A&A 373, 974-986 (2001)
DOI: 10.1051/0004-6361:20010580
Astrophysikalisches Institut Potsdam, An der Sternwarte 16, 14482 Potsdam, Germany
Received 21 February 2001 / Accepted 11 April 2001
Abstract
Time-resolved Doppler images of the rapidly rotating, but long-period
(25 days),
giant KUPegasi show several cool low-to-medium latitude spots as well as an
asymmetric polar feature. The average spot temperature is about 700K
below the photospheric temperature of 4700K. KUPeg is one of
the most massive, and currently the most evolved, late-type star
with a Doppler image. We obtained two independent
images from two consecutive stellar rotations covering 50 nights with a total
of 43 spectra.
From a cross-correlation analysis of the two maps, we detect systematic
longitudinal and latitudinal shifts that we tentatively interpret as latitude-dependent
differential rotation and local meridional flows, respectively.
The differential-rotation pattern is more complex than on the Sun, but on average
in the sense that the poles rotate slower than the stellar equator,
i.e. in the same direction and also of the same order than on the Sun.
The latitudinal shifts are of the order of 0.4day-1 towards
the stellar pole and occur at longitudes of around 40
and 330
.
The residual H
profiles show a stationary emission component at rest
wavelength and a blue-shifted absorption. The latter suggests an
outward pointed velocity field with a flow velocity of approximately 35 kms-1.
Key words: stars: activity - stars: atmospheres - stars: imaging - stars: individual: HD 218153 - stars: late-type
HD218153 = KUPeg (
,
,
2000.0, V = 7
61-7
72)
is a G9-K0 giant with strong
Ca IIH&K emission (Bidelman 1983),
a projected rotational velocity of
kms-1 (Fekel 1997),
and a photometric period of approximately 24 days
(Strassmeier & Hall 1988) which
was interpreted to be the stellar rotation period.
Despite its long rotational period the star
ranks near the top in ultraviolet emission-line strength compared to
other late-type giants (De Medeiros et al. 1992)
(only FKComae is significantly more active).
Its optical spectrum shows only a weak lithium feature in the 6707-Å range - Barrado y Navascues et al. (1998) report 16 mÅ and
Strassmeier et al. (2000) 13 mÅ, which suggests that the star is
in the expanding post-main-sequence evolutionary stage and not a contracting
pre-main-sequence star.
De Medeiros et al. (1992) detected small radial-velocity
variations and reported a preliminary spectroscopic orbit with a period of
1411 days (4 years). Such a long orbital period implies that the
star had evolved effectively like a single star similar to HD 51066 (=BM Cam;
Strassmeier et al. 1998, Paper VIII in this series) and thus may
be compared with a group of truly single, rapidly-rotating giants.
Such a group of evolved active stars was first identified by F. C. Fekel
and collaborators (see e.g. Fekel & Balachandran 1993) from
spectra with unusually strong lithium lines at 6708 Å, despite
their evolutionary status. At an age of several hundred Megayears one would
expect that the mixing in the outer convective envelope was thorough enough
so that most lithium was burned when it was getting to close to the bottom
of the convective envelope. How these rapidly-rotating giants maintained
a primordial level of lithium on their surface still remains unanswered.
Since KU Peg does not show significant lithium on its surface,
actually
,
but otherwise is a rapidly-rotating
extremely active giant, it may represent a border case
for stellar convection and light-element dredge-up models.
A brief discussion of the star's photometric history and our APT
(automatic photoelectric telescope) data prior
to 1996 was given by Strassmeier et al. (1999a).
The relatively small light and color amplitudes of
and
in 1996-98
suggested that either only a small fraction of the surface was covered
with spots or that the spot distribution was relatively symmetric.
Our line profile data show the typical bump-signatures from
cool spots and thus verify the existence of an asymmetric spot component
at that time.
In this paper, we present a series of 43 moderately high-resolution spectra
in the 6430-Å region taken over a two month period in late 1996 through early 1997.
photometry was gathered continuously from fall 1996 until now.
The combination of these data is used to study
the spatial surface distribution of the spots on KU Peg and its short-term
variations. The spectroscopic observations for Doppler imaging span a little
more than two stellar rotations and are used to derive two independent
surface maps. The instrumentation and the data reductions are described in
Sect. 2 while the stellar properties relevant for the
Doppler-imaging analysis are determined in Sect. 3. H
spectra from 1998 are also presented in this section. The Doppler maps from
three spectral regions and for the two stellar rotations are derived
in Sect. 4. Finally, we derive a latitude-dependent differential
rotation law from a cross correlation of the two consecutive maps and discuss
its relation with the Sun and other active stars.
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Figure 1: Three years of photometry of KU Peg. a) The 1996/97 APT data versus Julian date and versus phase calculated with the ephemeris in Eq. (1), b) and c) the same for the 1997/98 and 1998/99 seasons, and d) the periodogram for the combined 1996-99 APT data. |
HJD | Phase |
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![]() |
Rotation | Tele- |
(24+) | (![]() |
(kms-1) | cycle | scope | |
50391.770 | 90.4 | -77.5 | 1.5 | map 1 | NSO |
50392.661 | 103.3 | -82.0 | 2.7 | - | NSO |
50393.748 | 119.0 | -75.8 | 1.9 | map 1 | NSO |
50394.781 | 133.9 | -78.5 | 1.8 | map 1 | NSO |
50395.602 | 145.7 | -80.5 | 2.0 | map 1 | NSO |
50396.648 | 160.8 | -78.5 | 2.0 | map 1 | NSO |
50398.713 | 190.4 | -82.4 | 2.3 | - | NSO |
50400.629 | 218.2 | -81.1 | 1.6 | map 1 | NSO |
50401.630 | 232.6 | -79.6 | 2.4 | - | NSO |
50404.641 | 276.1 | -83.2 | 2.0 | map 1 | NSO |
50405.617 | 290.1 | -81.8 | 2.5 | map 1 | NSO |
50406.623 | 304.7 | -80.9 | 1.9 | map 1 | NSO |
50408.624 | 333.5 | -77.3 | 3.0 | map 1 | NSO |
50409.617 | 347.8 | -82.0 | 2.3 | map 1 | NSO |
50411.618 | 16.7 | -79.9 | 1.7 | map 1 | NSO |
50412.611 | 31.0 | -80.9 | 1.8 | map 1 | NSO |
50413.609 | 45.4 | -81.2 | 1.7 | map 1 | NSO |
50415.610 | 74.3 | -80.1 | 1.7 | map 1 | NSO |
50416.623 | 88.9 | -80.5 | 1.5 | map 1 | NSO |
50417.710 | 104.4 | -78.9 | 2.3 | - | NSO |
50418.683 | 118.6 | -80.4 | 1.2 | map 2 | NSO |
50419.597 | 131.8 | -80.2 | 1.6 | map 2 | NSO |
50420.607 | 146.4 | -80.3 | 1.8 | map 2 | NSO |
50421.690 | 162.0 | -79.4 | 2.1 | map 2 | NSO |
50422.602 | 175.1 | -79.6 | 1.3 | map 2 | NSO |
50423.609 | 189.7 | -80.6 | 1.5 | - | NSO |
50424.701 | 205.4 | -82.8 | 1.3 | map 2 | NSO |
50425.674 | 219.4 | -79.3 | 1.7 | map 2 | NSO |
50426.693 | 234.1 | -80.2 | 2.2 | map 2 | NSO |
50429.661 | 276.9 | -79.4 | 2.1 | map 2 | NSO |
50430.682 | 291.7 | -79.6 | 1.5 | map 2 | NSO |
50431.569 | 304.5 | -81.7 | 2.0 | map 2 | NSO |
50432.603 | 319.4 | -80.8 | 1.6 | map 2 | NSO |
50433.593 | 333.7 | -81.5 | 1.7 | - | NSO |
50434.599 | 348.2 | -81.4 | 2.1 | map 2 | NSO |
50436.645 | 17.6 | -80.9 | 1.7 | - | NSO |
50438.592 | 45.7 | -80.4 | 2.0 | map 2 | NSO |
50440.618 | 75.0 | -79.8 | 2.3 | map 2 | NSO |
50441.604 | 89.3 | -80.0 | 2.4 | - | NSO |
50445.594 | 146.9 | -80.0 | 2.2 | - | NSO |
50446.595 | 161.3 | -81.3 | 2.5 | - | NSO |
50450.682 | 220.0 | -81.3 | 1.9 | - | NSO |
50457.625 | 320.4 | -81.5 | 1.8 | - | NSO |
50821.638 | 170.3 | -78.7 | 1.1 | - | KPNO |
50826.670 | 243.0 | -80.0 | 1.0 | - | KPNO |
50828.628 | 271.4 | -78.5 | 1.6 | - | KPNO |
50918.011 | 120.6 | -79.0 | 1.0 | - | KPNO |
50921.003 | 163.4 | -80.3 | 0.7 | - | KPNO |
Spectroscopic observations were obtained at the National Solar Observatory
(NSO) with the McMath-Pierce main telescope using the stellar spectrograph
from October 31, 1996 to January 8, 1997. The 800800-pixel TI CCD
(TI-4 chip, 15
pixels) allowed for a resolving power of
40000
and a useful wavelength range of about 45 Å around 6430 Å.
Two additional spectra in the 6430-Å region plus one spectrum in the
H
region, along with corresponding reference spectra, were obtained
at Kitt Peak National Observatory (KPNO) with the coudé feed telescope in
December 1997/January 1998. A very similar 800
800 TI CCD (TI-5 chip,
15
pixels) was used in combination with grating A, camera 5, and the
long collimator. It allowed for a resolving power of 38000 and 22000 in the red
and the blue-wavelength regions, respectively, with a useful wavelength range
of between 50-80 Å. Two more spectra in the 6500-Å region included
H
and were obtained with the same instrumental set-up at the
coudé feed telescope in April 1998 but utilized the 3096
1024 CCD
(F3KB chip, 15
pixels). The resolving power was 28000, and the useful
wavelength range was 300 Å.
Table 1 is a summary of the spectroscopic observations.
Spectra taken after the second rotational cycle in early 1997 and spectra
with signal-to-noise ratios below 100:1 are not used for Doppler imaging
(these spectra are marked with a dash in Table 1).
All data were reduced using IRAF and included bias subtraction, flat
fielding, cosmic-ray removal, and an optimal aperture extraction. The
exposure level obtained with an integration time of 30 min corresponds
to a signal-to-noise ratio of approximately 150:1.
Usually twenty flat-field exposures with a Tungsten reference lamp were
taken at the beginning of the night and again at the end of the night. These
forty flat fields were co-added and used to remove the pixel-to-pixel
variations in the stellar spectra. Spectra of bright radial-velocity
standards were obtained several times throughout the night to ensure an
accurate wavelength calibration. Radial velocities for the 1996/97 NSO data
and the two KPNO spectra from JD 2450826 and JD 2450828 were derived
from cross correlations with spectra of the IAU velocity standard
Ari (K2III,
kms-1). The one spectrum from
JD 2450821 was cross correlated with
Gem (K0III,
kms-1). 16Vir (K0.5III,
kms-1) was used for the
remaining spectra.
The photometric data in this paper were obtained with the Amadeus 0.75-m
automatic photoelectric telescope (APT), part of the University of Vienna
twin APT at Washington Camp in southern Arizona (Strassmeier et al. 1997).
The observations were made differentially with respect to HD218610 as the
comparison star (
,
,
)
and
SAO91066 as the check star. The data cover the time interval
JD 2450395-459 (54 data points, season 1996/97),
JD 2450740-828 (80 data points, season 1997/98), and
JD 2450960-998 (47 data points, season 1998/99).
All photometry was transformed to match the Johnson-Cousins
system.
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Figure 2: A comparison of a spectrum of KUPeg (thick line) with a shifted and spun-up spectrum of the G9II-III standard star HD76294 (thin line). Identified are the spectral lines that are used in the Doppler-imaging analysis in Sect. 4 (a bold-face wavelength distinguishes the main mapping line from the blends included). |
We applied the multiple period search program PERIOD98 (Sperl 1998) in
single-period mode and use its Fourier option to determine the photometric
period of KUPeg (Fig. 1d). The application of this program to
photometry of spotted stars was recently discussed in Strassmeier et al. (1999a)
and we refer the reader to this paper. The fit with the smallest
was
obtained with a period of
for the 1996/97 data,
in 1997/98, and
for the 1998/99 data
(Figs. 1a-c).
By early 1997/98 the amplitude of the V-lightcurve had dropped from
0
07 in late 1996 to
0
03 in late 1997, and had risen again to
0
06 in 1998.
The Fourier analysis of the 1997/98 V data still shows a peak near the same
frequency as in 1996/97 but an alias at f=0.1 c/d appears with an even
stronger amplitude than the true photometric period.
Such a reduced rotational-modulation signal
suggests that the active regions that were visible in
1996/97 had significantly decreased in size and/or contrast by late 1997.
The lightcurves of the two seasons
(Figs. 1a and b) consistently show a minimum at around
phase 0.5 but the phased 1997/98-season data appear as a double-humped light
curve.
The final period is obtained from the combined data set and is of higher
precision due to the longer time range. We obtain
,
which we interpret to be the stellar rotation period and use it
to calculate the phases for all spectroscopic and photometric data
in this paper:
The effective temperature of KUPeg is estimated in several ways, first
from a comparison of the observed
color with synthetic colors
from the ATLAS-9 model atmospheres (Kurucz 1993). The bluest
value of
is observed at times of lightcurve maximum (Fig. 1) and is
adopted to best represent the "unspotted'' intrinsic color of the
photosphere. This value is different to the long-term average due to the
color variations from spots of up 0
03. Comparing this V-I value with the
grid of synthetic ATLAS-9 colors in the range
to 3.0 and metal
abundances of solar to 1.0 dex below solar, we find a temperature
of
K. This compares very well with the
4700K from the observed B-V of 1
13 and the calibration
for a bright giant from Bell & Gustafsson (1989).
Similar temperatures are obtained from the color-temperature
calibration of Flower (1996) (4622K), as well as
from the combined G9II, G9II-III and G9III temperatures listed in
Dyck et al. (1998) (
K), for a G9III star in
van Belle et al. (1999) (4679 K), and for a K0 giant
in Bell & Gustafsson (1989) (4820 K).
Another method for estimating the effective temperature is to use the temperature
dependence of spectral lines.
Using the calibration for giant stars from spectral line-depth ratios
from Strassmeier & Schordan (2000), we calculated the V-I colors and
the effective temperatures for 11 line-ratios. The resulting average effective temperature
was K, in good agreement with the values derived with other methods.
The average V-I-color of
is 0
11 smaller than the observed value,
which shows that the large interstellar absorption correction we apply in the next
section is indeed of the right order of magnitude.
In order to compare the spectrum of KUPeg numerically with spectra of
Morgan-Keenan standard stars, we make use of a computer program originally
designed by Barden (1985).
The standard-star spectra are Fourier transformed and
subtracted from a representative KUPeg spectrum. The resulting
difference spectra are iteratively minimized by changing the relative
continuum intensity, the rotational broadening, and the radial velocity
of the KU-Peg spectrum by means of a weighted
minimization.
With this procedure, we found the best fit with a spectral type of G9 to
K0III and a (preliminary)
rotational velocity of
kms-1. We
note that both standards,
Gem (K0III) and HD76294 (G9II-III),
fitted our KUPeg spectra equally well (the latter fit is shown in
Fig. 2).
To obtain an independent estimate for the luminosity class of KUPeg, we
use the luminosity-sensitive Sr II 4077-Å line and compare
KUPeg with Ari (K2III),
Gem (K0III),
Cnc (G8II), and 55Cam (G8II). Again, the program of
Barden (1985) was used to minimize the residual spectra.
Because the Sr II 4077 line of KUPeg is weaker than in the two
available G8II spectra, but only slightly stronger than in the
Gem
spectrum, we estimate KUPeg to have a luminosity class of II-III
and therefore a most likely spectral classification of G9-K0 II-III.
The rotational period and the projected rotational velocity determine the
minimum stellar radius to be
.
Using an
inclination angle of
- that will be derived later in
Sect. 4.2 - the most likely stellar radius of KU Peg is
18
.
This is not in agreement with the
for an average K0III
star according to van Belle et al. (1999), and just barely in
agreement with the
according to the Landolt-Börnstein tables (Schmidt-Kaler 1982)
but more likely indicates that the star is larger than a regular K0 giant.
Using above radius and an effective temperature of
K,
the bolometric magnitude of KU Peg is
.
The Hipparcos parallax of
milli-
(ESA 1997)
puts KUPeg in a distance of 188
+38-28 pc and, with the brightest
V magnitude observed so far, i.e.
in 1997 (see Fig. 1c),
results in an absolute visual brightness of 0
82
-0.40+0.35.
An interstellar extinction correction of
had been applied, which results
from
,
where E(B-V) is 0
13 using
(Guetter 1980;
Fernie 1983) and
(Schmidt-Kaler 1982).
This agrees with the
listed in Guarinos (1992).
With a bolometric correction of -0.48 (Flower 1996)
the bolometric magnitude is
,
which is
different by one magnitude to the above value from the combination of the minimum stellar
radius, the inclination from the Doppler-imaging analysis,
and the effective temperature from the spectral classification.
Adopting an average bolometric magnitude and a solar value of +4
64,
the luminosity of KUPeg is
.
A comparison of the position of KUPeg in the H-R diagram with the
evolutionary tracks of Schaller et al. (1992) interpolated
for 0.6 dex below solar metallicity (see Sect. 4.3), suggests a mass
of
and an age of approximately 870 Myrs.
De Medeiros et al. (1992) reported the detection of radial
velocity variations with a full amplitude of 6 kms-1 and calculated a preliminary orbit with the following
elements:
days,
,
,
kms-1,
kms-1,
and
.
The average from our radial velocity measurements from HJD 2450392-457 is
-80.3 kms-1, and agrees very well with the -velocity determined by
De Medeiros et al. (1992). Its root mean square is 1.8 kms-1 and
the average error of our individual velocities is also 1.8 kms-1. This values fit
the published orbit when a period of 1389 days is used (which is inside the
error bounds given by De Medeiros et al. 1992). The phase
coverage of our spectra, however, does not allow an independent period determination
but verifies the preliminary orbit.
Parameter | Value |
Spectral type | G9-K0II-III |
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7
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B-V |
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Rotation period
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Inclination i |
![]() ![]() |
Radius R | 18
![]() |
Equatorial velocity
![]() |
36.8 kms-1 |
Luminosity L |
![]() |
Mass M |
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Age | ![]() |
Microturbulence ![]() |
2.0 kms-1 |
Macroturbulence
![]() |
4.0 kms-1 |
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1.0dex below solar |
![]() |
0.6dex below solar |
Distance | 188 +38-28 pc |
Our KPNO spectra from April 1998 included the wavelength region around
the H
line, which is widely used as an activity and velocity
indicator for stellar chromospheres (see e.g. Bopp et al. 1988).
Our spectra show a broad and asymmetric
absorption line with a weak
blue-shifted emission component (Fig. 3). To extract the
H
contribution from the active part of the chromosphere, we
subtract the broadened spectrum of the inactive reference star 16Vir
(K0.5III). To first order, this removes the photospheric
and chromospheric flux that is unrelated to the magnetic activity.
This subtraction reveals a double-peaked emission profile with
unequally strong emission peaks which is typical for very active RS CVn
binaries.
We assume that the doubled peak results from one emission peak combined
with an off-centered absorption line. Using the Gaussian deblending routine
in the IRAF splot package, we find that the strong emission peak is
centered at the H
rest wavelength, while the absorption feature
appears blueshifted by
kms-1 in spectrum a
(JD 2450918) in Fig. 3 and
kms-1 in
spectrum b (JD 2450921).
The full width at half maximum (FWHM) of the Gaussian for the absorption
component of these fits is
Å,
which transforms to
kms-1 using the relation
described in Strassmeier et al. (1990), and thus seems to
originate from near the stellar surface.
Although we can not draw more firm conclusions from
these two H
spectra, we do note that the line-profile structure
and its asymmetry suggest a complex flow structure in the lower
chromosphere and possibly also in the outermost coronal layers, say, one
stellar radius above, where hydrogen can be trapped in coronal loops.
As for all previous papers in this series, the maps were calculated with the Doppler-imaging code TEMPMAP (Rice et al. 1989; Piskunov & Rice 1993; Rice 1996). See e.g. Paper X (Strassmeier et al. 1999b) for an updated description of the program and Paper VI (Strassmeier & Rice 1998) and Rice & Strassmeier (2000) for numerical tests with artificial data.
Because of the small wavelength coverage of our NSO spectra, we could only
use three main mapping lines: Fe I 6421, Fe I 6430 and
Ca I 6439 with
values of -2.23, -2.0 and +0.47 and
lower excitation potentials of 2.279, 2.176 and 2.526 eV, respectively.
Since all three lines are blended to a certain degree - the 6421-Å region seems to be especially vulnerable - the number of lines which had
to be synthesized was 6, 6, and 7 for the 6421, 6430 and 6439-line regions,
respectively.
All these blends are included in the inversion and treated simultaneously
with the primary mapping lines but only one spectral region can be handled
per solution. We employed a maximum-entropy regularization for the Doppler
maps presented in this paper, but the program also allows a Tikhonov regularizing
functional (for a comparison see, e.g., Piskunov & Rice 1993).
An appropriate set of model atmospheres was taken from the ATLAS-9 CD
(Kurucz 1993).
All radiation transfer calculations are done under the assumption of LTE.
Because the Doppler-imaging analysis is sensitive to the rotational velocity
and to the inclination of the stellar rotation axis, it can be used
to refine these two parameters with higher accuracy than with the method
described in Sect. 3 (see also, e.g., Unruh 1996).
Changing these two parameters one at a time, while all others are kept
constant, yields various values for the misfit between the data
and the model (). The value of the parameter corresponding to
the smallest
is the one we believe is
closest to the true value. The variation of
with the inclination i
and the rotational velocity
are plotted in
Figs. 4a and b, respectively. A minimum is seen in both cases:
for the inclination between
and
,
and for the projected
rotational velocity between 27 and 29 kms-1. We obtain
kms-1 and
as our final values,
according to the unweighted average minimum in Figs. 4a and b.
To determine approximate elemental abundances, we evaluate the run of the
from a series of inversion solutions with different initial
abundances. A straightforward computation of the local line profiles with
solar abundances and from model atmospheres with gravities between
,
and micro- and macroturbulences between 0.5-4.0 kms-1,
already indicated a relative underabundance of calcium and iron with respect
to the Sun.
Our test inversions were thus started
with solar abundances and decreased in steps of 0.1 dex
(Figs. 4c and d).
The transition probabilities, damping
constants, and laboratory wavelengths were kept constant at the values
specified in Sect. 4.1 and taken partially from the VALD database
(Kupka et al. 1999) and our previous papers in this series. We
then adopted the
abundances that resulted in the least
as the most probable values
for the surface abundances, i.e.
and
(on the
scale). Although an internally consistent set of parameters of
high precision, these abundances are of low accuracy
because test inversions with different sets of fixed parameters
(i.e. mostly different
and microturbulence) resulted
in similar Doppler maps but with individual abundances different by
up to
0.2 dex.
Figures 5a,b shows the average maps from three spectral regions and
for two consecutive stellar rotations, respectively.
One spectral region and two photometric bandpasses, V and ,
are used
simultaneously to produce a single map. Three such maps from Ca I 6439.08, Fe I 6430.84 and Fe I 6421.35 are then used to create an
average map by averaging the temperature in each pixel. Each map is given
equal weight because the overall
achieved is similar.
We caution though that the averaging may lead to an overemphasizing of the
strongest lines because of their larger intrinsic line width and thus a
more "smeared-out'' temperature distribution relative to a weaker line.
Therefore, we also present the individual maps
in Figs. 6 and 7 for the two stellar
rotations, respectively.
All computations are performed on a DEC-Alpha 500/400 workstation and require
between 20 to 30 min of CPU time depending on the number of spectral blends,
the number of input model atmospheres, and the number of photometric data points.
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Figure 7: Individual Doppler images for Rotation 2. Otherwise as in Fig. 6. |
Spot | Rotation 1 | Rotation 2 | ||||
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b |
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b |
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|
A | 20 | 35 | 4200 | 50 | 40 | 4050 |
B | 85 | 35 | 4250 | 110 | 35 | 4250 |
Ca | 160 | 45-75 | 3900 | 160 | 30-75 | 3850 |
D | 220 | 40 | 4050 | 230 | 35 | 4100 |
E | 285 | 40 | 4250 | 280 | 35 | 4100 |
F | 340 | 30 | 4150 | 330 | 45 | 4100 |
A visual comparison of the two consecutive maps in Figs. 5a and b
reveals a comparable
spot morphology as well as consistent absolute temperatures. A cool
polar spot with an appendage centered around a longitude of
and a latitude of
-
is seen in both maps (referred to as spot C in Table 3).
Its lower end reaches almost down to the equator in Rotation 1, but we
believe that this is partly due to latitude smearing in the image
reconstruction than a real single appendage and it is entirely possible
that the feature consists of a high-latitude appendage and a low-latitude spot at the
same longitude (the narrower Fe I6421 line shows actually two features
well separated). The Fe I6421 map for the second rotation does not
indicate the low-latitude spot anymore but instead show an enhanced polar appendage.
One possible explanation could be that the lower spot had merged with the polar
appendage within a single stellar rotation.
Additionally to the polar spot there seem to be five spots at low-to-medium
latitudes (named A, B, D, E, F from left to right in Fig. 5c). See
Table 3 for a summary of the spot positions on both maps.
Note that these numbers are measured off the maps by treating the Doppler
image as a CCD image in the CCD-photometry package digiphot in IRAF.
We define a box around the approximate location of a feature and then
use the IRAF routine phot to find a local (temperature) minimum within
this box. The resulting coordinates have sub-pixel precision but a significantly
worse accuracy
that is determined by a large variety of external sources (see Rice &
Strassmeier 2000). We estimate an average of 5
in latitude
and somewhat less than that in longitude.
The coolest of these spots has a temperature of about 800 K below the
effective photospheric temperature of 4700 K, and the warmest is
approximately 450 K cooler. Typical errors for spot temperatures from
TEMPMAP Doppler images are 70-100 K.
The equivalent widths of Ca I6439 Å and Fe I6430 Å are identical within the observational errors but the Fe I6421 line is approximately 20% weaker. It is severely affected by the nearby Fe I6419 line and its large number of blends that affect the blue wing of the Fe I6421 line. Its fit to the observed light curves of the second rotation is also noticeably different to the fits from the other two lines. Despite this inconsistency, all spectral lines indicate a cool polar cap. The main difference in the maps is that the polar spot from the Fe I6421 line appears 100-150 K cooler than in the other two maps while the low-latitude spots appear warmer by that amount and, consequently, the polar spot's enhanced contrast seems eye catching. It is a known problem of line-profile inversions that a polar spots' size and contrast depends strongly on the fit to the line wings (see Unruh et al. 1996) and thus the contrast of the 6421-line is not a big surprise.
Nevertheless, to disentangle such inconsistencies, we compute a rotation-averaged
map by combining all spectral data into one data set and invert it as usual. This
map is shown in Fig. 5c; it is already the average from all three
spectral regions. To quantify the differences between the fits of the individual
rotations and the rotation-averaged fit, we used the average map to calculate synthetic
line profiles in a forward solution and compare them with the inverse solutions
from the two rotations. The difference of their
is then an indirect measure
for the reality of surface changes from one rotation to the next.
Figure 8a summarizes these changes. Each individual point in this figure is the
squared minimum of the residuals between a single observation and the
rotation-averaged fit. We see, firstly, that the misfit is on average larger in
the second rotation and, secondly, that there are two phases in both rotations
with consistent
peaks from all three spectral regions
(at
100
and
300
). The plots of the difference maps
in Figs. 8b and c finally show the changes on the stellar surface. We find
that the asymmetry of the polar-spot appendage at
200
had vanished
in the second rotation and that two high-latitude regions appeared at
40
and
300
.
Two or maybe three of the lower-latitude
spots seemed to have migrated. We interpret these differences to
be due to real changes on the stellar surface and will investigate them further
in the next chapter. For the Fe I6421 line profiles, the
values in Fig. 8a are sometimes double than those for the two other lines.
We thus consider the severely blended Fe I6421-line map more
uncertain than the other two, albeit its overall
is comparable and
it recovers the features at approximately the same locations.
A further test of our Doppler images is to divide the spectra into "even'' and "odd'' data sets for each rotation, i.e. using first even-numbered spectra for an inverse solution and then the odd-numbered spectra, and then compare the resulting maps. The images from the rotation-averaged sets correlate very well, but the odd and even images for the individual rotations do not. The reason is that the phase coverage is just too sparse to derive results of the same quality as for the combined data set. Thus, the even-odd test is unfortunately not overly useful to verify the cross-correlation analysis applied in the next chapter.
By cross correlating longitudinal strips at successive latitudes from the
two maps in Figs. 5a, b, we can derive the differential surface rotation on
KUPegasi. Since the two maps are from consecutive rotations it is save
to assume that the surface features in the individual maps are still the
same (but see previous section). In such a case the spots can be used as
tracers for surface velocity fields, although the criterion
is not necessarily a stringent requirement if there are many spots with
the same general trend of migration. Of course, the interpretation is still
hampered by the possibility of a coincidental spot alignment that mimics a
latitude-dependent migration pattern. At this point, we simply caution the
reader that our data are just a snapshot and will be masked by spot evolution.
We applied the fxcor
routine of the IRAF package (for details see the IRAF
manual at iraf.noao.edu) to fit Gaussians to the
peak of the cross-correlation functions for the Ca I 6439 image, the
Fe I 6430 image, and the average of those two images.
The result in Fig. 9 shows a complex surface
differential-rotation
function: the shifts within
of the equator are tidily fitted
with a solar-like differential-rotation law proportional to
,
but
between +25
and +45
(and possibly also between -25
and -45
)
the function changes its sign and thus these regions appear to accelerate
again. Above
+50
the width of the cross-correlation
function increases rapidly due to the decreased surface resolution and
the longitudinal shifts cannot be measured very reliably there. Its error
bars from the Gaussian fits are
3 times larger than near the
equator. Despite this limitation there is some evidence though that the
rotation decelerates again above
+50
.
This is certainly
inconclusive from our two stellar rotations but should not go unnoticed.
We tried to fit the cross correlations with a simple
law
for the entire latitude range (full line in Fig. 9).
Such a fit is obviously not a good representation of the data (rms of 0.29)
but is intended to obtain a save lower limit for the
magnitude of the differential rotation.
We also did a
fit restricted to shifts within
25
latitude (dashed line in Fig. 9) and thereby also
obtain an estimate of the external error per latitude range. Its rms is
accordingly better, 0.11. The first fit leads to the following
differential rotation law for -40
to 65
![]() |
Figure 10:
a) Meridional changes on the surface of KU Peg. Plotted are the
latitudinal shifts per longitudinal bin from a cross correlation of the
hemisphere above the equator from Rotation 1 and 2.
The two bumps at 40![]() ![]() ![]() ![]() |
To quantify latitudinal changes on the stellar surface, we now
cross correlate the maps along meridional circles.
We just adopt the "northern'' hemisphere, i.e. the hemisphere that is fully
in view (all pixels with positive latitude), and cross correlate
its longitudinal strips of the Ca I 6439 images, the
Fe I 6430 images, and the average of those two images, respectively.
The result in Fig. 10a clearly shows that there were
two latitudinal shifts on the surface of KU Peg within one
stellar rotation that consistently appeared in both spectral lines.
One event at
phase 0.1 (40
)
and another at
0.9 (330
),
with an average magnitude of
(rms)
day-1 and
(rms)
day-1, respectively. The intermittent
longitudes on the opposite side of the visible pole may show a reversed
shift of magnitude -0.2
day-1, but this may be overinterpretation
given the large error bars and the inconsistency at phases from
approximately 135
to 180
.
The latter is most likely caused by
the well-known "north-south'' mirroring of the polar appendage at
.
As for the previous cross correlations the error bars per bin are estimated from a Gaussian fit of the FWHM of the cross correlation function. Both shifts have a positive sign and thus indicate a polar-directed change. We interpret these shifts as a meridional change of magnetic flux and, since it seems to be a local event on the stellar surface, tentatively suggest that magnetic reconnection and its associated plasma motions may be the underlying cause rather than a global meridional flow.
We detected a complex surface differential-rotation law for KU Peg with
acceleration along the equator and possibly also near a latitude of
50
,
and deceleration in between and above.
KU Peg's lap time
is
70 days for the
25
range
around the equator, but 260 days for the entire equator-to-pole range.
This differs from the solar case and from the findings
of Collier Cameron et al. (2001) for the three ultra-fast
rotators RXJ1508-4423 (G2), ABDor (K0/2V), and PZTel (K0IV/V)
with rotation periods of 0.31, 0.515, and 0.94 days, respectively, in that
these stars show a uniform differential rotation that follows a simple
solar-like
law. Their lap times
are 40, 110,
and 80 days, respectively, which are all shorter than the solar value of 120 days.
LO Peg (K5-7V,
days) on the other hand, does show signs for the
equator lagging behind (Lister et al. 1999). These authors argue
though that the absence of mid-latitude features together with the short time
between the two images (one day) possibly prevents the detection of significant
differential rotation.
Rice & Strassmeier (1996) detected differential rotation on the
weak-lined TTauri star V410Tau (
days) in the
same sense as on the Sun but with a laptime of 1800 days, a factor of 15
weaker, which is significantly different from the stars above. Either each
star with a given mass, rotation period, and evolutionary status has its
own distinct differential surface rotation, as e.g. indicated by the
theoretical models from Kitchatinov & Rüdiger (1999), or the
observations are masked by local and rapid magnetic-field
reconfigurations that just mimic a differential rotation law.
Differential rotation was reported for several other evolved stars: Weber
& Strassmeier (1998) found equatorial acceleration on the RSCVn
binary ILHya for latitudes below 45
and about a factor of 30 smaller
than on the Sun, as measured from two images taken
28
stellar rotations apart.
Strassmeier (1994) and Hatzes (1998) derived a
differential-rotation law for HUVir (also a RSCVn star) from the
comparison of two temperature maps taken 2 stellar rotations and 4 years
apart, respectively, and combined with contemporaneous photometric modeling
included in Strassmeier (1994). Both authors found differential
rotation of inverse behavior than on the Sun - polar regions rotate faster
than low-latitude regions - and a factor of 10 slower. However, the spot
features used in both studies were all at
and thus no
reliable information for the equatorial region was available.
Hatzes & Vogt (1992) found solar-like differential
rotation on the short-period (1.95 days) RSCVn binary EIEri,
i.e. equatorial acceleration and similarity in strength, but an inverse
behavior, and about a factor of 10 smaller than on the Sun, on the
6.4-days RSCVn binary UXAri (Vogt & Hatzes 1991).
So far, there is cumulative evidence that differential rotation profiles on
evolved stars (and possibly also on pre-main-sequence stars) appear to
be more complicated than on solar-type main-sequence or ZAMS stars.
A recent study of the RSCVn
binary HR1099 by Strassmeier & Bartus (2000) reveals
a general poleward spot migration of the order of 0.4day-1simultaneously to longitudinal spot migrations with both signs at the
same time, i.e. spots migrating faster and slower than the orbital
period but are located at approximately the same latitude. This is in
agreement with an earlier claim by Vogt et al. (1999)
based on 23 Doppler images taken throughout 11 years. KU Peg also shows
evidence, like HR1099, for poleward spot migration and even of the same amount.
There is also some similarity of our KU Peg result to the recently obtained rotation
profile for the rapidly-rotating long-period K0III binary
Gem
(
days, Kövári et al. 2001).
For
Gem, they found a differential rotation law in quadratic form
with acceleration in two latitudinal bands centered at approximately
40
around the equator, but deceleration along the equator
and near the one visible pole. We believe that all of these observations
hint toward a general dependence of differential rotation upon rotational
period. Giants seem to show a mixture of solar-like and anti-solar profiles
of various strengths, which seems to be partially in conflict with the
recent differential-rotation models of Kitchatinov & Rüdiger
(1999) who predict larger differential rotation in giants than
in dwarfs. It is also indicative that differential rotation is not the only
way to explain spot migrations and that the associated meridional flow may
play a stronger role on giants than, e.g., for the Sun.
The Sun, for comparison, has a very weak latitudinal flow pattern of
0.03
day-1 (Howard & Gilman 1986).
This flow transports magnetic flux from mid-latitude spots up to the rotation
poles where its opposite polarity causes the polarity reversal and the end of an
old magnetic cycle and the start of a new one.
So far, stellar observations of a poleward flow exist only for
stars with high-latitude active regions but are in agreement with the
picture first presented by Schüssler & Solanki (1992).
In that picture, the flux tubes
can arise at latitudes up to
60
if the star rotates
rapidly enough. However, an additional transportation mechanism is necessary
to move the spot towards the pole once it has emerged. This is different from
the very young stars, where a truly polar spot can emerge without the additional
need of a meridional flow (see Granzer et al. 2000 for a recent
discussion).
For some of the previously discussed stars, the time between the individual maps was usually many rotations, and thus spot changes with timescales less than a few rotations could not be determined. Only spots close to the poles seem to be persistent enough to be seen throughout many rotations and this may bias our meridional-flow detections. However, if some mechanism does transport active regions towards the pole, where they make up for a large torodial field that, in return, inhibits differential rotation, then the difficulty of detecting differential rotation on such stars is not a surprise. One such star is the 16-day G8II-III giant CMCam (Strassmeier et al. 1998), where cross correlations of Doppler images from four observing seasons with one year in between did not reveal a clear differential-rotation signal, despite that there is evidence for phase shifts on its surface. A very similar case is the single G5 giant HD199178 (Strassmeier et al. 1999b), where images taken one month and images taken one year apart were cross-correlated but no systematic migration pattern was found. Whether the time the magnetic field needs to reconfigure on these stars is too short to be detected or, whether an existing differential rotation pattern is simply masked by short-term field configuration changes, could not be answered in those two cases.
Acknowledgements
Thanks to Trudy Tilleman for operating the McMath telescope half of the time and for providing gourmet coffee all of the time, to the Austrian Fond zur Förderung der wissenschaftlichen Forschung (FWF) for support under grants S7301-AST and S7302-AST, and to the German Forschungsgemeinschaft (DFG) for grant HU 532/8. We thank the referee, Dr. J. R. De Medeiros, for his constructive criticism.