A&A 366, 178-196 (2001)
DOI: 10.1051/0004-6361:20000205
E. Rodríguez 1 - M. Breger 2
1 - Instituto de Astrofísica de Andalucía, CSIC,
PO Box 3004, 18080 Granada, Spain
2 -
Institut für Astronomie, Universität Wien, Türkenschanzstr. 17,
1180, Austria
Received 21 September 2000 / Accepted 3 November 2000
Abstract
We present a comprehensive analysis of the properties of the pulsating
Scuti and related variables based mainly on the content of the
recently published catalogue by Rodríguez et al. (2000a,
hereafter R00). In particular, the primary observational properties such as
visual amplitude, period and visual
magnitude and the contributions from the Hipparcos, OGLE and MACHO long-term
monitoring projects are examined. The membership of these variables in open
clusters and multiple systems is also analyzed, with special
attention given to the
Scuti pulsators situated in eclipsing binary
systems.
The location of the
Scuti variables in the H-R diagram
is discussed on the basis of HIPPARCOS parallaxes and uvby
photometry. New borders of the classical instability are presented.
In particular, the properties of the
Scuti pulsators
with nonsolar surface abundances (SX Phe,
Boo,
Pup,
Del and classical Am stars subgroups) are examined.
The Hipparcos parallaxes show that the available photometric
uvby
absolute magnitude calibrations by Crawford can be applied correctly
to
Scuti variables rotating faster than
100 kms-1 with normal
spectra. It is shown that systematic deviations
exist for the photometrically determined absolute magnitudes, which
correlate with
and
.
The photometric
calibrations are found to fit the
Boo stars,
but should not be used for the group of evolved metallic-line A stars.
The related
Dor variables and the pre-main-sequence
Scuti
variables are also discussed.
Finally, the variables catalogued with periods longer than 0
25 are
examined on a star-by-star basis in order to assign them to the proper
Scuti, RR Lyrae or
Dor class.
A search for massive, long-period
Scuti stars similar to
the triple-mode variable AC And is also carried out.
Key words: stars: variables: Scuti - stars: oscillations -
stars: fundamental parameters
Scuti-type variables are pulsating stars of short periods (<
)
located in the lower part of the Cepheid instability strip, with luminosities
ranging from the zero-age-main-sequence (ZAMS) to about 2 mag above the main
sequence with spectral types ranging from about A2 to F2. A handbook
with astrophysical reviews and discussions of these variables has now become
available (Breger & Montgomery 2000).
Recently, an updated catalogue of
Scuti stars has been
published (Rodríguez et al. 2000a) covering observational
information available up to January, 2000. This new catalogue
contains 636 variable stars, of which more than 50% have only been discovered
during the last six years. The majority of these new variables were found
in the data of long-term monitoring projects such as the Hipparcos
mission (ESA, 1997), OGLE (Udalski et al. 1994, 1995a,b,
1996, 1997) and MACHO (Alcock et al. 2000)
projects. Nevertheless, even ignoring the contributions from these three main
projects, a large number of new variables have been discovered since 1994 by
individual groups. Hence, an immense
amount of new information on
Scuti variables has been made
available during the last few years.
The main aim of this work is to present a comprehensive analysis on the present
status of the Scuti and related stars based mainly on the content of
the R00 catalogue. In Sect. 2, the observational properties of the stars in
the catalogue are presented together with an analysis of
the
Scuti variables which are members of double or multiple systems and
open clusters. In Sect. 3, we analyse the location of these variables in the
Hertzsprung-Russell (H-R) diagram on the basis of uvby
photometry
and parallaxes. In Sect. 4, several interesting groups of peculiar and/or
related variables are also studied.
The statistical distribution of the amplitudes, periods and apparent
magnitudes of the known Scuti stars is examined in
Figs. 1 to 3. Tables 1 and 2 compare the new average values
with those presented in the earlier catalogue on
Scuti stars
(Rodríguez et al. 1994, hereafter R94).
The histograms also show the distributions for the
variables discovered by the Hipparcos mission (ESA, 1997) and OGLE
(Udalski 1994, 1995a,b, 1996, 1997) and MACHO
(Alcock et al. 2000) projects are also studied.
The distribution of the amplitudes of the known Scuti stars
does not reflect the true distribution among the group: severe
selection effects exist, especially among the variables
discovered by the OGLE and MACHO projects, for which a limit of
0
1 are found (see Fig. 1). This is connected with extreme
faintness of these stars (see Fig. 3) leading to lower measurement
precision. The HIPPARCOS measurements present an intermediate case
between the OGLE/MACHO and the classical telescope surveys with millimag
precision.
Figure 1 shows that the
amplitudes are strongly dependent on the origin of the discovery. The
average amplitude of the variables discovered during high-accuracy variability
surveys is only about 0
01 (see Fig. 3 of Breger 1979). On
the other hand, nearly all the
Scuti variables discovered during
the OGLE and MACHO projects have amplitudes larger than 0
1. The
variables discovered during the Hipparcos mission have amplitudes between
those of the MACHO/OGLE projects and the "other'' sources. Our interpretation
that these differences are a reflection of the accuracy of the variability
measurements is confirmed by the brightness distribution of the stars, which
is shown in Fig. 3: the new
Scuti variables discovered by Hipparcos
satellite are much brighter than those discovered by OGLE and MACHO projects.
The higher accuracy of the measurements of the brighter stars makes it
possible to
discover smaller amplitudes. Consequently, the new MACHO/OGLE variables severely
distort the true statistics of the amplitude distribution among
Scuti
stars.
In spite of the great number of
Scuti-type pulsators with large
amplitudes discovered by the OGLE and MACHO projects, the majority of the
known
Scuti variables display small amplitudes.
In fact 51% of these variables have visual amplitudes smaller
than 0
05. Additionally, the number of low amplitude variables increases
nearly exponentially with decreasing amplitude. In particular, nearly 30%
of them show amplitudes smaller than 0
02. This percentage is about
45% if we ignore the contributions from Hipparcos, OGLE and MACHO.
The increased data used for Fig. 1 solves a puzzle from a similar
histogram of amplitudes (Fig. 2) in the R94 catalogue: the
apparent amplitude gap between 0
1 to 0
3.
In the R94 catalogue there were only 14 variables
in this interval (
), that is, 5% of the
full sample. However, in the R00 catalogue there are 94 variables (which
means, 15% of the total sample). Hence, it seems that there is not a strict
separation in two groups relative to the amplitude for the
Scuti-type
pulsators (low
amplitude variables with
and high amplitude variables with
). Note that the disappearance of the gap has been due to
the contributions from OGLE and MACHO projects. In fact, in both cases the
number of variables discovered, with visual amplitudes between 0
1 to
0
3 and larger than 0
3, is very similar (26 and 25 from OGLE; 41
and 42 from MACHO). Thus, it seems that the gap shown in the R94 catalogue
was due to a selection effect.
The period histogram (Fig. 2) shows that the majority of the
variables have short periods and that the number of variables decreases
with increasing period. This result is expected since the longer
period stars are more evolved because of the existence of a period-luminosity
relation.
Because of the relatively long life-time on the main sequence,
the probability of finding an evolved star is smaller than the probability
of detecting a main-sequence star. However, there exists another
reason as well. The majority of short-period stars on the main sequence
have very small amplitudes, which are more difficult to detect
in stars with longer periods because of the very high quality photometric data
required.
This selection effect had a
very strong influence in the
diagram of R94 catalogue.
However, our Table 1 (full sample) shows that this effect almost disappears
for the variables with periods longer than 0
05. This is caused
by the Hipparcos, OGLE and MACHO data, where the selection effect is not
present. However, in the other data (not containing the three data sets)
the selection effect still exists.
Figure 3 shows the distribution in visual amplitude.
Presently there exists a large number of Scuti stars with visual
apparent magnitudes greater than 16
0 mostly due to
the OGLE and MACHO contributions. In particular, 59
Scuti
variables fainter than 18
0 have been discovered during the last six years.
Table 2 shows that the selection effect "that the
amplitude is larger when the star is fainter'', still remains.
This effect was very pronounced in the older data used in the R94 catalog,
where the stars fainter than
had large amplitudes.
The explanation lies in the difficulty in detecting small amplitudes
in faint variables. For the new MACHO data, an amplitude jump to large
amplitude is evident only around
.
The increased astronomical
precision in the new data, which leads to the detection of small-amplitude
variability in faint stars, is also evident in the line marked "Other'' in
Table 2. This decrease is due to the very
low-amplitude
Scuti variables recently discovered during surveys
in open clusters using
high quality CCD photometry (Balona & Laney 1995;
Frandsen et al. 1996; Frandsen & Arentoft 1998a; Kim
et al. 1999).
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Figure 1:
Distribution of the variables in the catalogue (N) as a function of
the visual amplitude (![]() |
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Figure 2: Distribution of the variables in the catalogue (N) as a function of the period (P). The contributions from Hipparcos, OGLE and MACHO are also shown |
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0
![]() ![]() |
0
![]() ![]() |
0
![]() ![]() |
0.d20-0
![]() |
0
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|
Hipparcos | - | 0.098 | 0.066 | 0.062 | 0.074 | 0.065 |
OGLE | 0.160 | 0.294 | 0.317 | 0.320 | 0.343 | - |
MACHO | - | 0.317 | 0.303 | 0.226 | - | - |
Other | 0.024 | 0.084 | 0.213 | 0.226 | 0.352 | 0.238 |
Full sample | 0.025 | 0.151 | 0.202 | 0.186 | 0.278 | 0.214 |
R94 | 0.029 | 0.097 | 0.254 | 0.278 | 0.444 | 0.262 |
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Figure 3: Distribution of the variables in the catalogue (N) as a function of the visual magnitude (V). The contributions from Hipparcos, OGLE and MACHO are also shown |
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2![]() ![]() |
4![]() ![]() |
6![]() ![]() |
8![]() ![]() |
10![]() ![]() |
12![]() ![]() |
14![]() ![]() |
16![]() ![]() |
18![]() ![]() |
|
Hipparcos | - | 0.020 | 0.056 | 0.088 | 0.093 | - | - | - | - |
OGLE | - | - | - | - | - | - | - | 0.264 | 0.322 |
MACHO | - | - | - | - | - | 0.160 | 0.153 | 0.303 | 0.348 |
Other | 0.043 | 0.030 | 0.044 | 0.096 | 0.238 | 0.191 | 0.373 | 0.514 | 0.570 |
Full sample | 0.043 | 0.030 | 0.047 | 0.094 | 0.229 | 0.189 | 0.351 | 0.340 | 0.332 |
R94 | 0.043 | 0.032 | 0.052 | 0.120 | 0.281 | 0.436 | 0.544 | 0.594 |
Figure 4 shows the distribution in apparent magnitude
of the Scuti variables known to be
part of binary or multiple stellar systems. The R00 catalogue lists
86 such variables (62 with CCDM identification; Dommanget &
Nys 1994). This represents only 14% of the total sample of known
Scuti stars. Only five variables are fainter
than
:
three eclipsing binaries AB Cas (
,
,
Rodríguez et al. 1998), Y Cam (
,
,
Broglia & Conconi 1984) and V577 Oph (
,
,
Diethelm 1993) and the two recently discovered spectroscopic binaries
PL 43 (
,
d) and 29499-057 (
,
d; this case
is not completely confirmed yet) (Preston & Landolt 1999). Hence,
multiplicity is catalogued for 22% of all the
Scuti known up to
10
0. This percentage is very low because more than 50% of the stars
are expected to be members of multiple systems. It would probably be
incorrect to intepret the statistics to imply that multiplicity inhibits
pulsation since the observational bias against detecting multiplicity is very
high. One reason for this bias immediately comes to mind:
both pulsation and multiplicity lead to radial-velocity variability.
Long and accurate observations are required to separate the two effects.
These are generally not available.
Pulsating stars in eclipsing binaries are important for accurate
determinations of fundamental stellar parameters and the study of tidal
effects on the pulsations.
However, so far very few such systems with a Scuti variable as one
of its components have been discovered. Only 9 cases seem to be well
established and they are
listed in Table 3 and shown in Fig. 4.
Two of these variables have been very recently discovered:
AS Eri by Gamarova et al. (2000) and R CMa by Mkrtichian &
Gamarova (2000).
In all the cases listed in the table, the pulsation amplitude is
small. This, in combination with the large variations produced by the binarity,
makes it very difficult to detect
Scuti-like variations in such
systems.
One excellent example is the bright Algol-type eclipsing binary system RZ Cas.
This star was discovered to be variable by Muller in 1906. Later, Dugan
(1916)
obtained a complete light curve using a visual polarizing photometer and
derived the orbital elements for the system.
During the last two decades, unusual changes in the light curves have
been detected, leading to a number of different interpretations,
but the
Scuti-type variability of the primary component of this
system has only recently been discovered by Ohshima et al.
(1998). This demonstrates how difficult it is to detect low amplitude
Scuti-like variations in this type of binary systems: only
measurements of very high quality lead to successful results.
Besides the 9 stars listed in Table 3, two other
Scuti variables
(UX Mon and DL Uma) had been pointed out to belong to eclipsing binary
systems being listed in the R94 catalogue. However, these two
cases seems to be wrong. Olson & Etzel (1995) do not confirm the
Scuti-like variations claimed by earlier authors for the primary
component of the Algol-type system UX Mon. In the case of DL UMa, new
photometry carried out by Rodríguez et al. (2000c)
does not confirm the
existence of eclipsing binarity with an orbital period of 0
42 as given in
Kholopov et al. (1987).
Table 3 lists, for each system, the spectral types of both components
(ST1, ST2), visual magnitude V, period and amplitude of the pulsation
(
,
), orbital period and depths of the primary and
secondary minima (
,
,
)
and type of
binarity. In addition, the
Scuti component and the number (N) of
pulsation frequencies found are also listed. In all the cases, these 9 systems
are Algol-type binaries from an observational point of view as described in
the GCVS (Kholopov et al. 1985). However, from an evolutionary point
of view, only the six first cases listed in Table 3 are Algol-type systems
where the hotter component is an A-F type star, whereas the secondary
component is
much cooler and much less luminous. In these cases the
Scuti variable
is the primary component. On the other hand, AI Hya and RS Cha are early
F-type and
A-type double-lined eclipsing binaries, respectively, where the secondary
components display
Scuti-type variations. Finally, V577 Oph is
probably an A-type double-lined eclipsing system. In this case,
the
Scuti component would be the primary. In all the cases
only one or two reliable frequencies of pulsation have been determined, while
more frequencies of pulsation can probably still be found. In two cases
(RZ Cas and Y Cam),
multiperiodicity and nonradial pulsation are confirmed (Ohshima et al.
1998; Rodríguez et al. 2000b; Broglia & Conconi
1984). On the other hand,
monoperiodicity and radial pulsation have been found in AB Cas (Rodríguez
et al. 1998).
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Figure 4:
Distributions of ![]() |
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Star | ST1 | ST2 | V |
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Type | ![]() |
N | Source |
(mag) | (d) | (mag) | (d) | (mag) | (mag) | component | ||||||
AB Cas | A3V | K1V | 10.17 | 0.0583 | 0.05 | 1.3669 | 1.63 | 0.10 | Algol | 1 | 1 | 1 |
RZ Cas | A3V | K0IV | 6.26 | 0.0156 | 0.02 | 1.1953 | 1.50 | 0.07 | Algol | 1 | 2 | 2 |
WX Eri | A5 | K0V | 9.39 | 0.1645 | 0.04 | 0.8233 | 0.55 | 0.15 | Algol | 1 | 2 | 3 |
AS Eri | A3V | K0IV | 8.30 | 0.017 | 0.01 | 2.6641 | 0.70 | 0.10 | Algol | 1 | 1 | 4 |
R CMa | F0V | K1IV | 5.70 | 0.047 | 0.01 | 1.1359 | 0.58 | 0.08 | Algol | 1 | 1 | 5 |
Y Cam | A9IV | K1IV | 10.54 | 0.0665 | 0.04 | 3.3055 | 1.70 | 0.10 | Algol | 1 | 2 | 6, 7 |
AI Hya | F0 | F2 | 9.90 | 0.1380 | 0.02 | 8.2897 | 0.56 | 0.50 | F-type | 2 | 1 | 8 |
RS Cha | A8IV | A8IV | 6.05 | 0.08 | 0.01 | 1.6699 | 0.62 | 0.45 | A-type | 2 | 1 | 9 |
V577 Oph | A? | A? | 11.01 | 0.0695 | 0.05 | 6.0791 | 0.65 | 0.55 | A-type? | 1 | 1 | 10, 11 |
Pulsation provides an additional method to detect multiplicity through a study of the light-time effects in a binary system. This method generally favors high-amplitude variables with only one or two pulsation periods (which tend to be radial). Several decades of measurements are usually required to study these (O-C) residuals in the times of maxima.
A good example of such a system is the star SZ Lyn.
Moffett et al. (1988) found that the ephemeris of
this star can be well described with two terms: 1) a secular increase of the
period at a rate of 9 10-9 d/y and 2) an orbital period, as component of
a binary system, of
= 1118 d with a semiamplitude of 0
006.
Furthermore, five other high-amplitude
Scuti stars have been
suggested as
members of binary systems. They are listed in Table 4 together with their
probable orbital periods.
For a number of other variables, the interpretation of the (O-C)
residuals in terms of binarity is more controversial e.g., for
CY Aqr. The (O-C) variations are very complex with interpretations
in terms of period jumps (Rodríguez et al. 1995;
Breger & Pamyatnykh 1998). However, Zhou & Fu
(1998) propose that the period changes of this star
are a consequence of a continuously increasing period combined
with the light-time effect in a binary system with an orbital period of
years. Another controversial example VZ Cnc, where Arellano et
al. (1994) indicate binarity as one possibility to explain the
remaining residuals in their O-C analysis, while Fu & Jiang (1999b)
point out that multiperiodicity with more than two frequencies might be
the correct interpretation.
The question of binarity also plays an important role in the asteroseismological
interpretation of small-amplitude, nonradially pulsating Scuti stars
studied through large, multisite campaigns. The main frequencies
of the star
Tau (Breger et al. 1989) with
peak-to-peak amplitudes of 3 millimag or more show the (O-C)
values between -2 and +2 mins predicted from the 141-day orbit. The
companion is a less luminous star on the main sequence, also inside
the instability strip. The frequency spectrum of
2 Tau also
includes a number of very small-amplitude modes at values of approximately
twice that of the main frequencies. From an asteroseismological point of view,
it is important to determine whether these modes originate in the
companion or are part of a second excited frequency range of excited modes
in the primary star of
Tau.
Star | V |
![]() |
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Source |
(mag) | (d) | (mag) | (years) | ||
AD CMi | 9.39 | 0.1230 | 0.30 | 30.0 | 1 |
SZ Lyn | 9.44 | 0.1205 | 0.51 | 3.1 | 2 |
BE Lyn | 8.82 | 0.0959 | 0.39 | 6.4 | 3 |
KZ Hya | 9.96 | 0.0595 | 0.80 | 9.3 | 4 |
DY Her | 10.46 | 0.1486 | 0.51 | 43.0 | 5 |
BS Aqr | 9.40 | 0.1978 | 0.44 | 31.7 | 6 |
Cluster | log(age) | [Me/H] | EB-V | N | <V> | <P> |
![]() |
![]() |
Source |
(mag) | (mag) | (d) | (mag) | (km s-1) | |||||
![]() |
7.90 | 0.06 | 0.09 | 3 | 8.98 | 0.048 | 0.013 | 92 | 1, 9 |
Pleiades | 7.92 | 0.03 | 0.04 | 5 | 8.12 | 0.036 | 0.013 | 69 | 1, 9 |
Hyades | 8.80 | 0.18 | 0.01 | 8 | 4.81 | 0.084 | 0.016 | 114 | 1, 2 |
NGC 1817 | 8.87 | -0.27 | 0.30 | 7 | 13.72 | 0.049 | 0.009 | - | 1, 2 |
NGC 2264 | 6.99 | -0.15 | 0.06 | 2 | 10.03 | 0.150 | 0.040 | - | 1, 8 |
Melotte 71 | 9.0 | -0.29 | 0.20 | 4 | 13.61 | 0.098 | 0.028 | - | 5, 8 |
NGC 2516 | 7.79 | -0.07 | 0.10 | 1 | 10.70 | 0.060 | 0.020 | - | 1 |
Praesepe | 8.84 | 0.14 | 0.02 | 14 | 7.83 | 0.070 | 0.018 | 139 | 1, 2 |
NGC 2682 | 9.72 | 0.00 | 0.07 | 2 | 11.60 | 0.056 | 0.025 | 73 | 1, 2 |
NGC 3496 | 8.6 | - | 0.52 | 3 | 13.81 | 0.178 | 0.023 | - | 6 |
Coma Ber | 8.69 | -0.03 | 0.01 | 2 | 6.92 | 0.048 | 0.013 | 115 | 1, 8 |
NGC 5999 | 8.60 | - | 0.45 | 1 | 17.71 | 0.150 | 0.240 | - | 8 |
NGC 6134 | 8.84 | 0.28 | 0.36 | 7 | 12.93 | 0.096 | 0.014 | - | 4 |
NGC 6882 | 9.16 | -0.02 | 0.08 | 2 | 10.36 | 0.055 | 0.030 | - | 8 |
Melotte 227 | 8.57 | - | 0.04 | 1 | 8.01 | 0.055 | 0.010 | - | 8 |
NGC 7062 | 8.7 | - | 0.47 | 1 | 14.01 | 0.040 | 0.010 | - | 3 |
NGC 7245 | 8.5 | - | 0.40 | 2 | 14.80 | 0.098 | 0.015 | - | 3 |
NGC 7654 | 8.2 | - | 0.58 | 1 | 14.42 | 0.278 | 0.020 | - | 3 |
NGC 7789 | 9.3 | -0.26 | 0.24 | 1 | 14.06 | 0.087 | 0.030 | - | 7 |
Scuti-type pulsating stars in open clusters provide to the astronomer
an important tool to test stellar structure and evolution theory because the
same distance, age, initial chemical abundance and interstellar reddening can
be assumed for all the stars in the same cluster. This provides strong
constraints on the evolutionary status and the mode identification of the
observed frequencies. The classical studies of stars inside bright,
nearby clusters using photomultiplier detectors have recently been
extended to more clusters by utilizing CCD detectors. Especially for
stellar clusters, CCD cameras have important advantages because of the large
number of
cluster stars which can be monitored simultaneously on the same frame, higher
quantum efficiency and the ability to work under nonphotometric weather
conditions.
Some examples of such successful surveys are the open clusters
NGC 7789 (Jahn et al. 1995), NGC 3496 (Balona & Laney 1995),
NGC 6134 (Frandsen et al. 1996), NGC 7062, NGC 7245 and NGC 7654
(Viskum et al. 1997),
NGC 1817 (Frandsen & Arentoft 1998a) or Melotte 71 (Kim et al.
1999).
A difficulty with statistics of stars in clusters concerns their
membership. For the R00 list, the main source of information has been the
open clusters catalogue of Mermilliod (1995), although other
sources have also been consulted to determine whether or not a particular
Scuti star is a cluster member. In Table 5, the clusters are listed
together with their relevant parameters such as log(age),
[Me/H] and EB-V. In addition, the number (N) of
Scuti variables found to be members of each cluster together with the
mean values of the visual magnitude, period, visual amplitude and
are also listed. Altogether, Table 5 lists 67
Scuti stars distributed
in 19 open clusters. The great majority of these variables lie on the main
sequence together with the other cluster stars. Some variables have
already evolved off
the main sequence, i.e., the variables KW 204, KW 284 and KW 348 in
Praesepe (Hernández 1998) and IFA 161 in NGC 6134 (Frandsen et al.
1996; Bruntt et al. 1999).
The variables EX Cnc and EW Cnc in NGC 2682 (Gilliland & Brown
1992) and V10 in NGC 7789 (Jahn et al. 1995; Mochejska &
Kaluzny 1999) are blue stragglers; the variables W 2
and W 20 in NGC 2264 probably are in a pre-main sequence status (Breger &
Pamyatnykh 1998). For V16 in NGC 1817 (Frandsen & Arentoft
1998a) the location in the H-R diagram is uncertain, but the star
may be a nonmember.
As it can be seen from Table 5, the Scuti variables in open clusters
tend to show short periods and very low amplitudes. One exception from the
low amplitudes exists: the variable V32 in NGC 5999
(Pietrzynski et al. 1998) shows
,
but its membership is not fully confirmed. If the membership
is confirmed, then the star promises to be a very interesting object of study.
A few clusters contain
Scuti variables with
periods longer than 0
1 (NGC 2264, NGC 5999, NGC 3496, NGC 7654).
Special attention is drawn from the variable V10 in NGC 7654 (Viskum et al.
1997; Choi et al. 1999). This variable seems to be a
member of the cluster on the main sequence, but shows a very long period
(0
278), typical of an evolved star. If it is a
pulsating star, then g-modes are present.
Values of the projected rotational velocities,
,
are available for the
Scuti variables in only 6 clusters (in
which all the
Scuti pulsators have available
values). These rotational velocities are not unusual: inside the same cluster,
the
Scuti variables present a very
wide range of
values, e.g., in Pleiades,
=
10 kms-1
for TR 390 and 175 for TR 410; in Hyades, the range of values is also very
wide, from 30 to 205 kms-1 for 60 Tau and 69 Tau, respectively;
in
Persei,
the
values are between 50 kms-1 for H 606 to
150 kms-1 for H 906; in
Praesepe, the majority of the
Scuti stars present values higher than
100 kms-1 (the greatest is 200 kms-1 for KW 207), but a very low
value of
= 30 km s-1 corresponds to KW 284.
As it can also be seen from Table 5, the open clusters in which a
larger number of
Scuti variables have been found are Praesepe (14), Hyades (8),
NGC 1817 (7) and NGC 6134 (7). The two first clusters are bright and were
investigated during the decade of the seventies using single-channel
photometers while the two latter clusters are much fainter and the
corresponding surveys have been carried out recently (Frandsen et al.
1996; Frandsen & Arentoft 1998a) using CCD cameras. All the
four clusters have similar ages, about
.
There does
exist a detection bias in favor of a certain age range slightly below
1 Gyr: this corresponds to open clusters in which the isochrone turns upwards
near the hot border of the instability strip and
almost follows the instability strip. Younger clusters only contain
main-sequence stars inside the instability strip, which leads to a lower
detection probability
because of the very small amplitudes of the unevolved stars. Considerably
older clusters, on the other hand, no longer contain stars inside the
classical instability strip (except
possibly for blue stragglers or post-main-sequence objects). Frandsen &
Arentoft (1998b) give a list of open clusters suitable for CCD-camera
variability surveys of
Scuti variables.
Table 6 lists some open clusters in which variability surveys for
Scuti variables have not been successful in discovering these
variables. N refers to the number of cluster stars located inside the
instability strip which have been monitored for variability.
Column 3 indicates the type of photometry carried out:
CCD or simultaneous uvby photometry with photomultiplier detectors.
In fact, in NGC 7226, 68 potential
Scuti pulsators have been monitored
(Viskum et al. 1997;
), but none of them has been found
to be variable. While the interpretation can be argued to be uncertain for
individual stars
(e.g., contamination of the sample by less reddened field stars outside the
instability strip), the large number of nonvariable stars indicates that a
presently unknown physical parameter causes the lack of pulsation in these
stars. Follow-up studies of this interesting cluster seem warranted.
Cluster | N | Survey | Source | log(age) | Source |
(1) | (2) | ||||
NGC 654 | - | CCD | 1 | 7.08 | 6 |
NGC 663 | - | CCD | 2 | 7.13 | 6 |
Melotte 105 | ![]() |
CCD | 3 | 8.4 | 3 |
IC 4665 | ![]() |
uvby | 4 | 7.58 | 6 |
IC 4756 | ![]() |
uvby | 4 | 8.78 | 6 |
NGC 6633 | ![]() |
uvby | 4 | 8.66 | 6 |
NGC 7092 | ![]() |
uvby | 4 | 8.61 | 6 |
NGC 7226 | ![]() |
CCD | 5 | 8.7 | 5 |
In this section we will discuss the location of the
Scuti variables listed in the R00 catalogue in the H-R diagram on
the basis
of their uvby
Strömgren-Crawford photometry
and parallaxes determined by the Hipparcos satellite. The two sources
of absolute magnitudes allow us to compare the photometric and parallax
methods and to examine possible systematic errors in the photometric
calibrations used, especially for metallic-line stars.
The dereddened uvby
indices and photometric absolute magnitudes
(Mv(ph)) were calculated in the same way as in Rodríguez et al.
(1994), following the method described in Philip et al.
(1976) and using the reference lines of Philip & Egret
(1980).
The typical uncertainty for the derived Mv(ph) is about
.
For stars with available Hipparcos parallaxes, we have also
calculated absolute magnitudes (Mv(
))
by using
(
in arcsec).
Here, the error bars are determined from
where
/
.
Since the apparent magnitude
(but not the Hipparcos parallax) is affected by interstellar reddening,
we have applied a reddening correction determined from uvby
photometry by using the standard relation
V0=V-4.3 Eb-y
(Crawford & Mandwewala 1976). Additionally, when the components of a
binary system are resolved in the bibliography (mainly the Hipparcos catalogue
(ESA 1997)), corresponding corrections were applied.
We also have restricted the Hipparcos parallaxes to those stars with the
Hipparcos parallax uncertainty 0.20 of the parallax value,
.
This criterion corresponds to
.
Note this criterion also eliminates
the uncertain parallax of AD CMi, which made the star appear to lie below
the main sequence (Høg & Petersen 1997;
Antonello & Mantegazza 1997; Petersen & Høg 1998).
There are three stars which we have omitted because of the very uncertain
photometry and nature of variability: HD 60987,
HD 193084 and 2362-16.
Before the position of all the
Scuti stars in the catalogue
can be correctly placed in the H-R diagram, several subgroups need
to be investigated in more detail.
These four classes of stars are spectroscopically defined subclasses with
surface abundance anomalies. Although the stars Pup,
Del are
also pulsators, the groups named after these stars should not be regarded
as pulsation
subclasses of
Scuti stars. The abundance anomalies in these
stars affect the pulsation properties, e.g., the classical Am stars are
constant in light or show only small pulsation amplitudes (e.g., HD 1097,
Kurtz 1989).
These four groups are only a small selection from the bewildering
zoo of stars with unusual surface abundances
in this temperature region (e.g., see Kurtz 2000) and represent the
groups with Scuti pulsation.
Boo stars are metal-poor Pop. I objects with spectral types from
late-B to
mid-F. Details on the abundances of the different elements can be found in
Heiter (2000). These stars cover the whole main-sequence range between
the zero-age and the terminal-age main sequence. Both the
Boo and
SX Phe stars (see below) show metal-poor spectra and positive
m1indices in the uvby
system. This can make it difficult to
classify field stars
correctly if detailed spectroscopic or space velocity information are not
available. The status of BS Tuc (HD 6870) is now clear: it is not an SX Phe
variable, but a strong
Boo star (hF0A1V, Paunzen 2000).
The classical Am stars, which are on the main sequence,
have their evolved counterparts: the evolved Am,
Pup and
Del stars. Gray & Garrison (1989) make
a valiant attempt to provide a refined classification scheme for these
stars (e.g., they propose to abolish the
Del designation in
favor of spectroscopically more meaningful subgroups). A possible
reason for the present confusion may be our present lack of physical
understanding of these phenomena. Consequently, here we combine all these
groups together (albeit in an oversimplified manner) as
classical and evolved Am stars.
Variable | HD | P | V | ![]() |
Eb-y | (b-y)0 |
![]() |
Mv(ph) | ![]() |
![]() |
Mv(![]() |
(d) | (mag) | (mag) | (mag) | (mag) | (mag) | (mag) | (mas) | (mag) | |||
BS Tuc | 6870 | 0.065 | 7.49 | 0.02 | 0.000 | 0.170 | 0.060 | 1.97 |
![]() |
0.06 |
![]() |
BD Phe | 11413 | 0.0373 | 5.94 | 0.03 | 0.003 | 0.102 | 0.060 | 1.45 |
![]() |
0.05 |
![]() |
EX Eri | 30422 | 0.021 | 6.18 | 0.01 | 0.000 | 0.101 | 0.022 | 2.26 |
![]() |
0.04 |
![]() |
64491 | 0.049 | 6.23 | 0.007 | 0.000 | 0.195 | 0.049 | 2.69 |
![]() |
0.06 |
![]() |
|
HZ Vel | 75654 | 0.087 | 6.38 | 0.01 | 0.000 | 0.161 | 0.046 | 1.80 |
![]() |
0.05 |
![]() |
AK Ant | 83041 | 0.066 | 8.80 | 0.005 | 0.010 | 0.220 | 0.066 | 1.59 |
![]() |
0.24 | |
84948B | 0.078 | 8.14 | 0.01 | 0.000 | 0.196 | 0.036 | 1.27 |
![]() |
0.23 | ||
102541 | 0.050 | 7.94 | 0.02 | 0.021 | 0.142 | 0.055 | 2.47 |
![]() |
0.12 |
![]() |
|
105058 | 0.040 | 8.89 | 0.002 |
![]() |
0.20 |
![]() |
|||||
II Vir | 105759 | 0.0423 | 6.52 | 0.05 | 0.000 | 0.142 | 0.053 | 1.62 |
![]() |
0.10 |
![]() |
109738 | 0.033 | 8.28 | 0.02 | 0.011 | 0.154 | 0.062 | 1.78 | ||||
MO Hya | 111786 | 0.0322 | 6.14 | 0.02 | 0.003 | 0.159 | 0.062 | 2.20 |
![]() |
0.04 |
![]() |
120500 | 0.049 | 6.60 | 0.009 | 0.009 | 0.059 | 0.033 | 1.18 |
![]() |
0.13 |
![]() |
|
![]() |
125162 | 0.023 | 4.18 | 0.002 | 0.007 | 0.044 | 0.020 | 1.75 |
![]() |
0.02 |
![]() |
HR Lib | 142703 | 0.060 | 6.11 | 0.01 | 0.000 | 0.182 | 0.065 | 2.30 |
![]() |
0.04 |
![]() |
IN Lup | 142994 | 0.127 | 7.18 | 0.05 | 0.003 | 0.196 | 0.056 | 0.93 | |||
V346 Pav | 168740 | 0.036 | 6.13 | 0.01 | 0.001 | 0.135 | 0.064 | 1.88 |
![]() |
0.05 |
![]() |
V704 CrA | 168947 | 0.058 | 8.12 | 0.02 | 0.028 | 0.144 | 0.064 | 1.42 | |||
V1431 Aql | 183324 | 0.021 | 5.79 | 0.004 | 0.004 | 0.047 | 0.037 | 1.73 |
![]() |
0.05 |
![]() |
191850 | 0.074 | 9.62 | 0.03 | ||||||||
29 Cyg | 192640 | 0.0267 | 4.93 | 0.02 | 0.000 | 0.099 | 0.049 | 1.80 |
![]() |
0.02 |
![]() |
210111 | 0.036 | 6.37 | 0.02 | 0.000 | 0.136 | 0.046 | 1.73 |
![]() |
0.07 |
![]() |
|
V340 And | 221756 | 0.044 | 5.55 | 0.005 | 0.004 | 0.052 | 0.040 | 1.13 |
![]() |
0.05 |
![]() |
Both the Boo stars and the SX Phe variables
show metal-poor spectra and positive
m1indices in the uvby system. This can make it difficult to classify field stars
correctly if detailed spectroscopic or space velocity information is not
available. The status of BS Tuc (HD 6870) is now clear: it is not an SX Phe
variable, but a strong
Boo star (hF0A1V, Paunzen 2000).
Table 7 lists the Boo stars in the R00 catalogue. Only the star
29 Cyg have CCDM number, but it is not resolved by the Hipparcos catalogue
(ESA 1997). Hence, no corrections due to binarity is made for any
Boo star.
Inspection of the
table shows the good agreement between the photometric and Hipparcos absolute
magnitudes for
Boo stars. The position in the H-R diagram
of the pulsating
Boo stars from the catalogue is shown in Fig. 5.
![]() |
Figure 5:
Position of ![]() ![]() ![]() ![]() |
Open with DEXTER |
We now turn to the more difficult question of the classical and evolved
( Pup and
Del) metallic-line A stars, which are listed in
Table 8. Five of these stars have CCDM numbers (CP Oct,
Pup, V527 Car,
XZ Men and V1004 Ori), but corrections are available for only V527 Car
(0
09)
and XZ Men (0
74). The Hipparcos parallaxes for both stars were too
uncertain to be used (see Table 8), so that no corrections due to binarity
were applied to the listed Mv(
)
values.
Since metallic-line A stars were chosen, it is not surprising that the
index usually has a negative value. However, the value
of +0
040 for V527 Car, at first sight, appears incompatible with
the classification of A3m/A7/A9. However, the
index is an
unreliable indicator of metallicity for evolved stars and should only
be used with caution. Furthermore, the group of evolved Am stars
in the table contains some stars with abnormal element abundances,
which include both overabundances as well as underabundances of
different elements (e.g., the star
Del itself).
Table 8 clearly shows that for the group of classical and evolved
metallic-line A stars, the photometrically calculated absolute magnitudes
are seriously and systematically in error. In all cases, the
photometry underestimates the luminosity (i.e., the Balmer jump is
too small) with the error ranging from 0
4 to 3
0. A median value of
0
75 is found for the stars in Table 8. Further evidence of this
problem (and a confirmation of the correctness of the Hipparcos absolute
magnitudes) is provided by the length of the pulsation periods: the
period of 0
12 for CC Oct is compatible only with the Hipparcos parallax.
We conclude that the photometric absolute magnitude calibrations of the
uvby
system are not applicable to the group of classical and
evolved metallic-line A stars. These will not be used to determine the
location of these stars in the H-R diagram.
The middle panel of Fig. 5 shows the position of the pulsating evolved metallic-line
A stars in the H-R diagram. Note that the interesting classical Am star HD 1097
could not be plotted because of insufficient accuracy of the Hipparcos parallax.
Variable | HD | P | V | ![]() |
Eb-y | (b-y)0 |
![]() |
Mv(ph) | ![]() |
![]() |
Mv(![]() |
(d) | (mag) | (mag) | (mag) | (mag) | (mag) | (mag) | (mas) | (mag) | |||
CP Oct | 21190 | 0.1498 | 7.61 | 0.05 | 0.003 | 0.229 | -0.042 | 1.42 |
![]() |
0.15 |
![]() |
UY Col | 40765 | 0.1696 | 9.58 | 0.14 | 0.052 | 0.219 | -0.076 | 2.90 | |||
![]() |
67523 | 0.1409 | 2.83 | 0.09 | 0.040 | 0.220 | -0.049 | 2.02 |
![]() |
0.01 |
![]() |
V527 Car | 95321 | 0.2137 | 9.00 | 0.06 | 0.005 | 0.120 | 0.040 | 0.39 |
![]() |
1.37 | |
V388 Pav | 199434 | 0.1583 | 8.75 | 0.05 | 0.035 | 0.247 | -0.054 | 2.15 |
![]() |
0.30 | |
UZ Ret | 26892 | 0.1228 | 9.22 | 0.05 | 0.026 | 0.225 | -0.061 | 2.37 |
![]() |
0.29 | |
XZ Men | 31908 | 0.1083 | 7.85 | 0.03 | 0.050 | 0.162 | -0.021 | 1.63 |
![]() |
0.72 | |
V356 Aur | 37819 | 0.1893 | 8.07 | 0.08 | 0.108 | 0.266 | -0.095 | 2.02 | |||
V1004 Ori | 40372 | 0.061 | 5.89 | 0.02 | 0.000 | 0.122 | -0.007 | 1.32 |
![]() |
0.10 |
![]() |
V383 Car | 52788 | 0.12 | 8.39 | 0.04 | 0.035 | 0.207 | -0.029 | 1.47 |
![]() |
0.18 |
![]() |
BO Cir | 129494 | 0.1412 | 9.73 | 0.01 | 0.100 | 0.219 | -0.068 | 2.69 | |||
CC Oct | 188136 | 0.1249 | 8.01 | 0.06 | 0.007 | 0.265 | -0.182 | 4.62 |
![]() |
0.16 |
![]() |
![]() |
195961 | 0.1141 | 4.86 | 0.03 | 0.021 | 0.232 | -0.091 | 2.52 |
![]() |
0.04 |
![]() |
![]() |
197461 | 0.1568 | 4.43 | 0.07 | 0.004 | 0.187 | 0.018 | 1.31 |
![]() |
0.04 |
![]() |
EW Aqr | 201707 | 0.0966 | 6.47 | 0.04 | 0.009 | 0.168 | 0.000 | 1.17 |
![]() |
0.13 |
![]() |
AU Scl | 1097 | 0.0564 | 9.09 | 0.01 | 0.000 | 0.239 | -0.148 | 4.20 |
![]() |
0.29 |
The SX Phe variables are the
Scuti stars of Pop. II and
old disk population. Since such old stars at
8500 K should have
already evolved away and no longer exist in this part of the
Hertzsprung-Russell diagram,
they are also unusual from an evolutionary point of view. They probably
are in a post-giant branch stage of evolution and may be merged binary stars.
This would also explain the position of these stars on or near the main
sequence.
Verification of the existence of SX Phe stars as a group with an evolutionary
history different from that of the normal Pop. I
Scuti star comes
from the discoveries
of SX Phe stars in globular clusters such as
Cen, M 3 and M 55. The
recent paper by Rodríguez & López-González (2000)
discusses these
SX Phe stars in globular clusters in detail so that the much smaller sample
of SX Phe field stars does not reveal much additional astrophysical
information.
Table 9 lists the known field SX Phe stars. The absolute magnitudes
were calculated from the calibrations of uvby
photometry as
well as Hipparcos parallaxes. The table also shows that only one
SX Phe field star (SX Phe itself) has an accurate Hipparcos parallax
with
.
For SX Phe, the photometric and Hipparcos
absolute magnitudes differ by 0
8: this may be caused by systematic
errors in applying the Pop. I calibrations to this extremely metal-poor star.
Such an uncertainty in the photometric absolute magnitudes, though of
smaller size, may also apply to the other stars in the table.
The first 7 stars in Table 9 were
already classified as field SX Phe variables in the R94 catalogue. The other
6 stars (V879 Her, V2314 Oph, BQ Ind, PL 43, BQ Psc and 29499-057) are new and
were discovered by Wetterer et al. (1996), Martín &
Rodríguez (1995), Hipparcos catalogue (ESA, 1997),
Preston & Landolt (1998), Bernstein et al. (1995) and
Preston & Landolt (1999), respectively.
The position of the field SX Phe stars is also shown in Fig. 5, which
reveals the preference of these stars for the cooler half of the instability strip,
similar to that shown in globular clusters (Rodríguez &
López-González 2000).
Variable | HD | P | V | ![]() |
Eb-y | (b-y)0 |
![]() |
Mv(ph) | ![]() |
![]() |
Mv(![]() |
(d) | (mag) | (mag) | (mag) | (mag) | (mag) | (mag) | (mas) | (mag) | |||
BL Cam | 0.0391 | 13.10 | 0.33 | 0.211 | 0.149 | 0.098 | 2.00 | ||||
KZ Hya | 94033 | 0.0595 | 9.96 | 0.80 | 0.039 | 0.180 | 0.082 | 2.29 | |||
SU Crt | 100363 | 0.055 | 8.65 | 0.01 | 0.013 | 0.185 | 0.033 | 2.42 |
![]() |
0.32 | |
XX Cyg | 0.1349 | 11.87 | 0.80 | 0.065 | 0.156 | 0.056 | 1.31 |
![]() |
1.35 | ||
CY Aqr | 0.0610 | 10.93 | 0.71 | 0.036 | 0.164 | 0.053 | 1.83 |
![]() |
3.21 | ||
DY Peg | 218549 | 0.0729 | 10.36 | 0.54 | 0.055 | 0.155 | 0.048 | 1.78 |
![]() |
5.61 | |
SX Phe | 223065 | 0.0550 | 7.28 | 0.41 | 0.000 | 0.151 | 0.063 | 2.07 |
![]() |
0.06 |
![]() |
V879 Her | 0.0569 | 15.65 | 0.65 | ||||||||
V2314 Oph | 161223 | 0.144 | 7.43 | 0.05 | 0.093 | 0.147 | 0.059 | 0.97 | |||
BQ Ind | 198830 | 0.0820 | 9.87 | 0.25 |
![]() |
||||||
PL 43 | 0.0374 | 13.51 | 0.10 | ||||||||
BQ Psc | 0.0608 | 18.35 | 0.57 | ||||||||
29499-057 | 0.0417 | 13.8 | 0.04 |
In Fig. 6, we compare the results obtained for Mv(ph)
and Mv()
for the
Scuti variables
listed in the R00 catalogue with both uvby
photometry and
parallax available. Stars with less accurate Hipparcos parallaxes of
were omitted.
![]() |
Figure 6:
Comparison between absolute magnitudes of ![]() ![]() ![]() ![]() |
Open with DEXTER |
![]() |
Figure 7:
The dependence of
![]() ![]() ![]() ![]() |
Open with DEXTER |
The agreement between the Hippacos and photometric absolute magnitudes is
not satisfactory: the Hipparcos parallaxes indicate, on the average,
higher luminosities of
.
Figure 7
shows that the deviations originate in stars with low rotational
velocities and/or high metallicities (low
indices). (Note
that metallic-line stars are usually slow rotators so that it is
not easy to separate the two effects.
Since the Hipparcos parallaxes should not be effected by these
stellar properties, the photometric luminosity calibrations must
be in error for these stars.
The deviations of the photometric absolute magnitudes found by us
for Scuti stars have been noticed before in late A/early F stars.
Guthrie (1987) already proposed a correction based on
.
However, Domingo & Figueras (1999) have
shown that the Guthrie correction is not appropriate since
it leads to systematic errors for the normal A stars for which
the classical Crawford calibrations (as used by us)
are quite satisfactory. In their
comparison of the Hipparcos and photometric absolute magnitudes,
Domingo & Figueras (1999)
also noticed the rotational velocity and
correlations in the differences of the Hipparcos and photometric absolute
magnitudes. Considerable work is still required to improve the photometric
absolute magnitude calibrations and include the effects of rotational
velocity, metallicity and to provide a more accurately determined ZAMS relation.
The situation may be summarized as follows:
for Scuti stars with normal metals and rotational
velocities of 100 km s-1 the photometric calibrations are in
agreement with the Hipparcos parallaxes.
For the other stars, systematic corrections need to be
applied to the photometrically derived values of Mv, but the
size of these corrections still needs to be determined.
In the previous section, we have compared the results of Hipparcos
parallaxes with the absolute magnitudes determined from uvby
photometry. We have isolated one subgroup for which the calibrations
of the uvby
system are not applicable, viz., the group of
evolved and classical
metallic-line A stars. For these stars, the photometric calibrations
are not used in this section. We can now examine the
location of the
Scuti stars in the H-R diagram, using both these
methods. The following procedure was adopted:
(i) the colors were determined from the uvby
photometry. All the
stars were dereddened following the method for AF stars described in
Philip et al. (1976). A few hot stars can
also be dereddened as A-intermediate stars, but the results are very similar;
(ii) for those stars with accurate Hipparcos parallaxes (uncertainties
), only the Hipparcos absolute magnitudes were used,
after correcting the apparent magnitude, V, of the reddened
stars for interstellar extinction (determined in (i));
(iii) for those stars with Hipparcos parallaxes with uncertainties in
the absolute magnitudes between
and
(our previous limit),
an average absolute magnitude was calculated with weighting according
to the square of the standard deviation, which was assumed to
for the photometric absolute magnitudes;
(iv) for those stars with no (or very uncertain) Hipparcos parallaxes, only
the absolute magnitude calibrations of the uvby
system were used.
Note that our approach omits a few stars with poor absolute magnitude determinations. An example is the evolved metallic-line star, V527 Car, for which both the Hipparcos and photometric absolute magnitudes are extremely unreliable.
Figure 8 shows the location of the Scuti variables in the H-R
diagram.
Stars situated outside the instability strip borders have been marked with
their names and are discussed below:
(i) AB Cas: this star appears to lie a magnitude below the ZAMS, but the
location is unreliable. The star is an Algol-type binary system where the
primary component is a Scuti variable. A temperature of
K
and
was derived by Rodríguez et al. (1998)
for the
Scuti component, placing the star inside
the
Scuti region;
(ii) The four hottest stars Scuti stars (HD 191747, HD 214698,
HD 213272 and HD 97302) were already discussed by Rodríguez et al.
(1994). In all these four cases, very small variations have been
claimed to be present (
,
from peak to peak).
However, new observations of HD 191747 (Rolland et al. 2000)
and HD 213272 (Handler 1995) have not confirmed the variability.
More high-quality observations are needed to check whether or not these
stars are indeed
Scuti stars;
(iii) The evolved Am variable, HD 188136 (=CC Oct) shows the most extreme
overabundances of all the stars in our sample. Consequently, it is to be
expected that the photometric indices are unusual and that the
observed color is too red for the stellar temperature.
![]() |
Figure 8:
Location of 318 ![]() ![]() ![]() |
Open with DEXTER |
In this section we will discuss the current status of some interesting groups
of peculiar and/or related Scuti variables in order to clarify their
situation, connection with the
Scuti-type pulsators and inclusion in
the R00 catalogue.
During the last few years, a number of long period variables are being
discovered on or beyond the cool border of the instability strip
intersects with the main sequence. They are called Dor variables and
their periods are too long as compared with the corresponding ones for
Scuti variables. These stars are typically early F-type stars of
luminosity class V or IV, displaying low visual amplitudes of a few
hundredths of
a magnitude and periods ranging from 0
3 to 2
.
If one
considers the mean densities of these stars, the corresponding pulsation
modes must be nonradial gravity modes
of high order. Moreover, the radial velocity curves are nearly in phase with
the corresponding light curves, quite different to that occurring in
Scuti variables (Garrido 2000). Some
Dor stars coexist
with the
Scuti variables inside the instability strip, but the
majority
are located outside the cool border. Since Krisciunas (1993) proposed
these variables as a new group of pulsating
variables, a great effort has been made to confirm the pulsational nature of
these variables and to add new members to the sample. The pulsational nature
seems to be well established from observed multiperiodicity and interactions
between frequencies in some well studied objects. However, the
precise mechanism producing the pulsations is still subject of study.
The relatively long periods and very low amplitudes make the observational
detection
very difficult. The majority of the Dor stars have been discovered
by accident. Some were used as comparison stars for
differential photometry of previously known
Scuti variables, e.g.,
HD 108100 (Breger et al. 1997).
The connection between
Dor and
Scuti variables has been the
subject of several investigations, e.g., Breger & Beichbuchner (1996).
Multisite coordinated campaigns of observation have also been carried out for
some of these objects, e.g.
Dor (Balona et al. 1994,
1996), 9 Aur (Zerbi et al. 1997a), HD 108100 (Breger et al.
1997), HR 2740 (Poretti et al. 1997), HD 164615
(Zerbi et al. 1997b), HR 8799 (Zerbi et al. 1999),
HD 62454 and HD 68192 (Kaye et al. 1999b).
Systematic observations in open clusters have also been recently carried out,
e.g., NGC 2516 (Zerbi et al. 1998), Hyades (Krisciunas
et al. 1995; Martín et al. 2000), M 34 (Krisciunas &
Crowe 1997; Krisciunas & Patten 1999), Pleiades (Martín &
Rodríguez 2000), NGC 2301 (Kim et al. 2000).
Lists of bona fide members and candidates are available in
the literature (Krisciunas & Handler 1995; Handler & Krisciunas
1997). We caution that some of the variables are still uncertain.
The last revised list of Dor-type variables (Kaye et al.
1999a) contains only 13 objects, and one of them (HD 152569)
was later reclassified as a
Scuti variable (Kaye et al.
2000).
This star was, therefore, included into the R00 catalogue. One very
interesting object is HR 8799 which seems
to link three astrophysically interesting classes of
stars:
Dor,
Boo and Vega-like stars (Gray & Kaye
1999). Six new
Dor variables, discovered from systematic
observations in open clusters, have been added to this list (Martín et al.
2000). Five of them are members of open clusters.
Altogether, more than 100 candidates have been proposed as likely
Dor variables. Most of these were studied through the
Hipparcos mission (Aerts et al. 1998; Handler 1999a).
Further studies of these interesting candidates are
necessary in order to confirm their
Dor-type nature.
Studies of these variables in open clusters can give us some insight about the
constraints of the incidence of Dor pulsation.
Age was proposed by Krisciunas et al. (1995) as an important parameter in the sense that
pulsation would only be excited in young open clusters.
Handler (1999a) found that the
Dor stars to be
slightly deficient in metal abundances. In an unpublished study,
one of the authors (ER) finds a correlation between the Strömgren
colour index and temperature, in the sense
that
tends
to be positive (that is, deficient in metallicity) for the hotter stars and
is decreasing as the temperature decreases. For the coolest
stars this leads to
equal to zero (
,
Nissen 1988).
Hence, the metal abundance, rather than age,
seems to be a good parameter to characterize these stars.
This also
explains why the surveys have been unsuccessful in the Hyades cluster
(
,
Table 5). This conclusion is also in good agreement with the
results obtained by Martín et al. (2000) from systematic
surveys in
several open clusters. Thus, the clue to find new
Dor variables seems
to be to observe open clusters with
and a lot of main-sequence
stars in the region of interest. If we are dealing with field stars,
a good rule would be to observe in this region AF stars (that is,
and
)
being
slightly deficient in metal abundances (that is,
).
During the pre-main-sequence (PMS) contraction stage, many stars cross
the classical instability strip and Scuti-type pulsation can
take place. Recently, Marconi & Palla
(1998) have calculated the location of the instability strip for PMS
stars in the H-R diagram for the first three radial modes.
These authors show that, if the mass is larger than
1.5
,
the star crosses this region of pulsational instability as a
PMS object.
However, the typical time spent by a star within the boundaries of the
instability strip as PMS star is very small as compared with that spent as a MS
star. Hence, the probability of finding one star in this phase of evolution is
very small. However, a few PMS
Scuti-type pulsators have been already
identified (see Table 10) and a number of new PMS candidate pulsators have also
been proposed (Marconi & Palla 1998; Pigulski et al. 2000a).
On the other hand, the interior structure in a PMS star is not the same as that
of a classical Scuti variable evolving from the
MS. Consequently, the discovery of pulsations in PMS
stars is extremely important because it provides constraints on the internal
structure of young stars being a good method to test evolutionary models
(Marconi & Palla 1998). Moreover, as shown by Breger & Pamyatnykh
(1998),
the theoretically predicted evolutionary period changes are a factor of 10 to
100 larger than those for
Scuti variables in MS and with
different sign
(the period must be decreasing in PMS). Thus, the evolutionary changes might
not be hidden among other effects. Hence, the study of period changes
among the PMS
Scuti-type variables appears to be very promising in
contrast to MS and post-MS stars (Breger & Pamyatnykh 1998).
Table 10 lists the PMS Scuti variables known up to date in
the same order they were discovered. The first four variables are already
listed in the R00 catalogue, but the four latter ones have been discovered very
recently (Marconi et al. 2000; Pigulski et al. 2000b) and
they were not included. Additionally, Table 10 lists the visual
magnitude, visual amplitude, main period and number of pulsation frequencies
detected for each variable.
The first evidence for the existence of PMS
Scuti stars came from
the discovery of
Scuti-type pulsations in two members, V588 Mon and
V589 Mon, of the young open cluster NGC 2264 (Breger 1972) with main
periods of about 3.5 hours. New observations
were carried out by Kim (1996), who detected 3 and 4 frequencies,
respectively.
The existence of this subgroup was confirmed by the discovery
of
Scuti-like variations in
the well-known pre-main-sequence Herbig Ae star HR
5999 = V856 Sco by
Kurtz & Marang (1995). In this case, only one frequency was detected
corresponding to a period of 4.99 hours. However, the existence of additional
secondary frequencies can not be ruled out. The pulsation properties of this
star suggest, using the models of Marconi & Palla (1998), that the
mass of this star is about 4
with
second overtone pulsation. A new member of this kind of objects is the Herbig
Ae star HD 104237. The
Scuti-like variations were discovered by Donati
et al. (1997) and confirmed by Kurtz & Müller (1999).
These authors found
two very short and close periods of
and
.
This indicates nonradial pulsation.
PMS
Scuti-type pulsations have also been detected in the two Herbig
Ae stars HD 35929 and V351 Ori (Marconi et al. 2000) with quite
different pulsation periods
of 0
196 and 0
058, respectively. These periods suggest masses of about
3.6
for HD 35929 and 2
for V351 Ori (Marconi et al.
2000). In these two cases, only a few measurements were collected and
multiperiodicity can not be ruled out. Finally, two new PMS
Scuti
variables have been found in the young open cluster NGC 6823 (BL50 and HP57) by
Pigulski et al. (2000b). Over thirty nights of CCD photometry
revealed that both variables are double-mode pulsators.
Star | HD | V | ![]() |
P | N | Source |
(mag) | (mag) | (d) | ||||
V588 Mon | 261331 | 9.73 | 0.04 | 0.1400 | 3 | 1, 2 |
V589 Mon | 261446 | 10.32 | 0.04 | 0.1594 | 4 | 1, 2 |
V856 Sco | 144668 | 7.00 | 0.02 | 0.2078 | 1 | 3 |
104237 | 6.58 | 0.02 | 0.030 | 2 | 4, 5 | |
35929 | 8.2 | 0.02 | 0.196 | 1 | 6 | |
V351 Ori | 38238 | 8.9 | 0.04 | 0.058 | 1 | 6 |
BL50 | ![]() |
![]() |
0.072 | 2 | 7 | |
HP57 | ![]() |
![]() |
0.079 | 2 | 7 |
Only a few variables with periods ranging from
to 0
3, are
listed in the R00 catalogue. The true number may be larger, since
a selection effect probably exists for pulsators with such periods
because of the period overlap with other types of pulsators.
Pulsating variables with periods ranging within this interval are very
interesting because in this region the evolved Population I
Scuti
stars
and the Population II RRc-type variables can coexist. Since
Scuti
stars include all stars evolving from the main sequence to the giant region
(while inside the instability strip), even
Scuti stars with
periods longer than 0
3 are expected.
However, due to the rapid evolution of massive stars evolving towards the
giant region, the probability of finding a massive Scuti variable
in this region of the H-R diagram is small. Consequently, when a new pulsating
variable is discovered
with periods longer than 0
25, it is usually assumed that the star
is a RR Lyrae-type star situated on the horizontal branch.
Only when other parameters are
known such as metallicity, galactic location or spatial motions, is it
possible to distinguish between the two groups.
It is also possible to confuse a Scuti star with a relatively long
period
with the
Dor gravity-mode pulsators, especially among the cool stars
between spectral types F0 and F5. Since
Dor stars are on or near the
main
sequence and the long-period
Scuti stars are evolved, knowledge of the
luminosity allows us to distinguish between the two classes.
In this section, we will analyze star by star the variables catalogued in the
R00 list with periods longer than 0
25, in order to assign the appropriate
type (and, by inference, the evolutionary status). Some of
these variables have been included by earlier authors in lists of
Scuti variables and by other authors in lists of RR Lyr variables.
On the other hand, we are sure that some of these longer-period variables were
missed and not included in the R00 catalogue. A larger number of long-period
Scuti stars would allow us to improve the period-luminosity relations
(e.g., Petersen & Høg 1998 and references therein) by extending the
calibrating sample to include longer periods.
Table 11 lists, in order of increasing periods, the 14 variables listed in
the R00 catalogue with periods between 0
25 and 0
30 together with
their
relevant observed parameters. In this table, N is the number of detected
frequencies: this number is important since RR Lyrae stars are pulsating with
only one or two radial modes. The variable with the longest period is SS Psc
(
). Table 12 lists the
derived parameters, where the notation and limits are similar to
those used in earlier tables. The metallicity values, [Me/H],
have been calculated using the Nissen's
(1988) calibration for
or Smalley's
(1993) calibration for
.
None of these stars is known to contain peculiar abundances.
Only one star, DE Oct, is a known binary, which was not resolved by
the Hipparcos satellite (ESA 1997). Consequently, the computed
luminosities and colors should be accurate. There exists good agreement between
the luminosities derived photometrically and those derived from parallaxes.
Note that the photometric absolute magnitudes (Mv(ph)) have not been
determined for stars too luminous (
).
In the last column of Table 12, we list the most probable identification of
each star.
The variable 7654-501 is a member of the open cluster NGC 7654 being located
near to the blue edge of the instability strip. It was discovered by
Viskum et al. (1997) and confirmed by Choi et al. (1998,
1999). Hence, it probably is a Scuti variable.
There is very little information available on the variables II-52
and 4996-V5. II-52 was discovered by Yao et al. (1994)
in the direction of the globular cluster M 3, but its low amplitude
and light curve suggest this star to be
a field
Scuti variable. In the case of 4996-V5, the amplitude is large
and was discovered in the field of the young open cluster IC 4996 by
Pietrzynski (1996c). If this object is a member of the cluster, then it
would be a high amplitude
Scuti variable. However, these two latter
cases need to be confirmed.
The low amplitudes displayed by HN CMa, V388 Cep and S Eri suggest that
we are not dealing with RR Lyrae-type variables. This is also supported by the
large values of their rotational velocities. All three variables are
located in the upper part of the Scuti region. HN
2724 was
discovered as a multiperiodic variable by Baade & Stahl (1982).
This was confirmed by Breger et al. (1991) who found long term
amplitude variations and
four frequencies consistent with excited p-modes. This is similar to
V388
8851, discovered to be a multiperiodic
Scuti variable
by Hao & Huang (1993). Later, Hao et al. (1995)
identified 5 frequencies suggesting that 4 of them can be interpreted with
the same
-number (
)
and different m-values
(m=-2, -1, +1, +2). There is very
little information available about S
.
This variable was discovered by Millis
(1967) and later reobserved photometrically by Coates et al.
(1981). Coates et al. suggest that S Eri is a luminous
Scuti variable with a main
period of
.
The properties of this star are similar to those of
HN CMa and V388 Cep.
AD Ari and DE Oct were discovered by the Hipparcos mission (ESA 1997).
AD Ari is classified as Scuti variable in both the Hipparcos catalogue
and the list of Kazarovets et al. (1999). In the case of DE Oct, the
Hipparcos catalogue does not give an specific type of variability while
Kazarovets et al. (1999) list this star as RRc:. However, this star
was included in the R00 catalogue on basis of its visual amplitude and
location in the H-R diagram. In both cases, they are Population I
main sequence stars near the cool
edge of the instability strip. This means that nonradial g-modes pulsations of
high order are indicated by their periods. Thus, we are probably dealing with
two
Dor variables with the shortest periods known to date.
V1719 Cyg is a high-amplitude Scuti-type pulsating variable.
This star has been widely studied in the literature since its discovery by
González-Bedolla & Peña (1979), e.g., Gupta & Padalia
(1980, 1982), Padalia & Gupta (1982), Poretti
(1984), Joner & Johnson (1985). Johnson &
Joner (1986) made an exhaustive study using uvby
photometry
concluding that this star is a high amplitude multiperiodic
Scuti
variable
with a high metal abundances (
0.5). The star also exhibited an
unusual light curve of the m1-index.
The reverse shape of the m1 light curve was later explained by
the high metallicity of this star together with its position in the
H-R diagram (Rodríguez et al. 1991). In fact, a metallicity of
about
can be inferred from the observed m1-index curve.
This is also in good agreement with the metal abundances derived from
spectroscopy by Fernley & Barnes (1997). On the other
hand, Poretti & Antonello (1988) found two frequencies
f1=3.7412 c/d
(
)
and
f2=4.6775 c/d and the interactions f1+f2 and
f2-f1. These authors suggested that f1 and f2 correspond to the
first and second overtone of radial pulsation, respecticely.
In addition, these authors also
found that the light curves of V1719 Cyg are atypical in the sense that
descending branch is shorter than the ascending one. This is similar to
two other
Scuti variables: AN Lyn (
)
and
V798 Cyg (
)
(Rodríguez et al. 1997). Finally,
a value of
kms-1 has been determined by Solano & Fernley
(1997) whereas the RR Lyrae variables have no detectable rotation
(Peterson et al. (1996) estimated an upper limit of
10 kms-1 for the RR Lyr stars).
The luminosities and solar metallicity values of DE Lac and SS Psc suggest
that these stars are evolved Population I Scuti rather than
Dor or RR Lyrae variables. For DE Lac, the conclusion
is supported by the measured rotational velocity of
km s-1
(Solano & Fernley 1997). In addition, approximately solar abundances
were also derived
from spectroscopy by Fernley & Barnes (1997) for these two objects.
Moreover, in the case of SS Psc, simultaneous uvby
observations
were also collected by Rodríguez et al. (1993) leading to
,
in very good agreement with the value of
listed in
Table 12 (photometry from McNamara & Redcorn 1977). Furthermore,
is in very good agreement with the behaviour of the m1-index
curve observed in SS Psc (Rodríguez et al. 1993).
The low metallicity values and too high luminosities (c1
)
of DH Peg and UY Cam exclude a
Scuti and
a
Dor origin and suggests the membership in the RR Lyrae group.
(In principle, it could be argued that these stars are long-period SX Phe
stars which also are metal-poor. But such long-period SX Phe stars, which
are typically found in globular clusters, would be situated on the horizontal
branch, i.e., they would be RR Lyrae stars.)
Low metal abundances have also been spectroscopically determined by
Solano et al. (1997) (
for DH Peg) and Fernley &
Barnes (1997) (
for UY Cam). In the case of DH Peg,
the RR Lyrae interpretation is
also supported by the bump displayed just before the maximum of the light
curve (Lub 1977), which is common among RRc-type variables.
YZ Cap is also metal-poor (
,
Kemper
1982), typical of RR Lyr-type stars. Cacciari et al. (1989)
find this star showing typical properties of these variables.
In addition, this star also
presents a bump just before the light maximum (Lub 1977). Consequently,
we assign it to the RR Lyrae class.
Star | P | V | ![]() |
ST |
![]() |
b-y | m1 | c1 | ![]() |
![]() |
N |
(d) | (mag) | (mag) | (kms-1) | (mag) | (mag) | (mag) | (mag) | (mas) | |||
HN CMa | 0.2501 | 6.59 | 0.01 | A5IV-V | 155 | 0.117 | 0.152 | 1.088 | 2.784 |
![]() |
4 |
4996-V5 | 0.251 | 15.0 | 0.36 | 1 | |||||||
DE Lac | 0.2537 | 10.28 | 0.32 | F6 | 32 | 0.317 | 0.154 | 0.779 | 2.703 |
![]() |
1 |
II-52 | 0.2551 | 16.77 | 0.05 | 1 | |||||||
DH Peg | 0.2555 | 9.56 | 0.50 | A5-F0 | 0.246 | 0.075 | 1.116 | 2.772 |
![]() |
1 | |
UY Cam | 0.2670 | 11.44 | 0.34 | A3-6III | 0.149 | 0.110 | 1.140 | 2.754 |
![]() |
1 | |
V1719 Cyg | 0.2673 | 8.01 | 0.31 | F2III | 31 | 0.249 | 0.174 | 0.833 | 2.710 |
![]() |
2 |
AD Ari | 0.2699 | 7.43 | 0.06 | F0 | 0.191 | 0.171 | 0.727 | 2.747 |
![]() |
1 | |
V388 Cep | 0.2717 | 5.56 | 0.04 | A7V | 145 | 0.171 | 0.165 | 0.962 | 2.734 |
![]() |
5 |
S Eri | 0.273 | 4.78 | 0.03 | F0V | 195 | 0.165 | 0.170 | 1.007 | 2.750 |
![]() |
1 |
YZ Cap | 0.2735 | 11.38 | 0.49 | A6 |
![]() |
1 | |||||
DE Oct | 0.2778 | 9.15 | 0.07 | A9IV | 0.223 | 0.166 | 0.746 | 2.759 |
![]() |
1 | |
7654-501 | 0.278 | 14.42 | 0.02 | 1 | |||||||
SS Psc | 0.2878 | 10.99 | 0.39 | A9 | ![]() |
0.189 | 0.172 | 0.960 | 2.735 |
![]() |
1 |
Star | ![]() |
Mv(![]() |
Eb-y | (b-y)0 | m0 | c0 |
![]() |
![]() |
Mv(ph) | [Me/H] | Type |
(mag) | (mag) | (mag) | (mag) | (mag) | (mag) | (mag) | (mag) | ||||
HN CMa | 0.16 |
![]() |
0.000 | 0.117 | 0.152 | 1.088 | 0.340 | 0.045 | -0.39 | ![]() |
|
4996-V5 | ![]() |
||||||||||
DE Lac | 0.085 | 0.232 | 0.181 | 0.762 | 0.223 | -0.008 | 1.18 | 0.23 | ![]() |
||
II-52 | ![]() |
||||||||||
DH Peg | 9.47 | 0.108 | 0.138 | 0.110 | 1.094 | 0.370 | 0.083 | -0.80 | RR Lyrae | ||
UY Cam | 0.001 | 0.148 | 0.110 | 1.140 | 0.452 | 0.076 | -0.72 | RR Lyrae | |||
V1719 Cyg | 0.17 |
![]() |
0.028 | 0.221 | 0.183 | 0.827 | 0.269 | -0.009 | 0.73 | 0.25 | ![]() |
AD Ari | 0.07 |
![]() |
0.000 | 0.191 | 0.171 | 0.727 | 0.052 | 0.013 | 2.44 | -0.06 | ![]() |
V388 Cep | 0.06 |
![]() |
0.000 | 0.171 | 0.165 | 0.962 | 0.326 | 0.015 | -0.08 | ![]() |
|
S Eri | 0.07 |
![]() |
0.001 | 0.164 | 0.170 | 1.007 | 0.326 | 0.014 | -0.07 | ![]() |
|
YZ Cap | 0.85 | RR Lyrae | |||||||||
DE Oct | 0.23 |
![]() |
0.043 | 0.180 | 0.180 | 0.737 | 0.039 | 0.008 | 2.50 | 0.00 | ![]() |
7654-501 | ![]() |
||||||||||
SS Psc | 0.53 | 0.010 | 0.179 | 0.175 | 0.958 | 0.319 | 0.005 | 0.03 | ![]() |
Star | b | V | ![]() |
ST | [Me/H] | P1(mode) | P2(mode) | P3(mode) | Period ratios | Source | |
(![]() |
(mag) | (mag) | (d) | (d) | (d) | 1H/F | 2H/1H | ||||
AC And | -11 | 11.07 | 0.4 | F4 | -0.07 | 0.7112(F) | 0.5251(1H) | 0.4211(2H) | 0.738 | 0.592 | 1, 2, 3 |
GSC04018-01807 | +1 | 11.4 | 0.5 | - | - | 0.5126(1H) | 0.6689(F) | 0.4110(2H) | 0.766 | 0.614 | 4 |
V829 Aql | -11 | 10.28 | 0.3 | F5 | +0.17 | 0.2209(1H) | 0.2924(F) | 0.1767(2H) | 0.756 | 0.604 | 5, 6, 7 |
Stars with masses of 3 or more solar masses also cross the classical
instability strip as they evolve from the main sequence towards the
giant branch. The location inside the instability strip also excites
pulsation, but with periods in excess of the usual
or
limit. These stars with periods of pulsation typical of RR Lyrae
stars will nevertheless show the properties of
Scuti
stars because of the identical stage of evolution. In particular,
the rotational velocity,
,
can, in principle,
be as high as 200 kms-1, while RR Lyrae stars show
10 kms-1 (Peterson et al. 1996).
This, together with period ratios, allows us to recognize long-period
Scuti stars. Due to the rapid evolution across the so-called
Hertzsprung Gap, the probability of detecting such stars at a given
time is very small.
The most prominent example is the star AC And, which
shows three simultaneously excited pulsation periods, viz.,
,
and
,
together with a large number of harmonics
and combinations between the three main frequencies (Fitch & Szeidl
1976). The period ratios are different from those of RR Lyrae stars.
Fitch & Szeidl (1976) concluded that this object is not an
RR Lyr star but an evolved
Scuti variable with a mass of about
3
pulsating in the three first radial modes (fundamental,
first and second overtone, respectively).
The same conclusion was found by Kovacs & Buchler (1994) who
calculated a large number of radial pulsation models to fit the periods and
period ratios observed in AC And. This is also supported
by the Population I metallicity found for this star (Preston 1959, found
which is equivalent to
using the calibration of
Fernley & Barnes 1997) and its very low galactic latitude
(
). The properties of AC And and two other members
of the group are listed in Table 13, where
the observed periods, P1, P2 and P3, are listed in order
decreasing amplitude and F, 1H, 2H refer to the mode identification
of radial fundamental, first and second overtone, respectively.
Fernie (1994) showed that the interpretation of AC And as
a normal post-main-sequence star with 3
predicts
the correct period length based on his period-luminosity relation
(Fernie 1992). This suggests that this star is intermediate between
Scuti and classical Cepheids variables, bridging the gap between
these two classes of pulsators.
Another new three-mode variable similar to AC And is GSC04018-01807 (hereafter
GSC), discovered by Antipin (1997). For GSC, the fundamental
radial mode has a period of
.
V829 Aql, on the other hand, is
similar to these two variables, but has much shorter periods
(Diethelm 1997). This star was discussed by Handler
et al. (1998) as a radial triple-mode 2.1
pulsator with
as the fundamental radial mode. The identified evolutionary
status is supported by the Strömgren indices. Handler (1999b) finds
,
,
and
.
These colour indices place this star as an evolved
Scuti variable
near the cool egde with
and
.
Moreover, the metallicity inferred by
and
Nissen's (1988) calibration (
is typical of normal
Population I stars.
In addition, the Q value derived for the pulsation constant is of
in
very good agreement with
2209 as the first overtone. Thus, the two
variables AC And and GSC seem to be similar to V829 Aql, but with much longer
periods.
Two other possible examples of variables belonging to this part of the H-R
diagram are FW Lup (
,
,
,
ESA
1997) and ST Pic (
,
,
,
ESA
1997). Only single periods were observed for these stars so that the
radial period ratios cannot be used to assign the evolutionary status of these
stars. However, Eggen (1994) suggests on the
basis of a period-luminosity relation that these two stars resemble
first-overtone
Scuti stars rather than RR Lyrae variables. In both cases, the metal abundances
seems to be normal of Pop. I: Eggen (1994) deduced
and
-0.3, respectively.
Other metallicities determinations are available for
FW Lup, in very good agreement with the one derived by Eggen (1994):
by Solano et al. (1997) (from spectroscopic
observations) or
by Jurcsik & Kovack (1996) (from their
[Me/H] versus light curves relations).
Acknowledgements
This research was supported by the Junta de Andalucía and the Dirección General de Enseñanza Superiore Investigación Científica (DGESIC) under project PB98-0499, as well as the Austrian Fonds zur Förderung der wissenschaftlichen Forschung, project number S7304. This work has made use of the Simbad database, operated at CDS, Strasbourg, France.