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The Astrophysical Journal, 569:941–963, 2002 April 20
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WHAT IS HAPPENING AT SPECTRAL TYPE F5 IN HYADES F STARS?1
Erika Bo¨hm-Vitense
Astronomy Department, University of Washington, Seattle, WA 98195
Richard Robinson
Institute for Astrophysics and Computational Sciences, Department of Physics,
Catholic University of America, Washington, DC 20064
Kenneth Carpenter
LASP, NASA Goddard Space Flight Center, Code 681, Greenbelt, MD 20771
and
Jose Mena-Werth
University of Nebraska at Kearney, Physics Department, Kearney, NE 68849-1160
Received 2001 August 15; accepted 2002 January 1
ABSTRACT
Aiming at a better understanding of the mechanisms heating the chromospheres, transition regions, and
coronae of cool stars, we study ultraviolet, low-resolution Hubble Space Telescope/Space Telescope Imaging
Spectrograph spectra of Hyades main-sequence F stars. We study the BV dependence(s) of the chromospheric and transition layer emission line ﬂuxes and their dependences on rotational velocities. We ﬁnd that
the transition layer emission line ﬂuxes and also those of strong chromospheric lines decrease steeply between
BV ¼ 0:42 and 0.45, i.e., at spectral type F5, for which the rotational velocities also decrease steeply. The
magnitude of the line-ﬂux decrease increases for lines of ions with increasing degree of ionization. This shows
that the line-ﬂux decrease is not due to a change in the surface ﬁlling factor but rather due to a change of the
relative importance of diﬀerent heating mechanisms. For early F stars with BV < 0:42 we ﬁnd for the transition layer emission lines increasing ﬂuxes for increasing v sin i, indicating magnetohydrodynamic heating.
The v sin i dependence is strongest for the high-ionization lines. On the other hand, the low chromospheric
lines show no dependence on v sin i, indicating acoustic shock heating for these layers. This also contributes
to the heating of the transition layers. The Mg ii and Ca ii lines show decreasing ﬂuxes for increasing v sin i, as
long as v sin i is less than 40 km s1. The coronal X-ray emission also decreases for increasing v sin i, except
for v sin i larger than 100 km s1. We have at present no explanation for this behavior. For late F stars the
chromospheric lines show v sin i dependences similar to those observed for early F stars, again indicating
acoustic heating for these layers. We were unable to determine the v sin i dependence of the transition layer
lines because of too few single star targets. The decrease of emission line ﬂuxes at the spectral type F5, with
steeply decreasing v sin i, indicates, however, a decreasing contribution of magnetohydrodynamic heating for
the late F stars. The X-ray emission for the late F stars increases for increasing v sin i, indicating magnetohydrodynamic heating for the coronae of the late F stars, diﬀerent from the early F stars.
Subject headings: open clusters and associations: individual (Hyades) — stars: chromospheres —
stars: coronae — stars: rotation
number of solar ﬂares during solar activity maxima may
also support the suggested heating by ﬂares. The dependence of the emission measures on temperature in the upper
transition regions seems to support the assumption that at
least these layers are heated by conductive heat transport
from the coronae. In this case the question then remains:
what is heating the coronae?
The detailed solar observations of spatial structures and
time evolutions of chromospheric and transition layer emission and their relation to magnetic ﬁeld structures over the
last decades, together with the launch of solar observation
satellites SOHO and TRACE, have raised new general interest in the subject of the heating mechanisms of the outer
solar atmospheric layers. While these observations have
brought a large increase of detailed knowledge, they have
not yet revealed the governing heating mechanism(s). Stellar
observations, showing the dependences of the chromospheric, transition layer, and coronal emissions on the different physical parameters, like Teff , v sin i, and age will
have to be used in addition in order to identify the heating
1. INTRODUCTION
The identiﬁcation of the heating mechanism(s) for the
solar and stellar chromospheres, transition layers, and coronae has been a much discussed topic for several decades.
Acoustic waves (Schwarzschild 1948; Biermann 1946), magnetohydrodynamic waves (for instance, Vaiana & Rosner
1978), heat conduction from the coronae (Woolley & Allen
1948; Unso¨ld 1955), and many ﬂares (Sturrock et al. 1990)
have been discussed. Rosner, Tucker, & Vaiana (1978)
emphasized the importance of loops, which are especially
prominent in solar X-ray pictures. The observed increase of
solar chromospheric and transition layer emission during
increased solar activity has supported the suggestion of
heating by magnetohydrodynamic eﬀects. The increasing
1 Based on observations with the NASA/ESA Hubble Space Telescope
obtained at the Space Telescope Science Institute, which is operated by the
Association of Universities for Research in Astronomy, Incorporated,
under NASA contract NAS5-26555.
941
942
BO¨HM-VITENSE ET AL.
mechanism(s). From ground-based observations of mainly
ﬁeld stars with spectral types later than F5, the Ca ii line
emission was found to increase with increasing rotational
velocities (Wilson 1970). IUE observations of (mainly) ﬁeld
stars showed that the Mg ii and the transition layer lines for
stars with spectral types later than F5 also showed increased
ﬂuxes for stars with increasing rotational velocities (Walter
1983, 1986). Most of these stars also show rotational modulation of their Ca ii emission (Baliunas & Vaughan 1985),
attributed to star spots. On the other hand, the emission
lines of the early F ﬁeld stars did not show an increase of line
ﬂuxes with increasing v sin i and do not show rotational
modulation, although a few stars may show some very weak
modulation (Baliunas & Vaughan 1985). This led Simon &
Drake (1989) to suggest that for early F stars the heating of
the chromospheres and transition layers might be due to
acoustic wave damping, while for spectral types later than
F5 magnetohydrodynamic eﬀects would become important.
For larger v sin i the stellar dynamo might become more eﬃcient, leading to increased spot and magnetohydrodynamic
activity and rotational modulation of the line emission.
Skumanich (1972) studied ﬁeld and cluster G V stars of different ages and found decreasing activity for increasing
ages, as was discussed earlier by Wilson & Skumanich
(1964). This was attributed to the decreasing rotation with
increasing age. The dependences of activity on age and on
v sin i may, however, be confused in such studies since age
and rotation rate are correlated. We need observations of
stars with known ages and diﬀerent v sin i to determine
the dependences of the emission on rotation and on age
separately.
The diﬀerent suggested heating mechanisms indicate that
the heating may be dependent on Teff , vrot , age, chemical
composition, gravity, overall magnetic ﬁelds, and possibly
binarity (Bo¨hm-Vitense 1982 and references therein). In
order to study the dependence on one parameter, we have to
keep all other parameters constant as nearly as possible.
This can best be done for cluster stars, for which chemical
composition and age are the same. If we restrict our studies
to main-sequence stars of a small temperature range, then
gravity is also almost constant. We still have to deal with at
least three, if not four, parameters: Teff , v sin i, binarity, and
possibly magnetic ﬁelds. This means we have to observe a
number of stars in very narrow ranges for all parameters
except one, for instance Teff , to determine the dependence
on Teff , and similarly for the other parameters. This requires
a large number of stars to be observed. In order to reduce
the number of targets, we have observed only mainsequence Hyades F stars but as many as seemed feasible. F
stars appear to be especially important because of all the
abrupt changes observed around spectral type F5 and the
implied change in heating mechanism(s).
In a previous paper (Bo¨hm-Vitense et al. 2001, hereafter
Paper I), we used observations with the Hubble Space Telescope (HST) and the Space Telescope Imaging Spectrograph (STIS), together with IUE observations to study the
Mg ii lines of Hyades F stars. From these we found that,
contrary to the observations for the ﬁeld stars, for Hyades F
stars of a given Teff the Mg ii emission line ﬂuxes decrease
with increasing rotational velocities. This is true for all narrow intervals in BV as long as the rotational velocities are
not larger than about 50–60 km s1 and as long as for the
surface line ﬂux, F, we have log F > 6. For smaller ﬂuxes
and larger rotational velocities, another heating mechanism
seems to take over in the layers emitting the Mg ii lines. It
seems possible that the very eﬃcient heating mechanism
leading to large chromospheric line emission and increasing
line emission for decreasing v sin i only works for young
stars (perhaps owing to a fossil magnetic ﬁeld?), and therefore in the older ﬁeld stars we only see increasing surface line
ﬂuxes for increasing v sin i. In this paper we study the relation between Teff , rotation, binarity, and the emission line
ﬂuxes for other chromospheric and transition layer emission
lines of the same Hyades main-sequence F stars.
2. THE HST OBSERVATIONS
2.1. The HST Targets
We observed 18 Hyades F stars with the HST/STIS. For
these observations we used the G140L grating and the Large
Science Aperture to obtain spectra for the spectral region
˚ with a resolution R ¼ 2000. In this spectral
1150–1750 A
region many strong emission lines are found that originate
in the chromospheres and in the transition region between
the chromospheres and the coronae. In Table 1 we list our
HST targets and the basic properties of these stars. All photometric data were obtained from the Hipparcos catalogue
(ESA 1997) in order to have a homogeneous set of data. The
v sin i values were obtained from Uesugi & Fukuda (1982),
who collected the data from various sources in the literature, as referenced in their paper; ‘‘ vB ’’ indicates the van
Bueren numbers (van Bueren 1952).
The observing dates and exposure times for the HST
observations used here are given in Paper I. HD 27848 with
BV ¼ 0:450, the BV for which Ca ii and Mg ii emission
lines have the minimum ﬂux, was observed twice to see
whether any diﬀerences could be seen in the spectra, possibly owing to varying stellar activity. In Table 2 we give the
basic data for Hyades main-sequence stars that were previously observed with the IUE satellite.
In Figure 1 we show the color-magnitude diagram
(CMD) for our target Hyades main-sequence F stars. The
absolute magnitudes are plotted, using the parallaxes given
in the Hipparcos catalogue. Binary stars are indicated by
surrounding squares. More than 50% of our targets turned
out to be binaries (see Perryman et al. 1998). This gives us
an opportunity to study chromospheric and transition layer
emission line diﬀerences between binaries and single stars
but reduces the number of single stars for our statistics.
Fig. 1.—Positions of our target stars are shown in a CMD. Plus signs
indicate stars observed with IUE, crosses indicate stars observed with HST,
and surrounding squares indicate binaries.
TABLE 1
Basic Data for HST Hyades Program Stars
HD
vBa
mV
BV
v sin i
(km s1)
vr
(km s1)
Spectral Type
(mas)
MV
27561 ......
28736 ......
26345 ......
28568 ......
29225 ......
27848 ......
31845 ......
28406 ......
28483 ......
28608 ......
21847 ......
26784 ......
27808 ......
30738 ......
28205 ......
28033 ......
28237 ......
29419 ......
37
90
13
85
101
51
128
78
81
86
...
19
48
121
65
62
66
105
6.60
6.37
6.61
6.50
6.64
6.96
6.75
6.90
7.09
7.03
7.29
7.11
7.13
7.29
7.41
7.36
7.49
7.51
0.412
0.420
0.427
0.428
0.442
0.450
0.450
0.451
0.470
0.472
0.503
0.514
0.518
0.536
0.537
0.557
0.560
0.576
15
35
20
55
40
30
25
20
20
20
...
5
10d
10
10
5d
10
5d
39.2
40.3
33.0
43.6
33.7
43.2
44.1
38.6
38.0
41.4
...
38.5
38.9
42.7
39.3
38.8
40.2
39.9
F5
F5
F6
F2
F8
F8
F5
F8
F5
F5
F8
F8
F8
F8
G0
F8
F8
F5
19.46
23.13
23.22
24.28
22.99
18.74
23.09
21.59
19.94
22.96
20.45
21.08
24.47
19.30
21.83
21.54
21.18
22.60
3.046
3.191
3.439
3.426
3.448
3.324
3.567
3.571
3.589
3.835
3.843
3.729
4.073
3.718
4.105
4.026
4.120
4.281
Comment
Binaryb
Binary c
Binaryc
Binaryc
Binaryc
Binaryb
Binaryc
BY Dra variable?e
Binaryc
a
The van Bueren number.
Patience et al. 1998.
c Perryman et al. 1998.
d Upper limit.
e SIMBAD.
b
TABLE 2
Basic Data for IUE Observed Hyades Stars
a
b
HD
vB
mV
27397 ......
28294 ......
28677 ......
24357 ......
26462 ......
27901 ......
29169 ......
26015 ......
26911 ......
27429 ......
25102 ......
26737 ......
26345 ......
28568 ......
28911 ......
27534 ......
29225 ......
27731 ......
28483 ......
28608 ......
30869 ......
27991 ......
26784 ......
27808 ......
27691 ......
28394 ......
30810 ......
30311 ......
27859 ......
30
68
89
6
14
53
100
11
20
32
54
16
13
85
94
36
101
44
81
86
124
57
19
48
40
77
122
113
52
5.58
5.90
6.01
5.97
5.71
5.97
6.01
6.02
6.31
6.11
6.35
7.05
6.61
6.50
6.62
6.80
6.64
7.18
7.09
7.03
6.30
6.44
7.11
7.13
6.97
7.02
6.79
7.24
7.79
a
BV
b
0.283
0.325
0.338
0.357
0.360
0.378
0.380
0.397
0.400
0.404
0.417
0.424
0.427
0.428
0.429
0.441
0.442
0.462
0.470
0.472
0.502
0.509
0.514
0.518
0.518
0.528
0.543
0.560
0.599
v sin i
(km s1)
vr a
(km s1 )
30
102
109
59
6
150
80
30
53
132
54
70
18
55
40
12
40
30
18
20
25
18
12
12
10
25
6
7
5
42.0
44.2
36.0
35.0
36.6
36.6
43.3
36.4
36.9
42.0
39.6
38.4
33.0
43.6
35.0
37.0
33.7
34.1
38.0
41.4
38.8
36.4
38.5
38.9
42.7
39.9
42.4
42.5
38.8
The van Bueren number ‘‘ vB ’’ and vr are from Schwan 1991.
, BV, and mv from the Hipparcos catalogue.
c Spectral type from catalog 2000.
b
Spectral Type
b
(mas)
MV
Comment
F3
F0
F4
F4
F4
F4
F5
F3
F5
F3
F5
F5
F6
F2
F2
F5
F8
F5
F5
F5
F5
F7
F8
F8
G0
G0
F8
F5
G2
22.31
18.42
22.25
24.14
25.89
20.40
22.60
21.27
22.51
21.12
25.42
18.12
23.22
24.28
22.80
19.83
22.99
20.04
19.94
22.96
23.91
21.47
21.08
24.47
21.45
23.25
20.15
26.26
20.73
2.32
2.23
2.75
2.88
2.78
2.52
2.78
2.66
3.07
2.73
3.38
3.34
3.44
3.43
3.41
3.29
3.45
3.70
3.59
3.84
3.19
3.10
3.73
4.07
3.63
3.85
3.31
4.34
4.37
Binary
Binary
Binary
c
Binary
Variable?
Binary
Binary
Binary
Binary
Binary
Binary
Binary
Binary
Binary
Binary
Binary
Binary
Binary
Binary
Binary
BO¨HM-VITENSE ET AL.
944
2.2. The Spectra
In Figure 2 we show the spectra of the HST Hyades mainsequence single F stars. For HD 27848 we show the superposition of both spectra, taken at a time diﬀerence of 16
days, in order to demonstrate that no variations were seen
for this F5V star with BV ¼ 0:450.
For the HST spectra the signal-to-noise ratios are much
better, the resolution is higher, and the background noise is
lower than for the spectra of stars previously obtained with
the IUE. The accuracy of the HST observed emission line
ﬂuxes is therefore much better than the one for the IUE
measured emission line ﬂuxes. For the HST data we estimate the error bars for the strong lines to be less than 10%
or 0.04 dex. For the weak S i lines they may be 0.1 dex, and
for the very weak C i lines the errors for some stars may be
as large as 50% or 0.20 dex. The HST spectra were taken
through the long slit (0>5 by 5200 ). The geocoronal contribution and the background contributions to the measured
spectral ﬂuxes can thus be determined from pixels near but
outside of the spectrum proper and can be corrected for.
Vol. 569
For the IUE spectra of the Hyades late F stars (reduced
with the original, old IUESIPS reduction method), even the
usually strong lines appear weak and may have error bars
reaching 50% or 0.2 dex. Because of the strong geocoronal
contamination, Ly ﬂuxes cannot be measured for our IUE
spectra.
2.3. Error Discussion
In the following we will compare ﬂuxes of a given line for
diﬀerent temperature stars, or compare lines for stars with a
given temperature for diﬀerent rotation velocities. We are
mainly interested in relative surface ﬂuxes. Thus, for our
discussion mainly statistical errors are important. For reasonably strong lines these are mainly determined by the
ﬂux-measuring errors, which are mainly due to the personal
judgement of the level of the underlying apparent continuum, which may still be inﬂuenced by unrecognized, weak
emission lines. We measured the line ﬂuxes, f, independently
by three or four people and took the averages. The diﬀerent
measurements deviated for the strong HST lines by less than
Fig. 2.—Sample HST spectra are shown for Hyades F stars. For HD 27848 we show two spectra, the second one taken 16 days later than the ﬁrst one. No
changes larger than the noise are seen. For the major lines the emitting ions are indicated in the top panel. The spectra are ordered according to the BV colors
of the stars.
No. 2, 2002
SPECTRAL TYPE F5 IN HYADES F STARS
10% or 0.04 dex. For the very weak C i lines they may reach
50% or 0.2 dex for some stars.
The transformation of measured ﬂuxes, f, to surface line
ﬂuxes, F, discussed in x 3, makes use of equation (1), which
uses mðbolÞ, which we replaced by mv , because the bolometric corrections are close to zero for main-sequence F stars
(see, for instance, Perryman et al. 1998). When comparing
log F =f of F0 stars with F8 stars, this may cause an error of
about 0.01; mv is uncertain by about 0.01 and BV by about
0.001. The uncertainty in BV introduces an error of 0.15%
in the log F =f , and the uncertainty of mv leads to an uncertainty in log F =f of 0.004. These errors are negligible in
comparison with the measuring uncertainties in f.
945
3. THE MEASURED EMISSION LINE FLUXES
3.1. HST Observations
The measured chromospheric and transition layer emission line ﬂuxes, f, measurable on the HST/STIS spectra are
listed in Tables 3 and 4. In Tables 5A and 5B we give the surface ﬂuxes, F, for the chromospheric and transition layer
lines. They were calculated from the measured ﬂuxes, f,
using the relation between F and f given by Oranje (1986),
reproduced here in equation (1):
4
:
logðF =f Þ ¼ 0:328 þ 0:4mðbolÞ þ 4 log Teff
TABLE 3
HST Measured Chromospheric Emission Line Fluxes for Hyades F Stars
HD
Ly 1215
O i 1300
S i 1473
C i 1657
C i 1560
Si ii
1526+1533
27561 ......
28736 ......
26345 ......
28568 ......
29225 ......
27848 ......
31845 ......
28406 ......
28483 ......
28608 ......
21847 ......
26784 ......
27808 ......
30738 ......
28205 ......
28033 ......
28237 ......
29419 ......
73.36
43.30
29.20
58.27
45.28
19.83
32.62
20.90
31.69
27.29
11.87
32.80
25.57
24.43
22.60
14.31
24.38
19.53
7.46
4.64
2.85
5.12
4.23
1.65
2.61
1.84
2.68
2.28
1.40
2.65
2.60
1.82
1.86
1.05
2.03
1.68
0.16
0.22
0.22
0.64
0.13
0.12
0.08
0.12
0.21
0.06
0.10
0.21
0.12
0.16
0.08
0.12
0.09
0.08
0.85
...
...
...
...
0.18
0.18
0.36
1.30
0.63
0.57
1.37
1.86
1.17
1.23
0.49
1.60
1.42
0.37
0.17
...
...
0.24
0.11
0.24
0.12
0.28
0.45
0.29
0.53
0.39
0.23
...
...
0.41
0.29
0.64
0.59
0.38
0.70
0.64
...
0.45
0.12
0.62
0.53
0.33
0.53
0.68
0.56
0.45
0.37
0.45
0.42
Note.— Line ﬂuxes in units of 1014 ergs cm2 s 1.
TABLE 4
Measured Fluxes, f, for Transition Layer Emission Lines of HST Observed Hyades F Stars
HD
C ii 1335
C iii 1175
C iv 1550
Si iii 1206
Si iv 1393]
Si iv 1402
N v 1240
He ii 1640
27561 ......
28736 ......
26345 ......
28568 ......
29225 ......
27848 ......
31845 ......
28406 ......
28483 ......
28608 ......
21847 ......
26784 ......
27808 ......
30738 ......
28205 ......
28033 ......
28237 ......
29419 ......
19.31
10.36
6.67
13.28
9.18
3.15
5.47
3.08
4.47
4.16
2.22
3.81
4.13
3.22
3.19
0.71
3.18
2.42
10.97
5.61
3.33
7.05
4.24
1.39
2.60
1.25
2.44
1.76
0.77
1.72
1.80
1.48
2.26
0.84
1.40
0.94
28.19
15.51
7.97
19.97
11.18
3.34
6.41
3.59
6.21
4.28
2.40
4.45
4.19
3.36
3.26
0.67
3.61
2.83
12.07
7.39
3.33
7.45
4.15
1.36
2.52
1.63
2.63
1.99
1.59
1.68
2.09
2.13
2.25
0.53
1.98
1.37
7.79
4.39
2.17
4.48
3.51
1.16
2.03
1.06
1.93
1.41
0.98
1.84
1.94
1.45
1.49
0.26
1.51
1.32
4.79
2.96
1.55
3.21
2.35
0.65
1.28
0.78
1.20
0.84
0.61
1.13
1.14
0.97
0.87
0.17
0.88
0.83
3.87
1.91
1.05
3.65
1.47
0.51
1.11
0.53
0.95
0.55
0.35
0.65
0.66
0.54
0.77
0.28
0.54
0.46
3.04
2.77
1.64
1.65
2.64
0.84
2.45
1.02
2.02
1.53
0.50
2.10
1.66
0.92
1.53
Note.—Measured ﬂuxes in units of 1014 ergs cm 2 s1.
1.40
1.22
ð1Þ
BO¨HM-VITENSE ET AL.
946
Vol. 569
TABLE 5A
Log of Chromospheric Surface Line Fluxes F for HST Hyades Program Stars
HD
log F =f
Ly 1215
O i 1300
S i 1473
C i 1657
C i 1560
Si ii
1526+1533
27561 ......
28736 ......
26345 ......
28568 ......
29225 ......
27848 ......
31845 ......
28406 ......
28483 ......
28608 ......
21847 ......
26784 ......
27808 ......
30738 ......
28205 ......
28033 ......
28237 ......
29419 ......
18.283
18.180
18.269
18.219
18.257
18.380
18.293
18.350
18.402
18.376
18.455
18.350
18.350
18.400
18.442
18.402
18.451
18.433
6.15
5.82
5.73
5.98
5.91
5.68
5.81
5.67
5.90
5.81
5.63
5.87
5.79
5.79
5.80
5.56
5.84
5.72
5.16
4.85
4.72
4.93
4.88
4.60
4.71
4.61
4.83
4.73
4.60
4.77
4.76
4.66
4.71
4.42
4.76
4.66
3.50
3.52
3.60
4.03
3.39
3.47
3.22
3.44
3.73
3.17
3.46
3.68
3.41
3.60
3.33
3.50
3.39
3.34
4.21
...
...
...
...
3.63
3.54
3.91
4.52
4.17
4.21
4.49
4.62
4.47
4.53
4.09
4.65
4.58
3.85
3.40
...
...
3.63
3.34
3.67
3.41
3.84
4.02
3.91
4.07
3.94
3.76
...
...
4.06
3.90
4.09
3.95
3.84
4.07
4.06
...
3.94
3.44
4.20
4.10
3.98
4.07
4.18
4.15
4.09
3.97
4.11
4.05
Note.—Line ﬂuxes in units of ergs cm2 s1.
TABLE 5B
Log of Transition Layer Surface Line Fluxes for HST Hyades Program Stars
HD
C ii 1335
C iii 1175
C iv 1550
Si iii 1206
Si iv 1393
Si iv 1402
N v 1240
He ii 1640
27561 ......
28736 ......
26345 ......
28568 ......
29225 ......
27848 ......
31845 ......
28406 ......
28483 ......
28608 ......
21847 ......
26784 ......
27808 ......
30738 ......
28205 ......
28033 ......
28237 ......
29419 ......
5.57
5.20
5.09
5.34
5.22
4.88
5.03
4.84
5.08
4.99
4.80
4.93
4.97
4.91
4.95
4.25
4.95
4.82
5.32
4.93
4.79
5.07
4.88
4.52
4.71
4.45
4.79
4.62
4.34
4.59
4.60
4.57
4.80
4.33
4.60
4.40
5.73
5.37
5.17
5.52
5.31
4.90
5.10
4.90
5.20
5.01
4.84
5.00
4.97
4.93
4.95
4.23
5.01
4.89
5.36
5.05
4.79
5.09
4.88
4.51
4.69
4.56
4.82
4.67
4.66
4.58
4.67
4.73
4.79
4.13
4.75
4.57
5.17
4.82
4.61
4.87
4.80
4.44
4.60
4.37
4.69
4.53
4.45
4.61
4.64
4.56
4.62
3.82
4.63
4.55
4.96
4.65
4.46
4.73
4.63
4.19
4.40
4.24
4.48
4.30
4.24
4.40
4.41
4.38
4.38
3.62
4.40
4.35
4.87
4.46
4.29
4.78
4.42
4.08
4.34
4.07
4.38
4.12
4.00
4.16
4.17
4.13
4.33
3.85
4.19
4.10
4.77
4.62
4.48
4.44
4.68
4.30
4.68
4.36
4.71
4.56
4.16
4.67
4.57
4.36
4.63
...
4.60
4.52
Note.—Line ﬂuxes in units of ergs cm2 s1.
The constant 0.328 is determined to match the solar observations. The Teff were determined from the BV colors
according to Bo¨hm-Vitense (1981). The values for
F =f ¼ ðd=RÞ2 , where R is radius and d is distance, are given
in Table 5A.
For the Hyades F stars we could also have used the
known radii and distances to calculate F/f. However, many
of the main-sequence stars lie above the zero-age main
sequence and probably have lower mass companions, as
indicated by their smaller visual magnitudes. The eﬀective
radii are thus larger. We therefore considered it to be appropriate to use the apparent visual magnitude of each star in
equation (1) to estimate the eﬀective R/d. This is not an
accurate treatment of the contribution from the unknown
companion, but it seems better than ignoring it.
3.2. IUE Observed Hyades F stars
In Table 6 we give the IUE ‘‘ image ’’ numbers for the
short wavelength spectra and the observer’s name. The IUE
data are generally more noisy than the HST data. Often the
weak lines cannot be measured on these spectra. The uncertainty limits for strong lines on well-exposed IUE spectra
are usually about 25% and may reach a factor of 2 for the
weak lines. We have therefore decided to use here only the
strong C ii (1335) and C iv (1550) lines, which can be
measured most accurately. For the relatively faint Hyades
late F stars, even these IUE measured line ﬂuxes appear to
be rather uncertain. In Table 6 we list the measured ﬂuxes, f,
for these lines and also the surface ﬂuxes, F. Unfortunately,
Ly cannot be measured on the IUE spectra because of the
large geocoronal contamination.
No. 2, 2002
SPECTRAL TYPE F5 IN HYADES F STARS
947
TABLE 6
Measured Fluxes for IUE Observed Hyades Stars
HD
Image
(SWP)
Observera
log f þ 14
C ii
log f þ 14
C iv
log f þ 14
1735
log F =f
log F
C ii
log F
C iv
log F
1735
27397 ......
28294 ......
28677 ......
24357 ......
26462 ......
27901 ......
29169 ......
26015 ......
26911 ......
27429 ......
25102 ......
26737 ......
26345 ......
28568 ......
28911 ......
27534 ......
29225 ......
28483 ......
27731 ......
30869 ......
27991 ......
28608 ......
26784 ......
27808 ......
27691 ......
28394 ......
30810 ......
30311 ......
27859 ......
27874
10247
45934
45936
10204
45935
48879
40466
10205
10219
45928
48877
48878
9856
45917
48887
48888
52363
48860
45927
9877
48854
45941
48872
15320
15294
48846
45909
9854
E. B. V.
E. B. V.
E. B. V.
E. B. V.
E. B. V.
E. B. V.
E. B. V.
C. I.
E. B. V.
E. B. V.
E. B. V.
E. B. V.
E. B. V.
A. W.
E. B. V.
E. B. V.
E. B. V.
E. B. V.
E. B. V.
E. B. V.
A. W.
E. B. V.
E. B. V.
E. B. V.
A. W.
A. W.
E. B. V.
E. B. V.
A. W.
<0.67
1.09
1.37
1.28
1.30
1.36
1.40
1.50:
0.86:
1.27
1.16
1.03:
0.88:
1.10
1.03
> 0.56
0.91
0.37
0.61
1.23
0.88
0.71:
0.73
0.70
0.71
1.09
0.66
0.54
<0.28
<0.59
1.66
1.52
1.39
1.37
1.53
1.62
1.11
1.33
1.61
1.40
1.24
1.17:
1.20
1.09
1.16
1.07
0.90
0.93
1.41
0.90
0.69
?
0.99
0.83
1.66
0.79
0.72:
<1.02
0.87
?
1.50:
1.66
1.65
1.65:
?
1.54
1.30:
1.68
1.11
0.84:
0.93
1.00
0.91
0.75
0.95
0.46
0.48
1.07
0.69
0.25:
1.60:
0.23:
0.24
1.00
0.40
0.01
0.4
18.06
18.14
18.17
18.12
18.01
18.08
18.10
18.07
18.18
18.09
18.18
18.45
18.27
18.22
18.27
18.32
18.26
18.42
18.45
18.07
18.09
18.46
18.35
18.35
18.29
18.30
18.19
18.35
18.53
<4.73
5.23
5.54
5.40
5.31
5.44
5.50
5.57
5.04
5.36
5.34
5.48
5.15
5.32
5.30
> 4.9
5.17
4.79
5.06
5.30
4.97
5.17
5.08
5.05
5.00
5.39
4.85
4.89
<4.81
<4.65
5.80
5.73
5.55
5.38
5.61
5.72
5.18
5.51
5.71
5.58
5.69
5.44
5.42
5.36
5.48
5.33
5.33
5.38
5.48
4.99
5.15
?
5.34
5.12
5.96
4.98
5.07
<5.55
4.93
?
5.67
5.78
5.66
5.73
?
5.61
5.48
5.77
5.29
5.29
5.20
5.22
5.18
5.07
5.27
4.88
4.93
5.14
4.78
4.71
5.95
4.58
4.53
5.30
4.59
4.34
4.13
a
4.
E. B. V. = E. Bohm-Vitense; C. I. = C. Imhoﬀ; A. W. = A.Walker.
DEPENDENCE OF EMISSION LINE FLUXES
ON BV
4.1. Dependence of Chromospheric Emission Line Fluxes
on BV
Our HST observations are restricted to stars with
BV > 0:40, because earlier spectral types have been previously observed with IUE. But none of the chromospheric
lines except for the Mg ii lines can be accurately measured
on IUE spectra. We therefore have chromospheric data only
for stars with BV > 0:40.
The measured chromospheric emission line ﬂuxes, f, for
stars observed with the HST are given in Table 3. The ﬂux
ratios F/f and the surface line ﬂuxes, F, are given in Table
˚ are blends of three lines at
5A. The O i lines around 1300 A
˚
1302.2, 1304.9, and 1306.0 A, which cannot be separated on
our low-resolution spectra. The chromospheric C i (1657)
lines are also blends of several weak lines. The C i lines at
˚ are also quite weak and cannot always be measured.
1560 A
We did not apply corrections here for the interstellar (and
perhaps circumstellar) absorption, although three mediumresolution HST/STIS spectra (to be discussed in a later
paper) show that for Ly corrections of up to perhaps a factor of 2 should be applied. The O i lines also seem to be
aﬀected by interstellar absorptions, but the corrections cannot be determined from the low-resolution spectra. For the
Mg ii lines the interstellar medium (ISM) absorption corrections are on the order of 20%–30%. Since all our program
stars are in the Hyades, we guess that for a given line the
ISM corrections are rather similar for diﬀerent stars (see
Redﬁeld & Linsky 2001) and that the comparison of the
chromospheric emission line ﬂuxes for the diﬀerent stars,
measured on low-resolution spectra, does not depend critically on these corrections.
In Figure 3 we show the BV dependences of the surface
ﬂuxes for the chromospheric lines of single stars. For C i we
˚ blend. As
used the weak 1560 line rather than the 1657 A
to be expected, the weak S i (1473) and especially the very
weak C i (1560) lines show a large scatter owing to the
uncertain measurements. All other chromospheric lines
show a pronounced decrease in ﬂux between
0:42 < BV < 0:45, in agreement with the behavior of the
Mg ii lines reported previously (Paper I). This decrease is
parallel to a general decline in stellar rotation velocity as
shown in Figure 4, although for a given BV we see a large
spread in v sin i. The only exception to this behavior is in the
S i (1473) line, which shows no reliable variation. The ionization energy of S i is lower than that of the other atoms
and ions shown, and the S i line is likely to be formed in the
lowest chromospheric layers.
For BV > 0:46 the line ﬂuxes increase slightly by somewhat diﬀerent amounts for diﬀerent lines. For BV between
0.47 and 0.55, the ﬂuxes appear to be independent of BV.
For BV > 0:55, they may perhaps decrease again, but this
is indicated by only one star.
4.2. Dependence of Transition Layer Line Fluxes on BV
In Figure 5 we show for single stars the dependences of
the diﬀerent transition layer line ﬂuxes on BV. The lines
948
BO¨HM-VITENSE ET AL.
Fig. 3.—Surface line ﬂuxes for the chromospheric lines, shown as a function of BV for the single stars observed with HST. The lines are ordered
roughly according to the height of formation. The v sin i values (in km s1)
for the diﬀerent stars are given in the top panel. For the weak C i and S i
lines, the error bars are estimated to be 0.25 dex for C i and 0.15 for S i. For
the strong lines the measuring errors are less than 0.05 dex, but the ﬂuxes
are uncorrected for interstellar absorption.
Vol. 569
Fig. 5.—Log of the surface line ﬂuxes for the transition layer lines,
shown as a function of BV for the HST observed single stars. A steep
decrease is seen around BV ¼ 0:43. The v sin i values for the IUE measured stars are given in the C ii line panel. For the HST measured ﬂuxes the
error bars are less than 0.05 dex; for the IUE measured ﬂuxes they are
estimated to be 0.1 dex.
originating in the highest layers are shown at the top of the
ﬁgure, and those originating at the chromospheric boundary are plotted on the bottom. For the strong C ii and C iv
lines we can make use of the IUE measurements for
BV < 0:450. For the early F stars the ﬂuxes are independent of BV, as shown by the C ii and C iv lines. Again a
steep decline of the surface ﬂuxes is seen for BV between
0.42 and 0.45, as is seen for the chromospheric lines, but the
magnitude of the decrease is larger for the transition layer
lines than for the chromospheric lines. For larger BV, up
to BV 0:47, the line ﬂuxes for the single stars again
recover slightly. For BV > 0:47, they remain again essentially independent of BV.
4.3. Dependence of the BV 0:43 Flux Decrease on
Ionization Energy
Fig. 4.—Here log v sin i is shown as a function of BV. For the velocities,
two points of steep decreases in the distribution are recognized, one starting
at BV 0:41 and a smaller one probably at BV 0:56. The plus signs
refer to IUE observed stars and the crosses to HST observed ones. The
boxes around the symbols indicate binaries.
Comparing Figures 3 and 5 we realize that the magnitude
of the ﬂux decrease for BV between 0.42 and 0.45 is diﬀerent for lines originating at diﬀerent temperatures in the
chromospheres and transition layers. In order to quantify
this ﬂux decrease, we determined the ﬂux ratios of the diﬀer-
No. 2, 2002
SPECTRAL TYPE F5 IN HYADES F STARS
Fig. 6.—Log of the ﬂux ratios at the top and bottom of the steep ﬂux
drop at BV 0:43, shown as a function of the ionization energy for the
line-emitting ion. The relevant ions are indicated at the diﬀerent points.
Except for C i and S i the error limits are less than 0.1 dex. For S i the error
limit is about twice as large, and for C i it is 0.4 dex. The decrease in ﬂux
increases with increasing ionization energy, i.e., with increasing height in
the transition layers. He ii does not ﬁt in with the other ions. For the X-rays
the log of this ﬂux ratio is 0.6 (not shown).
ent lines in the single stars HD 27561 with BV ¼ 0:412
and HD 27848 with BV ¼ 0:450. In Figure 6 these ratios
are plotted as a function of the ionization energy v of the
observed ion, which is a measure of the temperature at
which the line is formed. While there is some scatter in the
data points (because the next lower ionization energy is also
important), there is convincing evidence that the discontinuity in the ﬂuxes increases for increasing ionization energies,
which means for increasing temperatures. If we tentatively
relate this to a possible change in the heating mechanisms,
we conclude that the importance of the very eﬃcient heating
mechanism(s) working for the early F stars decreases very
rapidly for stars with BV between 0.42 and 0.45, and the
decrease is larger the higher the temperature of the layer
under consideration. If another heating mechanism takes
over for BV slightly larger than 0.45 , then this one, as
compared to the early F stars, leads to relatively lower emission line ﬂuxes the higher the temperature is for the lineemitting region under consideration.
949
Fig. 7.—Log of the ﬂux ratio F(BV ¼ 0:516)/F(BV ¼ 0:450), shown
as a function of the ionization energy of the line-emitting ion for the HST
observed stars. The He ii value does not ﬁt in with the transition layer lines.
Error limits are the same as in Fig. 6.
BV < 0:45. In Figure 7 we show the relation between this
ﬂux ratio F(BV ¼ 0:516)/F(BV ¼ 0:450) and v, which
means the dependence of this ﬂux ratio on the temperature
of line formation. Even though at ﬁrst sight there appear to
be diﬀerences, this ratio is for most of the transition layer
lines rather independent of v, as shown in Figure 7. The
apparently large value for the very weak C i lines has a very
large uncertainty; we omit it in the discussion. The chromospheric Mg ii and Ca ii lines show an increase by about 70%
(see Paper I), but most of the chromospheric and transition
region lines show only a 25% increase. For the transition
layers, the Si iv lines with a 50% increase and the He ii lines
with a 100% increase are the exceptions. The behavior of the
He ii lines resembles those of the Mg ii and Ca ii lines.
A larger v dependence is seen for the ratios of
F(BV ¼ 0:412) to F(BV ¼ 0:516), as seen in Figure 8.
This, of course, reﬂects only the strong dependence of the
ﬂux ratio F(BV ¼ 0:412) to F(BV ¼ 0:450) and the
small variation seen for the ﬂux ratios of F(BV ¼ 0:45)/
F(BV ¼ 0:516). In this ﬁgure we also show the point for
the X-ray ﬂuxes, which was arbitrarily placed at ¼ 110 eV
and will be discussed later. Figure 8 shows rather clearly the
similarity of the behavior of the chromospheric line ﬂuxes,
the He ii line ﬂuxes, and the X-ray ﬂuxes.
4.4. The Flux Increase between BV ¼ 0:45 and
BV ¼ 0:516
4.5. The Line Flux Ratios
The ﬂux ratio for the diﬀerent line ﬂuxes for stars with
BV ¼ 0:45 and those with BV near 0.516 (we use the
average of the ﬂuxes for HD 26784 and HD 27808) is a
measure for the recovery of the ﬂuxes after the drop for
If the heating mechanisms change at spectral type F5, as
ﬁrst suggested by Simon & Drake (1985), we might at this
spectral type (which means around BV 0:45) also
expect a change of the temperature stratiﬁcations in the
950
BO¨HM-VITENSE ET AL.
Vol. 569
Fig. 10.—Same as Fig. 9, but for the silicon ions. Only HST data were
used.
Fig. 8.—Log of the ﬂux ratio F(BV ¼ 0:412)/F(BV ¼ 0:516), shown
as a function of the ionization energy of the line-emitting ion. The X-ray
value was arbitrarily placed at 110 eV. Neither the X-ray value nor the He ii
value ﬁt in with the high-energy transition layer line values. Error limits are
the same as in Fig. 6.
the C iii (1175) lines as a function of BV. In Figure 10 we
have plotted the line ﬂux ratios of the two Si iv (1393.8,
1403.8) lines to the Si iii (1206.5) line as a function of BV.
˚ are very weak and there(The Si ii lines at 1526 and 1533 A
fore rather uncertain.) From Figures 9 and 10 we see that all
line ﬂux ratios of the transition layer lines to the chromospheric ones decrease between BV ¼ 0:42 and
BV ¼ 0:45. For larger BV values, they increase very
slowly. The C ii lines originate at the bottom of the transition layer. The C iv/C ii line ﬂux ratio shows a rather small
decrease at BV 0:43, showing that for these lines the
BV dependence is nearly the same. On the other hand, the
C iv/C i ratio shows a very steep decrease at this BV.
Obviously, the BV dependence of the emission line ﬂuxes
for the transition layer lines is diﬀerent from the ones for the
chromospheric lines. The change apparently occurs in the
layer where the Si ii and C ii lines are formed.
chromospheres and transition layers. This should become
observable by comparing line ﬂuxes of given elements for
lines originating at diﬀerent heights, which means at diﬀerent temperatures, in these outer atmospheric layers.
In Figure 9 we show the line ﬂux ratios for the C iv
(1550) lines to the C i (1657, 1560), the C ii (1335), and
About 20 years ago (Bo¨hm-Vitense 1982) the question
was raised whether the line ﬂux increases observed for several F stars were due to an increase in rotation velocities or,
Fig. 9.—Logs of the line ﬂux ratios for the diﬀerent carbon ions, shown
as a function of BV for the line ﬂux ratios C iv/C i (1560) (squares),
C iv/C i (1657) (diamonds), C iv/C iii (1175) (triangles), and C iv/C ii
(1335) (crosses, HST; plus signs, IUE observations). The diﬀerence in the
ﬂux drop is largest for the C iv/C i lines, as also indicated by Fig. 6. (The
uncertainty for each point is 0.25 dex.) The diﬀerence is smallest for C iv/
C iii. For these ﬂux ratios the uncertainty for each HST point is less than
0.1 dex, and for the IUE points it is about 0.15 dex.
Fig. 11.—Logs of the Ly and O i surface line ﬂuxes, shown as a function
of BV. The plus signs refer to Ly, the crosses to the blend of the three
˚ . The boxes around the symbols indioxygen lines at 1302, 1304, and 1306 A
cate binaries. The average ﬂuxes of the binaries are not higher than for the
single stars, except for the region around BV ¼ 0:44. For this BV
region, the BV colors of the binaries are probably reddened owing to the
unresolved cooler companion.
5. THE INFLUENCE OF BINARITY ON THE
EMISSION LINE FLUXES
No. 2, 2002
SPECTRAL TYPE F5 IN HYADES F STARS
951
rather, due to the binarity of many F stars. We can now
answer this question, at least, for the Hyades F stars.
In Figure 11 we show again the dependence of the Ly
and O i emission line ﬂuxes on BV, but in these plots we
have included the binaries ( points in boxes). As is to be
expected, the points for the binaries scatter more than the
points for the single stars. We suspect that this is a result of
the slight color changes owing to the companions and the
somewhat uncertain procedure for the calculation of the
surface ﬂuxes for the binaries. Figure 11 clearly shows, however, that there is no indication of higher ﬂuxes for the
binaries; if anything, the binaries have lower ﬂuxes, except
for the stars with BV around 0.43, where the BV colors
of the binaries with cooler companions seem to be too red
by about D BV ¼ 0:02. A similar eﬀect is present for the
transition layer lines, as seen below.
6. DEPENDENCE OF EMISSION LINE FLUXES
ON v sin i
6.1. Uncertainties in v sin i
When discussing the v sin i dependence of the emission
line ﬂuxes, we also have to look at the accuracy of the v sin i
values that were taken from the data collection of Uesugi &
Fukuda (1982). They round the v sin i values to 0’s or 5’s at
the end, which means they consider the data to be uncertain
by at least 3 km s1. For F stars this corresponds to the photospheric turbulence. This is a large fraction of v sin i if
v sin i < 10 km s1. Values of 5 km s1 are often only upper
limits. For v sin i > 30 km s1 the uncertainties are probably
near 10%, as judged by the values given by diﬀerent authors.
We really would need to study the dependence of the emission line ﬂuxes on the rotational velocities, vrot . The uncertainty in sin i introduces an even larger uncertainty than the
one of v sin i, although statistically there are more large sin i
values than small ones, with the average sin i being about
0.8. When comparing diﬀerent lines of the same star, they
have, of course, the same v sin i and vrot .
Fig. 12.—Dependence of the log of C iv and C ii surface line ﬂuxes on
v sin i for stars with BV ¼ 0:427 0:003. For the C iv lines crosses indicate HST observations and plus signs IUE observations. For C ii the
squares indicate HST observations and the diamonds IUE observations.
The C iv line ﬂuxes depend more strongly on v sin i than the C ii line ﬂuxes.
This shows that the increase in ﬂuxes with v sin i is not due to larger ﬁlling
factors. The error bars for the HST measured ﬂuxes are less than 0.04; for
the IUE measured ﬂuxes they are about 0.1 dex.
than the chromospheric Mg ii lines. (For the star HD 26345
we have an IUE ﬂux measurement that is much larger than
the HST measured ﬂux. We suspect that this is a calibration
problem.)
For the C ii and C iv lines we ﬁnd
log F ðC iiÞ ¼ 4:99 þ 0:0068ðv sin iÞ ;
ð2Þ
log F ðC ivÞ ¼ 4:99 þ 0:0096ðv sin iÞ :
ð3Þ
For the chromospheric atomic lines we do not see the
strong temperature dependence for this BV range. Therefore, in Figure 13 we have combined the data for all stars
with 0:42 < BV 0:451. For the C i lines ( plus signs) and
6.2. The v sin i Dependence for Stars with
0:42 < BV 0:451
We saw above that for the BV range from 0.42 to 0.45,
the transition layer emission line strengths change steeply.
For the same range the maximum v sin i also decrease
steeply. In this range Teff and maximum v sin i are correlated, as seen in Figure 4. Because of the correlation of v sin i
with Teff we cannot uniquely distinguish between the
dependences on v sin i and Teff unless we look at each value
of Teff separately. This reduces the number of stars in each
bin to a very small number. As seen in Tables 1 and 2, there
are, however, four stars with 0:424 BV 0:429, i.e., in
a very small temperature range. In Figure 12 we have plotted the logarithms of the ﬂuxes for these stars as a function
of v sin i. The two straight lines are eye-ﬁtted linear regression lines for the C ii and the C iv lines, for which we also
have fairly reliable IUE measurements, although their
uncertainties are still at least twice as large as the HST measured ones. The C iv line ﬂuxes increase more rapidly with
v sin i than the C ii lines, which originate in deeper layers.
This is in agreement with the ﬁndings of Simon, Herbig, &
Boesgaard (1985), who also found that for the mainsequence G stars the high-excitation lines, like the C iv lines,
decrease faster with decreasing v sin i (or Rossby number)
Fig. 13.—Dependence of the log of surface line ﬂuxes for the atomic lines
on v sin i for stars with BV ¼ 0:427 0:003. Crosses are for the S i lines
and plus signs for the C i lines. The boxes around the symbols indicate
binaries. For v sin i < 35 km s1 all lines do not show a ﬂux increase for
increasing v sin i, but they do show an increase with increasing v sin i for
v sin i > 35 or 40 km s1. For Ly and O i the uncertainty for the ﬂux is 0.04
dex, for S i it is 0.1 dex, and for C i it is 0.2 dex.
952
BO¨HM-VITENSE ET AL.
S i lines (crosses), the data appear to be best represented by
a relation independent of v sin i for v sin i < 35 km s1 and
an increase for the higher two v sin i points (although this
latter suspicion rests on two binary stars only, it is strengthened when looking at the v sin i dependence of the transition
layer lines). It seems that for these chromospheres we are
dealing with the combined action of two heating mechanisms. The ﬁrst one, which we call heating mechanism I, is
independent of v sin i or even less eﬃcient for larger v sin i
than for lower ones. The second one, which we call mechanism II, is very ineﬃcient for small v sin i but becomes more
eﬃcient for larger v sin i. For the low chromospheric layers,
mechanism I is dominant for small rotation velocities, but
mechanism II appears to dominate for v sin i > 35 km s1.
However, at this point, we probably cannot completely
exclude that within the limits of error (about 0.05 dex for
Ly and the O i lines and about 0.1 dex for the S i and 0.2
dex for the C i lines), equations (2) and (3) may also hold for
these chromospheric lines.
Vol. 569
6.3. The Dependence of Emission Line Fluxes on v sin i for
Early F stars with BV 0:42
Fig. 14.—Dependence of the chromospheric Ly, O i (1300), and C i
(1657) surface line ﬂuxes on v sin i for the two HST observed stars with
BV < 0:42 (e.g., for the early F stars). They seem to decrease with
increasing v sin i, but not much can be concluded from these two stars.
Error bars are the same as Fig. 13.
For BV 0:42 we have only two HST observed single
stars with BV ¼ 0:412 and 0.420. The one with the smaller
v sin i has the larger Ly ﬂux. The same is true for the O i
and C i line ﬂuxes, as seen in Figure 14. This is in agreement
with the Mg ii and Ca ii line observations but carries little
weight because of the small number of stars and the
unknown sin i. (There is, of course, the slight possibility that
the star HD 27561, with v sin i ¼ 15 km s1 and a high line
ﬂux, may have a small sin i and actually a much higher vrot .)
For this BV range a fair number of stars were observed
with IUE. We can use these C ii and C iv emission line ﬂuxes
to study the v sin i dependence for the transition layer lines.
The results are seen in Figure 15. The IUE observed single
Fig. 15.—Logs of the surface line ﬂuxes for the C iv (1550) lines (top) and C ii (1335) lines ( bottom), shown as a function of v sin i for stars with
BV < 0:422. The left-hand plots are for the single stars, and the right-hand plots for binaries. The colons in the bottom right plot indicate uncertain measurements. For comparison, the straight lines in the top plots show the relations given by eq. (3) and the ones in the bottom plots those for eq. (2). For v sin i < 80
km s1 they match for the single stars the measured points rather well, although a somewhat less steep relations would match even better. The C ii lines of the
binaries do not follow this relation. The crosses stand for HST measured points and the plus signs for IUE measured values. For v sin i > 80 km s1 the line
ﬂuxes decrease, except for the C ii lines of the binaries.
No. 2, 2002
SPECTRAL TYPE F5 IN HYADES F STARS
953
early F stars all have BV < 0:40 and v sin i > 50 km s1.
For the single stars and for v sin i between 35 and about 100
km s1, an increase of the ﬂuxes is seen that for the C ii lines
is consistent with the gradient of equation (2) and for the
C iv lines consistent with the gradient of equation (3). For
larger v sin i > 80 km s1 the ﬂuxes seem to decrease for
increasing v sin i. This was also found for ﬁeld mainsequence stars by Rutten & Schrijver (1987). For G dwarfs
in young clusters Ayres (1999; Ayres et al. 1996) also ﬁnds
no further increase of the C iv line ﬂuxes for v sin i > 80 km
s1, for which he has, however, only one star in the Per
cluster. He attributes this to a ‘‘ saturation ’’ eﬀect, discussed
also by Vilhu & Rucinski (1983). We do not think that for
the Hyades early F stars we see a saturation eﬀect, because
we see decreasing ﬂuxes with increasing v sin i, which we do
not think can be explained as a saturation eﬀect. Rather, we
believe that we see for our stars a nonlinear relation between
the logs of the emission line ﬂuxes and v sin i. (We see no
compelling reason why the relation between the logs of the
transition layer line ﬂuxes and v sin i has to be linear.) Ayres
ﬁnds saturation for log½F ðC ivÞ=F ¼ 4:5. For our early F
Hyades stars the log of this ﬂux ratio for v sin i ¼ 80 km s1
is about 5.3, or about 0.8 dex lower than for the Per late
F star. It is quite possible that we see diﬀerent eﬀects.
For our binary targets the C iv line ﬂuxes also show a general increase for v sin i < 100 km s1, consistent with the
relation (3), plotted as a solid line in Figure 15 as a comparison. For the C ii lines we also plotted the line corresponding
to equation (2) for comparison; however, the C ii line ﬂuxes
for binaries are not represented by this line.
6.4. Chromospheric Emission Lines of Late F Stars with
BV > 0.46
We saw in x 5 that for BV > 0:46 the emission line
ﬂuxes for many of the lines are nearly independent of BV.
Therefore, we can combine these stars to study the dependence on v sin i. In Figure 16 we show the dependence of the
chromospheric line ﬂuxes on v sin i for stars with
BV > 0:46. For the single stars there does not seem to be
any dependence on v sin i for the Ly and the O i lines; however, within the limits of error, an increase with v sin i, as
expected from equations (2) and (3) probably cannot be
excluded. For the S i and C i lines there appears to be
decreasing ﬂux for increasing v sin i, but this conclusion is
based mainly on one star. We need more observations, especially of stars with a larger range in v sin i, in order to reach
deﬁnitive conclusions. For this BV range such stars are
not available in the Hyades. For binaries there seems to be
some ﬂux increase with increasing v sin i. For smaller separations perhaps rotation and the activity may be increased,
but presently we have no information about the orbital
parameters for the observed binaries.
6.5. Transition Layer Emission Lines of Late F Stars with
BV > 0.46
In Figure 17 we show the v sin i dependences for the transition layer lines. They are ordered according to the ionization energy for the line-emitting ion. Except for the C ii and
C iv lines we have reliable line ﬂuxes only from the HST
observations. The N v lines originate in the highest layers
studied here. The layer of origin for the He ii lines is not
clear. The high ionization plus excitation energy of the line
Fig. 16.—Log of the atomic line ﬂuxes, shown as a function of v sin i for
the single stars (top) and for the binaries (bottom) for the stars with BV
larger than 0.465. The asterisks refer to Ly, the triangles to the O i lines,
the squares to the C i lines, and the crosses to the S i lines. The line ﬂuxes for
the single stars appear to be independent of v sin i. The S i and C i line ﬂuxes
even seem to decrease with increasing v sin i, but we have too few targets to
say anything deﬁnitive. For the binaries an increase of the ﬂuxes with
increasing v sin i is likely.
suggests a high layer, but the relatively low ionization
energy prevents the existence of He ii at high temperatures.
The BV dependence of the line ﬂux is similar to the ones
for chromospheric lines. The ﬂux drop around BV ¼ 0:43
is also close to the one for chromospheric lines.
Figure 17 (left) shows that for single stars the transition
layer line ﬂuxes generally do not seem to increase with
increasing v sin i, but we have data only for v sin i up to 20
km s1. The scatter is large enough that at least for the carbon lines an increase with v sin i, as given by the gradients in
equations (2) and (3), cannot be excluded, as shown by the
solid lines drawn in for the C ii and C iv lines. For binaries
(Fig. 17, right) we see a trend of larger ﬂux increases for
increasing v sin i, but this trend depends mainly on one star.
The measurements are inconclusive because of the small
number of targets. In order to reach ﬁrm conclusions, we
need observations of more stars with a larger range in v sin i
as are available in the Pleiades.
954
BO¨HM-VITENSE ET AL.
Vol. 569
6.6. The Emission Measures
The emitted line ﬂuxes are proportional to the emission
measures
Z
Z
dh
EM ¼ n2e dh ¼ n2e
d log Te ;
ð4Þ
d log Te
where Te is the electron temperature. The integral has to be
extended over D log Te 0:30, which corresponds approximately to the temperature interval in which the ion is in the
state of ionization emitting the line under consideration.
We have calculated the emission measures for the single
Hyades F stars, as described earlier (Bo¨hm-Vitense &
Mena-Werth 1992). The results are given in Table 7.
Fig. 17.—Dependence of the log of the transition layer surface line ﬂuxes
on v sin i, for the single stars on the left-hand side and for binaries on the
right-hand side for the stars with BV > 0:465. The ﬁgures are ordered
according to the atmospheric height of the line formation region. The
straight lines for the C ii and C iv lines have the gradients given by eqs. (2)
and (3). For the N v, the He ii, and the Si iv lines, no increase with v sin i is
seen for the single stars, but within the limits of error a slight increase with
v sin i cannot be excluded. The number of observed targets is too small to be
sure. The binaries show generally increasing ﬂuxes for increasing v sin i.
Fig. 18.—Log of EM. plotted for single stars as a function of the log of
the electron temperatures Te for the strong transition layer lines of C ii,
C iv, Si iv, and N v originating approximately at log Te indicated at the top.
The diﬀerent symbols refer to diﬀerent stars with diﬀerent v sin i (in km s1).
The diﬀerent panels are for diﬀerent ranges of BV. The derived emission
measures for the Si iv lines were reduced by D log EM ¼ 0:35 to better
match the other emission measures. The straight lines are eye-ﬁtted linear
regression lines. The emission measure gradient is smaller for the early F
stars (top) in comparison with the other plots. For the slightly cooler F
stars, the highly ionized ions experience a larger ﬂux drop than the less ionized ones.
No. 2, 2002
SPECTRAL TYPE F5 IN HYADES F STARS
955
TABLE 7
Emission Measures for Single Stars
HD
BV
C ii 1334+35
C iv 1548+50
Si iv 1393+1402
N v 1238+1242
28.83
28.48
28.27
28.11
28.19
28.27
28.30
28.29
28.21
27.99
27.58
27.41
27.20
27.24
27.28
27.29
27.31
27.22
HST Targets
27561 ......
28736 ......
26345 ......
27848 ......
28608 ......
26784 ......
27808 ......
28237 ......
29419 ......
0.412
0.420
0.427
0.450
0.472
0.514
0.518
0.560
0.576
28.67
28.30
28.19
27.98
28.09
28.03
28.07
28.05
27.92
28.20
27.84
27.64
27.37
27.48
27.47
27.44
27.48
27.36
IUE Targets
28677 ......
24357 ......
27901 ......
29169 ......
26737 ......
28911 ......
26784 ......
0.338
0.357
0.378
0.380
0.424
0.429
0.514
28.64
28.50
28.46
28.60
28.57
28.40
28.16
Figure 18 shows for diﬀerent BV ranges the dependence
of the emission measures on the temperature in the lineforming layers. As discussed earlier (Bo¨hm-Vitense &
Mena-Werth 1992), the emission measures for the Si iv lines
always come out too high in comparison with those of the
carbon lines. Therefore, we have again reduced the log EM
(Si iv) for the plots by 0.35. (Formerly we had reduced them
by 0.5.) Figure 18 demonstrates that in the transition layers
the emission measures for the late F stars decrease faster
with increasing height than for the early F stars
(d log EM=d log Te 0:9 for the early F stars and
d log EM=d log Te 1:2 for the late F stars). This also
remains true if we consider the uncertainties of the EM. The
uncertainties in the atomic constants enter in the same way
for all plots and therefore cancel when we compare the gradients of the EM. The weak N v lines have the largest measuring uncertainties, which are, however, estimated to be less
than 0.1 dex. Also the diﬀerence in the gradients hardly
changes if we disregard the N v lines. The emitted energy is
lower for the late F stars as discussed above. The steeper
decrease of the emission measures probably means a relatively steeper temperature gradient in the higher layers,
reducing the line-emitting volume.
7. COMPARISON WITH X-RAY LUMINOSITIES
7.1. The BV Dependence of the X-Ray Luminosities
The X-ray luminosities of the Hyades stars have been
studied by Stern, Schmitt, & Kakabka (1995) by means of
the ROSAT satellite and by Micela et al. (1988). From the
paper by Stern et al. we have extracted the data for the
Hyades F stars. The v sin i values were taken from Uesugi &
Fukuda (1982) and the BV values from the Hipparcos
catalogue. In Table 8 we have collected the data. While most
stars with BV < 0:30 are not X-ray sources, we have
included a few ﬂux measurements for stars with
BV < 0:30. For the stars with BV < 0:30 we do not
expect intrinsic X-ray emission because convection and
28.16
27.98
28.08
28.19
28.16
27.83
related phenomena generally only occur for stars with BV
larger than 0.30. Since the Hyades X-ray sources with
BV < 0:30 are all binaries, we believe that the X-ray emission is due to the cooler companions, which are fainter in
the optical than the A9 or F0 stars but have measurable
X-ray ﬂuxes. Of course, for very few stars a white dwarf
companion could also be the source for the X-rays.
Stern et al. (1995) did not give the surface X-ray ﬂuxes,
but from the measured ﬂuxes they calculated the X-ray
luminosities that the stars would have if they all were at a
distance of 45 pc. We have collected the distances given in
the Hipparcos catalogue and determined the X-ray luminosities that the stars must have if they are at these distances.
They are also given in Table 8 and are called LX ðdÞ. [They
are diﬀerent from the LX ðdÞ listed by Stern et al. because
those authors used diﬀerent distances.] For distances larger
than 45 pc, the LX ðdÞ must be larger than the LX (45 pc)
given by Stern et al. (1995) in order to give the measured
X-ray ﬂuxes.
In Figure 19 we have plotted the log LX ðdÞ as a function
of BV. Diﬀerent symbols are used for single stars and
binaries. We also have distinguished spectroscopic binaries
and binaries indicated by an ‘‘ I ’’ in the catalog of Perryman
et al (1998). These binaries are probably, on average, wider
pairs because they were previously known as binaries, and
many of them were resolved. We therefore expect many of
them to behave almost like single stars, unless they are also
spectroscopic binaries. We will call these stars the ‘‘ wide ’’
binaries, although probably not all of them actually are.
Since for the spectroscopic binaries the companions are
bright enough to be seen in the spectra, the colors of these
stars may be distorted by the companions. Figure 19 looks
like a scatter diagram. The binaries show, on average, somewhat larger X-ray luminosities than the single stars (but
rarely more than a factor of 2), as already noticed from
Stern et al. (1995). We suspect that this is mainly due to the
additional ﬂux by the cooler companion.
The X-ray ﬂuxes show a strong dependence on temperature if we include the binaries with BV between 0.3 and
TABLE 8
Basic Data for Hyades X-Ray Sources
HD
vBa
Hipparcos
BV
log LX
(45 pc)
d(Hipparcos)
(pc)
log LX
(d )
v sin i
(km s1)
29388 ......
37147 ......
33254 ......
28546 ......
28052 ......
28556 ......
28226 ......
27176 ......
27397 ......
31236 ......
27749 ......
33204 ......
29375 ......
27628 ......
28294 ......
28677 ......
24357 ......
26462 ......
25570 ......
27901 ......
29169 ......
26015 ......
26911 ......
27429 ......
27561 ......
18404 ......
25102 ......
28736 ......
26737 ......
26345 ......
28568 ......
28911 ......
27524 ......
13871 ......
27534 ......
29225 ......
27848 ......
31845 ......
28406 ......
27483 ......
27731 ......
28483 ......
28608 ......
30869 ......
21847 ......
27383 ......
27991 ......
26784 ......
27691 ......
27808 ......
28394 ......
30809 ......
28363 ......
30738 ......
28205 ......
30810 ......
28033 ......
27406 ......
28237 ......
30311 ......
30676 ......
20430 ......
29419 ......
30589 ......
104
168
130
83
141
84
67
24
30
126
45
131
103
38
68
89
6
14
160
53
100
11
20
32
37
154
8
90
16
13
85
94
35
157
36
101
51
128
78
34
44
81
86
124
21589
26382
23983
21039
20713
21036
20842
20087
20219
22850
20484
24019
21588
20400
20873
21137
18170
19554
18975
20614
21459
19261
19877
20255
20357
13834
18658
21152
19789
19504
21053
21267
20349
10540
20350
21474
20567
23214
20948
20284
20491
21008
21066
22607
16517
20215
20661
19796
20440
20557
20935
22566
20916
22524
20815
22550
20712
20237
20826
22221
22496
15304
21637
22422
0.129
0.237
0.249
0.258
0.262
0.263
0.270
0.277
0.283
0.292
0.310
0.311
0.312
0.315
0.325
0.338
0.357
0.360
0.371
0.378
0.380
0.397
0.400
0.404
0.412
0.415
0.417
0.420
0.424
0.427
0.428
0.429
0.434
0.439
0.441
0.442
0.450
0.450
0.451
0.456
0.462
0.471
0.472
0.502
0.503
0.509
0.509
0.514
0.518
0.518
0.526
0.527
0.536
0.536
0.537
0.543
0.557
0.560
0.560
0.560
0.563
0.567
0.576
0.578
28.531
<27.954
<27.90
<27.95
30.122
28.968
<28.30
28.954
29.130
<28.43
<28.67
28.591
28.431
28.643
28.591
28.756
28.903
<28.52
29.360
28.959
28.785
<28.82
28.903
29.137
29.176
29.111
29.233
28.959
29.130
28.869
29.507
29.246
28.964
28.568
28.785
29.413
28.556
28.875
28.740
29.299
28.982
29.248
29.049
29.250
28.724
29.274
29.487
29.253
29.430
29.303
29.339
28.875
29.049
29.049
29.029
29.500
28.708
29.065
29.233
29.380
29.212
28.568
29.152
28.724
45.89
53.58
53.94
44.35
47.94
45.79
47.96
54.80
44.82
68.17
47.24
54.71
45.54
45.73
54.29
44.94
41.43
38.63
35.97
49.02
44.25
47.02
44.43
47.35
51.39
31.84
39.34
43.28
55.19
43.07
41.19
43.86
51.15
40.11
50.43
43.50
53.36
43.31
46.32
45.87
49.90
50.15
43.55
41.82
48.90
42.97
46.58
47.44
46.62
40.87
43.01
58.34
48.59
51.81
45.81
49.68
46.43
44.90
47.21
38.08
43.55
49.50
44.25
50.81
28.548
<28.11
<28.06
<27.94
30.177
28.984
<28.36
29.125
29.127
<28.79
<28.71
28.761
28.442
28.657
28.745
28.755
28.831
<28.39
29.165
29.033
28.771
<28.86
28.892
29.181
29.291
28.810
29.116
28.924
29.308
28.831
29.430
29.223
29.075
28.468
28.884
29.384
28.704
28.842
28.763
29.136
29.072
29.342
29.021
29.187
28.796
29.234
29.171
29.299
29.461
29.196
29.299
29.101
29.116
29.172
29.045
29.585
28.735
29.625
29.275
29.235
29.184
28.651
29.138
28.830
80
115
20
35
195
95
90
105
100
110
15
30
130
30
100
100
60
10
55
150
80
30
50
145
15
35
55
35
70
20
55
40
90
20
40
40
30
25
20
10
30
20
20
25
?
20
20
5
10
<10
25
?
<25
10
10
5
<5
10
10
5
15
<5
<5
<5
29
57
19
40
48
77
143
75
121
65
122
62
31
66
113
119
1
105
118
Binarity
I SB
I SB
I SB
SB
I
I
I SB
I
SB
I
I
I SB
B
SB
I SB
I
I
SB
I
I
Member?
SB
I
I
SB
SB
I SB
I SB
I SB
I SB
SB
I SB
SB
I SB
SB
SB
SB
SPECTRAL TYPE F5 IN HYADES F STARS
957
TABLE 8—Continued
HD
vBa
Hipparcos
BV
25825 ......
29310 ......
27859 ......
27836 ......
28344 ......
28992 ......
26767 ......
28068 ......
28099 ......
30246 ......
27685 ......
27989 ......
10
102
52
50
73
97
18
63
64
142
39
58
132
19148
21543
20577
20553
20899
21317
19786
20719
20741
22203
20441
20686
24020
0.593
0.597
0.599
0.604
0.609
0.631
0.640
0.651
0.664
0.665
0.677
0.680
0.707
a
log LX
(45 pc)
28.763
29.199
29.033
29.724
29.049
28.886
29.188
29.093
28.903
28.580
29.149
29.295
28.613
d(Hipparcos)
(pc)
46.71
42.48
48.24
44.94
47.42
43.12
45.07
45.96
46.69
51.50
37.09
43.33
54.71
log LX
(d )
29.265
29.149
29.093
29.723
29.095
28.193
29.189
29.112
28.935
28.697
28.981
29.262
28.782
v sin i
(km s1)
?
5
5
?
?
?
?
?
?
?
?
?
?
Binarity
SB
B
I
SB
SB
I SB
SB
I
The van Bueren number.
0.35. Since no single stars in this color range are observed to
be X-ray sources, it is rather likely that for these latter
binary stars, the X-rays are due to cooler companions. (The
low X-ray ﬂuxes require then that the companions have
BV > 0:7 or larger.) If we include these early F star
binaries, then the X-ray ﬂuxes start out low at BV ¼ 0:30
and increase up to BV about 0.41. This is quite diﬀerent
from the transition layer lines, which start out high immediately for BV > 0:32 (see Table 6). If, on the other hand,
the binary X-ray sources with BV between 0.30 and 0.35
owe their X-rays to the companions, then X-ray emission
for the F stars starts only at BV ¼ 0:357. It then also starts
out at about the same level as the stars with BV ¼ 0:40.
We must then ask why the onset of the X-ray emission is
delayed as compared to the transition layer emission.
Since the spectroscopic binaries probably have their colors distorted by the companions, we have in Figure 20 plotted only the X-ray ﬂuxes for the single stars (not known to
be binaries, although a few may still be) and the ‘‘ wide ’’
binaries, for which the BV colors are probably not so distorted. Figure 20 shows the true dependence of intrinsic X-
ray luminosities on BV much more clearly. In spite of the
large scatter, a steep decrease in ﬂux between BV ¼ 0:42
and BV ¼ 0:45 is clearly recognized, with a ﬂux decrease
of about 0.6 dex, i.e., 0.2 dex smaller than observed for N v.
A slow recovery of the X-ray ﬂuxes occurs for BV > 0:45,
which reaches the ﬂux observed for BV ¼ 0:42 again at
BV ¼ 0:52. This is similar to what is seen for the He ii line
ﬂuxes, which we have also plotted in Figure 20 for better
comparison. It is also similar to what is seen for the Ca ii
and Mg ii emission line ﬂuxes as shown in Paper I but quite
diﬀerent from the behavior of the transition layer ﬂuxes as
demonstrated earlier in Figure 8.
It has been suggested that the He ii line emission at 1640
˚ , originating from a high-excitation level, might be excited
A
by coronal X-rays. The similar BV dependences of the
X-ray ﬂuxes and the He ii line ﬂuxes argue in favor of this
suggestion, but, of course, this similarity is only a necessary
but not a suﬃcient condition to show that the excitation
of the He ii emission may be due to X-rays. (Other chromospheric lines, like the Mg ii lines, show similar BV
dependences.)
Fig. 19.—X-ray luminosities for Hyades F stars are plotted as a function of BV. The asterisks show single stars, the triangles presumably ‘‘ wide ’’ binaries,
and the squares spectroscopic binaries.
Fig. 20.—BV dependence of the X-ray ﬂuxes of single stars and ‘‘ wide ’’ binaries is compared with the BV dependence of the He ii (1640) line. The two
distributions look rather similar and are also similar to the ones observed for the Mg ii lines.
Fig. 21.—The v sin i dependence of the X-ray luminosities for single stars, shown for diﬀerent BV ranges. An approximate average value of BV is given
in each plot. Based on the top plots for BV < 0:42, the X-ray luminosities decrease for increasing v sin i, as long as v sin < 80 km s1, similar to the observations for the Mg ii lines. For very large v sin i they may perhaps increase as shown by one star only with v sin i ¼ 150 km s1. For the cooler stars with
BV > 0:42, we ﬁnd increasing luminosities for increasing v sin i. The straight lines in the plots for BV ¼ 0:425 and 0.44 have gradients corresponding to
eqs. (3) and (2), respectively. The diﬀerent v sin i dependences show that the heating mechanisms for the coronae are diﬀerent for stars with BV < 0:42 and
BV > 0:42.
SPECTRAL TYPE F5 IN HYADES F STARS
959
7.2. The Dependence of the X-Ray Fluxes on v sin i
Figure 20 shows a large scatter of the LX ðdÞ even for the
presumably single stars. Is this due to diﬀerent rotational
velocities? In Figure 21 we show the v sin i dependences of
LX ðdÞ for single stars only, within diﬀerent ranges of BV.
The values of BV given in the plots are approximate average values for the BV range used for the plot. The top two
plots suggest that for BV < 0:42 the ﬂuxes decrease for
increasing v sin i if v sin i < 80 km s1 but may increase for
larger v sin i, as possibly indicated by one star only. If so,
this would be similar to our ﬁndings for the Mg ii lines.
For stars with BV between 0:42 and 0.45 there seems to
be a real increase of LX ðdÞ with increasing v sin i. The gradient of the solid line in the plot for BV 0:425 agrees
with the gradient found in equation (3), and the gradient of
the line in the plot for BV 0:44 agrees with the gradient
given in equation (2). For the stars with BV 0:56 an
increase of the X-ray luminosities with increasing v sin i is
possible, but the v sin i are no larger than 10 km s1 with
uncertainties of 3–5 km s1. No ﬁrm conclusion can be
drawn from those data.
It seems important to notice that for increasing BV,
starting at BV ¼ 0:42, the X-rays show the same v sin i
dependence as the transition layer lines, which means they
increase with increasing v sin i, while for smaller BV they
decrease with increasing v sin i. The change occurs at the
same BV for which the decrease of the v sin i starts.
7.3. Relation between X-Ray Fluxes and C iv Line Fluxes
Ayres (1999) and Ayres et al. 1996 ﬁnd mainly for G stars
and some late F stars a tight correlation between
RX ¼ LX =L and RðC ivÞ ¼ F ðC ivÞ=F , namely, RX /
½RðC ivÞ1:7 . At the top of Figure 22 we have plotted for the
single F stars for which we know
both LX and F(C iv) the
4 , where T
log RX as a function of log F ðC ivÞ=Teff
eff was
determined from BV according to Bo¨hm-Vitense (1981).
The points for which F(C iv) was determined from HST
spectra are indicated by crosses, the ones with IUE measured C iv line ﬂuxes with plus signs. We ﬁnd a scatter diagram even if we omit the uncertain IUE measured points for
BV around 0.45. We can ﬁnd some systematics when we
look at the BV distribution. The points for stars with
BV 0:424 have been enclosed in boxes, the ones for
0:427 BV 0:472 in triangles. For these two groups we
have drawn in relations RX / RðC ivÞ1:7 , as shown by the
dashed and dashed-dotted lines, which are perhaps reasonable ﬁts for each group, although with large scatter. For
stars with BV > 0:50 the points are shown in the upper
left corner of Figure 22. These stars come close to the temperature range of the Ayres study. We do not have enough
points to determine any proportionality. Figure 22 clearly
shows that for the Hyades F stars the relation between LX
and F(C iv) is very temperature-dependent. To better understand this diagram, we have with the dotted line connected
the data points in succession of increasing BV.
In spite of the large scatter seen in the top half of Figure
22, we ﬁnd a surprisingly tight correlation between
log½RX =RðC ivÞ and BV, which we found puzzling but
which becomes understandable when we follow the dotted
line in the top ﬁgure. For a certain color range this line goes
more or less parallel to the dashed or dash-dotted lines, corresponding to the Ayres relation, and then crosses over to
the next color range. In the bottom ﬁgure we see for BV
Fig. 22.—Top: Log of the X-ray luminosities LX =L, plotted as a function
4 . The crosses show stars for which F(C iv) was
of the R(C iv)= F(C iv)/Teff
determined from our HST observations, and the plus signs are for stars for
which F(C iv) was determined from IUE observations. Symbols for stars
with BV < 0:424 are enclosed in boxes, those for stars with
0:424 < BV 0:472 in triangles. The remaining stars in the upper left
corner have BV > 0:50. We see a scatter diagram. The data are connected
by dotted lines in succession of increasing BV. For the stars with
BV 0:424, the dashed line gives the approximate relation
RX / RðC ivÞ1:7 . The dash-dotted line gives the similar relation for the stars
with 0:424 < BV 0:472. Bottom: Log of the ratio of RX =R(C iv) for single stars is plotted as a function of BV. The staircase curve demonstrates
again that diﬀerent relations between LX and F(C iv) hold for diﬀerent
BV. For BV < 0:42 and BV > 0:472, the gradients of the curve are
diﬀerent, indicating diﬀerent relations between the heating mechanisms for
the transition layers and the coronae.
between 0.42 and 0.472 a staircase, but for larger and
smaller BV values we ﬁnd smooth relations, which have,
however, diﬀerent gradients. For the whole range RX
increases faster than R(C iv), as also indicated by the Ayres
relation, meaning the emission of the coronae increases relative to the one for the transition layers.
In Figure 23 the log of the ratio of RX =RðC ivÞ1:7 for the
Hyades F stars is plotted as a function of BV. This ratio
should be equal to a constant, say, A, if the Ayres relation
would also hold for the Hyades F stars. However, we see
that between BV ¼ 0:42 and 0.47 this ratio increases
steeply because the decrease of the C iv line ﬂuxes is much
larger than the decrease in LX . For late F stars the gradient
is much smaller and may tend toward 0 for the spectral type
G0.
The variations of this ratio may also be interpreted in
another way. We write
RX ¼ ARðC ivÞc ;
ð5Þ
or in logarithmic form
log RX ¼ log A þ c log RðC ivÞ ;
ð6Þ
960
BO¨HM-VITENSE ET AL.
Vol. 569
nent c. We can calculate the values of c for stars with diﬀerent BV, which will reproduce the observed values of Rx/
R(C iv). We use log A ¼ 12:7, as obtained above for c ¼ 1:7
and BV ¼ 0:60. (In this way we are sure to obtain c ¼ 1:7
for BV ¼ 0:60.) The values of c required by the observations are shown in Figure 23 at the bottom. The range of c is
rather small, but a steep change is seen between
BV ¼ 0:40 and 0.45.
8. DISCUSSION
8.1. The Onset of Convection and Emission in the Outer
Stellar Atmospheres
It is well known that for ﬁeld stars chromospheric and
transition layer emission is observed for main-sequence
stars with BV > 0:29 or 0.30 (see Bo¨hm-Vitense &
Dettmann 1980). Our observations of the Hyades F stars
conﬁrm this: no measurable emission was seen for HD
27397 with BV ¼ 0:28, while HD 28294 with
BV ¼ 0:325 shows transition layer emission lines at the
same ﬂux level as the cooler stars. (We can exclude that the
emission lines from HD 28294 come from the companion.
Its high line ﬂux would require the companion to also be an
early F star because later F stars have less ﬂux. However,
the absolute magnitude of HD 28294 precludes such a luminous companion. See Fig. 1.) Since convection is believed to
be ultimately responsible for the chromospheric temperature increase, this shows that the younger age and the higher
v sin i for the Hyades early F stars, as compared to the average ﬁeld stars, do not change the BV for which eﬃcient
convection sets in abruptly. Some change in BV for the
onset might perhaps have been expected since rapid rotation
is expected to decrease convective eﬃciency (see Chandrasekhar 1961). Apparently, rotation velocities of 80 km s1 do
not yet delay the onset of convection and strong chromospheric and transition layer emission, although v sin i > 80
km s1 reduces the transition layer emission line ﬂuxes (see
Fig. 15).
It is, however, quite interesting to realize that for single
stars the X-ray emission only sets in at BV 0:357. It
seems that at this temperature the depth of the outer convection zone is large enough to create enough energy in heating
mechanism I, presumably acoustic shocks, to also heat the
corona.
Fig. 23.—Top: For single Hyades F stars, log½RX =RðC ivÞc ¼ A is
shown as a function of BV for c ¼ 1:7. If the Ayres relation would also
hold for Hyades F stars, the ratio of RX to RðC ivÞ1:7 should be constant.
Instead we see for these stars a steep gradient of A for BV between 0.42
and 0.47. Bottom: Attributing the diﬀerences from the Ayres relation to differences in the exponent c we calculate for a constant log A ¼ 12:7 the values of c required by the observed ratios RX =R(C iv). The values derived for
the exponent c are shown as a function of BV. A steep decline is seen
between BV ¼ 0:42 and 0.47.
where c ¼ 1:7 for the Ayres relation. If this equation would
hold with constant A and constant c for all F stars, we
should be able to represent the points in the top plot of Figure 22 by a straight line. Obviously, this is not possible.
With straight lines we can only represent points in a given
BV range with c ¼ 1:7 and use diﬀerent values of A for different ranges in BV. The necessary range in A is shown in
Figure 23 at the top. We can, of course, also attribute the
mismatch with the Ayres relation to diﬀerences in the expo-
8.2. Diﬀerent Heating Mechanisms in Diﬀerent
Atmospheric Heights
For low chromospheric emission lines (i.e., S i, C i, Mg ii,
Ca ii) we ﬁnd for all Hyades F stars with v sin i less than
about 40 or 50 km s1 either no v sin i dependence of the
emission line ﬂuxes (i.e., S i, C i) or we ﬁnd decreasing ﬂuxes
for increasing v sin i (Mg ii and Ca ii, as discussed in Paper
I). We attribute this to a heating mechanism that depends
only on the convective velocities, such as acoustic shock
heating. This should not depend on v sin i except if rotation
becomes fast enough such that Coriolis forces inﬂuence the
convective motions (Chandrasekhar 1961). In that case the
convective velocities slow down. It is not clear yet whether
the decreasing ﬂuxes for increasing v sin i, observed for the
Ca ii and Mg ii lines for moderate v sin i, and also for the
early F star coronae, can also be attributed to the same heating mechanism or whether an additional one is needed.
No. 2, 2002
SPECTRAL TYPE F5 IN HYADES F STARS
For the transition layer lines we ﬁnd for early F stars a
diﬀerent dependence on v sin i, namely, increasing emission
line ﬂuxes for increasing v sin i. This must be due to a diﬀerent heating mechanism, which also seems to contribute to
chromospheric heating if the rotation velocities become
very large. This can explain the large Mg ii emission line
ﬂuxes and X-ray ﬂuxes for very high v sin i.
The increasing eﬃciency of this heating mechanism for
increasing v sin i indicates that it is of magnetohydrodynamic origin. This is contrary to the belief that magnetohydrodynamic eﬀects become important only for spectral
types later than F5, as suggested by Simon & Drake (1989)
and also by Schrijver (1993). Mullan & Mathioudakis
(2000) emphasize, however, that the ﬂare observed for the
star HR 120 with spectral type F2, shows that magnetic
activity is already present for this spectral type. Here we ﬁnd
that magnetic activity is already present for the earliest F0V
star, for which transition layer emission is observed.
8.3. Diﬀerent Heating Mechanisms for Early and
Late F Stars?
For early and late F stars we see no diﬀerence in the v sin i
dependence of the chromospheric lines of S i and C i, as
shown here, and Mg ii and Ca ii, as shown in Paper I. We
also see no guaranteed diﬀerences in the v sin i dependence
of the transition layer emission line ﬂuxes for early and late
F stars, although the decreased ﬂuxes for spectral types later
than F5, attributable to the lower v sin i, indicate a lower
contribution of the magnetohydrodynamic heating mechanism. For the late F stars several line ﬂuxes (N v, He ii) also
appear to be independent of v sin i, but the relatively small
number of single stars and the small range of v sin i available
for the late Hyades F stars prevents a ﬁrm conclusion.
The diﬀerent v sin i dependences of the chromospheric
lines and the diﬀerent transition layer lines can be understood if we consider the height dependence of the two heating mechanisms. In the lower layers the v sin i independent
chromospheric heating mechanism (acoustic shocks?) is the
main heating mechanism. With increasing height, the magnetohydrodynamic mechanism with a strong v sin i dependence becomes increasingly more important, leading to the
increasing v sin i dependence of the emission line ﬂuxes. For
BV > 0:42 mainly the transition layer heating mechanism
changes, leading to reduced line emission. The chromospheric heating mechanism changes very little. For stars
with BV > 0:42 we thus expect initially smaller v sin i
dependence for the transition layer lines than for stars with
smaller BV, although this may change for larger BV.
The diﬀerent ﬂux decreases for diﬀerent transition layer
lines originating at diﬀerent heights clearly show that the
decrease is not (only) due to a reduction in the surface ﬁlling
factor for the emission, which means it is not due to a decrease
in the surface area covered by active regions.
The decrease of the transition layer line ﬂux level can
mainly be attributed to the decreasing v sin i between
BV ¼ 0:42 and 0.45, but the v sin i dependence does not
seem to be quite strong enough to explain the full amount of
ﬂux decrease. This is especially true for the chromospheric
Mg ii lines, for which we found increasing ﬂuxes for decreasing v sin i, yet there is a ﬂux decrease between BV ¼ 0:42
and 0.45. There must be an additional reason for the
decreasing ﬂuxes, which does not work, however, for the
low chromospheric lines. This looks like a third heating
961
mechanism may be involved, whose eﬃciency also decreases
around BV ¼ 0:42. It is also possible that the changes are
due to a change in damping length for the heating energy
ﬂux, perhaps due to changes in the magnetic ﬁeld or density.
We have to leave it to future theoretical studies to investigate this problem.
According to Figure 21 the v sin i dependence of the Xray emission for early F stars is very diﬀerent from the one
for late F stars, showing that diﬀerent heating mechanisms
are responsible for the heating. For late F stars the coronae
are probably heated by the same mechanism that heats the
transition layers, which does not seem to be the case for the
early F stars.
8.4. The v sin i Dependence of the Emission Line Fluxes and
the Rossby Numbers
Noyes et al. (1984) found for G and late F main-sequence
ﬁeld stars that the Ca ii emission line ﬂuxes F =Teff are
roughly proportional to the Rossby number Ro1 (Ro =
rotation period/convective turnaround time) if young and
old stars are combined and provided that the convective
turnaround times are computed for a ratio of mixing length
to pressure scale height of 2. Looking only at the young
stars, the F =Teff are approximately proportional to Rox
with x between 0.7 and 0.5. Comparing main-sequence stars
with a given temperature this means F / vrot for old stars or
F / vxrot for young stars. For the Hyades F5 stars we measure that the ﬂuxes are roughly proportional to v sin ix , where
x 0:67 for the C ii lines and x 0:9 for the C iv lines, for
which the data scatter much more. These exponents also
approximate the measured v sin i dependence of the transition layer lines for the early Hyades F stars and possibly also
the late F stars. It is in reasonable agreement with the result
of Noyes et al. for young G and late F stars. This result
seems to conﬁrm that the transition layer heating mechanisms for young G stars and F stars, including the early F
stars, are the same. The correlation with the Rossby number
also indicates that for these stars magnetic activity governs
the heating mechanism II, i.e., the heating mechanism for
all F star transition layers, and probably the late F star
coronae.
8.5. Inﬂuence of Binarity
We ﬁnd that the average emission line ﬂuxes for binaries
are the same as for single stars. Binarity generally does not
increase the emission line ﬂuxes more than might be
expected owing to the additional but lower emission from
the companion. There is, however, a trend for increasing
dependence on v sin i. This may ‘‘ explain ’’ why for close
ﬁeld star binaries with high-rotation velocities the emission
increases strongly, although we do not know the reason for
it.
9. SUMMARY
We found that at the spectral type F5 most chromospheric lines show a steep decrease of emission line ﬂuxes
for BV between 0.42 and 0.45, parallel to the decrease of
the rotational velocities. We found that for early F stars the
transition layer emission increases for increasing v sin i. We
have thus used the decrease of the rotational velocities as an
explanation for the major part of the ﬂux decreases between
BV ¼ 0:42 and 0.45. This ﬂux decrease is larger for lines
962
BO¨HM-VITENSE ET AL.
of ions with higher degrees of ionization, indicating that the
rotation is less inﬂuential in the lower transition layers,
showing that another heating mechanism contributes also
to the heating of these lower layers. It is reasonable to
equate this one with the chromospheric heating mechanism,
presumably acoustic shock heating, which is not expected to
increase with increasing v sin i.
We have taken the observed strong v sin i dependence of
the early F star transition layer lines as an indication that
these transition layer lines are heated by magnetohydrodynamic eﬀects, contrary to the often expressed belief (for
instance, Schrijver & Zwaan 2000; Simon & Drake 1989)
that for the early F stars the transition layers are heated by
acoustic shocks. We have to remember that the absence of
rotational modulation does not necessarily show the
absence of magnetic activity. We will not detect rotational
modulation for small-scale ﬂux diﬀerences distributed
homogeneously over the whole stellar surface. We speculate
that for the thin convection zones of the early F stars with
BV < 0:42 dynamo activity is present, but magnetic ﬁelds
and structures are small and nearly evenly distributed.
The line ﬂux decreases around BV ¼ 0:43, attributed to
the decrease in vrot , show that for the late F stars the contribution of the magnetohydrodynamic heating is reduced. On
the other hand, star spot activity shows large-scale magnetic
activity and explains the observed rotational modulation
observed for late F stars. For cooler stars the photospheric
densities increase, and the convection zones become thicker.
This leads to larger dynamo activity and larger magnetic
ﬁelds in larger, but fewer spots, which lead to rotational
modulation. The average line ﬂux is lower than for the early
F stars.
We still need to explain why the rotational velocities
decrease for BV around 0.43. It is generally assumed that
this is due to stellar winds that carry away angular momentum (see, for instance, Schrijver & Zwaan 2000). But why do
stellar winds suddenly become so eﬀective for braking? They
must have been even more eﬀective for the cooler mainsequence stars because those were braked much faster. We
would expect that strong stellar winds of the Parker type
(Parker 1958), as seen for the Sun and perhaps expected for
F5 stars, require high coronal temperatures, which lead to
high coronal emission. We do not see that for the Hyades F
stars with 0:42 BV 0:45. On the contrary, for single
stars the X-ray emission also decreases between
BV ¼ 0:42 and 0.45. If strong winds do the braking, there
must be additional or other properties responsible for the
strong braking.
In this context it seems important that for BV 0:42,
the X-ray emission shows a positive dependence on v sin i, in
Vol. 569
agreement with the transition layer lines, while for smaller
BV the X-ray luminosities decrease for increasing v sin i.
This means that for the early F stars magnetohydrodynamic
eﬀects do not reach into the corona; the magnetic structures
are small (see also Giampapa & Rosner 1984). Only when
the magnetic structures become large enough to reach into
the corona do the magnetic ﬁelds in the winds increase the
braking eﬃciency of the winds (see, for instance, Schrijver &
Zwaan 2000, chapter 13.3) to cause the steep drop in rotation velocities. While for F0 stars the hydrogen and helium
convection zones are separated, they merge for slightly
cooler stars, and the depth for the outer convection zone
suddenly becomes much larger. At which BV this happens
depends on the eﬃciency of the convective energy transport.
(See Bo¨hm-Vitense & Nelson 1976.) The abrupt changes
observed for BV around 0.43 tell us that perhaps it happens at this BV. We expect then that because of the sudden
increase of the depth of the convection zone, the magnetic
ﬁelds and the geometric scales of the magnetic structures
increase, star spots form, and open and closed loops reach
into the corona. Braking by the winds then increases rather
abruptly because the wind is frozen in the magnetic ﬁeld out
to the Alfve´n point. Any wind then acquires and carries
away much more angular momentum than for stars with
BV 0:42. The deeper the convection zone, the larger the
magnetic ﬁelds and structures, and the stronger the braking
will be.
If indeed the dynamo action and the size of magnetic
structures increase for BV 0:42, should we not expect
an increase in chromospheric and transition layer emission
line ﬂuxes and also coronal emission in comparison with the
early F stars, rather than less emission and a dip in the Xray emission? If our arguments are correct, then the geometric scale of the magnetic structures is mainly important for
the braking, while for the emission line ﬂuxes the eﬃciency
of the magnetohydrodynamic heating is important. It then
seems that the scale of the magnetic structures and the heating eﬃciency for the transition layers and coronae are not
necessarily correlated.
We are very much indebted to the staﬀ of the Space
Telescope Science Institute for their continuous help and
support to obtain the observations. Without their dedicated
help this research would not have been possible. We also
proﬁted from a long and informative discussion with Tom
Ayres. We are also very grateful for ﬁnancial support by
NASA grant GO-07389-01-98A to E. B.-V. and NASA
grant GO-07389.02-98A to K. G. C., without which this
study could not have been completed.
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