ASTRONOMY AND ASTROPHYSICS Doppler imaging of stellar surface structure

Transcription

1 Astron. Astrophys. 357, (2000) ASTRONOMY AND ASTROPHYSICS Doppler imaging of stellar surface structure XIII. The flaring RS CVn-binary HD = V1355 Orionis K.G. Strassmeier Institut für Astronomie, Universität Wien, Türkenschanzstrasse 17, 1180 Wien, Austria (strassmeier@astro.univie.ac.at) Received 19 January 2000 / Accepted 9 March 2000 Abstract. We present the first Doppler images of the newly discovered RS CVn binary HD The star is a synchronuously rotating subgiant in a 4-day spectroscopic binary with an unseen companion and is particularly interesting because it is one of the few evolved stars that show large flares. Three consecutive years of moderate-resolution KPNO spectra and continuous VI and by photometry with our APTs reveal a large, cool, and long-living asymmetric polar spot. Its temperature is approximately 1100 K below the average photospheric temperature. Several low latitude and even equatorial spots were evident at the same time. We interpret this latitudinal spot bimodality to be due to a combination of poleward meridional circulation and more frequent magnetic reconnections near the pole than near the equator. During the 1998 Doppler-imaging observations, a strong flare in Hα was detected which coincided with a meridian passage of the most asymmetric part of the polar spot. We estimate a total flare energy of erg typical for flares on very active RS CVn systems. No photospheric heating is evident in our surface maps from optically thin spectral lines of various temperature sensitivities. Key words: stars: activity stars: imaging stars: individual: HD stars: late-type stars: starspots 1. Introduction HD (=V1355 Ori, K0-2IV, V=9 ṃ 0 9 ṃ 3) was detected as a bright EUV source in the ROSAT WFC all-sky survey (Pounds et al. 1993, Pye et al. 1995) which suggested, in combination with its late spectral type, the occurence of chromospheric and photospheric magnetic activity. Further attention was drawn to it when Cutispoto et al. (1995) discovered its light variability with a period of 3.82 days. At the time of their observations in late 1993, HD showed one of the largest photometric amplitudes ever recorded for a spotted star (0 ṃ 37 in the V bandpass). This prompted Kazarovets & Samus (1997) to add HD to the 73rd name list of variable stars. The star s Visiting Astronomer, Kitt Peak National Observatory, operated by the Association of Universities for Research in Astronomy, Inc. under contract with the National Science Foundation mean colors and optical spectrum are consistent with a composite K1-2IV+G2V classification but also with a single K2-3 dwarf classification (Cutispoto et al. 1995). A K0-2IV classification was suggested from a single, high-resolution, red-wavelength spectrum presented and analyzed by Osten & Saar (1998). HD was observed during the Vienna-KPNO search for Ca ii H&K emission stars in and was detected to exhibit very strong emission lines above the local continuum (Strassmeier et al. 2000). A single Hα spectrum showed a weak Hα emission line which usually indicates a very active atmosphere even in comparison with other RS CVn-type binaries. Finally, our initial photometric data from the observing season 1996/97 caught the star with an amplitude of up to 0 ṃ 25 in V and with a period of 3.87 days (Strassmeier et al. 1999), thus confirming Cutispoto s period from a comparably short observing interval in time. The spectroscopic observations in the present paper show the star to be a single-lined spectroscopic binary with an orbital period similar to the photometric period. In the following sections, we present and analyse continuous photometric and spectroscopic observations from the years and derive Doppler images for April 1997, April 1998, and February Additionally, we report the detection of a large Hα flare in April 1998 and discuss its relation with the photospheric surface maps. 2. Observations 2.1. Moderate-resolution spectroscopy All spectra were obtained at Kitt Peak National Observatory (KPNO) with the 0.9-m coudé feed telescope during three runs in April 1997, April 1998, and February The Ford 3k- CCD detector was employed together with grating A, camera 5, the long collimator, and a 280-µm slit to give a resolving power of 27,000 at 6500 Å, i.e Å as judged from the full width at half maximum (FWHM) of unblended Th-Ar comparison-lamp lines. The useful wavelength range was 300 Å and the exposure time was set to 4000 sec in 1997, 5400 sec in 1998, and 3600 sec in The shorter exposure time in 1999 was still sufficient because the telescope optics were cleaned and aluminized just a few nights prior to our run. The analog-to-digital units in the continuum correspond to signal-to-noise ratios of around

2 K.G. Strassmeier: Doppler imaging of stellar surface structure. XIII 609 Fig. 1a c. Photometric observations of HD a Light and color curves from All data were obtained with the Vienna Wolfgang- Amadeus APT in by (plusses) and V(I) C (dots). The horizontal bars indicate the epochs of our Doppler maps. b Window function for the period analysis and c periodogram from the combined V and y data. The best-fit period of days is indicated :1. Fourty nightly flat fields were co-added and used to remove the pixel-to-pixel variations in the stellar spectra. The Ford CCD (F3KB) shows no signs of fringing around Å and no attempts were made to correct for it other than the standard flat-field division. Continuum fitting with a very loworder polynomial was sufficient to find an excellent continuum solution Automated photometry Continuous Johnson-Cousins V(I) C and Strömgren by photometry of HD is being obtained with the two Vienna automatic photoelectric telescopes (APTs; Strassmeier et al. 1997b) at Fairborn Observatory in southern Arizona. Observations started in early 1997 and are still ongoing. All measures are obtained differentially with respect to HD (V=7 ṃ 04, B V=0.97, V I c =0.95) as the comparison star and HD as the check star. Both stars are found to be constant to within the expected standard deviation. Three measures in VI C as well as by per night were taken and added up to a total of 916 data points over the course of three years (Fig. 1). One data point consists of three 20-second (60-sec in by ) integrations on the variable, four integrations on the comparison star, two integrations on the check star, and two integrations on the sky. A 30 diaphragm was used with both APTs. The standard error of a nightly mean from the overall seasonal mean was for the Wolfgang APT 0 ṃ 0023 and for the Amadeus APT 0 ṃ 0035 in V and 0 ṃ 005 in I C. For further details we refer to Strassmeier et al. (1997a, 1997b). 3. Astrophysical parameters of HD Orbital elements We obtained 43 radial velocities from our red-wavelength spectra and find the star to be a single-lined spectroscopic binary. One additional velocity of 19: km s 1 was obtained by Osten & Saar (1998) under the assumption that the star has no composite spectrum. This velocity deviates by 15 km s 1 from our smallest value ( 4 km s 1 ) and is therefore not used in the orbit computation. A single velocity from a blue-wavelength Ca ii H&K spectrum from April 1998 (Strassmeier et al. 2000) was only used for the period search but not in the orbit determination. Initially, all velocities were given equal weight for a preliminary orbital solution. Final orbital elements were derived with the differential-correction program of Barker et al. (1967) as described in the update by Fekel et al. (1999). The solution with preliminary elements converged at an eccentricity of 0.001±0.01 so that a formal zero-eccentricity solution was adopted. The standard error of an observation of unit weight was 1.5 km s 1 but two O-C residuals were as large as 4 5 km s 1 and were given half weight in the final solution. As noted by the referee, such a deviation could be introduced by the presence of low-latitude spots and would even be of the correct amount for a spot filling factor that is representative for HD (see Saar & Donahue 1997 and Hatzes 1999). The elements are listed in Table 2 and the computed velocity curve is plotted in Fig. 2 along with the observations.

3 610 K.G. Strassmeier: Doppler imaging of stellar surface structure. XIII Table 1. Spectroscopic observing log and radial velocities HJD Phase v r σ vr O-C Reference (245+) (Eq. 1) (km s 1 ) star β Gem β Gem β Gem β Gem β Gem β Gem β Gem β Gem β Gem β Gem Vir Vir Vir β Gem Vir Vir Vir Vir Vir Vir Vir Vir β Gem Vir Vir Vir Vir Vir β Oph α Ari α Ari α Ari α Ari α Ari α Ari α Ari α Ari α Ari α Ari α Ari α Ari α Ari α Ari Throughout this paper, photometric and spectroscopic phases are always computed from a time of maximum positive radial velocity, T 0, and with the orbital period: HJD =2, 450, E (1) 3.2. The stellar rotation period Continuous photometry of a spotted star allows the determination of a precise stellar rotation period because its brightness Table 2. Orbital elements for HD Orbital element Value P (days) ± T 0 (HJD) 2,450, ±0.005 γ (km s 1 ) +35.7±0.26 K 1 (km s 1 ) 38.6±0.36 e 0.0 a 1 sin i (km) 2.055± f (m) (M ) ± Standard error of an observation of unit weight (km s 1 ) 1.5 Fig. 2. Radial velocity curve and orbit. The circles are the measurements and the line is the fit. Half weight was given to the two deviant points at phase and will be rotationally modulated by the inhomogeneous surface. This principle is being applied to solar-type stars as well as to very active RS CVn-type binaries and already led to a large set of precise stellar rotation periods (e.g. Lockwood et al. 1997, Strassmeier et al. 1997a, Cutispoto 1998). We separate our photometric data in Fig. 1 into the three consecutive observing seasons and perform a period search on each of them. The Fourier search range included a large number of frequencies up to the Nyquist frequency with a frequency spacing optimized for each data set. The frequency with the highest Fourier amplitude was adopted and a non-linear leastsquares minimization of the residuals performed which led to the best-fit photometric periods. Our period search program (Sperl 1998) is optimized for use with sinusoidal-like data and thus well suited for targets with rotational modulation. Our good time coverage furthermore allows the unique separation of the true period and its aliases, as shown in Fig. 1. The periods obtained were 3.862±0.005, ±0.0005, and ± days for 1997, 1998, and 1999, respectively. These periods are at most ±0.4% different from the orbital period which confirms that the primary star is a synchronous rotator. Seasonal deviations of this order are commonly explained with a differentially rotating stellar surface. The best-fit photometric period from the combined three-year interval is ± days and is assumed to be the rotation period of HD

4 Table 3. Stellar parameters of HD K.G. Strassmeier: Doppler imaging of stellar surface structure. XIII 611 Parameter Adopted Spectral type K0-2IV B V 0 ṃ 94 V I C 1 ṃ 03 T phot 4750±120 K log g 4.0±0.3 v sin i 41.0±0.5 km s 1 Inclination i 50 ±15 Rotation period ± days Radius R 4.1 R Luminosity L 6.4 L Mass M 1.3 M Micro turbulence ξ 2.0 km s 1 Macro turbulence ζ R = ζ T 4.0 km s 1 log[ca] abundance 0.2 dex solar log[fe] abundance 0.1 dex solar 3.3. Effective temperature, spectral type, gravity, stellar radius, luminosity, and mass No trigonometric parallax is available for HD but Cutispoto et al. (1995) estimated a distance between pc from their UBVRI photometry and suggested a K2-3V or K1-2IV+G2V classification. The average B V color is 0 ṃ 93 which leads to an effective temperature of approximately 4980 K with the transformation from Flower (1996) and Bell & Gustafsson (1989). The unspotted B V and V I C colors from our photometry in Fig. 1 are 0 ṃ 94 and 1 ṃ 03, respectively, and suggest effective temperatures within a few tens of degrees as with above transformations but give 4750 K according to the calibration for normal giants from Bessell (1979). As a further constraint, we compared several line ratios in the 6400-Å wavelength region. The particularly well-defined ratio of Fe i 6430 to Fe ii 6432 Å of 3.24±0.03 (among others) suggest a classification of K2±1 if the luminosity class is IV or III and if using the calibration from Strassmeier & Fekel (1990). Osten & Saar (1998) found a good match of the optical spectrum of HD with an inactive template star of spectral type K0IV, thus confining the range of most-likely spectral types to K0 to K2. However, an even larger range of possible spectral subclasses results if the luminosity class adopted is, say, III-IV. The corresponding effective temperature would then be 4750±160 K according to Bell & Gustafsson (1989). Various tables of absolute stellar parameters, e.g. Gray (1992), list an absolute magnitude for a K0-2IV star of +3 ṃ 2to+3 ṃ 5 and thus a distance of approximately pc for m V =8 ṃ 98. Our value for the surface gravity in Table 3 is based on the extent of the line wings of the strong Ca i 6439-Å line and on the Fe i/fe ii ratio. We computed a series of model spectra and fit them to the particularly sensitive line regions. The fit is done by a simple trial-and-error comparison. The best fit is achieved with a fixed T eff of 4750 K and log g =4.0 ± 0.3. The iron and calcium abundances are less than solar by 0.2 dex and 0.1 dex, respectively. These abundances are to be taken with care because Fig. 3. Observed and synthetic spectra for the 6430-Å region of HD The synthesized lines are identified with their element, ionization state, wavelength, and logarithmic transition probability. The thick line is the observed spectrum and the thin lines are the synthetic spectra with T eff =4750 K and log g of 3.5, 4.0, and 4.5 (log g=4.5 has the strongest lines except for Fe ii 6432 Å where it is the weaker of the three). they are just based on a 15-Å wide wavelength range. A plot of the 6430-Å line region and three representative model spectra from ATLAS-9 atmospheres are shown in Fig. 3. A similar procedure but for three different temperatures (4500, 4750, and 5000 K) and fixed log g of 4.0 confines the effective temperature to 4750±120 K (under the assumption of unaltered abundances with respect to the results in Fig. 3). Thereby, we also obtain a (preliminary) value for the projected rotational velocity, v sin i, of 41 km s 1 that is used as a starting value in the Dopplerimaging analysis. The final value will be determined later in Sect. 4.2 to 41.0 km s 1. The rotation period and the projected rotational velocity determine together the lower limit of the stellar radius to R sin i=3.13±0.04 R. With T eff =4750 K and the most likely inclination of i=50, the bolometric magnitude is M bol =+2 ṃ 3 (or M bol =+2 ṃ 6 if i=60 ) according to the solar bolometric magnitude adopted by Schmidt-Kaler (1982). Interstellar absorption was neglected in this case. The unprojected stellar radius is then 4.1 R (R=3.6 R if i=60 ). With the bolometric correction of (Flower 1996), the absolute visual brightness is M V =+2 ṃ 6 (or +2 ṃ 9ifi=60 ) and its luminosity in solar units 6.4 L (or 6.6 L if i=60 ). This absolute magnitude is still brighter by several tenths of a magnitude when compared to the expected luminosity for an early K subgiant as outlined previously in this chapter but, taking into account the uncertainties in the absence of a trigonometric parallax and the fact that HD is heavily spotted, is consistent with our log g determination of 4.0±0.3 from the 6430-Å spectra and thus being a subgiant. The position of HD relative to the evolutionary tracks of Schaller et al. (1992) for solar metallicity suggests then a mass of 1.3 M with a formal uncertainty of ±0.1 M. Table 3 summarizes the relevant stellar parameters.

5 612 K.G. Strassmeier: Doppler imaging of stellar surface structure. XIII 3.4. Lithium abundance Eighteen of our spectra in 1999 include the wavelength region around the neutral lithium line at 6708 Å and we use them to measure its average equivalent width. A double Gaussian fit to Li i Å and the nearby Fe i Å line yields an average equivalent width for the lithium line of 65±7 må (rms). Although this value is still the combined equivalent width for the close blend from the 6 Li and the 7 Li isotopes, it explicitely excludes the Fe i line ( 10 må). No phase-dependent variations are detected which suggests that the bulk of the line equivalent width comes from surface regions that do not participate in the rotational modulation, e.g. the polar region. The curves of growth from Pavlenko & Magazzú (1996) convert the 65 må to a logarithmic lithium abundance of 1.60±0.05 or 1.80±0.05 for their LTE and non-lte versions, respectively (on the regular log n(h) = scale for hydrogen). This is approximately 0.6 dex above solar. Such an abundance is not unusual for active stars, not even for evolved active stars, despite that HD should have had enough time to deplete its surface lithium (Fekel & Balachandran 1993). The newly derived space motions are (U, V, W )=( 31.4, 15.8, 6.0) km s 1 (in a right-handed coordinate system with respect to the Sun) because HD turned out to be a spectroscopic binary and supersede the velocities given in Strassmeier et al. (2000). According to the UVplane definitions of Eggen (1989), HD belongs to the old disk stars Limits on the inclination of the stellar rotation axis The mass function of ±0.0006, together with the expected mass of around 1.3 M, sets an upper limit for the inclination angle of 69 ±1. This is because a cos i must be less than R 1 + R 2 in the absense of eclipses (R 2 was assumed to be 1 R according to the G2V spectral type suggested by Cutispoto et al. 1995). The small mass function implies a lowmass secondary star with a mass of between 0.48 M for i =70 and 1.07 M for i =30. Because we do not see the secondary spectrum in our red-wavelength spectra nor in the one bluewavelength spectrum shown in Strassmeier et al. (2000), the secondary must be fainter by at least 2 ṃ 5 (no ultraviolet spectrum of HD exists). Basically, this excludes any subgiant companion as well as all main-sequence stars hotter than G0 and cooler than M1. Thus, the mass ratio primary/secondary must be in the range , which sets a lower limit to the inclination of Doppler imaging 4.1. The TempMap code All maps in this paper were generated with the Doppler-imaging code TempMap of Rice et al. (1989). For a brief description, we refer to previous papers in this series (e.g. Rice & Strassmeier 1998, Strassmeier & Rice 1998, Strassmeier et al. 1999, Strassmeier 1999, Strassmeier & Bartus 2000). In the present paper, we employ a maximum-entropy regularization for all inversions. A grid of 9 model atmospheres with temperatures between 3500 and 5750 K and log g =4.0 were taken from the ATLAS-9 CDs (Kurucz 1993). Maps in 1997 and 1998 were computed for five spectral-line regions of width 3 4 Å centered on the following main spectral lines (the numbers in parenthesis are the number of blends included and the total equivalent width in må measured): Ca i Å (4, 300), Fe i Å (8, 217), Fe i Å (8, 248), Fe i Å (6, 366), and Fe i Å (6, 283). The adopted log gf values and lower excitation potentials (in ev) are +0.47/2.53, 1.85/2.18, 0.42/3.65, 0.52/3.60, and 1.62/2.43, respectively. Maps in 1999 were computed from Ca i Å (7, 213), Fe i Å (3, 227), Fe i Å (2, 206), Ni i Å (1, 173), and Fe i Å (5, 200), with 0.85/2.71, 1.47/2.69, 2.48/2.42, 2.30/1.67, and 1.65/2.76, respectively. For each season we compute an average map from the five individual maps together with their standard deviations and compare them in order to estimate the stability of our solutions. Such a procedure ensures that inconsistencies from a particular spectral line appear suppressed in the final map while consistent detail appears emphasized (see also Beryugina et al for a similar approach) Refined rotational velocity and inclination Starting values were determined from independent methods as described in Sect. 3.3 and 3.5 and are refined in the course of the Doppler-imaging analysis. We tried inclination angles between 5 and 85 for the line profile inversion and found the largest reduction of the sum of the squares of the residuals (at maximized entropy) when inclinations between were used. The final value for v sin i is 41.0±0.5 km s 1 and for the inclination 50 ±15. The most likely secondary star is then a 0.56 M red dwarf star of spectral type M0. Note that our maps would not change significantly even if they were reconstructed with inclinations of 30 or Doppler maps for 1997, 1998, and 1999 The annual images from the individual spectral regions are shown in Fig. 4 through Fig. 6. These plots allow a comparison of the individual maps and show the consistency (and inconsistency) with which lines of different strength, excitation potential, and amount of blending can be modelled. The misfits to the observed profiles allow some judgement of the quality of the data and the model. The length of a plotbar for a profile point indicates the ±1-σ deviation based on the signal-to-noise (S/N) ratio per pixel. This is what we call the internal error and is an input parameter for our Doppler-imaging code while the more severe external errors like flat-fielding errors, spektrograph flexure during an observing run or scattered light are basically unknown. Rice & Strassmeier (2000) made extensive numerical simulations to test the impact of such external errors and found that our Doppler-imaging code is very robust against phase-unrelated inconsistencies in the line profiles, e.g. an unnoticed weak cosmic-ray hit in one or even several profiles, but

6 K.G. Strassmeier: Doppler imaging of stellar surface structure. XIII 613 Fig. 4a g. Doppler images of HD from April Each panel presents the results from one spectral region and shows a pseudo-mercator projection of the reconstructed surface temperature (top), the observed line profiles and their fits (middle), and the light curve in Johnson V (5500 Å) and Cousins I (7900 Å) and their respective fits (bottom). Rotational phase is indicated to the right of each line profile and also marked in the maps. The main mapping line is identified at the top of each map: a Ca i Å, b Fe i Å, c Fe i Å, d Fe i Å, and e Fe i Å. Panel f shows the average map and g the standard-deviation map. The individual maps were given equal weight. significant global artifacts can be introduced once the external errors vary (pseudo)periodically in size and in time like, e.g., a flat-field artifact due to a hot pixel of changing sensitivity. The latter is not a significant problem in case of a binary like HD because the spectral line shifts due to the orbital motion. Nevertheless, some less-than-perfect fits of otherwise high-s/n data are attributed to these external errors (for a more detailed discussion on external errors, we refer to previous papers in this series, e.g. Strassmeier & Bartus 2000). The photometric data are treated as if they had infinite S/N ratio. During the first rounds of line-profile fits the VI photometry is included without its zeropoints, i.e. only the shape of the light curve in the two bandpasses is used to constrain the map. The final rounds are then computed with the zeropoints added. At this point, we note that the inclusion of the zeropoints does not significantly change the average temperature level because the temperature information primarily comes from the line-profile equivalent widths and from using lines of different excitation potentials and depend only to a lesser degree on

7 614 K.G. Strassmeier: Doppler imaging of stellar surface structure. XIII Fig. 5a g. Doppler images from April Otherwise as in Fig. 4. the photometry. The surface geometry recovered remains the same with or without zeropoints but the advantage of having a color zeropoint, mostly V I C in our case, is that it is sensitive to metallicity, i.e. the chemical abundance of the main mapping line s chemical spezies. The latter is obtained from a series of inversions with a large range of trial abundances and is fixed for the final solutions. The fits and the observations in each of the V and I C bandpass are compared in the respective lower panels in Figs April 1997 For 1997, 9 spectra and 40 VI C light-curve points are available. The photometry was combined from a 20-day interval centered around the mid time of the spectroscopic observations which covered a total of 14 consecutive nights. Fig. 4a f shows the data and the results. Four of the five mapping lines recover a cool polar cap-like spot with an average temperature difference to the unspotted photosphere of T = 1000±80 K (rms). The most blended of the five lines (Fe i 6411 Å, W λ = 248 ± 10 må) recovers a circumpolar band up to a latitude of Based on the overall quality of the fit, we can not decide whether this is a real physical effect in the stellar atmosphere, e.g. a height dispersion of the magnetic filling factor or an artifact due to several small errors in the atomic data for the line blends. Since only a small region directly at the rotation pole is affected, we tend to believe that the effect is not real. The average map emphasizes two moderate-latitude spots at a longitude of l 100 (latitude +40 ) and l 210 (latitude

8 K.G. Strassmeier: Doppler imaging of stellar surface structure. XIII 615 Fig. 6a g. Doppler images from February The spectral lines mapped are now different than in the other seasons. a Ca i 6717 Å, b Fe i 6677 Å, c Fe i 6663 Å, d Ni i 6643 Å, and e Fe i 6546 Å. Panel f is again the average map and g the standard-deviation map. +50 ), respectively that appear at the same positions than the two polar-spot appendages in the following 1998 map and probably are the precursors of that features. The spot at +40 is also the main cause for the light curve minimum at phase (l 90 ). The coolest feature in the 1997 image is the part of the polar spot between longitude with T 1100 K and coincides with the light-curve maximum. Against a common assumption in light-curve spot modelling it shows that the brightest hemisphere is not necessarily the least spotted April 1998 For 1998, altogether 18 spectra and 40 VI C light-curve points are available. The photometry was combined from the same 20 consecutive nights as the spectroscopic observations. Fig. 5a f shows the data and the maps. The 1998 maps are again dominated by a large and asymmetric polar spot of T = 1100±50 K (rms). The longitudinal location of the most asymmetric part is 120 and reaches down to a latitude of almost +30, thus spanning a significant part of the visible hemisphere. A second appendage is evident at a longitude of 210 and down to +45 that appears connected to the larger appendage. Except maybe for the Fe i 6411-Å map, the two features are well separated at their lower rims but appear merged at latitudes closer to the pole. These two appendages are also the cause of the single but asymmetric light-curve minimum. As we will discuss later in Sect. 5.2, a large Hα flare coincided in time with the faintest part of the light-curve mini-

9 616 K.G. Strassmeier: Doppler imaging of stellar surface structure. XIII Fig. 7a g. Annual average maps in spherical projection and their temperature profiles. The maps in the top row are for phase 90, the lower maps for 270. The plot to the right of each map is the latitudinal temperature profile (repeated in the lower maps), the plot below a map is the longitudinal profile of the map shown. The dotted and dashed lines are the effective and average temperatures, respectively. The grey tones represent temperature according to the scale in Figs mum and we suspect that it was related with this large polar spot at a longitude of l 120. A third appendage, though weaker, is evident at l 300. A set of two equatorial features each consisting of three separated and latitudinally elongated spots appear in the average map. These group of features is reconstructed consistently from Ca i 6439 and Fe i 6430 and partly also from the other three lines. Although we are usually sceptical about fine details near the equator because of less favorable mapping conditions at low latitudes but, since these features repeat in several lines and the phase coverage is very good, we have no reason not to believe in their reality February 1999 For 1999, 15 spectra and 40 VI C light-curve points are available except for the Fe i 6677-Å and Fe i 6546-Å lines where one profile each was severely affected by cosmic-ray hits and were excluded from the analysis. The photometry was combined from the same 17 consecutive nights as the spectroscopic observations. Fig. 6a f plots again the data, the fits, and the maps. All maps recover a large polar-spot appendage at l 135 and at an average latitude of +70. Its appearance is very similar to and at the same location as in 1998 and we strongly believe that it is the same feature as seen one year earlier. Its approximate size of 4.5% of the total sphere also remained constant. A significant change, however, occured at equatorial latitudes where a spot at l 315 appeared and caused a second lightcurve minimum. Fig. 1a nicely demonstrates that the change was a gradual process from a 0 ṃ 3-amplitude single-humped light curve in 1998 to a 0 ṃ 06-amplitude double-humped minimum in The third polar appendage seen in 1998 at l 300 appears now as a separate high-latitude spot at l 295 and b +55. Finally, Fig. 7 compares the annual average maps in spherical projection. It also shows the temperature profiles from the individual maps (thin lines) and the average maps (thick lines) along binned longitudes and latitudes. All maps indicate a general temperature gradient from the pole to the lower rim at 50 from 3900 K at the pole to the photospheric temperature of 4750 K at the equator. We emphasize though that the regions below the equator are not well defined and no reliable information

10 K.G. Strassmeier: Doppler imaging of stellar surface structure. XIII 617 Fig. 8. Observed Hα-line profiles as a function of phase for a February 1999, b April 1998, and c April Phase is indicated in each subpanel. Notice the flare in 1998 marked with t 0. The spectrum on the night prior to this flare is marked t 1, the two spectra in the following two nights are marked t +1 and t +2, respectively. The dotted vertical line is the Hα rest wavelength with respect to the photospheric lines. The tick marks on the wavelength axes are in steps of 5 Å and indicate a range of 30 Å in each panel. at all is obtained for latitudes below, say, 30. This is simply because the time a surface feature is in view gets increasingly shorter the lower the latitude and could, with the given phase coverage, be missed entirely. Therefore, we would not expect to recover the mirrored temperature gradient in the opposite hemisphere, despite that the inclination of the stellar rotation axis is just 50. We interpret the temperature gradient thus to be a real gradient from the pole to the equator and primarily due to a distinct polar spot. 5. Hα line variations 5.1. Chromospheric inhomogeneities Fig. 8 shows the Hα profile variations as a function of orbital phase for the three observing epochs in 1997, 1998, and The Hα profile appears quite complex, ranging from extremely strong emission due to flaring (phase 0. p 37 in 1998), emission with a central sharp absorption (phase 0. p 21 in 1998), partial emission with a red-shifted absorption component (phase 0. p 82 in 1999) to pure but very narrow absorption (phase 0. p 86 in 1998). Such a behavior is not unexpected for an active RS CVntype star. As magnetic activity increases, the Hα-line depth first increases and then fills in and turns into pure emission. The latter is observed only in the most active stars and is in agreement with theory (e.g. Cram & Mullan 1979, Houdebine et al. 1995). Additionally, surface plages will add the asymmetry to the profile shape and also cause the rotationally modulated line variations as observed, e.g., for the very active RS CVn binary V711 Tau (=HR 1099, Buzasi et al. 1991). However, chromospheric dynamics as well as the presence of circumstellar material may dilute or even suppress the rotationally modulated plage signature. Evidence for global mass inflow due to prominences and/or Hα flaring was detected in several active late-type stars, e.g. for BD (Eibe et al. 1999), GK Hya and TY Pyx (Gunn et al. 1997), AB Dor (Collier Cameron & Robinson 1989), and the many RS CVn binaries observed by the PennState group (e.g. Hall & Ramsey 1992).

11 618 K.G. Strassmeier: Doppler imaging of stellar surface structure. XIII We tried to quantify the line variations in Fig. 8 by fitting Gaussians to the emission and absorption component of the individual Hα profiles. The temporal evolution of the fit parameters is shown in Fig. 9a d. Panel a plots the velocity of the emissionline center, v c, after subtraction of a reference spectrum of β Gem (for 1997 and 1998) and α Ari (for 1999). Panel b plots the full width at continuum, FWAC, panel c shows the centroid velocity of the sharp absorption line, v abs, and panel d plots the velocity of the blue wing of the emission line at continuum level. The latter indicates a more or less constant value of 100 km s 1 over the course of the three years just interrupted by the large chromospheric outflow velocity during flaring. In 1998, the velocities from the sharp absorption line show some evidence of a sinusoidal variation with an amplitude of 20 km s 1 and a maximum at phase 0. p 5 that sometimes still exceed the ±v sin i surface limit (the dotted lines in Fig. 9c). In the years 1997 and 1999 no clear trend is obvious. It is likely that the absorption is due to some variable hydrogen absorbing region in the stellar corona driven by an inward pointed flow pattern. Such an inhomogeneous antiwind in Hα has been detected for the K0III- IV binary component of HU Virginis (Strassmeier 1994, Hatzes 1998) or for the single K5 dwarf BD (Eibe et al. 1999) and seems to be a common feature in the outer atmospheres of very active stars. A purely circumstellar, or even circumbinary, nature of this absorption line seems possible but is unlikely to be explained with the transit of cool co-rotating prominences across the projected stellar disk. For that, the radial-velocity dispersion is too small, say by a few hundred per cent when compared to the prominences on AB Dor (Collier Cameron & Robinson 1989), and the absorption observed for HD appears rather stable with a net redshift most of the time. The chromospheric velocity signatures in Fig. 9a,b, and c indicate variations that are mostly anticorrelated with the orbital velocity variations from photospheric lines. This is most obvious for the April 1998 data set (the filled dots in Fig. 9). E.g., the full width of the emission line at continuum level had a maximum at the time around the maximum negative photospheric velocity, i.e. at one of the two times of quadrature. Because the inner Lagrangian point is seen at both quadratures and the behavior does not repeat with 1/2 P orb, we may exclude it as the location of the Hα emission. Fig. 9b also shows that the bulk of the broad emission lines with FWAC levels in excess of 150 km s 1 appeared after the occasion of the large flare at phase 0. p 37 (see Sect. 5.2). This, in turn, suggests that all subsequent Hα profiles with large FWAC s were affected by that flare. Another similar, but less dramatic event may have occurred at phase 0. p p 06 in We adopted the calibration from Hall (1996) to convert the relative flux in the central 1-Å band of the line core into an absolute surface flux. This flux is plotted in Fig. 10 versus rotational phase and for the years 1997, 1998, and A clear phase-dependent variation is only seen in the 1998 data. From this we conclude that a significant part of the Hα emission must be due to chromospheric plages, otherwise no rotational modulation of the flux would be expected. This is confirmed by the close analogy to the photospheric light curve. It showed a full Fig. 9a d. Hα-line parameters for 1999 (plusses), 1998 (filled circles), and 1997 (open circles). a The radial velocity of the emission profile. b Its full width at continuum level (FWAC). c The radial velocity of the central sharp absorption line. d The radial velocity of the blue wing of the emission profile at continuum level. The line in panels a and c is the photospheric orbital velocity and the two dotted lines indicate the range of ±v sin i. The small arrow highlights the value during flare peak in the 1998 set. Fig. 10. Hα line-flux variation in April 1998 (filled circles). The line is an overplot of the contemporaneous V -light curve in 1998 and shows the anti-correlation of the continuum variation with the Hα flux. Note that the flare occurred exactly during light-curve minimum. The Hα fluxes during the observations in 1999 (plusses) and 1997 (open circles) remained mostly constant. amplitude of 0 ṃ 3inV in 1998, thus twice as large as in 1997 and even three times as large as in Its shape anticorrelates with the Hα flux which suggests that the bright plages are spatially related to the cool spots. That plages overlay spots is a common solar phenomenon and is also seen in many active RS CVn stars, e.g. HU Vir (K0III-IV; Strassmeier 1994), CM Cam (G8II-III, Strassmeier et al. 1998), IL Hya (K0III-IV, Weber & Strassmeier 1998), and many more stars summarized by Catalano et al. (1996) The large Hα flare at JD 2,450,909.6 Fig. 8 also shows the Hα profile during the large flare in April 1998 (marked with t 0 ). The Hα equivalent width during the flare was Å, i.e. about 70 times larger than the day before (that spectrum is marked with t 1 in Fig. 8) and still 10 times larger than the day after the flare (marked with t +1 ). When plotted versus phase as in Fig. 8 there appears an increase of the full width of the emission as well as of the equivalent width

12 K.G. Strassmeier: Doppler imaging of stellar surface structure. XIII 619 at the phases immediately following the flare. Notice that the equivalent width and the profile shape were apparently back to normal two days after the event as shown by the spectrum marked with t +2. It then appeared very similar to the profile observed one day prior to the flare. The emission peak during the flare is very sharp and is redshifted by +21±3kms 1. The entire profile appears asymmetric having a base velocity at continuum level of 326 km s 1 in the blue wing and +428 km s 1 in the red wing (its FWAC is offscale in Fig. 9b and is indicated by an arrow). A similar red excess was seen in a large Hα flare on the RS CVn binary HR 1099 (Foing et al. 1994) and was interpreted as the infalling loop structure after flare peak. The photospheric absorption lines in the flare spectrum seem to be entirely unaffected by the flare. This indicates that the location of the flare must have been very high up in the stellar chromosphere and corona that left the photosphere practically unaffected. Recent flare models for X-ray light curves of active stars suggest loop heights of the order of one stellar radius or even higher (e.g. Kürster & Schmitt 1996, Endl et al. 1997), and are in agreement with the qualitative behavior of HD The escape velocity at the surface of HD is 350 km s 1 and only 250 km s 1 one stellar radius above the surface. We may thus expect that the star looses some mass due to flaring. The total energy released by the flare is estimated by integrating the absolute flux over the duration of the flare. The non-flaring Hα flux at the time of the flare is interpolated from Fig. 10 and subtracted from the peak emission. For the flare duration we have only an upper limit of approximately 1 2 days ( 0. p p 5) but no information on the decay flux curve. Assuming a one-day duration at constant absolute flux, we obtain a crude estimate of the total flare energy released in Hα of erg. This compares well with the erg released byahα flare on HR 1099 (Foing et al. 1994) but falls short by two orders of magnitude when compared to a long-duration Hα flare on HK Lac (Catalano & Frasca 1994) Does the flare relate to the Doppler image? The time of the flare coincided with a time when the dominating, high-latitude spot seen in the maps in Fig. 5 crossed the central meridian. Although possibly coincidental, we suggest that a spatial correlation exists and that the most likely position of the flare is related to this high-latitude spot at a latitude of around +60 and a longitude of 130 (0. p 36). However, the absence of eclipses in the HD system does not allow to determine more direct constraints on the actual flare position. 6. Summary and conclusions We present Doppler images of HD from three observing intervals in 1997, 1998, and All images are characterized by large and cool high-latitude spots that appear to be mostly appendages of a polar cap-like feature. Several low-latitude spots are present at the same time and such a bimodal spot distribution in latitude still remains a challenge for the flux-tube evolution model of Schüssler et al. (1996). This is because the tubes are diverted to higher latitudes the more rapidly the star rotates and the deeper its convective envelope. The bimodality of starspots, i.e. flux tubes that are emerging predominantly at high and low latitudes but not in the middle, was proposed in the numerical T Tauri star models by Granzer et al. (2000a) due to the existence of a mid-latitude stability island of the field strengths in the overshoot region. We computed a flux-tube model following the procedure described in Granzer et al. (2000a) for the stellar parameters of HD and find that no flux tubes emerge near the stellar surface for latitudes below approximately +30 and above Yet our Doppler images recover many low-latitude and even equatorial features together with a polar spot (or at least very high-latitude spots). A similar morphology was observed on other spotted stars in various evolutionary stages (a summary was given in Granzer et al. 2000b, but see also previous papers in this series). It seems that the latitudinal spot bimodality is now observationally well established. Vogt & Hatzes (1996) and Vogt et al. (1999) discovered a poleward migration of spots on the RS CVn binary HR 1099 and Strassmeier & Bartus (2000) were able to quantify their latitudinal migration rate to /day. If interpreted due to meridional circulation, we may expect that magnetic flux tubes are starting to significantly float towards the pole once they got close to the surface where the numerical simulations must be terminated because the thin flux-tube approximation breaks down (i.e. approximately in the outer 2% of the stellar radius). Because it is likely that magnetic reconnection occurs very frequently and all the time (see, e.g., Thomas & Montesinos 1997), not all flux tubes will float much further in latitude from their points of emergence but on average the flux tubes should make it up to the polar regions in approximately 100 days (i.e. in a full observing season) to create the polar spot and its appendages. On the Sun, sunspots migrate to lower and higher latitudes and can sometimes even cross the equator (e.g. Howard & Gilman 1986). Such a bimodal meridional migration may also exist on very active late-type stars like HD and would force lowto-mid latitude spots to migrate to the equator while the high latitude spots migrate toward the poles. Such a scenario could be verified with the technique of Doppler imaging if the time resolution between consecutive maps could be significantly increased. The evolution of the two 4200-K cool high-latitude features in 1997 into two 3700-K polar appendages in 1998 suggests that significant changes take place after flux-tube emergence. The temperature decrease of 500 K is likely explained by a mass exchange along closed field lines of different polarity and from different active regions. The cool polar region is presumably filled with magnetic field and suggests that reconnections occur more frequently near the pole than closer to the equator, as indirectly observed on HD During the 1998 observations a large flare was detected in Hα. It coincided with a time when the polar-spot appendage crossed the central meridian and we believe that the flare was spatially related to this cool photospheric spot. If interpreted

13 620 K.G. Strassmeier: Doppler imaging of stellar surface structure. XIII due to violent magnetic reconnections as in a two-ribbon flare (Martens & Kuin 1989), its pure existence lends further evidence to our explanation of the latitudinal spot bimodality with a more frequent reconnections at higher latitudes than near the equator. The flare released a total energy of around erg and was thus very similar to a well-observed flare on the bright RS CVn binary HR 1099 (Foing et al. 1994). The (optically thin) photospheric absorption lines showed no sign of a general brightening of the flare hemisphere or the appearance of a bright spot and suggest that the flare electrons were restricted to the overlaying stellar chromosphere and corona. Acknowledgements. I am grateful to the Austrian Science Foundation (FWF) for their support through grant S7301-AST for the operation of the two APTs and S7302-AST for Doppler imaging. 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