log (L/M) 4), as expected, since L T 4 . collected from the literature. Data points of our sample with large open A.1),is similar although to∼ the it is value greater found than by measurements Lefever∝ eteff al. derived (2007,from see Table the outliers.from the literature. Data points of our sample with large open circles are A large scatter1 is observed around the stars that have outliers.circles are outliers. UVA.1), lines although (725 km it iss− greater, Prinja than et al. measurements 1990). derived from the large log gcor ( 2.5);1 in particular HD 42087, HD 52089, and UV lines (725 km s− , Prinja et al. 1990). ≥ 4 HD 64760 do not follow the relation. Two of these stars, BSG sample 4 4.2.HD Global 42087 and properties HD 64760, of stellar are close and to wind the parameters TAMS so we believe BSG sample they4.2. Globalhave just properties left the main of stellar sequence. and HDwind 64760 parameters is a high rotator 3.5 Althoughand, as a the consequence, stars present we spectral should observe variability, higher it is apparent still possible bolo- 3.5 metricAlthough luminosities the stars present and lower spectral effective variability, temperatures it is still than possible those HD 52089 to describe their global properties. Figure 6 shows a comparison 3 HD 64760 HD 52089 ofexpectedtoT describeeff of BSGs for their their with global non-rotating the spectral properties. stellar subtype. Figure counterparts The 6 shows sample a(Frémat iscomparison a collec- et al. 3 HD 64760 tion2005of T ofeff; dataZorecof BSGs from et al. with the 2005 literature the), spectral giving and asubtype. this later work. evolutionary The With sample the stage exceptionis a collec- of the object.tion of data from the literatureT and this work. With the exception 2.5 HD 42087 of those stars that have eff based on the BCD spectrophotomet- Log(g) T 2.5 HD 42087 ricof calibration thoseFrom stars Table (Zorec that3 we have can et al. distinguisheff 2009),based the on starsthe effective BCD showing temperaturespectrophotomet- two different was Log(g) obtainedwindric calibration regimes. via adjustments One(Zorec group et al. of gathers 2009), the line stars the profiles e withffective steep of temperature He velocityi, He ii, gradi- and was 2 lineents,obtained intensityβ via1, ratio inadjustments good of Si agreementiv/Si ofiii theor linewith Si iii profiles/ theSi ii, standard as ofin this He windi work., He theory,ii, and 2 ≤ whileline intensity the other ratio group of Si showsiv/Si iii valuesor Si iii of/Siβiibetween, as in this 1.5 work. and 3.3, 1.5 We performed a fit to the observed Teff versus spectral sub- similar to the results reported by Crowther et al.(2006), Searle 1.5 typeWe relation performed using a a fit third-order to the observed polynomialTeff versus regression spectral (a sub-= et al.(2008), and Markova et al.(2008). Stars with β between 2 type54.6 relation7.7, b using= 973 a.4 third-order107.5, c polynomial= 6054. regression3 387.9, (anda = 1 −and 2.5± are the most common± ones. High− values± of β are often 4 4.05 4.1 4.15 4.2 4.25 4.3 4.35 4.4 4.45 4.5 d 54= .276 3167.7, b331).= 973 The.4 derived107.5, coec ffi=cients6054 agree.3 very387.9 well, and 1 Log(T ) − ± ± − ± 4 4.05 4.1 4.15 4.2 4.25eff 4.3 4.35 4.4 4.45 4.5 foundd = 27 among 316± the331). late B-type The derived stars. coefficients agree very well Log(T ) with the fit obtained by Markova & Puls (2008). We illustrate1 eff withThe the fit terminal obtained± velocities by Markova range & Puls from (2008). 155 We km illustrate s− to Linear relation (log T –log g ) of our sample of B Galac- our fit in Fig.1 6 with a grey wide band that has a dispersion Fig. 7. eff cor of1700our 1 290 fit km in K. s Fig.− The,6 while supergiants with mass-lossa grey of wide our rates sample, band range that with has generally the a exception dispersion from Fig.tic supergiants7. Linear relation (dashed (log line).Teff The- log surfacegcor) of gravity our sample was corrected of B Galac- by the 6 1 6 1 tic supergiantsLinear (dashed relation line). (log T Theff - surface log g gravity) of our was sample corrected of B by Galac- the of0.08of HD 1 290 52089,10 K.− The HDM 53138,supergiantsyr− to and 0.7 HD of our 4724010− sample,M (indicatedyr with− (four the in the exception stars figure fall Fig.centrifugal 7. acceleration. e cor outside× this mass-loss range). × centrifugaltic supergiants acceleration. (dashed line). The surface gravity was corrected by the withof HD large 52089, open HD circles), 53138, fall and inside HD the 47240 traced (indicated relation. in Moreover, the figure centrifugal acceleration. wewith canIn large observe addition, open that circles), we the calculate departures fall inside the of the average the traced effective values relation. temperatures for Moreover, the total large circles) can be fitted by the following relationship: formechanicalwe the can early observe BSGsmomentum that range the flow departures roughly of the from wind,of the 0 KD eff tomomective 5, 000 as temperaturesa K function and tend of the luminosity of the star. This is known as the modified WLR 2005; Zorec et al. 2005), giving a later evolutionary stage of the tofor be the lower early the BSGs later rangethe B spectral roughly subtype. from 0 K This to 5 same 000 K tendency and tend log Dmom = 1.96( 0.28) log L/L + 17.98( 1.43), (Kudritzki et al. 1995), given by object.2005; Zorec et al.± 2005), giving a later evolutionary± stage of the isto also be lower observed the in later Markova the B spectral et al. (2008). subtype. This same tendency for B0 B1.5. (2) is also observed in Markova et al. (2008). object.From− Table 3 we can distinguish stars showing two different Figure˙ 7 displays0.5 a1 linear/αeff correlation between log Teff and Dmom = M v R L , (1) windFrom regimes. Table One 3 we group can gathers distinguish stars stars with showing steep velocity two di gradi-fferent log gFigurecor (the 7 surface∞ displays∝ gravity a linear corrected correlation by the between centrifugal log T accel-eff and The second WLR was traced using BSG stars with spectral log g (the surface gravity corrected by the centrifugal accel- ents,windβ regimes.1, in One good group agreement gathers with stars the with standard steep velocity wind theory, gradi- eration,wherecorα see= Tableα δ A.1).(see Thisdetails kind in Kudritzki of correlation et al. was1999 previously; Puls et al. types B2-B9≤ (see Fig.8, large triangles). The values of the wind eration,e seeff Table A.1). This kind of correlation was previously whileents, theβ other1, in group good showsagreement values with of theβ between standard 1.5 wind and theory, 3.3, reported1996); α byand Searleδ are− theet al. force (2008). multipliers We performed related to a the linear line fit opacity and momentum≤ rates of these stars are clearly lower than those of the reported by Searle et. al. (2008).. We performed a linear fit and similarwhile theto the other results group reported shows by values Crowther of β between et al. (2006); 1.5 and Searle 3.3, obtainedand wind a ionization, slope of 3 respectively.98 0 31. This indicates that most of early-type objects. Furthermore, the WLR displays a different theobtained stars in a our slope sample of 3. have98± a0 similar.31. This luminosity-to-mass indicates that most ratio of etsimilar al. (2008), to the and results Markova reported et al. by (2008). Crowther Stars et with al. (2006);β between Searle 2 Our sample splits into± two different groups (see Fig.8): the slope (less steep) with luminosity. For these objectsβ we find a (logthe ( starsL/M) in our4), assample expected, have since a similarL T luminosity-to-mass4 . ratio andet al. 2.5 (2008), are the and most Markova common et al.ones. (2008). High Stars values with of β betweenare often 2 early BSGs with spectral types betweene B0ff and B1.5 (blue dia- linear regresion of the form (log (L/M)∼ 4), as expected, since L∝ T 4 . foundand 2.5 among are the the most late B-type common stars. ones. High values of β are often monds)A large and∼ scattermid- to late-typeis observed BSGs around with∝ spectrale theff stars types that from have B2 foundThe among terminal the velocitieslate B-type range stars. from 155 km s 1 to 1 700 largeto B9A log (green largegcor scatter and( red2.5); is diamonds). observed in particular These around HD two the42087, groups stars HD seem that 52089, to have be log Dmom = 1.43( 0.42) log L/L + 19.94( 2.23),− ≥ The1 terminal± velocities range from 155± km s 1 to 1 7006 andlarge HD log 64760gcor do( not2.5); follow in particularthe relation. HD Two 42087, of these HD 52089, stars, km s− , while mass-loss rates range generally from− 0.08 10− clearly separated.≥ However, among the early B supergiants we for B211 B9. 6 1 × (3)6 and HD 64760 do not follow the relation. Two of these stars, Mkmyr s−− ,−to while 0.7 mass-loss10− M yr rates− (four range stars generally fall outside from this 0.08 mass-10− HDfound 42087 three and mid HD B 64760, stars (HD are close41117, to HD the 42087, TAMS and so thewebelieve variable 1 × 6 1 × theyHD have 42087 just and left HD the 64760, main sequence. are close HDto the 64760 TAMS is a so high we rotator believe lossM range).yr− to 0.7 10− M yr− (four stars fall outside this mass- star HD 99953) with wind momentum rates comparable to the lossThe range).observed offset× between the WLRs increases from the early- and,earlythey as have ones. a consequence, just left the main we should sequence. observe HD higher 64760 apparentis a high rotator bolo- toIn the addition, mid-type we BSGs. calculate Despite the average the large values errors for of the these total quanti- me- metricand,From as luminosities a consequence, a linear regression,and welower should e weffective observe found temperatures that higher the apparent sample than those of bolo- the chanicalties,In the addition, momentum tendency we is calculate oppositeflow of the to the averagewind, resultsD valuesmom observed, as for athe function by Kudritzki total me- of expectedearlymetric BSGs luminosities for their plus non-rotating the and three lower mentioned stellar effective counterparts mid temperatures B stars (Frémat (see than Fig. et those al.8, thechanicalet luminosityal.(1999 momentum) and of predicted the star. flow This by of Vink the is known wind,(2000).D asmom the, asmodified a function WLR of expected for their non-rotating stellar counterparts (Frémat et al. the luminosity of the star. This is known as the modified WLR Article number, page 12 of 28 A91,Article page number, 12 of 28 page 12 of 28 Haucke et al.: Wind Properties of Variable B Supergiants M. Haucke et al.: Wind properties of variable B supergiants A&A proofs: manuscript no. 31678corr_last

31 log D mom7 (Fig. 8, black dots) were considered to fit the regres- Early (B0 − B1.5) L to stars. We think that the values we estimated in section §4.1 Mid (B2 − B3) sion lineM of the WLR for each group, i.e. early B or mid/late Late (B4 − B9) E are quite reliable. Uncertainties in the stellar radius yield an in- B stars. 6 In the next section, we return to the issue of mass-loss secure luminosity and a mass-loss discrepancy. This last issue is 30 variations. discussed in the next subsection. It 5 is interesting to stress that the wind regime of the stars defining each WLR is different. From Table 3 we found that the ) 29 early-type 4 stars have mostly β . 2 and terminal velocities greater 5.2. Spectral and wind variability esc mom than 500 km s 1, while the mid- and late-type stars have β 2 / V −

∞ 1 ≥ Variable B supergiants often show spectroscopic variability in

and terminalV 3 velocities lower than 500 km s− . Log(D 28 the optical and UV spectral region that can be attributed to the On the other hand, the parameter α , which is the inverse eff large-scale wind structures (Prinja et al. 2002). Variations linked of the 2 slope of the WLR, changes from 0.5 for the early BSGs to rotational modulation have been reported, for example, in 27 / α to 0.7 for the mid late ones. An increase in eff might suggest HD 14134 and HD 64760 (Morel et al. 2004; Prinja et al. 2002, highly 1 structured ionized outflows that lead to a weakening of respectively), while those associated with pulsation activity were the radiation force. Under this condition, different mechanisms found in HD 50064 and 55 Cyg (Aerts et al. 2010; Kraus et al. 26 0 4.5 5 5.5 6 6.5 (e.g. pulsations, 3.8 clumping) 4 might 4.2 also contribute 4.4 to 4.6 driving the 4.8 2015, respectively). Nevertheless, variability related to the pres- Log(L/L ) Sun wind. Log Teff ence of weak magnetic fields has also been proposed (Henrichs Finally, Fig. 9 shows the ratios of V /Vesc as a function of et al. 2003; Morel et al. 2004). Fig. 9. Ratios of V /Vesc as a function of log∞ Teff . Early-type (E), mid- Fig. 8. Wind-momentum luminosity relationships (dotted lines). Blue, logFig.T ff 9.. InRatios the of figureV∞ /V weesc as present a function our of results log Teff . together Early-type with (E), the mid- Wind-momentum luminosity relationships (dotted lines). Blue, e ∞ Most of the stars in our sample show Hα line variations in Fig.green, 8. and red diamonds represent early-, mid-, and late-type BSGs, typetype (M), (M), and and late-type late-type (L) stars are are indicated indicated by by blue, blue, green, green, and and red red green, and red diamonds represent early-, mid-, and late-type BSGs, two temperature regions defined by Markova & Puls (2008), respectively. The upper fit line was derived using data points in large diamonds,diamonds, respectively. respectively. Stars with wind variability variability are are indicated indicated with with both shape and intensity. These variations can take place on a respectively. The upper fit line was derived using data points in large which show significantly different V /Vesc ratios: 3.3 for effec- circles, while the lower fit line corresponds to data points with large largelarge open open triangles triangles and and connected with with∞ a a solid solid line. line. The The largest largest vari- vari- scale of a few days, but daily and even hourly variability has also circles, while the lower fit line corresponds to data points with large tive temperatures above 23 000 K, and 1.3 below 18 000 K. Er- triangles. Diamond symbols connected with solid lines correspond to abilityability is is shown shown by by HD HD 99953, which is located located in in the the region region limited limited been reported (e.g. Morel et al. 2004; Kraus et al. 2012, 2015; triangles. Diamond symbols connected with solid lines correspond to rors of 33% for the cooler and 43% for the hotter regions are observations of the same star in different epoch (black dots represent byby vertical vertical black lines lines (the (the bi-stability bi-stability jump). jump). The The shaded shaded areas areas refer refer to Tomic´ et al. 2015). Among our sample, HD 75149, HD 53138, ff observations of the same star in di erent epoch (black dots represent showntotwo two zones as zones shaded with with di areasfferent different (forV / detailsVV :/V 3.3esc on: for 3.3 the eff forective error effective temperatures determinations, temperatures above and HD 111973 display emission-line episodes or dramatic line their average values). esc∞ their average values). seeabove23 Markova 000 23 K, 000 and & K, Puls 1.3 and below 2008). 1.3 below 18∞ 000 18 K,000 with K, with errors errors of 43% of 43% and 33%, and 33%, re- profile variations with a diversity of alternating shapes (from a In Fig.8, observations of the same star at different epochs respectively,spectively,The region which which in between, were were identified identified delimited by by Markova Markova by vertical & & Puls Puls (2008).black(2008 lines,). is pure absorption to a P Cygni profile and vice versa). called the bi-stability jump. Markova & Puls (2008) showed that (Kudritzkiare connected et al. 1995), by solid given lines. by For these stars, average values of Regarding the amplitude of the mass-loss variation in indi- thevariation stars located for this there star show is observed a gradual in decreasethe WLR. in WithV / theVesc exception. Over- vidual objects, Prinja & Howarth (1986) measured a variability log Dmom (Fig.8, black dots) were considered to fit the regres- radiative transfer problems in the co-moving frame,∞ based on 0.5 1/αeff all,of our HD results 80077, show which a similar has a huge behaviour mass-loss to that rate, reported we do not by observe these at 10% level on timescales of a day or longer, while Lefever et al. Dmomsion= lineM˙ v ofR the WLRL for, each group, i.e. early B or mid/late(1) NLTE line-blanketed spherically symmetric wind models. ∞ authors. In particular,V we/V found many stars in our sample located (2007) and Kraus et al. (2015) reported on ratios between maxi- B stars. In the next∝ section, we return to the issue of mass-loss thatFor a decrease most of in the starsesc inis our accompanied sample, we by found an increase good agree- in the insidemass-loss the bi-stability rate, as predicted∞ region. byMoreover, Vink et al. as(1999 some). stars In general, are vari- mid mum and minimum mass-loss rates ranging from 1.05 up to 2.5. wherevariations.αeff = α δ (see details in Kudritzki et al. 1999; Puls et al. T − ablement we alsobetween plotted our ine Fig.ff estimates 9 the various and˙ those positions obtained of by those Crowther stars 1996);Itα isand interestingδ are the force to stress multipliers that the related wind toregime the line of opacity the stars andet al. late (2006, BSGs six present stars in values common), of M Lefeverthat areet similar al. (2007, to or eight lower In this way, based on a carefully study of the variation in the observed during different epochs. These positions are connected anddefining wind ionization, each WLR respectively. is different. From Table3 we found that the thanstars those in common), of the stars and Searle located et on al. the (2008, hot five side stars of the incommon), bi-stability Hα line among the stars we observed at different epochs (mod- with a solid black line. ˙ early-typeOur sample stars splits have into mostly twoβ di <ff2erentand terminal groups (see velocities Fig. 8): greater the jump.within error margins of 1500 K. Good agreements were also elled with the code FASTWIND), the largest variations in M (a 1 The variation in V /Vesc±of those stars on the cool side of the factor of 1.5 to 2.7) are detected in HD 75149, HD 99953, and earlythan BSGs 500 km with s− spectral, while thetypes mid- between and late-type B0 and B1.5stars have(blueβ dia- 2 found with the values∞ derived by Fraser et al. (2010, twelve 1 ≥ bi-stability jump seems to be small. Instead, HD 99953 (green α monds)and terminal and mid- velocities to late-type lower BSGs than with500 km spectral s− . types from B2 stars in common), who used TLUSTY. Among all the pre- HD 111973. As the H line is very sensitive to changes in the diamonds5. Discussion with triangles), which is located inside the bi-stability mass density and in the velocity gradients at the base of the wind to B9On (green theother and red hand, diamonds). the parameter Theseα twoeff, groups which is seem the inverseto be vious mentioned authors, the largest discrepancies in Teff are clearlyof the separated. slope of the However, WLR, changes among the from early 0.5 B for supergiants the early BSGs we region,found presents for six stars, such HD a large 42087, wind HD 52089, variability HD 53138, that it HD seems 64740, to (Cidale & Ringuelet 1993), a connection between line variability switch5.1. Photosphericfrom a slow to parameters a fast regime. The same variation for this and changes in the wind structure due to photospheric instabili- foundto 0.7 three for mid the mid/late B stars (HD ones. 41117, An increase HD 42087, in α andeff might the variable suggest HD 92964, and HD 99953, which show appreciable differences star is observed in the WLR. With the exception of HD 80077, ties can be expected. starhighly HD 99953) structured with ionized wind momentum outflows that rates lead comparable to a weakening to the of For(∆Teff some> 2 000 stars K). different Two of them estimates (HD 53 of 138T andeff HDand 64740)log g which has a huge mass-loss rate, we do not observe that a de- earlytheones. radiation force. Under this condition, different mechanisms (have∆Teff shown> 1500 eitherK, and some∆ log degreeg 0 of.2) ion are stratification found in the or literature ioniza- To gain insight into the origin of the variability of the BSG / ∼ (e.g.From pulsations, a linear clumping)regression, might we found also contributethat the sample to driving of the the crease(seetionin Table changesV A.1Vesc in).is the accompanied UV lines (see by an details increase in Prinja in the et mass-loss al. 2002). winds, we gathered information from the literature for all our rate, as predicted∞ by Vink et al. (1999). In general, mid and late objects on periods of light and spectroscopic variations (quoted earlywind. BSGs plus the three mentioned mid B stars (see Fig. 8, TheIn star relation HD 52089 with has these a magneticTeff estimates, field and it is X-ray important emission to point(see BSGs present values of M˙ that are similar to or lower than those in Table 1), and on previous determinations of the wind param- large circles)Finally, can Fig. be9 shows fitted by the the ratios following of V relationship:/Vesc as a function of outsection that 4.1) this and parameter like HD is 99953 derived is located using inside different the methodsbi-stability or ∞ of the stars located on the hot side of the bi-stability jump. eters derived from individual modelling attempts. A comparison log Teff. In the figure we present our results together with the two approaches.region. This Quite might often suggest the that atmospheric these stars parameters show intrinsic are obtainedtemper- temperature regions defined by Markova & Puls(2008), which byature fitting variations the observed since we line do notprofiles observe with a tendency a synthetic regarding spectrum a of the mass-loss estimates listed in Table A.1 is of particular in- log Dmom = 1.96( 0.28) log L/L +17.98( 1.43), for B0 B1.5. terest as they display ratios between maximum and minimum show significantly± different V / Vesc ratios:± 3.3 for effective− tem- calculatedparticular method from NLTE used line-blanketedto derive their atmospheric plane-parallel parameters. hydrostatic peratures above 23 000 K, and∞ 1.3 below 18 000 K. Errors of(2) 33% 5.model DiscussionWe atmospheres, also found that using the theline codes spectra TLUSTY do not show and He SYNSPEC forbid- values that range from 1 to 7. These discrepancies are not com- ff for the cooler and 43% for the hotter regions are shown as shaded (denHubeny components & Lanzsuggesting 1995; Lanz that & the Hubeny He abundance 2007). This is close procedure to the pletely true because these models were computed using di erent The second WLR was traced using BSG stars with spectral 5.1. Photospheric parameters ff areas (for details on the error determinations, see Markova & wassolar mainly value.carried This is in out accordance by Fraser with et al. what(2010 Kraus), McErlean et al. (2015) et al. stellar radii, and the di erences in mass-loss rates could be much types B2-B9 (see Fig. 8, large triangles). The values of the wind lower. For instance, for HD 34085 we found a discrepancy in M˙ Puls 2008). For(observed1999 some), Kudritzki stars in 55 di Cyg.fferent et In al. estimates addition,(1999), while Crowther of T eCrowtherff and et al. log (2006) etg al.(∆(T2006 dideff not>), momentum rates of these stars are clearly lower than those of the 1 500report K, and any∆ discrepancieslog g 0.2) on are the found Si and in Hethe abundances. literature (see Table of 1.9 when compared with the result given by Markova & Puls The region in between, delimited by vertical black lines, is Lefever et al.(2007∼), Markova & Puls(2008), Searle et al. ˙ 6 early-type objects. Furthermore, the WLR displays a different We did not find any correlation between V and T ff, as (2008), but it could be only 1.3 times higher (M = 0.44 10− called the bi-stability jump. Markova & Puls(2008) showed that A.1).(2008) used codes like FASTWIND or CMFGENmacro (Pulse et al. 1 × slope (less steep) with luminosity. For these objects we find a M yr− ) if a value of R? = 115 R is adopted (keeping all the the stars located there show a gradual decrease in V /V . Over- 2005wasIn relation reported; Hillier with by & Fraser Miller these etT 1998e al.ff estimates, (2010).) that In solve general, it is multi-level important we needed to non-LTE point larger linear regresion of the form esc 1 other parameters constant). all, our results show a similar behaviour to that reported∞ by these outradiativevalues that this for transfer parameterVmacro (mostly problems is derived around in usingthe 50 co-moving60 diff kmerent s− methods frame,for all the based or stars) ap- on − Therefore, to compare the values obtained by different au- authors. In particular, we found many stars in our sample located proaches.NLTEthan the line-blanketed Quite ones often given the in spherically theatmospheric literature symmetric parameters to be able wind to are models. fit obtained the pho- log Dmom = 1.43( 0.42) log L/L +19.94( 2.23), for B2 B9. bytospheric fitting the line observed widths. line In some profiles cases with it awas synthetic even necessary spectrum to thors we should discuss the scaling properties of the correspond- inside the bi-stability± region. Moreover, ± as some stars are− vari- For most of the stars in our sample, we found good agree- calculatedslightly fromincrease NLTEV sin line-blanketed i as well. This plane-parallel could be due tohydrostatic our use of ing models since line fits are not unique (Puls et al. 2008). Thus able we also plotted in Fig.9 the various positions of those stars(3) ment between our Teff estimates and those obtained by Crowther ˙ 1.5 a mean value for Vmacro to model all the lines and, particularly, we will adopt the optical-depth invariant Qr = M/R? , instead observed during different epochs. These positions are connected modelet al. atmospheres,(2006, six stars using in common), the codes TLUSTYLefever et and al.(2007 SYNSPEC, eight 1.5 the line-forming region for He might extend into the base of the of the traditional parameter Q = M˙ /(R? V ) (Puls et al. 1996) withThe a observed solid black off line.set between the WLRs increases from the (Hubenystars in & common), Lanz 1995; and Lanz Searle & Hubeny et al.(2008 2007)., five This stars procedure in com- ∞ wind where higher turbulent velocities can be expected. since we assume that V can change due to the wind variability. early-The to the variation mid-type in BSGs.V /Vesc Despiteof those the stars large on errors the of cool these side wasmon), mainly within carried error out margins by Fraser of 1500 et al. (2010);K. Good McErlean agreements et wereal. ∞ quantities, the tendency is∞ opposite to the results observed by (1999);The Kudritzki major discrepancy et al. (1999), found± while among Crowther the stellar et al. parameters (2006); The corresponding calculations and errors for log Qr are given of the bi-stability jump seems to be small. Instead, HD 99953 alsowith found previous with works the values is related derived to the by determination Fraser et al.( of2010 the, twelvestellar in Table A.1. An overview of these values shows that the differ- Kudritzki(green diamonds et al. (1999) with and triangles), predictedby which Vink is (2000). located inside the Lefeverstars in et common), al. (2007); who Markova used& TLUSTY. Puls (2008); Among Searle all the et al. pre- radius, due to the different methods used to calculate distances ences in log Qr (for the stars we observed at different epochs) bi-stabilityIn Fig. 8, region,observations presents of the such same a large starwind at diff variabilityerent epochs that (2008)vious used mentioned codes like authors, FASTWIND the largest or discrepanciesCMFGEN (Puls in T eteff al.are areit connected seems to byswitch solid from lines. a For slow these to a stars, fast average regime. values The same of 2005;foundArticle Hillier for number, six & stars, pageMiller HD 14 of1998) 42087, 28 that HDsolve 52089, multi-level HD 53138, non-LTE HD 64740,

Article number,A91, page page 13 13 of of28 28 A&A 614, A91 (2018)

HD 92964, and HD 99953, which show appreciable differences wind (Cidale & Ringuelet 1993), a connection between line vari- (∆Teff > 2000 K). Two of them (HD 53 138 and HD 64740) have ability and changes in the wind structure due to photospheric shown either some degree of ion stratification or ionization instabilities can be expected. changes in the UV lines (see details in Prinja et al. 2002). The To gain insight into the origin of the variability of the star HD 52089 has a magnetic field and X-ray emission (see BSG winds, we gathered information from the literature for Sect. 4.1) and like HD 99953 is located inside the bi-stability all our objects on periods of light and spectroscopic varia- region. This might suggest that these stars show intrinsic tem- tions (quoted in Table1), and on previous determinations of the perature variations since we do not observe a tendency regarding wind parameters derived from individual modelling attempts. a particular method used to derive their atmospheric parame- A comparison of the mass-loss estimates listed in Table A.1 ters. is of particular interest as they display ratios between max- We also found that the line spectra do not show He forbid- imum and minimum values that range from 1 to 7. These den components suggesting that the He abundance is close to the discrepancies are not completely true because these models were solar value. This is in accordance with what Kraus et al.(2015) computed using different stellar radii, and the differences in observed in 55 Cyg. In addition, Crowther et al.(2006) did not mass-loss rates could be much lower. For instance, for HD 34085 report any discrepancies on the Si and He abundances. we found a discrepancy in M˙ of 1.9 when compared with We did not find any correlation between Vmacro and Teff, as the result given by Markova & Puls(2008), but it could be 6 1 was reported by Fraser et al.(2010). In general, we needed larger only 1.3 times higher (M˙ = 0.44 10− M yr− ) if a value 1 × values for Vmacro (mostly around 50 60 km s− for all the stars) of R? = 115 R is adopted (keeping all the other parameters than the ones given in the literature− to be able to fit the pho- constant). tospheric line widths. In some cases it was even necessary to Therefore, to compare the values obtained by different slightly increase V sin i as well. This could be due to our use of authors we should discuss the scaling properties of the cor- a mean value for Vmacro to model all the lines and, particularly, responding models since line fits are not unique (Puls et al. the line-forming region for He might extend into the base of the 2008). Thus we will adopt the optical-depth invariant Qr = 1.5 1.5 wind where higher turbulent velocities can be expected. M˙ /R? , instead of the traditional parameter Q = M˙ /(R? V ) The major discrepancy found among the stellar parameters (Puls et al. 1996) since we assume that V can change∞ due with previous works is related to the determination of the stellar to the wind variability. The corresponding∞ calculations and radius, due to the different methods used to calculate distances errors for log Qr are given in Table A.1. An overview of these to stars. We think that the values we estimated in Sect. 4.1 are values shows that the differences in log Qr (for the stars we quite reliable. Uncertainties in the stellar radius yield an inse- observed at different epochs) are two or three times larger cure luminosity and a mass-loss discrepancy. This last issue is than their corresponding error bars suggesting real variations discussed in the next subsection. in the wind parameters. These differences can also be observed among the Qr values we calculated from data given by different 5.2. Spectral and wind variability authors. It is interesting to note that we found five stars in the domain Variable B supergiants often show spectroscopic variability in of the bi-stability region. HD 99953 shows pronounced Hα vari- 1 the optical and UV spectral region that can be attributed to the ations from which different estimations of V (250 km s− , 1 1 ∞ 6 1 large-scale wind structures (Prinja et al. 2002). Variations linked 500 km s− , and 700 km s− ) and M˙ (0.08 10− M yr− , 6 1 6 ×1 to rotational modulation have been reported, for example, in 0.13 10− M yr− , and 0.22 10− M yr− ) were obtained. HD 14134 and HD 64760 (Morel et al. 2004; Prinja et al. 2002, This implies× variations in V and×M˙ of a factor of about 2.8. This respectively), while those associated with pulsation activity were result goes against the theoretical∞ prediction made by Vink et al. found in HD 50064 and 55 Cyg (Aerts et al. 2010; Kraus et al. (1999) who argue that the jump in mass loss is accompanied by 2015, respectively). Nevertheless, variability related to the pres- a steep decrease in the ratio V /Vesc (from 2.6 to 1.3), which is ence of weak magnetic fields has also been proposed (Henrichs close to the observed bi-stability∞ jump in the terminal velocity. et al. 2003; Morel et al. 2004). Moreover, these authors proposed that if the wind momentum Most of the stars in our sample show Hα line variations in MV˙ were about constant across the bi-stability jump, the mass- both shape and intensity. These variations can take place on a loss∞ rate would have increased by a factor of two from stars with scale of a few days, but daily and even hourly variability has also spectral types earlier than B1 to later than B1. been reported (e.g. Morel et al. 2004; Kraus et al. 2012, 2015; In the particular case of HD 99953, our result indicates that Tomic´ et al. 2015). Among our sample, HD 75149, HD 53138, MV˙ is not constant (see Fig.8) nor does the mass loss increases ∞ and HD 111973 display emission-line episodes or dramatic line when the ratio V /Vesc decreases, even when the observed profile variations with a diversity of alternating shapes (from a changes in the variables∞ are of the order of the expected values pure absorption to a P Cygni profile and vice versa). ( 2.8). On the other hand, the mass loss derived for HD 99953 ∼ 1 Regarding the amplitude of the mass-loss variation in indi- (log M˙ = 6.66, when V = 700 km s− ) agrees with the value − ∞ vidual objects, Prinja & Howarth(1986) measured a variability computed by Vink et al.(1999, 6.54 for V /Vesc = 2.6) in at 10% level on timescales of a day or longer, while Lefever Table 1; however, the terminal velocity− we derived∞ is the half of et al.(2007) and Kraus et al.(2015) reported on ratios between the theoretical value found by these authors. We do not believe maximum and minimum mass-loss rates ranging from 1.05 up to that the discrepancy could be due to the luminosity of the star 2.5. (log L/L = 4.87 0.37) since the models were computed for a In this way, based on a carefully study of the variation in log L/L = 5. We± also found that this is not related to the velocity the Hα line among the stars we observed at different epochs law adopted for the model since HD 47240 does not fit either. (modelled with the code FASTWIND), the largest variations in Looking for other observed BSGs with log L/L 5 (see M˙ (a factor of 1.5 to 2.7) are detected in HD 75149, HD 99953, Table1), we found that the wind parameters of the stars∼ on and HD 111973. As the Hα line is very sensitive to changes in the hot side of the bi-stability jump, like HD 38771 (with 1 the mass density and in the velocity gradients at the base of the log M˙ = 6.85 and V = 1500 km s− ) and HD 52382 (with − ∞ A91, page 14 of 28 HauckeM. Haucke et al.: et Wind al.: Wind Properties properties of Variable of variable B Supergiants B supergiants are two or three times larger than their corresponding1 error bars 31 log M˙ = 6.62 and V = 1000 km s− ), resemble those listed E suggesting− real variations∞ in the wind parameters. These differ- M in Table1 by Vink et al.(1999). In contrast, the wind param- L enceseters can of the also stars be observed on the cool among side the ofQ ther values bi-stability we calculated jump, like from data given by different authors. 30 HD 34085 (with log M˙ = 6.64 and V /Vesc = 0.8), HD 58350 It is interesting to note that− we found∞ five stars in the domain (with log M˙ = 6.82 and V /Vesc = 0.9), and HD 111973 (with of the bi-stability− region. HD 99953∞ shows pronounced Hα vari-

˙ ) log M between 6.7 and 6.85, and V /Vesc = 1.4) do1 not ﬁt 29 ations from which− diﬀerent− estimations of∞V (250 km s− , 500

any model predictions. They often have lower terminal veloci- mom 1 1 ˙ 6 ∞ 1 6 kmties s− and, and mass-loss 700 km s rates− ) and thanM (0 values.08 10 expected− M yr from− , 0 the.13 models.10− M yr 1, and 0.22 10 6 M yr 1) were× obtained. This implies× − − − Log(D 28 This means that× a decrease in the ratio V /Vesc, both inside variations in V and M˙ of a factor of about 2.8.∞ This result goes and on the cool∞ side of the bi-stability jump, is not accompa- againstnied by the an theoretical increase prediction in the mass-loss made by rate. Vink This et al. result (1999) supports who arguethe conclusions that the jump drawn in mass by Markova loss is accompanied & Puls(2008 by) who a steep used de- the 27 crease in the ratio V /Vesc (from 2.6 to 1.3), which is close to modified optical depth∞ invariant and found that both M˙ and V theare observed decreasing bi-stability in parallel. jump in the terminal velocity. More-∞ ˙ 26 over,We these believe authors that proposed the decrease that if in the both wind the momentum mass-loss rateMV and 4.5 5 5.5 6 6.5 ∞ Log(L/L ) werethe terminalabout constant velocity, across on the the cool bi-stability side of jump, the bi-stability the mass-loss jump, Sun ratecould would be explained have increased with the byδ a-slow factor hydrodynamic of two from solutionstars with for spectral types earlier than B1 to later than B1. Fig. 10. Comparison of various empirical WLR: this work (solid and radiation line-driven winds (Curé et al. 2011). Venero et al. Fig.dashed 10. Comparison black lines of are various for early empirical BSGs WLR:and mid- this or work late-type (solid BSGs, and (2016In the) found particular that a case change of HD in 99953,the ionization our result of the indicates wind, that char- dashedrespectively), black lines Herrero are for et early al.(2002 BSGs, solid and green mid- line), or late-type Searle etBSGs, al.(2008 re- , MV˙ is not constant (see Figure 8) nor does the mass loss in- spectively), Herrero et al. (2002, solid green line), Searle et al. (2008, acterized∞ by the parameter δ of the line-force multiplier (Abbott solid cyan line), and Mokiem et al.(2007, magenta solid line). The the- creases1982), when defines the two ratio differentV /Vesc stationarydecreases, wind even regimes whenthe (fast ob- and solidoretical cyanWLR line), and of Vink Mokiem(2000 et) is al. in (2007, blue, magenta solid and solid dashed line). lines The are the- for ∞ oretical WLR of Vink (2000) is in blue, solid and dashed lines are for servedslow). changes The fast in regime the variables is always are present of the order at high of the effective expected tem- Teff > 27 500 K and Teff < 22 500 K, respectively. The figure includes values ( 2>.8). On the other hand, the mass loss derived for Teffthe> sample27 500 of K stars and T showneff < 22 in 500 Fig.8 K,. respectively. The figure includes peratures∼ ( 25 000 K), while fast or slow regimes1 could develop HD 99953 (log M˙ = 6.66, when V = 700 km s− ) agrees with the sample of stars shown in Fig. 8. at low effective temperatures.− Both∞ wind regimes are separated theby value an instability computed zone by Vink where et non-stationary al. (1999, 6.54 flow for regimesV /Vesc exist.= − ∞ and X = 1.35 0.14 and D = 21.08 0.73 for the mid BSGs, 2.Venero6) in Table et al. 1;(2016 however,) proposed the terminal that anvelocity initial density we derived perturbation is the 0 samewhich work both the agree authors± with show our values. preliminary± computations of wind halfcould of the trigger theoretical a switch value in the found wind by regime these between authors. fast We and do not slow believe that the discrepancy could be due to the luminosity of the hydrodynamics using time-dependent wind solutions that lead solutions (and vice versa) in the B mid-type supergiants. In the to the formation of kink velocity field structures which might be starsame (log workL/L the= authors4.87 0 show.37) since preliminary the models computations were computed of wind 5.4. The pulsation–mass-loss relationship for a log L/L = 5. We± also found that this is not related to the related to the evolution of absorption components like the ones hydrodynamics using time-dependent wind solutions that lead observed in the UV line spectrum (DACs). In addition, at tem- velocityto the formationlaw adopted of for kink the velocity model since field HD structures 47240 doeswhich not might fit Saio(2011) demonstrated that one or more radial strange modes either. peratures < 17 000 K the models predict a decrease in both ter- be related to the evolution of absorption components like the minalare expected velocity toand be mass-loss excited in rate evolved (see BSGs.Figure These4 in Venero authors et com- al. Looking for other observed BSGs with log L/L 5 (see puted the periods of excited modes during the evolution of only ones observed in the UV line spectrum (DACs). In addition,∼ at 2016). Futher research on the slow hydrodynamics solution is in Tabletemperatures 1), we found<17 000that Kthe the wind models parameters predict of a the decrease stars on in the both two stellar models: stars with initial masses of 20 M and 25 M . ˙ progress (Venero et al. 2017). hotterminal side of velocity the bi-stability and mass-loss jump, like rate HD (see 38771 Fig. 4 (with in Venero log M et= al. It appears that the periods of these radial strange modes tend 1 ˙ 62016.85). and FutherV = research1500 km ons the− ) slow and hydrodynamics HD 52382 (with solution log M is= in to decrease when the effective temperature of the star increases − ∞ 1 6.62 and V = 1000 km s− ), resemble those listed in Table 5.3.during The the wind blue-ward momentum–luminosity evolution. relationship −progress (Venero∞ et al. 2017). 1 by Vink et al. (1999). In contrast, the wind parameters of the As radial strange modes are considered suitable to trigger stars on the cool side of the bi-stability jump, like HD 34085 The empirical WLR of massive stars is generally represented by 5.3. The wind momentum–luminosity relationship mass loss, we first checked which of the known (photometric (with log M˙ = 6.64 and V /Vesc = 0.8), HD 58350 (with a linearand spectroscopic) regression of the periods form (see Table1) might be considered − ∞ log M˙ = 6.82 and V /Vesc = 0.9), and HD 111973 (with as radial strange modes, and then opted for a comparison of The empirical− WLR of∞ massive stars is generally represented by log M˙ between 6.7 and 6.85, and V /Vesc = 1.4) do not fit logtheseDmom periods= X log withL/L the+ D amplitude0. of the mass-loss rates(4) (i.e. a linear regression− of the− form ∞ any model predictions. They often have lower terminal velocities ratios between maximum and minimum values) using the Qr Values for the coefficients X = 1/α and D reported in the andlog mass-lossDmom = X rateslog thanL/L values+ D0. expected from the models. This(4) parameters given in Table A.1. eff 0 means that a decrease in the ratio V /Vesc, both inside and on the literatureAccording are given to in the Table models 4, together of Saio with et al. the(2013 parameter), the radialαeff ∞ cool sideValues of the for bi-stability the coefficients jump,X is= not1/α accompaniedeff and D0 reported by an inin andmodes, the corresponding especially in late-type errors. It stars, is worth display mentioning the longest that periods. pre- creasethe literature in the mass-loss are given rate. in This Table result4, together supports with theconclusions the parame- viousTherefore, investigations for each on star the we WLR searched for Galactic for correlations BSGs (see with Table the drawnter α byeff Markovaand the corresponding & Puls (2008) who errors. used It the is worth modified mentioning optical 4)longest show that known wind period. momentum As shown rates derivedin Fig. 11 by, means we found of line- a lin- depththat previousinvariant investigations and found that on both theM WLR˙ and forV Galacticare decreasing BSGs (see in blanketedear correlation analyses for are stars systematically mostly represented higher than by those mid/late reported BSGs ∞ parallel.Table4) show that wind momentum rates derived by means by(i.e. Kudritzki spectral et type al. (1999). from B2 to B9) with P > 6 days. The longer ofWe line-blanketed believe that analyses the decrease are systematically in both the mass-loss higher than rate andthose theThe period coeffi thecients lower (X, theDo) amplitude derived in of this the work parameter for theQ earlyr. We thereported terminal byvelocity, Kudritzki on et the al.(1999 cool). side of the bi-stability jump, BSGsadded are to close our to sample those obtainedtwo well-studied by Herrero variable et al. (2002) stars having and couldThe be explained coefficients with (X, theD )δ derivedslow inhydrodynamic this work for solution the early Mokiemmeasurements et al. (2007), of their who stellar used and a wind sample parameters of stars consisting (Deneb and 0 − forBSGs radiation are close line-driven to those winds obtained (Curé by et Herrero al. 2011). etVenero al.(2002 et) and al. mostly55 Cyg, of Scuderi OSGs. Both et al. coe 1992fficients; Chesneau also agree et al. with2010 the; Kraus theoret- et al. (2016)Mokiem found et al.that(2007 a change), who in used the ionization a sample of the stars wind, consisting char- ically2015 predicted). Three values mid-type of Vink BSGs (2000), with measuredbut disagree values with the of P val- and acterizedmostly of by OSGs. the parameter Both coefficientsδ of the line-force also agree multiplier with the (Abbott theoret- uesM˙ obtainedwere rejected by Kudritzki from the et al. sample: (1999) HD and 53138, Searle HDet al. 75149, (2008). and 1982),ically defines predicted two values different of Vink stationary(2000), wind but regimes disagree (fast with and the TheseHD 99953 relations (Fig. are 11 plotted, green in dots Fig. without 10. triangles). Two of these slow).values The obtained fast regime by Kudritzki is always et present al.(1999 at) high and e ffSearleective et tem- al. starsParticularly, HD 75149 we and find HD that 99953 the havedifference been poorly between studied our (X and, D newo) peratures(2008). These(> 25 000 relations K), while are plotted fast or slowin Fig. regimes 10. could develop valuesperiods and should those of probably Searle et be al. determined. (2008) could The be former, due to thein partic- fact at lowParticularly, effective temperatures. we find that Boththe difference wind regimes between are our separated (X, D0) thatular, the might sample have of stars inaccurate selected minimum by these values authors of consideredM˙ since two a byvalues an instability and those zone of Searle where et non-stationary al.(2008) could flow be regimes due to the exist. fact mixof of our early observed and mid profiles BSG types. are in In absorption. fact, if we Thesplit star their HD sample, 111973 Venerothat the etal. sample (2016) of proposed stars selected that an by initial these density authors perturbation considered a weis obtain not includedX = 1.85 since0.36 we and haveD0 = observations18.47 2.00 during for early an BSGs, interval ± ± couldmix oftrigger early a and switch mid in BSG the types. wind regime In fact, between if we split fast their and sample, slow andofX time= 1 of.35 two0. days14 and whichD0 = is21 less.08 than0.73 the for reported the mid period BSGs, of ± ± solutionswe obtain (andX = vice1.85 versa)0.36 inand theD B mid-type= 18.47 supergiants.2.00 for early In BSGs, the whichvariation. both agree with our values. ± 0 ± Article number,A91, page page 15 15 of of 28 28 A&A 614, A91 (2018) A&A proofs: manuscript no. 31678corr_last Table 4. Coefficient for the WLR derived by different authors. Table 4. Coefficient for the WLR derived by different authors. Reference XD0 αeff sample used Reference XD0 αeff sample used Kudritzki et al. (1999) 1.34 0.25 21.24 1.38Kudritzki et al.(1999 E ) 1.34 0.25 21.24 1.38 E ± ± Kudritzki et al.(1999) 1.95 ± 0.2 17.07 ± 1.05 M Kudritzki et al. (1999) 1.95 0.2 17.07 1.05 M ± ± ± ± Vink(2000) 1.826 0.044 18.68 0.26 Teff > 22.5 kK from theory Vink (2000) 1.826 0.044 18.68 0.26 Teff > 22.5 kK from theory± ± ± ± Vink(2000) 1.914 0.043 18.52 0.23 Teff < 27.5 kK from theory Vink (2000) 1.914 0.043 18.52 0.23 Teff < 27.5 kK from theory± ± Herrero et al. (2002) 2.18 ±0.21 16.81 ± 1.16 0.46 0.04Herrero et al. mostly(2002) OSGs2.18 0.21 16.81 1.16 0.46 0.04 mostly OSGs ± ± ± Mokiem et al.(2007) 1.84 ± 0.17 18.87 ± 0.98 0.54 ± 0.05 mostly OSGs Mokiem et al. (2007) 1.84 0.17 18.87 0.98 0.54 0.05 mostly OSGs ± ± ± Searle et al. (2008) 1.59 ± 19.86 ± 0.78 0.63 ± 0.06Searle et al.(2008 E &) M 1.59 19.86 0.78 0.63 0.06 E&M ± ± ± This work 1.96± 0.28 17.98 ± 1.43 0.51 ± 0.07 E This work 1.96 0.28 17.98 1.43 0.51 0.07 E ± ± ± This work 1.43 ± 0.42 19.94 ± 2.23 0.70 ± 0.21This work M & L 1.43 0.42 19.94 2.23 0.70 0.21 M&L ± ± ± ± ± ± Notes. E, M, and L refer to early-, mid-, and late-type B supergiants. Notes. E, M, and L refer to early-, mid-, and late-type B supergiants. Haucke et al.: Wind Properties of Variable B Supergiants

100 oretical or empirical values found by Vink (2000); Herrero et al. 5.4. The pulsation–mass-loss relationship E M L 55 Cyg Linear ﬃ Late 6.5 HD 80077 (2002); Mokiem et al. (2007). As well, our coe cients would 90 Mid Saio (2011) demonstrated that one or more radial strange modes Early agree with data from Searle et al. (2008) if we separated their 55 CYG are expected to be excited in evolved BSGs. These authors com- 80 Deneb sample of stars into two groups: early and mid BSGs. 6 puted the periods of excited modes during the evolution of only Furthermore, we observe that a decrease in the ratio V /Vesc 70 ∞ two stellar models: stars with initial masses of 20 M and 25 M . 60 in the bi-stability region, and on its cool side, is accompanied by

It appears that the periods of these radial strange modes tend to 60 Sun 5.5 a decrease in mass-loss rate. This observation is contrary to the decrease when the eﬀective temperature of the star increases dur- 40 theoretical result obtained by Vink et al. (1999), based on a fast (days) 50 32 max ing the blue-ward evolution. Log L/L wind regime. Instead, this behaviour seems to be consistent with

P 5 40 HD 99953 As radial strange modes are considered suitable to trigger HD 41117 predictions from highly ionized stellar winds (δ slow solution, − mass loss, we first checked which of the known (photometric 30 20 HD 64760 Venero et al. 2016). The possible switch between fast and δ 4.5 HD 52089 − and spectroscopic) periods (see Table 1) might be considered as slow solutions should be explored in detail. radial strange modes, and then opted for a comparison of these 20 15 Finally, we found an empirical correlation (period–mass-loss 4 periods with the amplitude of the mass-loss rates (i.e. ratios be- 10 12 relationship) that associates the amplitude of the mass-loss vari- 4.8 4.6 4.4 4.2 4 3.8 tween maximum and minimum values) using the Qr parameters ations with long-term photometric and/or spectroscopic variabil- Log Teff given in Table A.1. -1.2 -1 -0.8 -0.6 -0.4 -0.2 0 0.2 ity. This relation is expressed in terms of the Qr parameter. The log[Qr /Qr ] According to the models of Saio et al. (2013), the radial min max found period–mass-loss relationship could indicate that radial Fig.Fig. 12. 12.Position of the sample of of stars stars in in the the HR HR diagram. diagram. Evolutionary Evolution- (or strange) pulsation modes with periods > 6 days are related modes, especially in late-type stars, display the longest periods. Fig. 11. RatiosRatios between between minimum minimum and and maximum maximumQQr valuesr values as asa function a func- ary tracks (taken from Ekström et al. 2012, black solid lines) correspond ˙ 1.5 tracks (taken from Ekström et al. 2012, black solid lines) correspond to the wind variability. To confirm this trend, it is necessary to Therefore, for each star we searched for correlations with the tionof the of longest the longest photometric/spectroscopic photometric/spectroscopic period, period, where whereQr = MQ/rR?=. to non-rotating stars with M? = 12, 15, 20, 32, 40, and 60 M . The 1.5 to non-rotating stars with M? = 12, 15, 20, 32, 40, and 60 M . The longest known period. As shown in Fig. 11, we found a linear MThe˙ /R linear. The correlation linear correlation was obtained was obtainedusing the using up-pointing the up-pointing triangle. tri- perform long-term high-resolution spectroscopic campaigns to ? colouredcoloured curves curves define define instability instability boundaries boundaries of of various various modes modes ( (SaioSaio correlation for stars mostly represented by mid/late BSGs (i.e. angle. 2011): a) stability boundary for low-order radial and non-radial p modes properly measure oscillation periods and simultaneously derive 2011): a) stability boundary for low-order radial and non-radial p modes homogeneous sets of mass-loss rates. spectral type from B2 to B9) with P > 6 days. The longer the (blue(blue solid solid line); b) b) non-radial non-radial g g modes modes with withl=1l and= 1 2 and (red 2 short-dashed (red short- The relationship we obtained is P = 100.27 period the lower the amplitude of the parameter Qr. We added to max dashedand solid and lines, solid respectively); lines, respectively); c) oscillatory c) oscillatory convection convection modes (cyan modes dot- Acknowledgements. We thanks our anonymous referee for the helpful comments mechanism( 7.64) X + that83. drives97( 4. the47) wind., where However, X = log pulsationQr / modesQr , may and (cyan dot-dashed lines); d) monotonically unstable radial modes (above our sample two well-studied variable stars having measurements ± ± min max dashed lines); d) monotonically unstable radial modes (above the ma- and suggestions. We also want to acknowledge J. Puls for allowing us to use the of their stellar and wind parameters (Deneb and 55 Cyg, Scud- alsowas play derived a significant using the role data in (quasi-)periodically enclosed in triangles lifting (the the used ma- thegenta magenta dotted-line, dotted-line, where where the hypergiant the hypergiant star HD star 80077 HD 80077 is found). is found). The FASTWIND code, for his help and advice with the code, and for the fruitful sug- The blue, green, red, and magenta diamonds correspond to early (E), gestions. This research has also made use of the SIMBAD database, operated eri et al. 1992; Chesneau et al. 2010; Kraus et al. 2015). Three terial,values which of Qrmin willand thenQrmax be acceleratedare listed in outwards bold in Table by radiation. A.1). The In blue, green, red, and magenta diamonds correspond to early (E), mid differences in the selected values are two or three times> larger mid (M), and late (L) BSGs, and 55 Cyg, respectively. at CDS, Strasbourg, France, and the International Variable Star Index (VSX) mid-type BSGs with measured values of P and M˙ were rejected this way, radial strange-mode pulsation with periods 6 days (M), and late (L) BSGs, and 55 Cyg, respectively. database, operated at AAVSO, Cambridge, Massachusetts, USA. L.C. and M.C. from the sample: HD 53138, HD 75149, and HD 99953 (Fig. mightthan their be responsible corresponding for error the wind bars, variability suggesting in real mid- variations and late- in acknowledge support from the project CONICYT + PAI/Atracción de capital 11, green dots without triangles). Two of these stars HD 75149 typethe mass-loss BSGs. rates. humano avanzado del extranjero (folio PAI80160057). L.C. also acknowledges boundaries for the different pulsation modes, as defined by Saio and HD 99953 have been poorly studied and new periods should The group that comprises the early BSGs (HD 38771, the ated with strong shocks originating in the interface between slow financial support from CONICET (PIP 0177), La Agencia Nacional de Promo- et al.(2013). The stability boundary for low-order radial and ción Científica y Tecnológica (PICT 2016-1971) and the Programa de Incentivos probably be determined. The former, in particular, might have binary star HD 47420, HD 52089, HD 64760, and HD 148688, and fast wind streams, such as in a model of co-rotating interac- 5.5. Location of the BSGs in the HR diagram. non-radialtion regionsp modes (Prinjais etrepresented al. 2002), and by its the WISE blue solid image line. reveals This a (G11/137) of the Universidad Nacional de La Plata (UNLP), Argentina. R.V. is inaccurate minimum values of M˙ since two of our observed pro- blue diamonds) does not show a clear correlation with their peri- grateful for the finantial support from the UNLP under programme PPID/G003. bendskind of and wind-lobe becomes structure. horizontal We due propose to strange-mode that this star instability could be files are in absorption. The star HD 111973 is not included since ods. Some of these stars show differences in their Qr parameters M.C., S.K., and C.A. acknowledge support from Centro de Astrofísica de Val- Figure 12 displays the position of our sample stars in the HR dia- which occurs when log L/M > 4. Non-radial modes (low-degree we have observations during an interval of time of two days of the order of the errors suggesting that they do not show strong evolving towards the RSG stage. paraíso. C.A. also thanks BECAS DE DOCTORADO NACIONAL CONICYT gram, using data from Table 3, together with the stability bound- high-order g-modes) can be excited in the region delimited/ by ˇ which is less than the reported period of variation. wind variability. In the same HR diagram we can also see that the mid- late- 2016-2017. M.K. acknowledges support from GACR (17-02337S). The Astro- aries for the different pulsation modes, as defined by Saio et al. redtype lines, stars whilst are located monotonically in the region unstable where modes oscillatory exist convection above the nomical Institute Ondrejovˇ is supported by the project RVO:67985815. This The relationship we obtained is P = 97.92( 7.86) X + We can then conclude that as the observed line profiles can max (2013). The stability boundary for low-order radial and non- magentamodes would dotted be line observable in the most (cyan luminous lines, Saio part 2011; of the Saio HR et dia- al. work was partly supported by the European Union European Regional Devel- 81.99( 4.56), where X = log Q / Q , and was derived± using be matched using a velocity β-law, radiation pressure is the opment Fund, project “Benefits for Estonian Society from Space Research and rmin rmax radial p modes is represented by the blue solid line. This bends gram2013), (as with is the the caseexception for the of hypergiantHD 41117 and star HD HD 99953, 80077). which On the are the data± enclosed in triangles (the used values of Q and Q main mechanism that drives the wind. However, pulsation modes Application” (KOMEET, 2014 - 2020. 4. 01. 16 - 0029), and by the institutional rmin rmax and becomes horizontal due to strange-mode instability which otherlocated hand, together two stars with in the our early sample B types. (HD The 52089 star and 55 HD Cyg 64760, (taken are listed in bold in Table A.1). The differences in the selected may also play a significant role in (quasi-)periodically lifting research funding IUT40 1 of the Estonian Ministry of Education and Research. occurs when log L/M > 4. Non-radial modes (low-degree high- withfromlog KrausL/L et al.< 2015)4.6, log wasL also/M included.< 4) are located inside the Financial support for International− Cooperation of the Czech Republic and Ar- values are two or three times larger than their corresponding er- the material, which will then be accelerated outwards by radi- order g-modes) can be excited in the region delimited by red instabilityIt is worth strip mentioning of the β Cep, that foundWISE byimages Saio reveal et al.( signs2013) of and pre- gentina (AVCR-CONICET/14/003) is acknowledged. ror bars, suggesting real variations in the mass-loss rates. lines,ation. whilst In this monotonically way, radial strange-mode unstable modes pulsation exist above with the periods ma- > Georgyvious strong et al. wind-ISM(2014), near interaction the cool in border three ofobjects that region.in our sample Par- The group that comprises the early BSGs (HD 38771, the genta6 days dotted might line be in responsible the most luminous for the wind part of variability the HR diagram in mid- and late-type BSGs. ticularly,(HD 52382, HD HD 64740 99953, presents and HD in 115842) the UV pointing spectral out region that clear these binary star HD 47420, HD 52089, HD 64760, and HD 148688, (as is the case for the hypergiant star HD 80077). On the other signsstars could of phase be evolving bowing towards and ionization the blue-loop changes after in a the post-RSG wind blue diamonds) does not show a clear correlation with their peri- hand, two stars in our sample (HD 52089 and HD 64760, with References 5.5. Location of the BSGs in the HR diagram thatphase. could This be same associated also holds with true strong for 55 shocks Cyg (Kraus originating et al. in2015). the ods. Some of these stars show differences in their Qr parameters log L/L < 4.6, log L/M < 4) are located inside the instabil- Abbott, D. C. 1982, ApJ, 259, 282 interface between slow and fast wind streams, such as in a of the order of the errors suggesting that they do not show strong ityFigure strip 12 of displays the β Cep, the found position by of Saio our et sample al. (2013) stars and inGeorgy the HR model of co-rotating interaction regions (Prinja et al. 2002), Aerts, C., Lefever, K., Baglin, A., et al. 2010, A&A, 513, L11 wind variability. et al. (2014), near the cool border of that region. Particularly, Aerts, C., Simón-Díaz, S., Catala, C., et al. 2013, A&A, 557, A114 diagram, using data from Table3, together with the stability and6. Conclusions its WISE image reveals a kind of wind-lobe structure. Astraatmadja, T. L. & Bailer-Jones, C. A. L. 2016, ApJ, 833, 119 We can then conclude that as the observed line profiles can HD 64740 presents in the UV spectral region clear signs of phase Barlow, M. J. & Cohen, M. 1977, ApJ, 213, 737 be matched using a velocity β-law, radiation pressure is the main bowingA91, page and 16 ionizationof 28 changes in the wind that could be associ- We studied a sample of 19 Galactic BSGs by fitting theoretical Benaglia, P., Vink, J. S., Martí, J., et al. 2007, A&A, 467, 1265 line profiles of H, He, and Si to the observed profiles. The syn- Bouret, J.-C., Lanz, T., & Hillier, D. J. 2005, A&A, 438, 301 Article number, page 16 of 28 thetic line profiles were computed with the NLTE atmosphere Carpay, J., de Jager, C., & Nieuwenhuijzen, H. 1991, A&A, 248, 475 code FASTWIND to obtain new estimates for their stellar and Carpay, J., de Jager, C., Nieuwenhuijzen, H., & Moffat, A. 1989, A&A, 216, 143 Castor, J. I., Abbott, D. C., & Klein, R. I. 1975, ApJ, 195, 157 wind parameters. The mass-loss rates for HD 74371, HD 99953, Chesneau, O., Dessart, L., Mourard, D., et al. 2010, A&A, 521, A5 and HD 111973 have been derived for the first time. In addi- Chesneau, O., Kaufer, A., Stahl, O., et al. 2014, A&A, 566, A125 tion, some BSGs in our sample show short-term variability in Chini, R., Hoffmeister, V. H., Nasseri, A., Stahl, O., & Zinnecker, H. 2012, MN- the Hα and UV resonance lines, indicating important changes in RAS, 424, 1925 the wind structure. Cidale, L. S. & Ringuelet, A. E. 1993, ApJ, 411, 874 Cohen, D. H., Cooper, R. G., Macfarlane, J. J., et al. 1996, ApJ, 460, 506 In relation with the global properties of the stellar winds of Crowther, P. A., Lennon, D. J., & Walborn, N. R. 2006, A&A, 446, 279 our sample, we find two different behaviours in the WLR. The Curé, M., Cidale, L., & Granada, A. 2011, ApJ, 737, 18 early and mid B-type stars have mostly β < 2 and terminal veloc- Ebbets, D. 1982, ApJS, 48, 399 1 Ekström, S., Georgy, C., Eggenberger, P., et al. 2012, A&A, 537, A146 ities greater than 500 km s− , while the other group is comprised mostly by mid and late B-type stars with β 2 and terminal ve- Evans, C. J., Lennon, D. J., Trundle, C., Heap, S. R., & Lindler, D. J. 2004, ApJ, 1 ≥ 607, 451 locities lower than 500 km s− . In addition, we find that the co- Feldmeier, A. 1998, A&A, 332, 245 efficients for the linear regression of early BSGs agree with the- Firnstein, M. & Przybilla, N. 2012, A&A, 543, A80

Article number, page 17 of 28 M. Haucke et al.: Wind properties of variable B supergiants

We propose that this star could be evolving towards the RSG grateful for the finantial support from the UNLP under programme PPID/G003. stage. M.C., S.K., and C.A. acknowledge support from Centro de Astrofísica de In the same HR diagram we can also see that the mid-/late- Valparaíso. C.A. also thanks BECAS DE DOCTORADO NACIONAL CONICYT 2016-2017. M.K. acknowledges support from GACRˇ (17-02337S). type stars are located in the region where oscillatory convection The Astronomical Institute Ondrejovˇ is supported by the project RVO:67985815. modes would be observable (cyan lines, Saio 2011; Saio et al. This work was partly supported by the European Union European Regional 2013), with the exception of HD 41117 and HD 99953, which are Development Fund, project “Benefits for Estonian Society from Space Research located together with the early B types. The star 55 Cyg (taken and Application” (KOMEET, 2014 – 2020. 4. 01. 16 - 0029), and by the institutional research funding IUT40 1 of the Estonian Ministry of Education and from Kraus et al. 2015) was also included. Research. Financial support for International− Cooperation of the Czech Republic It is worth mentioning that WISE images reveal signs of pre- and Argentina (AVCR-CONICET/14/003) is acknowledged. vious strong wind-ISM interaction in three objects in our sample (HD 52382, HD 99953, and HD 115842) pointing out that these stars could be evolving towards the blue-loop after a post-RSG References phase. This same also holds true for 55 Cyg (Kraus et al. 2015). Abbott, D. C. 1982, ApJ, 259, 282 Aerts, C., Lefever, K., Baglin, A., et al. 2010, A&A, 513, L11 Aerts, C., Simón-Díaz, S., Catala, C., et al. 2013, A&A, 557, A114 6. Conclusions Astraatmadja, T. L., & Bailer-Jones, C. A. L. 2016, ApJ, 833, 119 Barlow, M. J., & Cohen, M. 1977, ApJ, 213, 737 We studied a sample of 19 Galactic BSGs by fitting theoretical Benaglia, P., Vink, J. S., Martí, J., et al. 2007, A&A, 467, 1265 line profiles of H, He, and Si to the observed profiles. The syn- Bouret, J.-C., Lanz, T., & Hillier, D. J. 2005, A&A, 438, 301 thetic line profiles were computed with the NLTE atmosphere Carpay, J., de Jager, C., Nieuwenhuijzen, H., & Moffat, A. 1989, A&A, 216, 143 Carpay, J., de Jager, C., & Nieuwenhuijzen, H. 1991, A&A, 248, 475 code FASTWIND to obtain new estimates for their stellar and Castor, J. I., Abbott, D. C., & Klein, R. I. 1975, ApJ, 195, 157 wind parameters. The mass-loss rates for HD 74371, HD 99953, Chesneau, O., Dessart, L., Mourard, D., et al. 2010, A&A, 521, A5 and HD 111973 have been derived for the first time. In addi- Chesneau, O., Kaufer, A., Stahl, O., et al. 2014, A&A, 566, A125 tion, some BSGs in our sample show short-term variability in Chini, R., Hoffmeister, V. H., Nasseri, A., Stahl, O., & Zinnecker, H. 2012, the Hα and UV resonance lines, indicating important changes in MNRAS, 424, 1925 Cidale, L. S., & Ringuelet, A. E. 1993, ApJ, 411, 874 the wind structure. Cohen, D. H., Cooper, R. G., Macfarlane, J. J., et al. 1996, ApJ, 460, 506 In relation with the global properties of the stellar winds of Crowther, P. A., Lennon, D. J., & Walborn, N. R. 2006, A&A, 446, 279 our sample, we find two different behaviours in the WLR. The Curé, M., Cidale, L., & Granada, A. 2011, ApJ, 737, 18 early and mid B-type stars have mostly β < 2 and terminal veloc- Ebbets, D. 1982, ApJS, 48, 399 1 Ekström, S., Georgy, C., Eggenberger, P., et al. 2012, A&A, 537, A146 ities greater than 500 km s− , while the other group is comprised Evans, C. J., Lennon, D. J., Trundle, C., Heap, S. R., & Lindler, D. J. 2004, ApJ, mostly by mid and late B-type stars with β 2 and terminal 607, 451 1 ≥ velocities lower than 500 km s− . 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Appendix A: Photospheric line-fitting and FASTWIND and the parameters quoted in Table3. These compiled data from the literature parameters are also given in Table A.1 together with a set of stellar and wind parameters gathered from the The following figures present a comparison between observed literature. and synthetic photospheric line profiles calculated with

A91, page 19 of 28 A&A 614, A91 (2018) ˇ cka &) ( 2001 Kubát Crowther et al. ( 2006 ) Lefever et) al. ( 2007 McErlean et) al. ( 1999 Crowther et al. ( 2006 ) Kudritzki et) al. ( 1999 ) ( 1980 Nerney This work This work This work Chesneau et) al. ( 2014 Firnstein &) Przybilla ( 2012 Przybilla et al. ( 2006 ) McErlean et) al. ( 1999 This work Crowther et al. ( 2006 ) This work Morel et al. ( 2004 ) Fraser et) al. ( 2010 Lefever et) al. ( 2010 This work Lefever et) al. ( 2010 Krti Fraser et) al. ( 2010 Searle et al. ( 2008 ) McErlean et) al. ( 1999 Gathier et) al. ( 1981 Prugniel et) al. ( 2011 McErlean et) al. ( 1999 Benaglia et) al. ( 2007 McErlean et) al. ( 1999 Morel et al. ( 2008 ) References Israelian et) al. ( 1997 Scuderi et) al. ( 1998 25 11 . . 0.25 Lefever et) al. ( 2007 0 0 ± ± ± 2 10 06 12 11 13 17 14 11 . 13 17 ...... 8.77 9.05 ) ( 1980 Nerney 0 0.17 Markova et( 2008 ) al. 0.11 This work 0.17 Searle et( 2008 ) al. 0.13 Kudritzki et) al. ( 1999 0.13 Kudritzki et) al. ( 1999 0.13 Kudritzki et) al. ( 1999 r 8 9 9 9 0.20 This work 0.33 Searle et al. ( 2008 ) 0 0 0 0 0 0 − − − − − − Q ± ± ± 16 07 73 82 44 ± ± ± ± ± ± ± ± ± ± ± ± . . . . . / / / / / 8.43 ) ( 1980 Nerney 8.36 ) ( 1980 Nerney 9 8 8 8 8 26 − − − − − − 24 42 81 85 . 60 95 log 37 25 21 11 ...... 26 . . . . 9 . 8.76 8.07 9.63 8.84 9.42 9.28 8.52 8.65 8 9 8 8 8 8 0.25 / 8 0 9 0 9 − − − − − − − − − − − − − − − − ± − ± ± 23 19 . . 8.92 9 9 − − − 09 83 08 27 14 09 ...... 13 12 13 13 13 13 8 5 Q 57 84 94 71 29 78 78 79 . 12 87 12 58 82 69 10 00 20 86 . − − − − − − ...... / / / / / / 12 12 · · ······· − ············ ······························ 13 12 12 12 12 13 13 12 13 12 12 13 13 12 13 13 13 12 − 68 93 35 42 35 36 − − − − − − − − − − − − − − − − − ...... 13 12 13 13 13 13 − − − − ) log

L 4 / 22 04 34 34 9 68 35 45 56 87 48 52 70 11 08 30 96 30 09 78 84 65 31 72 03 93 49 94 27 31 L ...... dex · · · · · · − · · · − · · · − · · · − · · · − 5 4 4 5.55 5 4 log ( 61 5 0 4 2 5 0 5 7 5 6 5 2 5 7 5 9 5 0 4 4 ...... ?

R 58 ··· ··· ························ ············ ··· ························ 59 i R 1 − sin 6055 39 58 65 5 80 83 22 82 30 3625 115 5 91 27 81 28 5 40 61 71 36 6060 35 58 54 72 61 71 26 33 135 36 63 ······ ··············· ······ km s V 1 − 40 22 36 109 5 65 3540 30 129 5 40 macro ··· ··· ··· ··· ··· ··· ····················· ··· km s V 1 5 − . 8 17 17 22 18 23 35 1516 2020 32 53 56 micro ··· ········· ······ ····················· ······ ······ ··· ······ ······ ··············· km s V 1 350 3450 600 450 10 60 40 46 5 − ∞ / / / / ··············· 1870 km s 400 1 − 4 530 24 1000 15 55 94 27 4 94 300 1 1900 31 490 15 45 38 50 5 . . . . . 24 600 yr 1 . 0 0 1 0

0 7 580 62 510 1000 / ˙ . 36 865 20 27 1350 10 80 13 34 230 90 1525 12 20 49 1200 20 1390 15 85 50050 20 11 650 735 15 40 45 500 20 23 155 1014 52 30 1500 72 1317 5 6090 510 80 510 13 1057 4 10 6520 700 4024 23 735 1502 4 80 450 80 90024 10 55 60 8 5 1000 95 65 10 30 10 65 5 11 55 4 21 4 . . / / / − 095 620 ...... / MV M 1 . 1 2 0 0 0 0 7 6 . 2 86 76 21 . . − 17 . 0 . . 0 0 10 5 0 0 0 0 0 5 0 5 0 5 0 1 1 0 0 2 0 0 0 2 0 6 0 5 0 00 0 0 0 0 00 0 0 2 0 0 0 ...... β 2 1 ······ ········· 1 1 1 3 1 cor g 05 2 90 1 94 48 1 71 16 2 72 2 76 1 32 2 35 2 47 2 47 1 00 1 48 2 26 2 ...... ··· ·················· ·················· ·················· ························ ··············· ·········································· ···································· ··· ········· ············································· ········· 2 log ﬀ e g 30 35 75 2 75 00 00 05 15 20 75 70 75 00 70 2 07 2 40 25 50 30 50 55 40 2 20 22 25 15 2 70 1 60 70 2 3010 2 2 45 2 40 2 00 3 45 2 25 2 ...... ········· ··· ········· ·················· log ﬀ e K dex dex ············ ············ ············ ············ T 18 50018 500 2 2 15 500 1 12 000 1 27 50027 500 3 3 20 100 3 15 400 2 19 500 2 12 000 12 10012 500 1 1 13 000 1 26 00026 500 3 2 26 000 3 20 000 2 19 500 2 18 00019 000 2 2 20 50020 500 2 2 19 000 2 22 00022 200 3 3 23 000 20 800 16 50017 000 2 2 13 000 1 19 000 2 18 500 17 150 12 700 1 25 000 2 19 000 2 16 500 2 19 000 2 23 000 3 21 500 2 18 000 2 Stellar and wind parameters. Highlighted values were used to make Fig. . 11 Name HD 34085 HD 38771 HD 41117 HD 42087 HD 47240 HD 52089 HD 52382 HD 53138 Table A.1. Notes.

A91, page 20 of 28 M. Haucke et al.: Wind properties of variable B supergiants ˇ ˇ cka &) ( 2001 Kubát cka &) ( 2001 Kubát Fraser et) al. ( 2010 This work Fraser et) al. ( 2010 Prinja &) Massa ( 2010 References This work Fraser et) al. ( 2010 This work Fraser et) al. ( 2010 Fraser et) al. ( 2010 Fraser et) al. ( 2010 This work Fraser et) al. ( 2010 Fraser et) al. ( 2010 Prinja &) Massa ( 2010 Lefever et) al. ( 2007 This work Fraser et) al. ( 2010 Searle et al. ( 2008 ) Lefever et) al. ( 2007 Benaglia et) al. ( 2007 Krti ) ( 1984 Underhill Fraser et) al. ( 2010 McErlean et) al. ( 1999 This work 09 13 . . 0.15 This work 0.11 This work 0 0 ± ± ± ± 20 42 25 09 09 15 17 25 09 11 ...... 9.28 8.75 0.11 This work 0.13 This work r 9 9 0.65 Carpay et) al. ( 1989 0.25 Lefever et) ( 2007 al. 0.25 Lefever et) al. ( 2007 0.25 Lefever et) al. ( 2007 0 0.13 This work 0.23 Searle et al. ( 2008 ) 0 0 0 0 0 0 0 − − − − Q ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± ± / / 8.10 Crowther et al. ( 2006 ) 9.71 Krti 8.00 Crowther et al. ( 2006 ) − − − 29 08 08 18 01 39 87 35 log . 09 13 ...... 8.70 8.06 8 9.57 9.07 8.62 8.45 0.11 / 0.15 / 8.00 9 9 8 8 9 8 9 8.79 0 0 − − − − − − − − − − − − − − − − ± ± ± ± 38 53 . . 9.19 9.72 9 9 − − − − 1 02 02 . 19 07 01 . . . . . 97 15 . . 13 13 13 13 14 13 12 12 8 8 − Q 76 98 93 . 82 85 55 16 17 . 91 83 61 43 39 − − − − − ...... / − / / / / / / / 12 12 ······ ······ 12 12 12 12 12 11 12 13 12 12 12 13 13 97 63 9 − − . 79 19 . . 14 79 90 − − − − − − − − − − − − . . . . . 11 13 12 12 13 13 13 12 − − − − − − − − ) log

24 L . / 5 3 4 57 33 67 16 23 53 87 68 35 . 48 65 33 49 95 . 45 18 23 10 23 36 - L ...... dex · · ·5 − ········· ········· 5 13 . log ( 3 5 2 5 4 7 5 3 5 . . . . . ?

R 68 ············ ············ ···························································· ············ ············ 162 6 62 i R 1 − sin 84 34 72 36 50 57 45 81 ······ 255 265 23 230 24 5 km s V 1 70 40 54 5 60 40 61 5 − 60 35 48 5 / / 40 20 30 31 macro ··· ··· ··· ··· ··· ··· ∼ km s V 1 − 17 55 / 15 22 19 28 36 14 36 39 22 2822 36 37 49 14 6616 39 44 48 20 34 30 18 25 32 micro ····················· ········· ······························ ····················· km s V 1 350 9 200 10 60 10700 195 18 6 50 50 25 4 807 − 233 11 50 ∞ / / / / - 520 km s 1 − 15 175 25 400 22 250 21 350 12 28 520 15 50 31 4840 5 1200 15 40 50 42 5 ...... yr 0 0 0 0 0 1

4 150 8 1700 10 120 70 35 5 0 1180 10 - ˙ 42 1500 15 10040 230 12 400 4 . 11 5349 40 370 61 11 5 40. 45 7015 5 1200 11 65 50 31 5 10 1600 00 140 . 80 435 11 530 10 50066 15 450 6070 30 140 39 4 75 725 15 70 320 20 42 1400 14 14 250 1228 40 200 37 10 65 60 5 30 73 5 / / / / / ...... MV M ··· 5 0 0 0 1 6 09 08 14 25 10 − 116 ...... 10 2 5 . 3 5 0 5 0 3 0 0 0 0 0 0 0 5 1 5 1 0 0 5 0 1 0 0 5 2 0 1 5 0 0 1 0 0 8 0 5 0 ...... β 0 1 ········· 1 − cor 5 3 g 09 0 11 2 02 3 20 3 21 2 3311 2 2 76 2 47 2 02 3 10 3 85 1 . 06 2 60 2 27 0 77 2 81 2 0 ...... ············ ··· ······ ············ ················································ ·················· ··············· ··············· ··············· log ﬀ e g 0 2 0 90 3 10 2 . 20 2 20 2 30 2 10 2 75 2 45 2 00 2 25 38 30 00 75 . 00 10 2 15 65 2 45 50 2 20 05 2 25 40 2 10 13 20 3 75 1 80 1 90 ...... ······ ········· ········· ········· ····················· log 16 000 2 ﬀ e K dex dex − T 16 000 2.10 23 000 2 16 00015 800 2 2 17 70018 000 2 2 19 00016 500 2 25 500 2 21 000 2 2 26 000 3 28 000 3 16 00016 500 2 17 000 2 24 800 2 15 10015 000 2 15 60018 00017 400 2 16 800 2 2 25 50020 70021 000 2 2 2 15 90016 000 2 2 13 600 18 500 2 22 000 2 14 50015 000 2 2 24 000 3 13 900 13 500 1 13 700 1 13 400 1 15 000 continued. Name HD 58350 HD 64760 HD 75149 HD 79186 HD 80077 HD 92964 HD 99953 HD 111973 HD 115842 HD 148688 HD 74371 Table A.1.

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Fig. A.1. HD 34085, HD 38771, HD 41117, and HD 47240: Line model ﬁttings to observations.

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Fig. A.2. HD 42087, HD 52089, HD 52382, and HD 53138: Line model ﬁttings to observations.

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Fig. A.3. HD 53138, HD 58350, and HD 64760: Line model ﬁttings to observations. . A91, page 24 of 28 M. Haucke et al.: Wind properties of variable B supergiants

Fig. A.4. HD 74371 and HD 75149: Line model ﬁttings to observations.

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Fig. A.5.Fig.HD A.5. 75149,HD HD 75149, 79186, HD HD 79186, 92964, HD and 92964, HD and80077: HD Line 80077: model Line ﬁttings model to ﬁttings observations. to observations. . .

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Fig. A.6.Fig.HD A.6. 99953HD and99953 HD and 111973: HD 111973: Line model Line ﬁttings model toﬁttings observations. to observations.

ArticleA91, number, page 27 page of 28 25 of 28 A&AA&A proofs: 614,manuscript A91 (2018) no. 31678corr_last

Fig. A.7. HD 111973, HD 115842, and HD 148688: Line model ﬁttings to observations.

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