Pulsating
Pre-Main Sequence Stars
In Young Open Clusters
Dissertation
eingereicht von
Maga . Konstanze Zwintz
zur Erlangung des akademischen Grades
Doktorin der Naturwissenschaften
an der Fakultät für Geowissenschaften, Geographie und Astronomie
der Universität Wien
Institut für Astronomie
Türkenschanzstraße 17
A-1180 Wien, Österreich
Wien, im Oktober 2005
Meinen Eltern, Dr. Edgar und Mag. Sigrid Zwintz, gewidmet.
Abstract
Asteroseismology of pulsating pre-main sequence (PMS) stars has the potential of
testing the validity of current models of PMS structure and evolution. As a first
step a sufficiently large sample of pulsating PMS stars has to be established which
allows to select candidates optimally suited for a detailed asteroseismological analysis
based on, e.g., COROT, MOST or ground based network data. In a second step, the
parameter space for pulsation has to be determined as an analogon to the classical
instability strip.
At the beginning of this study the known PMS pulsators were limited to only
eight. A search for pulsating pre-main sequence stars was therefore performed in the
young open clusters NGC 6383, IC 4996 and NGC 6530 using CCD time series
photometry in the Johnson B and V filters. All three clusters are younger than
106 years and their members with spectral types later than B9 are still contracting
towards the ZAMS. Hence, they were ideal candidates for the investigation of PMS
pulsation among A and F type stars, which cover the classical instability region and
even beyond. For in total 593 stars detailed frequency analyses in both filters have
been performed.
These analyses resulted in the discovery of 15 new pulsating PMS cluster stars:
ten bona fide PMS δ Scuti-type pulsators, three PMS δ Scuti-type candidates and
two γ Doradus-like candidates. Hence, compared to the situation at the beginning
of this work, where only eight members of this group have been known, the total
number of detected pre-main sequence pulsating stars and candidates has significantly increased to 37. This allowed for the first time to probe the instability strip
for pre-main sequence stars in the Hertzsprung-Russel diagram observationally and
compare it, both with the theoretical PMS instability strip and with the classical δ Scuti and γ Doradus instability regions of the corresponding post- and main
sequence counterparts.
Pre-main sequence stars differ from their evolved counterparts of same temperature and luminosity only in their interior structure, whereas their global envelope
properties are quite similar. Therefore, the determination of the evolutionary stage
of a field star may be ambiguous. The study of pulsation in young stars that are still
in their deuterium burning phase and contract towards the zero-age main sequence
provides the unique chance to distiguish between pre- and post-main sequence stars
and hence leads to a better fundamental understanding of stellar structure and evolution.
1
Moreover, the discovery of the potential new class of PMS pulsating objects, the
PMS γ Doradus stars, is specifically interesting for the study of stellar structure
and evolution. As the mechanism driving γ Doradus pulsation in post- and main
sequence stars is currently suggested to be related to convection, first the existence
of young objects showing a similar type of pulsation seems very likely. Secondly, the
study of the pulsational properties of PMS γ Doradus stars could help to solve the
problem of the driving mechanism of γ Doradus stars in general.
Contents
Abstract
1
1 Early Stellar Evolution
1.1 Star formation . . . . . . . . . . . . . . . . . . .
1.2 Evolution of pre-main sequence stars . . . . . . .
1.2.1 The birthline for low-mass stars . . . . . .
1.2.2 The birthline for intermediate-mass stars
1.3 Evolutionary tracks . . . . . . . . . . . . . . . . .
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2 Pre-main sequence stars
2.1 T Tauri stars . . . . . . . . . .
2.1.1 Classes of T Tauri stars
2.1.2 Variability . . . . . . . .
2.2 Herbig Ae/Be stars . . . . . . .
2.2.1 Variability . . . . . . . .
2.2.2 Evolutionary stage . . .
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3 Asteroseismology
3.1 Introduction . . . . . . . . . .
3.2 Pulsation . . . . . . . . . . .
3.2.1 δ Scuti stars . . . . . .
3.2.2 γ Doradus stars . . . .
3.3 The classical instability strip
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4 Pulsation in PMS stars
4.1 Historical background . . . . . . . . . . . . . . . .
4.2 The PMS instability strip . . . . . . . . . . . . . .
4.2.1 Theoretical investigations . . . . . . . . . .
4.2.2 Comparison with observations (status 2000)
4.3 Seismology of PMS stars . . . . . . . . . . . . . . .
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5 Young open clusters
5.1 Basic definitions . . . . . . . . . . . . . . . . . . . . . . . . . . . . .
5.1.1 Embedded and exposed clusters . . . . . . . . . . . . . . . . .
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4
CONTENTS
5.2
5.3
5.4
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7 Observational Results
7.1 Pulsating PMS stars in NGC 6383 . . . . . . . .
7.1.1 NGC 6383 170 . . . . . . . . . . . . . . .
7.1.2 NGC 6383 198 . . . . . . . . . . . . . . .
7.1.3 NGC 6383 152 . . . . . . . . . . . . . . .
7.1.4 Summary of PMS pulsators in NGC 6383
7.2 Pulsating PMS stars in IC 4996 . . . . . . . . . .
7.2.1 IC 4996 37 . . . . . . . . . . . . . . . . .
7.2.2 IC 4996 40 . . . . . . . . . . . . . . . . .
7.2.3 IC 4996 106 . . . . . . . . . . . . . . . . .
7.2.4 IC 4996 46 . . . . . . . . . . . . . . . . .
7.2.5 Summary of PMS pulsators in IC 4996 . .
7.3 Pulsating PMS stars in NGC 6530 . . . . . . . .
7.3.1 NGC 6530 5 . . . . . . . . . . . . . . . . .
7.3.2 NGC 6530 82 . . . . . . . . . . . . . . . .
7.3.3 NGC 6530 85 . . . . . . . . . . . . . . . .
7.3.4 NGC 6530 263 . . . . . . . . . . . . . . .
7.3.5 NGC 6530 265 . . . . . . . . . . . . . . .
7.3.6 NGC 6530 278 . . . . . . . . . . . . . . .
7.3.7 NGC 6530 281 . . . . . . . . . . . . . . .
7.3.8 NGC 6530 288 . . . . . . . . . . . . . . .
7.3.9 Summary of PMS pulsators in NGC 6530
7.4 Other variables . . . . . . . . . . . . . . . . . . .
7.4.1 Variable stars in NGC 6383 . . . . . . . .
7.4.2 Variable stars in IC 4996 . . . . . . . . .
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5.5
5.6
Sequential star formation . . . . . . . . .
PMS stars in young open clusters . . . . .
NGC 6383 . . . . . . . . . . . . . . . . . .
5.4.1 Historical background . . . . . . .
IC 4996 . . . . . . . . . . . . . . . . . . .
5.5.1 Historical background . . . . . . .
NGC 6530 . . . . . . . . . . . . . . . . . .
5.6.1 Historical background . . . . . . .
5.6.2 Cloud collapse and star formation
5.6.3 Proper motion studies . . . . . . .
6 Observations and data reduction
6.1 NGC 6383 . . . . . . . . . . . . . .
6.1.1 Bias level variations . . . .
6.1.2 Color dependent extinction
6.2 IC 4996 . . . . . . . . . . . . . . .
6.2.1 SigSpec . . . . . . . . . . .
6.3 NGC 6530 . . . . . . . . . . . . . .
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CONTENTS
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8 Modelling pulsation
8.1 Pulsation constants . . . . . . . . . . . . .
8.2 Pulsation models . . . . . . . . . . . . . .
8.2.1 Discussion of observed frequencies
8.2.2 PMS γ Doradus type pulsators . .
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7.5
7.4.3 Variable stars in NGC 6530
Summary of cluster properties . . .
7.5.1 NGC 6383 . . . . . . . . . .
7.5.2 IC 4996 . . . . . . . . . . .
7.5.3 NGC 6530 . . . . . . . . . .
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9 The empirical PMS instability strip
124
9.1 All known pulsating PMS stars . . . . . . . . . . . . . . . . . . . . . 124
9.2 The new PMS instability strip . . . . . . . . . . . . . . . . . . . . . 124
10 Conclusions
129
A Photometric data
A.1 Stars in the field of NGC 6383 . . . . . . . . . . . . . . . . . . . . .
A.2 Stars in the field of IC 4996 . . . . . . . . . . . . . . . . . . . . . . .
A.3 Stars in the fields of NGC 6530 . . . . . . . . . . . . . . . . . . . . .
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Abbreviations
147
Bibliography
148
Curriculum Vitae
151
Publications
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Danksagungen
157
Chapter 1
Early Stellar Evolution
The study of the first stages in the formation of stars is one of the currently most
active research fields in stellar astronomy. The relatively short time span between
the formation of stars from interstellar clouds and the core burning of hydrogen in
stars is called the pre-main sequence (PMS) phase.
Star formation is taking place in two very distinct regimes: massive stars can
only be formed in giant molecular clouds, while low-mass star formation can occur in
giant molecular as well as in less massive dark-clouds. The evolution of intermediatemass PMS stars is qualitatively different from that of lower- and higher-mass stars
owing to the differences in stellar and circumstellar processes, as well as in time
scales.
1.1
Star formation
Stellar evolution theory mostly addresses stars that are in hydrostatic equilibrium
(i.e. gas pressure and gravity are balanced), where the motion and inertia of the gas
are neglected. The main problem to be solved is to determine the initial conditions
of stellar evolution (i.e. masses, radii and internal structure) at the moment when
the young stars become mainly hydrostatic for the first time. After these initial
stages, the stars contract towards the zero-age main sequence (ZAMS), where the
energy radiated from the photosphere is equal to the nuclear energy production in
the interior.
The question of the initial conditions of the stars at the beginning of their premain sequence evolutionary tracks in the Hertzsprung-Russell (HR-) diagram is still
unanswered. This is due to the fact that such young stars still accrete mass from their
circumstellar surroundings. The hydrostatic star and its photosphere are directly
connected to the moving hydrodynamic circumstellar material that is accreted. This
has to be taken into account by theoretical flux calculations, as well as the equations of radiation hydrodynamics and convection. The latest progress in computer
technology and the improvement of convection models make it possible to calculate
the pre-main sequence evolution from the initial molecular cloud conditions, follow6
1.2. Evolution of pre-main sequence stars
7
ing the protostellar collapse until mass accretion stops and the stellar photospheres
become visible for the first time (e.g. Wuchterl 1999).
1.2
Evolution of pre-main sequence stars
An interstellar cloud begins its dynamical collapse at the density for which selfgravity begins to overwhelm the cloud’s internal pressure support. The cloud collapses nonhomologously and quickly establishes a characteristic hydrostatic core
surrounded by an optically thick dust envelope, which hides the core from view.
The structure and evolution of the core is dependent on the mass accretion rate
during the free collapse of the envelope onto the core. Collapse calculations predict
a mass accretion rate of the order of 10−5 M¯ yr−1 during the main accretion phase.
After accretion of the envelope to the core, the core begins quasi-static contraction
along a convective Hayashi track. At that moment, the star is no longer hidden by
its dusty envelope and becomes optically visible.
Together with protostar theory it is possible to predict the locus in the HRdiagram where pre-main sequence stars of various masses should first appear as
visible objects. This is called the birthline. Observationally the birthline forms the
upper boundary of the distribution of pre-main sequence stars in the diagrams of
very young clusters.
1.2.1
The birthline for low-mass stars
The birthline for low-mass stars, i.e. in the mass range of 0.2 M¯ ≤ M ≤ 1 M¯
was calculated by Stahler (1983) based on a spherically symmetric collapse of a
Jeans unstable parent cloud neglecting the possible influence of magnetic fields,
rotation or turbulent motion. Although the assumptions have been quite simple,
the computed birthline is in excellent agreement with observations of low-mass T
Tauri stars. There is a sensitivity of the location of the birthline to the collapse rate
of the parent cloud. This implies that the clouds from which low-mass stars form
cannot have been strongly affected by other forces than thermal pressure prior to
their collapse.
Once a low-mass star becomes optically visible and contracts along its Hayashi
track, it presumably becomes a T Tauri star. The low-mass stars in star forming
regions such as Taurus-Auriga, Orion or Ophiuchus seem to cluster below the theoretical birthline. This is explained by the fact that almost all T Tauri stars began
contracting from the birthline.
1.2.2
The birthline for intermediate-mass stars
In the case of intermediate-mass stars the imprint of the previous accretion history
persists much longer and their evolution is much closer tied to the protostellar conditions than for low-mass stars. The star’s surface luminosity increases sharply early
during contraction. Also, the star inherits a thick, subsurface mantle of deuterium,
8
1. Early Stellar Evolution
which must ignite in a shell and fuse to helium during the subsequent approach to
the ZAMS. The fusion of deuterium to helium, which plays a dominant role in the
evolution of low-mass protostars, is less significant for stars with higher masses.
Some stars initially expand once they become optically visible, while others skip
their early convective phase. Stars more massive than ∼10M¯ never have a premain sequence phase at all and can only be observed in IR as accreting protostars
or in their later evolutionary stages.
To derive the birthline for intermediate mass stars, Palla & Stahler (1990) combined pre-main sequence evolutionary tracks with a theoretical mass-radius relationship for accreting protostars. As in the low-mass case, the burning of interstellar
deuterium plays a dominant role in determining the protostar radius. Four main
stages in the burning process can be distinguished (Figure 1.1):
Figure 1.1: Deuterium burning in protostars (taken from Palla & Stahler 1990)
(a) For a star with ∼1M¯ , deuterium burns near the center and keeps the star fully
convective. The freshly accreted deuterium is quickly transported to the center by
convective eddies and a situation of steady-state burning is maintained.
(b) With growing mass the interior temperature of the star slowly rises causing a
concurrent drop in opacity. At a given point the deuterium burning cannot keep
the star fully convective. The transition to the radiative stability is first manifested
by the appearance of an internal radiative barrier: a localized region first becomes
stable against convection and prevents the accreted deuterium from reaching the
center.
(c) The portion of the star inside the barrier quickly becomes radiatively stable,
whereas the outer layers are still too cold to ignite the freshly accreted deuterium.
(d) With further increasing mass, the temperature is rising just outside the radiative
1.3. Evolutionary tracks
9
barrier. If it reaches 106 K, deuterium ignites in a shell and maintains convection in
the outer layers.
For stars more massive than 2.5M¯ theory predicts that the star is radiatively
stable after accretion ends. The star will first appear directly on the radiative portion
of its evolutionary track. Not only does the theoretical birthline coincide well with
observations of Herbig Ae/Be (HAEBE) stars, but there is also good agreement for
the intersection of the birthline with the ZAMS: the observed stars seem to be on
the ZAMS for log Teff ≥ 4.4 corresponding to masses ≥10M¯ .
The observed intermediate mass (2 ≤ M? /M¯ ≤ 10) pre-main sequence stars
indeed show an upper envelope that is close to the theoretical predictions by Palla
& Stahler (1990).
1.3
Evolutionary tracks
Several groups have calculated pre-main sequence evolutionary tracks, which differ
mostly in the constitutive physics (equation of state, convection, atmospheric opacities etc.), but also in the treatment of the surface boundary conditions. Figure 1.2
gives an example of two different sets of frequently used PMS evolutionary tracks for
intermediate mass stars by Palla & Stahler (1993, black solid lines) and D’Antona
& Mazzitelli (1994, red dashed lines).
Figure 1.2: PMS evolutionary tracks by Palla & Stahler (1993, black solid lines)
and D’Antona & Mazzitelli (1994, red dashed lines) for 1.5, 2.0, 2.5 and 3.0 M¯
illustrating the differences due to different input physics.
10
1. Early Stellar Evolution
The evolutionary tracks by Palla & Stahler (1993) occupy a much smaller portion of the HR-diagram, which is a consequence of the initial conditions used by
them, specifically the modest radii attained by each star during its accretion period.
D’Antona & Mazzitelli (1994) included several updates in the input physics, among
them two different sets of recent low temperature opacities and two different treatments of overadiabatic convection (mixing length theory and the Canuto-Mazzitelli
model). But generally, the different sets of evolutionary tracks for pre-main sequence
stars do not differ too much especially in the region of the instability strip, which is
important for this work.
Chapter 2
Pre-main sequence stars
Pre-main sequence (PMS) stars lie between the birthline and the ZAMS in the HRdiagram. They interact with the circumstellar environment in which they are still
embedded; hence they are characterized by a large degree of activity, strong near- or
far-IR excesses and very often by emission lines. It can be distinguished between two
major groups: T Tauri and Herbig Ae/Be objects. Members of both groups show
photometric and spectroscopic variabilities on time scales from minutes to years,
indicating that stellar activity begins in the earliest phases of stellar evolution, prior
to the arrival on the main sequence. The fact that stars move across the instability
region during their evolution to the main sequence suggests that at least part of
their activity can also be due to pulsations.
2.1
T Tauri stars
T Tauri stars are newly formed low-mass stars that have recently become visible in
the optical range. They were discovered by Joy (1942; 1945; 1949) in the TaurusAuriga dark cloud and named after their brightest member, T Tauri. They appeared
worth studying at that time because they were variable stars. They display irregular
and large light variations and are always associated with dark or bright nebulae.
T Tauri stars are primarily of spectral types G, K or M; objects as early as
A type stars are in principle also included, although they are not very numerous.
T Tauri stars have apparently normal photospheres with overlying continuum and
line-emission characteristics of a hotter (say 7 000 K to 10 000 K) envelope. Several
studies (e.g. Joy 1945 & 1949; Herbig 1962) established beyond doubt that these
stars are in their pre-main sequence phase of evolution. Numerous investigations
focused on the nature of the envelope. Many attributes including the H and K lines
and the IR triplet of Ca II, Hα and other Balmer series lines, Fe I lines, a variety of
emission lines in the UV, continuum emission in the far-blue and near UV and in
the near-IR, resemble features seen in the solar chromosphere and other active stars.
This can be explained by a deep chromosphere model: the overlying characteristics
are generated in an extended chromosphere just above the stellar photosphere, anal11
12
2. Pre-main sequence stars
ogous to the solar chromosphere. Features which cannot be explained yet include
the strong Hα , the far-IR emission and the forbidden lines in some stars. These
indicate the presence of gas and dust in a more extended region around the stars.
2.1.1
Classes of T Tauri stars
Classical T Tauri stars (CTTSs) were discovered from Hα surveys. Optical spectroscopic criteria that define a CTTS, according to Herbig (1962), are the following:
(a) Hydrogen Balmer lines and Ca II H and K lines are in emission. (b) Anomalous
emission of Fe I at λ = 4063 and 4132 Å is often observed. (c) Forbidden emission
of O I and S II is observed in many CTTSs. (d) Li I at λ = 6707 Å absorption is
conspicuously strong.
Stars showing equivalent widths less than 5 Å are called naked or weak-lined T
Tauri stars (WTTSs) showing weaker Hα emission (Bertout 1989). Herbst (1986)
suggests the reason could be that some T Tauri stars at a given mass and age will be
in a very active state (CTTSs) and some others in less active states (WTTSs). As
a star ages it might spend less time in the more active states, giving rise to a larger
number of weak-emission T Tauri stars far from parental clouds. WTTSs are X-ray
sources with an optical counterpart showing pre-main sequence characteristics. In
particular, Li I at λ = 6707 Å is present with equivalent widths larger than 100 mÅ,
and stellar radial velocity is consistent with membership in the associated molecular
cloud (Bertout 1989).
2.1.2
Variability
T Tauri stars can vary on time scales ranging from minutes to decades and the
variations can be different at different wavelengths. Initial attempts to understand
the variability of T Tauri stars were led by Parenago (1954) who established an own
classification scheme. Some T Tauri stars vary in periodic or quasi-periodic fashion
at least occasionally. The first really convincing discovery of periodic behaviour was
made by Rydgren & Vrba (1983) in the WTTSs V 410 Tau and HD 283447. The stars
show sinusoidal light variations with periods ranging from 1.9 to 4.1 days, which is
a result of rotational modulation of an inhomogeneous photosphere. Hence, the
rotation period of the star and parameters for the spots causing the light variations
may be derived, including their temperature and size relative to the photosphere.
Hot and cool spots are required to explain all the behaviour observed.
Periodic signals coming from some T Tauri stars are not the rule, but the exception, and in some cases the periodic component is buried in a much larger quasiperiodic or aperiodic variation. Periods can only be found in relatively quiescent
stars because it would be difficult to find periodicities in otherwise irregular light
variations with amplitudes of 1 to 2 magnitudes.
However, as T Tauri stars are of later spectral types and generally do not fall in
the instability region of the HR-diagram, they have not been primary candidates to
search for pulsations. They are described here for completeness.
2.2. Herbig Ae/Be stars
2.2
13
Herbig Ae/Be stars
Herbig Ae/Be (HAEBE) stars are the more massive counterparts of the T Tauri
stars, and hence possess masses between 2 and 10M¯ . The lower limit corresponds
to the mass above which stars are radiatively stable when they begin their quasistatic contraction. The upper limit corresponds to the mass above which stars
start burning hydrogen before they emerge from their contracting envelope, i.e. it
occurs where the stellar birthline (Stahler 1983) intersects the ZAMS. Higher-mass
PMS stars are therefore not expected to be optically visible before they reach the
ZAMS.
HAEBE stars were first mentioned as a group by Herbig (1960) who studied Ae
and Be stars associated with nebulosity and defined empirically three criteria for his
new class of objects: (a) the stars have spectral types A or B, (b) they are located
in an obscured region and (c) they illuminate reflection nebulae in their vicinity.
Herbig (1960) also identified 26 stars showing these properties. Additions to this
original list have been made by Finkenzeller & Mundt (1984) and Herbig & Bell
(1988), a new catalog of HAEBE stars was generated by Thé et al. (1994) and
HAEBE candidate stars were investigated by Vieira et al. (2003).
A slight modification of the definition for HAEBE stars had to be made, because
stars were discovered that are not associated with any nebulosity. So, currently
HAEBE stars are identified according to the following characteristics (Waters &
Waelkens 1998): They are of spectral types A or B and show emission lines, they
possess an IR excess due to hot or cool circumstellar dust or both and they have
luminosity classes III to V.
Spectral energy distributions (SED) of HAEBE stars are characterized by the
presence of sometimes very large amounts of circumstellar matter, which can dominate the SED in the IR and contribute also to the continuum in the UV. This
clearly illustrates that the circumstellar material has a wide range of temperatures
and densities, which is significantly above and below the stellar effective temperature. Sometimes it can be hard to distinguish between properties of the stellar
photosphere and effects of the circumstellar matter. The extinction law of HAEBE
stars can deviate significantly from the average extinction derived for the interstellar
medium, because these stars are often found in star forming regions and can have
substantial circumstellar extinction. Sometimes a UV excess is present in HAEBE
stars, which is caused by accretion with rates on the order of 10−7 M¯ /yr.
The difference between HAEBE and normal main sequence stars is the presence
of emission lines and the complex variability of the emission and absorption features.
Very prominent in HAEBE stars is Hα emission, but emission is also observed in
other atoms and ions, such as O I, Ca II, Si II, Mg II or Fe II.
HAEBE stars rotate with typical v · sin i values between 60 and 200 km/s and
lack slow rotators.
14
2.2.1
2. Pre-main sequence stars
Variability
HAEBE stars display regular and irregular light variations on very different time
scales due to several reasons.
The well-studied phenomenon of sudden drops in brightness of up to three magnitudes in V accompanied by an increased reddening and degree of polarization
and followed by a slow recovery lasting weeks is characteristic for UX Orionis type
variables, which are named after their prototype UX Ori. Those large drops in
brightness are observed only in stars of spectral types A0 and later. It is suggested
that the lack of strongly variable Herbig Be stars is due to the fact that these stars
are optically invisible for most of their pre-main sequence phase (Waters & Waelkens
1998).
Another type of variability is characterized by long-term fading or brightening
over time scales up to decades. This is connected with FU Orionis type outbursts
or with gradual changes in the degree of circumstellar extinction.
On time scales of weeks the reason for photometric variability is variable extinction due to circumstellar dust. Clumped accretion or chromospheric activity may
be responsible for variations between hours and days.
Variability on time scales longer than approximately a day have been studied
frequently, but for most of the HAEBE stars no information on light variations with
periods shorter than that is available. If the observed periodicities lie between half
an hour and few hours and if the star is crossing the region of instability in the
HR-diagram, the origin of stellar variability is pulsation. The amplitudes expected
for this phenomenon are at the millimagnitude level.
Hence, pre-main sequence field and cluster stars with HAEBE type characteristics are primary candidates to search for pulsation, where in this work the focus was
on the investigation of cluster members.
2.2.2
Evolutionary stage
Pre-main sequence stars differ from their more evolved counterparts of same temperature and luminosity only in their interior structure, whereas their envelope properties are quite similar (Marconi & Palla 1998). As the evolutionary tracks for preand post-main sequence stars intersect each other several times (see Figure 2.1, preand post-main sequence evolutionary tracks are taken from D’Antona & Mazzitelli
(1994) and Breger & Pamyatnykh (1998), respectively), the determination of the
evolutionary stage of a field star may be ambiguous. Additional information, like
the age or distance of the star, is needed to decide on this ambiguity.
2.2. Herbig Ae/Be stars
15
Figure 2.1: Intersecting pre- and post-main sequence evolutionary tracks for 1.6,
2.0 and 2.5 M¯ and the boundaries of the classical instability strip, where REobs
denotes the empirical red edge, BE the blue edge for the radial overtones and BEF
the blue edge for the fundamental mode (Breger & Pamyatnykh 1998).
Chapter 3
Asteroseismology
3.1
Introduction
Seismology on the Earth is the study of earthquakes and related phenomena including the measurement of speeds at which the seismic waves travel through the Earth.
Similarly helioseismology applies seismic methods very successfully to the Sun. In
the Sun a huge number of modes is excited simultaneously, where on the order of 107
modes possess amplitudes large enough for observation. Each mode carries information from the Sun’s interior and helps to investigate the solar structure. Many stars
other than the Sun support pulsations with similar properties and asteroseismology
allows to put constraints on the stellar interiors by studying them.
3.2
Pulsation
Pulsation is characterized by the nature of the restoring force that is responsible
for the oscillatory behaviour. For acoustic (p) modes pressure is the restoring force;
such modes can be found in the Sun and in many types of pulsating stars, e.g in
δ Scuti stars. Gravity (g) modes, for which the restoring force is buoyancy, can be
found in white dwarf pulsators, for example.
The pulsation eigenmodes can be described as the product of a function of radius
and a spherical harmonic assuming that the stars can be described as spheres. The
spatial and temporal variation of a perturbation to the star’s mean state are given
as (e.g. Brown & Gilliland 1994):
ξnlm (r, θ, φ, t) = ξnl (r)Y l m (θ, φ)e−iωnlm t
(3.1)
ξ is any scalar perturbation associated with the mode; r, θ, φ and t are the radial
coordinate, the colatitude, the longitude and the time, respectively. The radial order
n specifies the number of nodes between the center of the star and its surface. The
angular degree l is a product of stellar radius and the horizontal wavenumber of
the modes. A high number of l means that the sign along the hemisphere changes
very often. The azimuthal order m can be described as the projection of l on to the
16
3.2. Pulsation
17
equator, so it never can be larger than l, i.e. |m| ≤ |l|. p-modes may be purely radial
(l = 0), but g-modes - that are driven by buoyancy - always show a variation in the
horizontal coordinates and hence have l ≥ 1. The mode frequency ωnlm depends on
n and l, hence on the restoring force and the structure of the star. The results of
observations are often written using the circular frequency νnlm ≡ ωnlm /2π.
Figure 3.1 shows three examples for non-radial pulsation patterns, where l denotes the total number and m the number of longitudinal node lines, i.e. those
crossing the equator, on the stellar surface. Pulsation with l = 1 and m = 0 (on the
left), l = 4 and m = 2 (in the middle) and l = 4 and m = 4 (on the right) are shown.
Figure 3.1: left: pulsation with l = 1 and m = 0; middle: pulsation with l = 4 and
m = 2; right: pulsation with l = 4 and m = 4
Stellar pulsations can be observed by measuring photometric intensitites or radial
velocities including also the determination of amplitudes and line-widths.
Several types of stellar pulsations driven by different mechanisms can be found
across the HR-diagram: from the hot β Cephei and slowly pulsating B (SPB) stars
into the region of the classical instability strip, which is populated by Cepheids, RR
Lyrae, δ Scuti and rapidly oscillating Ap (roAp) stars, to the cooler γ Doradus stars.
Figure 3.2 shows the location of these pulsators in the HR-diagram and the modes,
in which they pulsate. These different types of pulsations have been discovered for
numerous main sequence or slightly more evolved stars.
The detection of pulsation in pre-main sequence stars is an important test for
stellar evolution models and helps to investigate the interiors of such young objects.
PMS pulsators are searched among the young A and F type stars, as the more
massive B type stars do not have a pre-main sequence phase at all and stars of later
spectral types are still deeply embedded in their protostellar material. Hence, δ
Scuti- and γ Doradus-like pulsations can be expected in PMS stars.
3.2.1
δ Scuti stars
δ Scuti stars possess spectral types in the range A - F occupying a position in the
HR-diagram close to and slightly above the main sequence. Their pulsation periods
lie between ∼ 30 minutes and 6.5 hours and their pulsation amplitudes range from a
few millimagnitudes to several tenths of a magnitude. The source of energy for the
pulsation is an instability driven by the κ mechanism in the He II ionization zone
near 48 000 K.
18
3. Asteroseismology
Figure 3.2: Location of different types of pulsating stars across the HR-diagram.
Some δ Scuti stars pulsate purely radial, but most of them show a large number
of nonradial p-modes simultaneously. Photometrically mostly low-degree (l ≤ 3)
and low-order (n = 0 ... 7) p-modes can be measured, while spectroscopically highdegree nonradial modes with l up to 20 can be detected (Kennelly et al. 1998).
When the fundamental and first overtone modes are present, their period ratio can
be used to test models of the structure of δ Scuti stars.
Once the stars have evolved significantly off the main sequence towards the
giant branch the pulsations become more complicated than simple p-modes. With
an increasing helium content in the core, an important gradient of molecular weight
develops in the stellar interior, causing a sharp increase in the buoyancy frequency
in those regions. The latter is dominating the pulsational response of the stellar
core and global pulsations can develop a dual character. The so-called phenomenon
of avoided crossing exists between two decoupled oscillators: one is a gravity wave
(g-mode) showing high amplitudes in the stellar interior, the other is an acoustic
wave (p-mode) centered in the outer envelope of the star. As they interfere the mode
is a p-mode in the envelope and a g-mode in the interior.
3.2. Pulsation
19
Pulsation constant
For δ Scuti stars the pulsation constant, Q, can be calculated to distinguish if the
observed period is a radial fundamental or higher overtone mode. The relation
Q = P (ρ/ρ¯ )1/2
(3.2)
can be also written as (Breger 1979):
log Q = −6.454 + log P + 0.5 log g + 0.1 Mbol + log Teff
(3.3)
For the fundamental mode in δ Scuti stars Q evaluates to 0.033 d, for the first
overtone mode Q is 0.025 d, for the second harmonic it is 0.021 d and for the third
0.017 d. The smaller the Q values the higher is the corresponding radial overtone
pulsation mode for δ Scuti stars.
3.2.2
γ Doradus stars
γ Doradus stars possess a convective core, a radiative envelope and a small outer
convective zone close to the photosphere (Kaye et al. 1999). The relationship
between evolved γ Doradus and δ Scuti stars is not yet clear. Both share a similar
parameter space in the HR-diagram even with overlapping zones (see Figure 3.3).
γ Doradus stars are high radial order n and low spherical degree l, g-mode pulsators
(Kaye et al. 1999), while classical δ Scuti stars mostly pulsate with low radial order
p-modes. Hence, the excitation mechanisms are also different. While pulsation in
δ Scuti stars is driven by the κ mechanism, the only presently suggested mechanism
for γ Doradus type pulsation is similar to convective blocking in the relatively thin
convective envelopes of these stars (Guzik et al. 2000).
The longest pulsation periods of δ Scuti stars listed in the catalogue of Rodriguez
et al. (2000) are ∼ 6.5 hours and the shortest pulsation periods of γ Doradus stars
(Handler & Shobbrook 2002) are ∼ 7.5 hours. It may be suspected that there is
an overlap in the pulsational behaviour of those two classes of pulsators. However,
the 54 known γ Doradus stars are so far only found on the main sequence, and
the long-period δ Scuti stars all seem to be evolved. So, this overlap seems to be
not a physical one. Using the pulsation constant Q, the ambiguity is removed,
because with typical Q values larger than 0.23 days (Handler & Shobbrook 2002),
the γ Doradus stars are well separated from the δ Scuti stars. However, γ Doradus
stars often have multiple photometric periods of up to three days and sinusoidal
light curves with amplitudes of a few millimagnitudes. Radial velocity variations of
2-4 km/s and changing spectroscopic line profiles have also been observed in some
stars.
γ Doradus stars are often confused with ellipsoidal variables and rotationallymodulated chemically peculiar objects, but can be separated from them (Handler &
Shobbrook 2002): both other variables will only show one or two dominant periods
in a frequency analysis; and if there are two frequencies, they will be harmonically
related. The study of the relative amplitudes and phases of the measured signals in
20
3. Asteroseismology
Figure 3.3: The HR-diagram with the location of all known γ Doradus stars, the
γ Doradus instability strip and the ZAMS (thick black line) taken from Handler
(2005). The open star symbol marks HD 209295, a γ Doradus star with most likely
tidally excited pulsation, the cross relates to HD 221866, a binary with a γ Doradus
component, and the filled star symbol corresponds to HD 8801, which shows both
γ Doradus and δ Scuti pulsations.
different photometric filters can also be used to tackle this question: little color modulation with B/V amplitude ratios less than 1.05 is present in ellipsoidal variables
and eclipsing binaries, because their light variations are dominated by geometrical
effects. Quite large color variations together with large phase shifts between the
filters are seen in the light curves of rotationally-modulated Ap stars. Color amplitude ratios of γ Doradus stars are quite similar to those of δ Scuti stars. Typical
B/V amplitude ratios for γ Doradus stars pulsating with photometrically detectable
modes would be between 1.2 and 1.35 according to model calculations, which is also
expected for δ Scuti type pulsation (Handler & Shobbrook 2002).
3.3
The classical instability strip
The theoretical borders of the classical instability strip are strongly affected by the
choice of the global input parameters, such as initial chemical composition, opacity
data, treatment of convection etc. The instability domain for δ Scuti stars calculated
by Pamyatnykh (2000) was computed with OPAL opacities (Iglesias & Rogers 1996),
without taking into account effects of rotation and convective overshooting (Figure
3.4). An initial hydrogen abundance X = 0.70 and metallicity Z = 0.02 are assumed.
Near the ZAMS the radial fundamental mode (p1 ) and seven overtones (p2 - p8 ) can
be excited. The blue edge for the fundamental mode lies in the center of the strip
and the blue edges of unstable regions are hotter for higher overtones. Towards the
3.3. The classical instability strip
21
blue edges of higher overtone modes the fundamental and lower overtones remain
stable in the models because their amplitudes are larger in the interior resulting
in stronger damping below the main driving region. The general blue edge is the
hottest envelope of all unstable stellar models. The modes higher than the seventh
overtone remain stable due to the damping region above the He II ionization zone
and the according short periods. The empirical red edge (REobs in Figure 3.4) is
determined using a transformation of observational data into the theoretical HRdiagram. The location of the general blue edge of the classical instability strip is
also affected by the available opacity data and amount of metallicity: the blue edge
shifts towards the cool side if the opacities, and hence the metallicities, increase.
Figure 3.4: Theoretical blue edges of the classical δ Scuti instability strip for radial
pulsations (taken from Pamyatnykh 2000): the blue edge for the fundamental mode
is marked with p1 , the blue edges for the according overtones are marked with p2
- p8 , respectively. REobs indicates the location of the corresponding empirical red
edge.
The cooler δ Scuti stars possess considerable outer convection zones, making it
impossible to calculate the position of the red edge of the instability strip using
linear nonadiabatic pulsation models with a simple assumption about interaction
between convection and pulsation, namely that the convective flux is constant during an oscillation cycle. Pulsationally induced fluctuations of the turbulent fluxes
become important for the selection mechanism of modes with observable amplitudes.
22
3. Asteroseismology
Houdek (2000) found that with increasing effective temperature the turbulent pressure becomes larger in the upper convective layers and eventually dominates over
the gas pressure. The return to stability at the cool border of the instability domain
is exlusively due to the fluctuations of the turbulent pressure and without including the latter, the pulsation calculations fail to produce the red edge of the δ Scuti
instability strip.
Comparison with observations allows to put constraints on the treatment of
convection adopted in the stellar models and on the interaction between convection
and pulsation. The black dots in Figure 3.4 mark the positions of the post- and main
sequence δ Scuti stars in the theoretical HR-diagram. ∼ 25% are hotter than the
blue edge for the fundamental mode, hence pulsate only in overtones. For the couple
of stars that seem to be even hotter than the general blue edge or are located below
the ZAMS (Pamyatnykh 2000), a systematic reobservation is required. Standard
photometric calibrations may result in wrong fundamental parameters if applied
to non-normal stars, like chemically peculiar stars, as the calibrations have been
derived from normal stars.
Chapter 4
Pulsation in PMS stars
During their evolution to the main sequence, young stars move across the instability
region in the HR-diagram, which suggests that part of their activity is due to stellar
pulsations. PMS stars differ from their counterparts with same effective temperature
and luminosity, but which have already evolved off the main sequence, mostly in the
inner regions, while their atmospheres are quite similar (Marconi & Palla 1998).
The discovery of pulsating PMS stars is extremely important, because it allows to
constrain the internal structure of young stars and to test evolutionary models.
4.1
Historical background
The first two pre-main sequence pulsators were discovered by Breger (1972) in the
young open cluster NGC 2264. The position of V588 Mon (HD 261331, NGC
2264 2) and V589 Mon (HD 261446, NGC 2264 20) in the HR-diagram agreed
with that of the post- and main sequence δ Scuti stars. The A7 III-IV type star
V588 Mon showed a period of 2.6 hours, while the slightly cooler F2 III star V589
Mon was variable with a period of 3.0 hours, both determined from three nights of
observations (see Figure 4.1).
It took more than 20 years until the next PMS pulsating star was discovered. The
pre-main sequence pulsator, HR 5999 (HD 144668, V856 Sco), detected by Kurtz
& Marang (1995), is a Herbig A7 III-IVe star that was intensively studied in the
years before (e.g. Praderie et al. 1991). But none of these studies was trying to
detect δ Scuti-like pulsation because it was not expected at that time. HR 5999
is a fast rotator with v · sin i = 180 ± 50 kms−1 , its mass was estimated to 3 M¯
and its radius to 6.9 R¯ . Its effective temperature Teff is ∼ 7800 K and the surface
gravity log g ∼ 3.5 − 4.0. The star is physically associated with the peculiar late B
star HR 6000, which is separated from it by 44 arcseconds. HR 6000 is a He weak,
variable CP star with a period of about 2 days (Kurtz & Marang 1995). Both stars
are embedded in an obscured region of Scorpius, which also includes many T Tauri
stars in the associated dark cloud.
Previous photometric studies have shown that HR 5999 varies irregularly from
23
24
4. Pulsation in PMS stars
Figure 4.1: Original light curves of the two first known PMS pulsators V588 Mon
(left panels) and V589 Mon (right panels) taken from Breger (1972).
a maximum brightness of V ∼ 6.8 mag down to V > 8 mag on timescales between 48
days and 301 days caused by the obscuring medium. As the star lies within the
classical δ Scuti instability strip, Kurtz & Marang (1995) carried out observations
in order to search for pulsation. They detected a peak-to-peak pulsation amplitude
of about 0.013 mag in Johnson V and a period of ∼ 5.0 hours in the presence of
0.35 mag background variability (Figure 4.2). The pulsation is most likely caused
by one or more p-modes.
This result made it possible to examine for the first time the internal structure
of a pre-main sequence star and to put constraints on the models. Marconi & Palla
(1998) tried to reproduce the pulsation period of ∼ 5.0 hours of HR 5999 using linear
non-adiabatic pulsation models for the first three radial modes, and hence performed
the first asteroseismic investigation for a PMS pulsator. The most plausible model
gave 4.0 M¯ and pulsation in the second overtone mode.
The HAEBE type star HD 104237 was found to pulsate with a period of only
37 minutes by Donati et al. (1997) during an investigation of potential magnetic
fields in HAEBE and T Tauri stars using high-precision spectropolarimetry. Kurtz
& Müller (1999) performed photometric observations to confirm this period, but
find a frequency of their highest amplitude mode of 33.29 c/d corresponding to a
period of 43 minutes. Assuming that - within the measurement errors of Donati et
al. (1997) - the observed modes are the same, it can be concluded from calculations
of the pulsation constant that HD 104237 must be a high overtone PMS pulsator.
4.2. The PMS instability strip
25
Figure 4.2: Part of the original light curve of HR 5999 taken from Kurtz & Marang
(1995). On top of the long-term light variation with high amplitude the pulsation
with a period of ∼ 5 hours is clearly visible.
The pulsation of V351 Ori and HD 35929 was discovered by Marconi et al.
(2000) during a search for δ Scuti type variability in seven Herbig Ae stars using
Strömgren uvby time series photometry. At that time it was only possible to derive
a single, ‘cycle-count’ period for each of the two stars, namely 1.4 hours for V351
Ori and 4.7 hours for HD 35929.
NGC 6823 HP 57 and NGC 6823 BL 50 were found to be pulsating PMS
stars by Pigulski et al. (2000) within a search for variable stars in the young cluster
NGC 6823 using U BV (RI)C CCD time-series photometry. Both stars show two
periods simultaneously: BL 50 with 1.7 hours (IC amplitude = 18 mmag) and with
2.4 hours (IC amplitude = 6 mmag); HP 57 with 1.9 hours (IC amplitude = 27
mmag) and with 1.5 hours (IC amplitude = 20 mmag). The authors considered the
two stars as PMS δ Scuti like members of the cluster.
A list of the eight known δ Scuti like pulsating PMS stars as of 2000, corresponding to the beginning of this work, is given in Table 4.1. Note that - except for the
two stars in NGC 6823 - for each δ Scuti type PMS pulsator only a single frequency
could be detected at that time.
4.2
The PMS instability strip
The comparison of the hot and cool border of the classical instability strip with
observations has been an important test for stellar structure and evolution codes.
The determination of these borders by dedicated observations of PMS stars will be
comparably important for the theory.
4.2.1
Theoretical investigations
Marconi & Palla (1998) studied the theoretical instability properties for PMS stars
for the first time and investigated whether a PMS star can indeed pulsate. PMS
26
4. Pulsation in PMS stars
Name
V589 Mon
V588 Mon
NGC 6823 HP57
NGC 6823 BL50
V351 Ori
HR 5999
HD 35929
HD 104237
RA (2000.0)
[hh:mm:ss]
06:39:28.46
06:39:05.9
19:43:06.78
19:43:09.07
05:44:18.79
16:08:34.29
05:27:42.79
12:00:05.08
DEC (2000.0)
[dd:mm:ss]
+09:42:4.1
+09:41:3.4
+23:16:37.8
+23:17:49.6
+00:08:40.4
-39:06:18.3
-08:19:38.4
-78:11:34.6
sp
F2 III
A7 III/IV
A7 IIIe
A7 III/IVe
F0 IIIe
A4 V
V
[mag]
10.32
9.73
14.60
14.50
8.90
6.98
8.20
6.60
log Teff
log L/L¯
3.85
3.90
3.86
3.86
3.87
3.85
3.86
3.93
1.51
2.05
1.25
1.60
1.15
2.12
1.92
1.50
f
#
1
1
2
2
1
1
1
1
Table 4.1: Parameters of the eight known PMS pulsators in the year 2000: name,
right ascension (RA) and declination (DEC) at the epoch 2000.0, spectral type (sp) if
available, V magnitude, effective temperature (log Teff ) and luminosity (log L/L¯ ).
The last column specifies the number of frequencies found in each star as of 2000.
evolutionary models were computed for low- and intermediate-mass stars starting
at the birthline determined by the protostellar accretion phase (Palla & Stahler
1990 & 1993). Several sequences of linear non-adiabatic radial pulsation models at
fixed mass covering a wide range of luminosities and effective temperatures were
used by the authors to provide information about periods and modal stability in
PMS stars. The study of Marconi & Palla (1998) is limited to the first three radial
modes of pulsation and was especially developed for the case of HR 5999. Marconi
& Palla (1998) estimate the location of the blue boundaries of the theoretical PMS
instability strip for each mode, but no information about the theoretical definition
of the red boundary is given because the strong effects of convection are not taken
into account. However, the authors found their red edge to lie between 6500 ≤ Teff
≤ 7100 K and the blue edge between 7100 ≤ Teff ≤ 7500 K.
The κ and γ mechanisms in the hydrogen and helium ionization zones are assumed to drive the pulsation in such young stars. The time PMS stars spend in the
instability region is typically 5–10% of the total PMS contraction time, the KelvinHelmholtz time scale. For a 1.5 M¯ star this accounts to ∼ 106 yr and for a 4.0
M¯ star it is only 8 × 104 yr. Although this phase lasts relatively short, a number
of PMS stars have the right combination of effective temperature and luminosity to
become pulsationally unstable.
4.2.2
Comparison with observations (status 2000)
Figure 4.3 shows the HR-diagram with PMS evolutionary tracks from D’Antona &
Mazzitelli (1994, solid lines) for 1.5, 2.0, 2.5 and 3.0 solar masses and the location
of the - in the year 2000 - known eight PMS pulsators (coloured symbols). Also,
the borders of the classical δ Scuti instability strip for the theoretical, fundamental
blue edge (BEF , thin solid line), the theoretical, general blue edge (BE, thick solid
line) and the empirical red edge (RE, dotted line) are drawn (Breger & Pamyatnykh
1998). The PMS instability strip (Marconi & Palla 1998) for the first three radial
modes is marked as dot-dashed blue lines. It can be seen that the PMS blue edge
4.3. Seismology of PMS stars
27
for the second overtone mode matches well with the post-main sequence blue edge
for the fundamental mode. Out of the eight stars known at that time, six fall inside
the theoretical PMS instability region and the other two stars have been thought to
be rather the exception than the rule. At that time it was believed that pre-main
sequence stars pulsate rather monoperiodically and purely radial. The two stars
located close to the general blue edge of the classical instability region gave a hint
that PMS stars could indeed pulsate with higher overtone modes. But such stars
just had not been discovered at that time. As the statistics with only eight stars
was very poor, new detections of PMS pulsators were urgently needed.
Figure 4.3: HR-diagram with the location of the eight known PMS pulsators in
the year 2000, the borders of the classical and the PMS instability strips and PMS
evolutionary tracks (see text for additional information).
4.3
Seismology of PMS stars
The unstable modes in pulsating PMS stars known so far are the same as those for
classical δ Scuti stars, namely low radial order g- and p-modes. Frequencies of l = 0
modes computed with same radial order are nearly identical for pre- and post main
28
4. Pulsation in PMS stars
sequence stars (Suran et al. 2001), because stars of both evolutionary stages have
similar mean density and outer layers. For nonradial modes (l > 0) the patterns
are more complicated due to evolutionary changes in the stellar interior.
Avoided crossings exist only for post-main sequence stars, because nuclear reactions in the stellar interior of the more evolved stars produce the internal structure
responsible for such a phenomenon. The inner parts of pre-main sequence stars are
quite homologous without the presence of nuclear reactions, which is the reason for a
lack of avoided crossing. This is a very interesting difference to post-main sequence
stars and can be tested using asteroseismology.
The theoretical pulsation frequency spectra of pre- and main sequence stars
with same masses, effective temperatures and luminosities look quite similar at first
glance (Figure 4.4, Pamyatnykh, private communication). In the case of only a few
observed frequencies it always will be possible to find a pre- and main sequence
model which fits the observations within the errors equally well. However, if a larger
part of a frequency spectrum is available, frequency spacing allows to distinguish the
models. Of course, longer time series obtained with multi-site campaigns or using
space telescopes are needed to derive a dense enough pulsation frequency spectrum.
Hence, it would be possible to discriminate between different evolutionary stages of
stars located in the same region of the HR-diagram from analysis of their oscillation
frequency distributions (Suran et al. 2001).
Figure 4.4: Differences of non radial pulsation frequencies for a two solar mass star
with same Teff = 7900 K and L/L¯ = 1.35 in the pre- (dots) and main sequence
(circles) phases for l = 0, 1 and 2 (Pamyatnykh, private communication).
Chapter 5
Young open clusters
Open clusters appear to be continuously forming in the galactic disk, and, in principle, direct studies of the physical processes leading to their formation are possible.
These studies have been seriously complicated by the fact that galactic clusters
form in giant molecular clouds and during their formation and earliest phases of
evolution they are completely embedded in molecular gas and dust, and are thus
obscured from view. Hence, observations are extremely difficult in the optical range
and the situation improved only because of technical developments of IR astronomy
and detectors.
These new observations revealed that embedded clusters are quite numerous and
that the vast majority of stars may form in such systems. Furthermore, open clusters
span a wide range of stellar mass within a relatively small volume of space. Hence,
their study can directly address a number of fundamental astrophysical questions
concerning the origin and early evolution of stars and planetary systems.
Young clusters are most suitable to search for pulsating PMS stars because all
members have the same age and distance, hence confusion with more evolved objects
can be reduced. Those members which have not yet evolved to the ZAMS therefore
can be most probably identified as PMS stars.
5.1
Basic definitions
A cluster is defined as a group of stars that are physically related and whose observed
stellar mass volume density is large enough to stabilize the group against tidal
disruption by the galaxy (Lada & Lada 2003).
A cluster consists of enough members to ensure that its evaporation time1 is
greater than 108 yr, the typical lifetime of open clusters in the field. Hence, a stellar
cluster normally has more than 35 members allowing to distinguish between multiple
systems with less than six members and stellar associations being loosely grouped,
physically related stars.
1
i.e., the time it takes for internal stellar encounters to eject all its members.
29
30
5. Young open clusters
5.1.1
Embedded and exposed clusters
It can be distinguished between two environmental classes depending on their association with the interstellar matter:
• Exposed clusters possess little or no interstellar matter within their boundaries. Almost all clusters found in standard open cluster catalogs (e.g. Lynga
1987) fall into this category, e.g. the 5 Myr old NGC 2362.
• Embedded clusters are fully or partially embedded in interstellar dust and gas.
They are frequently completely invisible at optical wavelengths and best detected in the IR. These are the youngest known stellar systems and can also be
considered protoclusters because upon emergence from molecular clouds they
will become exposed clusters. Known members of this group are, for example:
NGC 2264, the Trifid nebula, NGC 6611 and NGC 6530, the latter studied in
this work.
The embedded phase of cluster evolution appears to last 2 – 3 Myr. Clusters with
ages larger than 5 Myr are rarely associated with molecular gas (Leishawitz et al.
1989).
5.2
Sequential star formation
Young open clusters provide important information concerning star formation processes. Most massive stars show a relatively small age spread. Therefore the formation of massive stars in young clusters is nearly coeval, whereas low-mass cluster
members have longer pre-main sequence lifetimes and are still in their pre-main
sequence stage.
The age and mass of a pre-main sequence star can be estimated using PMS
evolutionary models. This allows to gain important information on star formation
history as well as an initial mass function (IMF) of the cluster. Before that the
crucial question of membership criteria especially for the low-mass stars in the PMS
stage has to be settled.
5.3
PMS stars in young open clusters
In a cluster that is only a few million years old, fainter members are still in the
process of gravitational contraction from the prestellar medium to the ZAMS. As
the contraction rate is higher for more massive stars, the CMDs for young clusters
consist of a normal main sequence for the brightest stars, which extends to some
point depending on the age of the cluster. Fainter stars of later spectral types have
not reached the ZAMS yet, hence are still in their pre-main sequence evolutionary
phase.
The lack of a complete cluster main sequence makes it difficult to obtain reliable
estimates of the cluster distances and ages, which is the reason why the ages of
5.3. PMS stars in young open clusters
31
such young clusters typically have relatively large error bars on the order of up to
the cluster age itself. However, this fact emphasizes the relative youth of the corresponding cluster making it impossible that its A to F type members have already
evolved off the ZAMS.
As all cluster members have the same age and distance, confusion with more
evolved objects can be avoided. The number of cluster stars showing the spectral
types of interest is typically around 15 to 20, hence providing a good sample of
candidates for the search for pre-main sequence pulsators.
The clusters selected for the search for pulsating PMS stars had to meet the
following criteria and have been selected accordingly:
• Their ages are lower than 10 Myr.
• They possess a normal main sequence down to spectral types of about B9/A0,
while fainter stars of later spectral types are still gravitationally contracting
towards the ZAMS, hence are in their pre-main sequence evolutionary phase.
• A significant number (say N ≥ 10) of cluster members possess the spectral
types of interest.
• Previous determination of the positions, magnitudes and colors of cluster members are available from the literature. For many of the extremely young clusters, not even the position and magnitudes of its members have been studied.
One reason may be that the cluster is located in a highly obscured region still
hiding its members from view in the visual spectral range. For some clusters it is also difficult to determine its dimension on the sky or to identify its
particular members. But to be able to conduct an investigation of the pulsational behaviour of cluster stars, at least the basic information of position and
magnitude are necessary.
NGC 6383, IC 4996 and NGC 6530 have been chosen to conduct the search of
pulsating PMS cluster members using the criteria mentioned above.
32
5.4
5. Young open clusters
NGC 6383
The young open cluster NGC 6383 belongs to the Sgr OB1 association together with
NGC 6611, NGC 6530 and NGC 6531 and is centered around the bright spectroscopic
binary HD 159176. A number of authors has studied the cluster photometrically
and spectroscopically in the past, but no search for variability and/or pulsation has
been performed before. An overview of the studies available in the literature is given
below.
Figure 5.1: False color image of the region of NGC 6383 with a field of view of 20’
× 20’ (taken from the First Digitized Sky Survey).
NGC 6383 can be compared to three other well-studied clusters of similar age,
NGC 2264, NGC 6530 and the Orion nebula region. The absence of a dense nebula
with much dust is significant. NGC 6383 may be a case in which the formation of
smaller mass stars ceased prematurely after the formation of the central cluster of
massive stars, resulting both in a lack of faint stars and in the absence of T Tauri
stars as bright as those found in regions where star formation has continued.
α2000
δ2000
age
diameter
distance
17h 34.8m
−32◦ 340
1.7 ± 0.4 Myr
200
1.5 ± 0.2 kpc
Table 5.1: Main cluster properties for NGC 6383.
5.4. NGC 6383
5.4.1
33
Historical background
The cluster was first observed photoelectrically by Eggen (1961), who found that its
CMD resembles that of the very young cluster NGC 2264 studied by Walker (1956).
It consists of a normal main sequence to a spectral type of about A0 and stars
beyond were considered in the state of contraction. While NGC 2264 is embedded
in bright and dark nebular matter, Eggen did not find any nebulosity associated
with NGC 6383. He also determined the distance modulus for the cluster to be
+10.5 mag and the color excess E(B − V ) = 0.30 mag.
Thé (1965) observed photographically a total number of 99 stars in NGC 6383
down to V = 13.8 mag located in a circular area with a radius of about 12 arcminutes
and confirmed Eggen’s result.
Fitzgerald et al. (1978) characterized NGC 6383 as a young open cluster with
a strong central core and a possible extended halo. They performed photoelectric
U BV photometry and MK spectroscopy of 25 stars within 2 arcminutes of the
center of the cluster and confirmed the pre-main sequence nature of stars redder
than (B − V )0 ' 0.0. The authors also claim that their star #3 is a foreground B9
IV star with a faint, very close companion and that the F type star #21 is also not a
cluster member. Star #24 has a spectral type of B8 Vn and showed emission at Hβ
during one night of their observations. It is interpreted by the authors as an early
type flare star undergoing the final stages of pre-main sequence contraction. Star
#10 is probably a variable star (∆ V = 0.3 mag) in its pre-main sequence stage of
evolution and still surrounded by the remnants of its protostellar cloud. Fitzgerald
et al. (1978) think that the central star, HD 159176, is significantly older than
the rest of the cluster core. This massive binary system may have stimulated the
formation of the cluster core stars and maybe of the stars in the outer regions as
well.
Lloyd Evans (1978) confirmed that the fainter stars in NGC 6383 fall above the
main sequence and are presumed to be still contracting to it. He obtained U BV
photometry of 86 stars down to V = 18.1 mag and 33 spectra of 16 stars and
discussed the interstellar reddening in the field of the cluster as well as the cluster
membership. The author found several variable stars – six of them seem to be premain sequence variables – and determined the upper age limit of NGC 6383 to be
5 · 106 yr.
Thé et al. (1985) discussed the spectral energy distribution of stars above the
ZAMS in the central part of NGC 6383 using spectroscopic and photometric observations in the red and near-IR. For the photometry they used the Walraven W U LBV
system supplemented by V RI (Cousins’ system) and JHKL(M ) measurements.
Their most interesting result is that three stars were found to have excess IR radiation most probably due to thermal emission of circumstellar dust grains, indicating
that they are pre-main sequence objects. The central star, HD 159176, is found to
be a double-lined binary with an effective temperature of ∼38 500 K for which a
mass loss rate of ∼ 10−6 M¯ yr−1 can be expected. It has no excess infrared radiation up to ∼ 5 µm. Their star #2, HDE 317847, has a mass loss rate smaller than
34
5. Young open clusters
10−8 M¯ yr−1 , furthermore also no excess near-IR radiation could be found. Stars
#20 and #24 show excess near-IR radiation, which is in agreement with the fact
that the observations shortward of the Balmer jump show substantial UV excess.
Assuming appropriate physical parameters of the extended gaseous shell, which is
responsible for the above mentioned emission, the observed near-IR excess can be
explained (see Schild et al. 1974, Gehrz et al. 1974). Star #4 in the study by Thé
et al. (1985) shows variability of ∼ 0.1 mag, which is in agreement with the V RI
measurements. Also their star #5 is suspected to be variable. E(B − V ) for these
stars is higher than the mean foreground color excess which is caused by circumstellar dust shells. The authors examined the temperatures, masses and distances
of the dust shells from the central star in detail. Star #6 of Thé et al. (1985) is an
early A-type star showing strong IR-excess. Stars #4, #5 and #6 are located above
the main sequence in the HR-diagram and show near-IR excess caused by thermal
re-emission of a heated dust shell which supports the idea that they are genuine
pre-main sequence stars. Star #3 is a spectroscopic variable where asymmetric hydrogen lines show the presence of a gas shell. Its location in the CMD indicates
that it is (probably) a pre-main sequence giant. #T54 might be a foreground object
because its color excess E(B − V ) is much smaller than the average value of the
cluster.
14 stars mostly located in the core of the cluster were selected by van den Ancker
et al. (2000) from the publication by Thé et al. (1985). Low resolution CCD spectra
of these stars were obtained and new spectral classifications were performed. No
deviation from a normal interstellar extinction law (i.e. Rv = 3.1) could be found.
Stars classified with luminosity classes III and IV seem to be located to the right
of the main sequence and therefore are probably true pre-main sequence stars. But
no strong correlation between the position in the HR-diagram and the presence of
an infrared excess seems to be present within their sample. Star #4 seems to be a
new Herbig Ae/Be type star because it shows a large IR excess, Hα in emission and
some indications for the presence of circumstellar gas in the spectrum.
Rauw et al. (2003) report the detection of a number of X-ray sources associated
with the cluster using observations performed with XMM-Newton.
5.5. IC 4996
5.5
35
IC 4996
IC 4996 is located in the direction of Cygnus, 40 pc above the galactic plane, and
is part of a large region with active star formation that contains other young open
clusters and Wolf-Rayet stars. An IRAS map of the region (Lozinskaya & Repin
1990) shows the presence of a dusty shell that surrounds the cluster. The age
and distance estimates from different authors agree with each other: the cluster is
slightly younger than 107 years and is located ∼ 1.7 kpc from the Sun (see Table 5.2).
In particular, the inferred age indicates the likely existence of pre-main sequence
members in the range of spectral types A and F.
Figure 5.2: False color image of the region of IC 4996 with a field of view of 10’ ×
10’ (taken from the First Digitized Sky Survey).
α2000
δ2000
age
diameter
distance
20h 16m 30s
37◦ 380 000
8.87 Myr
60
1.732 kpc
Table 5.2: Properties of IC 4996 taken from the WEBDA database.
5.5.1
Historical background
Alfaro et al. (1985) obtained uvby and Hβ observations for 15 stars brighter than
V = 12 mag with spectral types earlier than A0 and report a mean distance modulus
of 11.43 ± 0.31. The age of IC 4996 was estimated by them to be 7.5 · 106 years.
36
5. Young open clusters
Vansevic̆ius et al. (1996) performed CCD observations in the BVRI system of
126 stars in the central part of the cluster. It seems that IC 4996 lies in an area
where interstellar extinction is large and variable across the cluster. The authors
fitted theoretical isochrones to the observed CMDs and determined the age of the
cluster to be 9 ± 1 · 106 years.
A total of 1120 stars was measured by Delgado et al. (1998) in a field of ∼7’ × 7’
of IC 4996. They found the average values of the colour excess and true distance
modulus to be E(B − V ) = 0.71 ± 0.08 mag and (V0 − MV ) = 11.9 ± 0.1 mag. There
seems to be an indication that the cluster had two episodes of star formation: the
existing lower main sequence stars were formed first, and the presumed PMS members are the result of a second episode of star formation. Using isochrone fitting to
the upper part of the sequences in the CMDs the authors derived a cluster age of
7 ± 3 · 106 years, which corresponds well with the other values from the literature.
Delgado et al. (1999) performed spectroscopic observations with the aim to
estimate radial velocities and spectral types for 16 proposed PMS stars in order to
confirm or reject their cluster membership. They also searched for possible spectral
features indicative of PMS nature. The heliocentric radial velocity of cluster stars of
–12 ± 5 kms−1 is in good agreement with published values of other young clusters
which are also located in the Cygnus star forming region.
A spread in the color-color diagram is detected (Delgado et al. 1999), which is
probably due to the diversity of actual reddening features and ages. Regardless of
this effect the observed stars span a range in colors and spectral types that nicely
links the coolest HAEBE with the hottest T Tauri stars. It has to be noted that the
PMS objects can mix with fore- or background field stars in the CMD and cannot
be unambiguously separated from each other.
In any case, the presence of a pre-main sequence in IC 4996 covering a range in
spectral types from A to early G is strongly confirmed by different authors.
5.6. NGC 6530
5.6
37
NGC 6530
NGC 6530 is located in the central part of the HII region M8, the Lagoon nebula (see
Figure 5.3). Since the first study performed by Trumpler already in 1930, several
investigations have been devoted to study this cluster and to estimate its parameters.
A review of the publications on NGC 6530 that are considered important for this
work is given below, the main cluster properties taken from the literature are listed
in Table 5.3.
Figure 5.3: Mosaic image of M8, the Lagoon nebula (Copyright by Robert Gendler;
www.robgendlerastropics.com).
Author
Walker (1957)
Kilambi (1977)
Chini & Neckel (1981)
Mc Call (1990)
Sung (2000)
V0 − MV
[mag]
E(B − V )
[mag]
10.7
age
[Myr]
3.0
1.0–3.0
distance
[kpc]
0.39±0.09
11.35±0.08
11.25±0.1
1.86±0.07
0.35
1.5
Table 5.3: Astrophysical properties of the PMS pulsators in NGC 6530.
5.6.1
Historical background
Walker (1957) concluded from UBV photoelectric observations of 118 stars that
the CMD of this cluster consists of a normal main sequence extending from O5
to about A0 with stars of later spectral type still contracting towards the ZAMS.
This was confirmed later by several authors (e.g. Kilambi 1977; van den Ancker
et al. 1997; Sung et al. 2000). Walker (1957) already found that for stars fainter
38
5. Young open clusters
than V = 14 mag the effect of the nebulosity surrounding the stars is causing large
irregular variations in the brightness.
Lada et al. (1976) performed the first millimeter-wave observations of this cluster
and took high quality optical interference-filter photographs toward the NGC 6530M8 star forming region. They find the bright O7 star Herschel 36 to be a newly
born star surrounded by circumstellar dust visible in the IR.
Kilambi (1977) obtained U BV photographic photometry of NGC 6530 and found
that all stars fainter than V = 12.0 mag have not reached the ZAMS yet. They
also determined the cloud temperature to lie between 5 and 10 K and encountered
deviations from a normal galactic reddening, which seem to occur mainly in the
regions surrounded by nebulous material.
van den Ancker et al. (1997) obtained Walraven WULBV, Johnson/Cousins
UBV(RI) and near-IR JHK photometric data and performed spectroscopy of NGC
6530 on different sites. They report about a good agreement between spectral classifications from photometry and from spectroscopy, indicating that the assumption
of a normal extinction law when obtaining the classifications from photometry is not
too far off. The authors found that the cluster contains a mixture of normal main
sequence stars, young stars still contracting towards the ZAMS, as well as older stars
evolving off the main sequence. Hence, they conclude that star formation must have
started a few times 107 years ago and probably is continuing up to now.
37 pre-main sequence stars with Hα emission were detected in NGC 6530 by
Sung et al. (2000) using U BV RI and Hα photometry. They also derived the cluster
age to be 1.5 million years with an age spread of about 5 million years. Moreover the
authors confirm the presence of a small amount of differential reddening across the
cluster. Low-mass pre-main sequence stars being at relatively earlier stages of their
PMS evolution are more likely to be obscured by circumstellar disks than relatively
more evolved PMS stars.
119 X-ray point sources in the Lagoon Nebula region have been recently detected
by Rauw et al. (2002) in a 20 ks XMM-Newton observation. They found that
most of the X-ray sources are associated with pre-main sequence stars of low and
intermediate mass. A larger list of X-ray point sources with a much better spatial
resolution was obtained by Damiani et al. (2004) using Chandra ACIS-I X-ray
data. One of the most important features in their CMDs is the well defined blue
envelope of the CMDs; it is due to the presence of the giant molecular cloud, which
prevents us from seeing field stars (mostly main-sequence) more distant than the
cloud. Therefore, the well defined blue envelope of the CMDs is populated by mainsequence field stars at the distance of the cloud. Background field stars are highly
obscured by the cloud and therefore they would be visible at magnitudes and colors
much fainter and much redder than their intrinsic values.
5.6.2
Cloud collapse and star formation
NGC 6530 is embedded within ionized gas, where at least six O stars contribute to
the ionization of the region. The flattened appearance of the zone of massive star
5.6. NGC 6530
39
formation, which is manifested by similarly elongated distributions of molecular and
ionized gas and heated dust, suggests that cloud collapse was not symmetrical.
A phenomenological model for the evolutionary history and structure of the M8
region was given by Lada et al. (1976): The stars were born about 2 · 106 years
ago at the edge of a massive molecular cloud. Since then, the stars have moved
away from and/or have severely disrupted the portion of the cloud in which they
were born. The remnants of the hole left by these stars in the cloud are visible as
the outermost bright-rim structure and low-surface brightness Hα emission observed
toward M8. This hole allows to see deeper into the molecular cloud, where, possibly,
more recent star-forming activity has taken place.
The stars which lie above the ZAMS show a great scatter which might be due to
variations of initial formation conditions in the cloud, due to an age spread in the
formation of stars or to activity of the shell structure itself (Kilambi 1977). There are
at least 10 stars, which lie below the main sequence between −0.25 ≤ MV ≤ +1.75
(Kilambi 1977). The location of such stars in NGC 6530 and other young clusters
such as NGC 2264 has found no natural explanation in the context of standard premain sequence evolutionary tracks. If the location of pre-main sequence stars in the
CMD is affected by the presence of circumstellar shells, it seems logical to assume
that stars below the main sequence reflect extreme shell phenomena. These stars
may be surrounded by gas and dust shells, which are optically thick in the visible
making the stars less luminous. At the same time the incipient emission from the
gas shell makes them display uncommonly negative color indices. The combined
effects of both gas and dust will place these stars below the ZAMS.
5.6.3
Proper motion studies
For 363 stars brighter than 13.05 mag proper motion distribution parameters have
been determined by van Altena & Jones (1972). As a large difference in accuracy of
the motions between the bright and faint stars was noted, the other and fainter 135
stars of their sample have not been included into their absolute parameter solution.
But anyway van Altena & Jones (1972) tried to compute membership probabilities
for such faint stars using the parameters for the bright stars, even if they are not
strictly applicable. The errors of the proper motions for the faint stars are about
twice as large as for the brighter stars. The validity of the authors’s analysis is
problematic as low membership probabilities were assigned to many of the early
type stars (Sung 2000).
Chapter 6
Observations and data reduction
6.1
NGC 6383
For NGC 6383, the first cluster investigated in this study, CCD photometric time
series in Johnson B & V filters were obtained with the 0.9m telescope (f /13.5) at
the Cerro Tololo Interamerican Observatory (CTIO), Chile. Between Aug. 11 and
Aug. 24, 2001, NGC 6383 was observed using the 2048 x 2046 SITe CCD chip,
which provides a field of view (fov) of 13.5’ × 13.5’ (see Figure 6.1) with a scale of
0.396”/px. In total, 53.25 hours of time-series photometry could be acquired within
8 clear out of 14 granted nights (Table 6.1).
Figure 6.1: False-color image of the observed field of NGC 6383 (fov ∼ 13.5’ × 13.5’,
South is at the top and East is to the left).
40
6.1. NGC 6383
41
Figure 6.2: Raw image of the observations of NGC 6383 read out by four amplifiers.
The CCD chip was read out in quad mode by four amplifiers providing a readout
time of 32 sec, which was chosen due to the high time resolution needed. An example
of the original image is shown in Figure 6.2, illustrating the overscan strip lying in
the center of the frame as well as the slightly different electrical offsets (i.e. bias
levels) of the four quadrants of the CCD. On the right side a bad column is situated,
but its presence affected the data acquisition only marginally: it had to be assured
that no important star is located close to this zone. Although sky flats in both filters
were obtained every evening, the 10 dome flats per filter (with exposure times of 90
sec in B and 50 sec in V ) turned out to be better for flat-fielding the images.
The basic reductions (bias subtraction, flat-fielding) were performed using the
IRAF ared.quad1 package. The Multi Object Multi Frame (MOMF) software developed by Kjeldsen & Frandsen (1992) was used to extract the photometric signal.
MOMF is optimized to analyze photometric time series (i.e. a large amount of
CCD frames per night) of semi-crowded fields by combining point-spread function
(PSF) fitting and aperture photometry. The reduction with MOMF relies on the
selection of 10 stars at the beginning, which are used to compute the PSF that will
be applied to all stars on the frames. Each of the 10 stars is used as reference for
compensation of tracking errors of the telescope and for the reduction itself. MOMF
determines absolute and relative magnitudes of each star identified on the frames
and their corresponding standard deviations. The absolute values are raw, uncorrected, instrumental magnitudes, whereas the relative light curves are determined by
subtracting a weighted mean of all stars on the frame. Variable and non-variable,
1
The IRAF ared.quad package was especially developed by NOAO for reduction of CCDs used
at CTIO and KPNO observatories read out in quad mode.
42
6. Observations and data reduction
extremely red or blue stars are used to determine the weighted mean, requiring
colour-dependent extinction corrections.
On the observed images 286 non-saturated stars have been identified (see Figure
6.3), for which light curves using the optimum aperture producing minimal pointto-point scatter were generated. Nightly means were subtracted to correct for zeropoint changes and long-term irregular light variations, which most likely are due to
variable extinction by circumstellar dust.
2000
1800
1600
1400
px
1200
1000
800
600
400
200
0
0
200
400
600
800 1000 1200 1400 1600 1800 2000
px
Figure 6.3: Schematic map of the observed field of NGC 6383 (fov ∼ 13.5’ × 13.5’,
South is at the top and East is to the left) with all stars measured in Johnson B &
V, where 1 pixel (px) corresponds to 0.33 arcseconds.
For all 286 stars, a detailed frequency analysis was performed in both filters using
the Fourier Analysis program Period98 (Sperl 1998) which is based on the Discrete
Fourier Transformation (DFT, Deeming 1975) and provides a multi-sine fit option.
A signal was considered to be significant, if it exceeds four times the noise level in
the amplitude spectrum (Breger et al. 1993, Kuschnig et al. 1997). The errors of
amplitudes, σ(A), and frequencies, σ(f ), were calculated using the relations given
by Montgomery (1999):
r
2
· σ(m)
N
(6.1)
6
1
σ(m)
·
·
,
N πT
A
(6.2)
σ(A) =
r
σ(f ) =
6.1. NGC 6383
43
where σ(m) is the rms magnitude of the data set, A the corresponding amplitude,
N is the number of data points and T is the time base of the observations.
Our own star numbers are used, cross references with the literature are given
according to the publications by Fitzgerald et al. (1978), e.g. F 4, by Thé (1965),
e.g. T 47, and Lloyd Evans (1978), e.g. EV 281. All photometric measurements for
the stars in NGC 6383 including the cross references, where available, are listed in
the Appendix.
6.1.1
Bias level variations
During the reduction an interesting effect was encountered: the bias level changed
from night to night with the ambient temperature. It increases at lower temperatures
and decreases at higher outside temperatures (see Figure 6.4; note the different scales
on the y-axes!). As a consequence the bias correction had to be performed separately
for each night.
Figure 6.4: Changing bias level with ambient temperature of the four quadrants of
the CCD which are read out by four different amplifiers (Note the different scaling
between top, AMP11 and AMP12, and bottom, AMP21 and AMP22!).
44
6.1.2
6. Observations and data reduction
Color dependent extinction
A systematic effect was encountered for some of the light curves. Towards the end
of the nights some stars became continuously brighter, but others fainter. The
corresponding Bouguer plots (i.e. magnitude vs. airmass) showed that the different
colors of the stars were the explanation. Hence, the extinction correction had to
include also the color-dependent coefficient k 0 (Sterken & Manfroid 1992):
m = m0 − (k 0 + k 0 · CI) · X,
(6.3)
where m0 is the uncorrected magnitude, X is the airmass and k 0 the principal
extinction coefficient. In our case, the color index CI was taken as (B − V ).
4
0.02
0
slope k
(B-V) literature
3
2
-0.02
1
-0.04
0
0
1
2
(B-V) instrumental
3
4
-0.06
0
0.5
1
1.5
2
2.5
(B-V)
Figure 6.5: Modelling the color-dependent extinction effect: left: Determination of
the (B − V )trans of all stars using an inverse second-order polynomial (solid line),
which describes the relation between instrumental and literature (B − V ) values for
97 of 286 stars. right: Dependence of the slope of the Bouguer plot, k, on (B−V )trans
and weighted linear regression. As an example data from the 8th night are shown
in this figure. The symbol areas correspond to the weights of the individual data
points, where larger symbols are related to higher weights.
For only 97 stars (B − V ) values were available in the literature and they show a
clear correlation with the slope, k, of the Bouguer plots. However, it was necessary
to transform the instrumental (B − V ) values for all observed stars to the standard
system to be able to correct for the color dependent extinction effect. Hence, the
relation between the 97 stars with (B − V ) from the literature and the instrumental (B − V ) values from our observations is modelled by an inverse second-order
polynomial (solid line in Figure 6.5). The three polynomial coefficients evaluate to
6.2. IC 4996
45
a0 = −0.712 ± 0.052, a1 = +0.692 ± 0.084 and a2 = +0.177 ± 0.029. (B − V )instr
values for all stars could then be transformed to the standard system according to:
s
(B − V )trans =
a21
((B − V )instr − a0 )
a1
+
−
a2
2 a2
4 a22
(6.4)
where (B −V )trans are the transformed indices and (B −V )instr are our instrumental
values (see Fig. 6.5).
6.2
IC 4996
Figure 6.6: False-color image of the observed field of IC 4996 (fov ∼ 6’ × 6’).
IC 4996 was observed with the 1.5m telescope (f /8) at Sierra Nevada Observatory, Spain, between Sep. 2 and Sep. 15, 2002, in Johnson B & V filters using a 1k x
1k CCD chip with a scale of 0.33”/px providing a field of view of 6’ × 6’ (see Figure
6.6). In total, 69.76 hours of time-series photometry have been obtained in 12 out of
14 granted nights (see Table 6.1). Only the first ten nights (corresponding to 62.38
hours of observations) were used for the analysis because during the last two nights
the weather conditions were too bad resulting in an extremely high scatter in the
light curves. As the quality would have decreased significantly, the according data
have been rejected for the frequency analysis.
The images were already bias- and dark-corrected by the standard procedure
used at the observatory. The flat field images showed ring-shaped structures which
looked different in the V and B filters indicating the presence of dust particles on
the filter (see Figure 6.7). Also, the flat field images did not have the required “flat”
46
6. Observations and data reduction
shape. Hence, the creation of the combined “super” flat field images was performed
using an own code written in the IDL language.
Figure 6.7: Original flat field image of the observations performed at OSN. Clearly
visible are the ring-shaped structures.
The actual flat field correction was performed within the reduction software
PODEX written by Kallinger (2005). PODEX allows to extract the photometric signal
of CCD time series photometry using a combination of aperture photometry and
point-spread function fitting. For each selected star the light curve is computed.
The mean value of the comparison light curve is then subtracted, where the stars
used for the comparison light curve can be selected arbitrarily, optionally flat field
and (color-dependent) extinction corrections can be applied. The output of the
program not only contains the photometric signal at each integration, but also the
according value of the airmass and the signal of the comparison light curve. The
advantage of PODEX is that variable stars can be identified easily and immediately
deselected for the computation of the comparison light curve.
The 113 observed stars lie in the range between V = 11 − 18 mag (see Figure
6.8).
Unfortunately - due to observations performed in service observing mode - the
exposure times were 90 sec per default for each image. This was not enough for
such faint stars. Fortunately, always two frames per filter were taken after each
other. This allowed us to sum subsequent images to improve the signal-to-noise
ratio considering that the time resolution decreased.
Again the resulting light curves had to be corrected for zero-point changes and
long-term irregular light variations coming from the surrounding nebulous area by
subtraction of nightly means. Also, the effect of color-dependent extinction had to
be taken into account similar as in the case of NGC 6383.
For all 113 identified stars a detailed frequency analysis was performed using
primarily again Period98, but also the new software SigSpec developed by Reegen
6.2. IC 4996
47
1000
900
800
700
px
600
500
400
300
200
100
0
0
100 200 300 400 500 600 700 800 900 1000
px
Figure 6.8: Schematic map of the observed field of IC 4996 (fov ∼ 6’ × 6’) with all
stars measured in Johnson B & V, where 1 pixel (px) corresponds to 0.33 acrseconds.
(2005). Only if frequencies appeared significant using both methods they are believed to be intrinsic and not artefacts of the reduction.
Our own star numbers are used, but cross references with numbers given by
Delgado et al. (1998 & 1999), e.g. D 32, and Purgathofer (1964), e.g. P 66, are
listed. The photometric measurements of all stars including cross references are
listed in the Appendix.
6.2.1
SigSpec
SigSpec computes significance levels for amplitude spectra of time-series with arbitrarily given sampling. The probability density function (PDF) of a given amplitude
level is solved analytically including dependence on frequency and phase. A detailed
description of this concept is given by Reegen (2005). For a given time-series dataset
SigSpec calculates both an amplitude and a significance spectrum and the frequency,
amplitude and phase at maximum significance in the considered frequency range.
Optionally consecutive prewhitening is provided to perform multi-frequency analysis.
The significance of an amplitude (A) is calculated using:
sig(A) := −lg[ΦFA (A)],
(6.5)
where ΦFA is the false alarm probability. A significance of 8, for example, means that
in one out of 108 cases the according amplitude is due to noise. A signal-to-noise
(S/N) ratio of 4 (Breger et al. 1993) corresponds to a significance of 5.46.
48
6. Observations and data reduction
In this work, a signal was considered to be significant, if it exceeded a S/N of 4
and yielded significances higher than 5.5.
6.3
NGC 6530
Figure 6.9: Composed false-color image of the two observed fields of NGC 6530.
The overlapping region is clearly visible. Each frame has a fov of ∼ 13.5’ × 13.5’,
South is at the top and East is to the left.
For NGC 6530 CCD photometric time series in Johnson B & V filters were
obtained again with the 0.9m telescope (f /13.5) at the Cerro Tololo Interamerican
Observatory (CTIO), Chile. Between Aug. 1–7 and Aug. 9–15, 2002, NGC 6530
was observed using again the 2048 x 2046 SITe CCD chip, which provides a field
of view of 13’ × 13’ with a scale of 0.396”/px. In total, 80.16 hours of time-series
photometry could be acquired within 14 nights. As the cluster was slightly too large
to be observed on a single frame, two overlapping regions have been chosen for which
the observing time was split. Out of 3437 scientific frames (see Table 6.1), 1601 were
observed for field 1 and 1336 for field 2 (see Figure 6.9).
6.3. NGC 6530
49
Bias subtraction and the creation of the superflat-images was again performed
using the IRAF ared.quad package. Flat field correction was performed within the
reduction software Podex, which was also used to extract the photometric signal,
similar as for the data of IC 4996.
In total 194 stars have been identified in both fields where 43 stars lie in the overlapping region, 79 only in field 1 and 72 only in field 2 (see Figures 6.10 and 6.11).
Nightly means were subtracted from the light curves in order to correct for zero-point
changes and long-term irregular light variations caused by variable extinction due to
circumstellar material. Again the effect of color-dependent extinction was encountered and corrected using the same method as for the other two clusters. It turned
out that the effect is stronger in the B than in the V filter. Frequency analysis was
performed using Period98 and SigSpec. Only if frequencies appeared significant
using both methods they are believed to be intrinsic and not to be artefacts of the
reduction.
2000
1800
1600
1400
px
1200
1000
800
600
400
200
0
0
200
400
600
800 1000 1200 1400 1600 1800 2000
px
Figure 6.10: Schematic map of the observed field 1 of NGC 6530 (fov ∼ 13.5’ × 13.5’),
South is at the top and East is to the left) with all stars measured in Johnson B &
V, where 1 pixel (px) corresponds to 0.369 arcseconds. The area marked with red
lines corresponds to the overlapping region.
50
6. Observations and data reduction
2000
1800
1600
1400
px
1200
1000
800
600
400
200
0
0
200
400
600
800 1000 1200 1400 1600 1800 2000
px
Figure 6.11: Schematic map of the observed field 2 of NGC 6530 (fov ∼ 13.5’ × 13.5’),
South is at the top and East is to the left) with all stars measured in Johnson B &
V, where 1 pixel (px) corresponds to 0.369 arcseconds. The area marked with red
lines corresponds to the overlapping region.
cluster
NGC 6383
IC 4996
NGC 6530
nights
#
8 (14)
10 (14)
12 (14)
time base
[h]
53.25
62.37
80.16
science frames
#
2434
1990
3437
Table 6.1: Observing Statistics for NGC 6383, IC 4996 and NGC 6530. Numbers in
braces correspond to the total observing time granted.
Chapter 7
Observational Results
7.1
Pulsating PMS stars in NGC 6383
In NGC 6383 two new (NGC 6383 170 and 198) and one new suspected (NGC
6383 152) pulsating pre-main sequence stars were discovered. Their astrophysical
parameters are listed in Table 7.1, while the measured significant frequencies and
amplitudes for the two bona fide PMS pulsators are shown in Table 7.2.
7.1.1
NGC 6383 170
For NGC 6383 170 (V = 12.61 mag), Thé et al. (1985) found Hα in emission and
a large amount of excess radiation in the near-IR typical for HAEBE stars. These
findings are confirmed by the results of low resolution CCD spectroscopy (van den
Ancker et al. 2000) that clearly show an emission feature in the Hα line and a
spectral type of A5 IIIe for this star. Together with a confirmed membership to
NGC 6383, and a position above the ZAMS, star 170 is an ideal target to search for
PMS pulsation.
Five frequencies between 8.2 c/d and 19.5 c/d, spanning a period range between
1.2 and 2.9 hours have been found to be significant with amplitudes from 16 mmag
down to 7.6 mmag (Figure 7.2). A simultaneous multi-sine fit to the data (solid line
in Figure 7.1) was performed. Comparison of the fit with the shape of the observed
light curve clearly illustrates that there is strong indication for additional frequencies
buried in the noise. The five frequencies and Johnson B & V amplitudes are listed
in Table 7.2.
7.1.2
NGC 6383 198
Thé (1965) found for NGC 6383 198 - his star number 55 - a V magnitude of
12.60 mag and (B − V ) = 0.36 mag. No spectral classification is available in the
literature. But from its position in the HR-diagram it is most likely to be a cluster
member falling into the region of the classical instability strip, which made it an
ideal candidate to search for pulsation.
51
52
7. Observational Results
Figure 7.1: Differential light curves of NGC 6383 170; top: V filter, bottom: B filter
(shifted for better visibility). The solid line represents the multi-sine fit with the
five significant frequencies.
Only one frequency of 19.024 c/d, corresponding to a period of ∼ 1.26 hours,
with amplitudes of 26.4 mmag in B and 20.8 mmag in V , is significant in both filters
(see Table 7.2 and Figure 7.4). The comparison of the corresponding sine-fit with
the shape of the light curve (Figure 7.3) indicates multi-periodicity and the presence
of additional frequencies that were not discovered in these observations.
7.1. Pulsating PMS stars in NGC 6383
53
Figure 7.2: Amplitude spectra of NGC 6383 170 in V (top panel) and B (bottom
panel) filters; the identified frequencies are marked with arrows.
54
7. Observational Results
Figure 7.3: Differential light curves of NGC 6383 198; top: V filter, bottom: B filter
(shifted for better visibility). The solid line represents a sine-fit with a frequency of
19.024 c/d.
star
WEBDA
170
198
152
27
55
54
Other No.
#
F4
T 55
T 54
V
[mag]
12.61
12.90
12.45
(B − V )
[mag]
0.60
0.36
0.70
(U − B)
[mag]
0.36
-
sp
A5 IIIe
B-A ?
Table 7.1: Astrophysical parameters of the two bona fide and one suspected PMS
pulsators in NGC 6383 (other numbers taken from F ... Fitzgerald et al. 1978, T ...
Thé 1965).
7.1. Pulsating PMS stars in NGC 6383
55
Figure 7.4: Amplitude spectra of NGC 6383 198 in V (top panel) and B (bottom
panel) filters; ‘f1’ marks the identified frequency.
56
7.1.3
7. Observational Results
NGC 6383 152
NGC 6383 152 has a (B−V ) = 0.57 mag and, hence, lies in the region of the classical
instability strip. But only one significant frequency of 2.55 c/d with a peak-to-peak
amplitude of ∼30 mmag appears in the B data. Unfortunately, the V filter data
are of poor quality, where the noise is so dominant that no peak in the amplitude
spectrum exceeds four times the noise level.
Although this star has been one of the primary targets in NGC 6383 to search
for pulsation, it remains inconclusive in our data and is flagged as candidate PMS
pulsator. Longer time series of better quality have to be obtained to unambiguously
decide on its variability.
7.1.4
Summary of PMS pulsators in NGC 6383
Two bona fide, δ Scuti-like pre-main sequence stars have been found in NGC 6383.
Their frequencies and amplitudes are given in Table 7.2. All errors were calculated
using equations 6.1 and 6.2 (see Chapter 6).
star
no
170
f1
f2
f3
f4
f5
f1
198
frequency
[c/d]
14.376(3)
19.436(4)
13.766(4)
8.295(5)
17.653(6)
19.024(2)
V amp.
[mmag]
12.5(8)
11.3(8)
9.8(8)
8.6(8)
7.6(8)
20.8(7)
B amp.
[mmag]
16.0(8)
14.9(8)
12.3(8)
11.1(8)
9.8(8)
26.4(7)
Table 7.2: Frequencies and amplitudes determined for the two PMS pulsators NGC
6383 170 and NGC 6383 198, where the errors in the last digits of the corresponding
quantities are given in parentheses.
7.2. Pulsating PMS stars in IC 4996
7.2
57
Pulsating PMS stars in IC 4996
In IC 4996 two new (IC 4996 37 and 40) and one suspected (IC 4996 46) pulsating
pre-main sequence stars have been discovered as well as one new potential γ Doradus
type PMS pulsator (IC 4996 106). The astrophysical parameters of these stars are
given in Table 7.3 and their corresponding frequencies and amplitudes are listed in
Table 7.4.
star
WEBDA
37
40
106
46
201
171
1095
1085
Delgado No.
#
32
30
95
85
V
[mag]
15.30
15.03
15.71
15.30
(B − V )
[mag]
0.80
0.75
0.90
0.72
(U − B)
[mag]
0.44
0.43
0.39
0.36
sp
A5
A4
-
Table 7.3: Astrophysical parameters of the PMS pulsators in IC 4996.
7.2.1
IC 4996 37
For IC 4996 37, Delgado et al. (1998) measured V = 15.30 mag, (B − V ) = 0.8 mag
and (U − B) = 0.44 mag using CCD photometry. This coincides well with the observations performed by Vansevicius et al. (1996) who give V = 15.207 mag and
(B − V ) = 0.784 mag, but also (V − R) = 0.428 mag and V − I = 0.921 mag. Delgado
et al. (1999) obtained long-slit spectra for 16 stars of IC 4996, among those also
star 37, which was found to have a spectral type of A5. Hence, it is located in the
region of the HR-diagram, where pulsation can be expected, and this makes it an
ideal target to search for pulsation.
The frequency analysis with Period98 yielded a single intrinsic significant frequency in both filters at 31.875 c/d corresponding to a period of 45 minutes. This
frequency was also detected using SigSpec with significances 1 of 14.7 in V and 5.9
in B. In the amplitude spectra of both filters the peak at the frequency of 31.875 c/d
can clearly be noticed (Figure 7.5). Several alias frequencies related to the one-day
alias appear to be significant in the amplitude spectra of both filters, but were omitted in the analysis. As the data quality is rather poor the light curve of IC 4996 37
and the sine fit with this frequency are not really convincing and also the scatter is
very high in the B filter data. But for consistency reasons the light curves are also
shown as well (Figure 7.7). For better visibility, also the phase plots of the data in
both filters are given (Figure 7.6).
1
For the definition of SigSpec significances see Chapter 6.
58
7. Observational Results
Figure 7.5: Amplitude spectra of IC 4996 37 in V (top panel) and B (bottom panel)
filters; the identified frequencies are marked with arrows.
7.2. Pulsating PMS stars in IC 4996
59
Figure 7.6: Phase plots with a period of 45.18 minutes for IC 4996 37 in V (top)
and B (bottom) filters.
60
7. Observational Results
Figure 7.7: Differential light curve of IC 4996 37 in V (top) and B filters (bottom,
shifted for better visibility). The solid lines show the sine-fit with a frequency of
31.875 c/d.
7.2. Pulsating PMS stars in IC 4996
7.2.2
61
IC 4996 40
For IC 4996 40, Delgado et al. (1998) give V = 15.03 mag, (B − V ) = 0.75 mag and
(U − B) = 0.43 mag. A spectral type of A4 was determined by Delgado et al. (1999),
which places the star inside the instability region in the HR-diagram.
The frequency analysis with Period98 resulted in the detection of a single significant frequency in both filters at 33.569 c/d corresponding to a period of 43 minutes.
The same frequency was also detected using SigSpec with significances 2 of 33.1 in
V and 22.2 in B. The amplitude spectra clearly show the frequency at 33.569 c/d
(Figure 7.10). Again a number of frequencies related to the one-day alias appear in
the amplitude spectra and have been omitted in the analysis.
As the data quality is rather poor the sine-fit to the light curve of IC 4996 40 is
not really convincing and the scatter in the B light curve is much higher than for
V . However, the phase plots (Figure 7.8) visualize the variability in a better way.
For completeness, also the light curve with the sine fit is plotted (Figure 7.9).
Figure 7.8: Phase plots with a period of 42.89 minutes for IC 4996 40 in V (top)
and B (bottom) filters.
2
For the definition of SigSpec significances see Chapter 6.
62
7. Observational Results
Figure 7.9: Differential light curve of IC 4996 40 in V (top) and B filters (bottom,
shifted for better visibility). The solid lines show the sine-fit with a frequency of
33.569 c/d.
7.2. Pulsating PMS stars in IC 4996
63
Figure 7.10: Amplitude spectra of IC 4996 40 in V (top panel) and B (bottom
panel) filters; the identified frequencies are marked with arrows.
7.2.3
IC 4996 106
For IC 4996 106, Delgado et al. (1998) find V = 15.71 mag, (B − V ) = 0.90 mag and
(U − B) = 0.39 mag. No information on the spectral type of the star is available in
the literature, but from its position in the HR-diagram it is most likely a cluster
member lying right at the red edge of the instability region.
A frequency of 2.74 c/d appears to be significant in both filters corresponding
to a period of 8.76 hours (see Figure 7.12). Such a long period is rather difficult
to explain as δ Scuti-like pulsation. For post- and main sequence stars in this
64
7. Observational Results
part of the HR-diagram the classical instability strip overlaps with the γ Doradus
instability region. As the mechanism responsible for γ Doradus pulsation is believed
to be strongly related to convection and as convection is quite strong in such young
stars, IC 4996 106 is suspected to be the first PMS γ Doradus pulsator. To illustrate
this periodicity, the phase plots in both filters are given (Figure 7.11). The peak
at 2.74 c/d is also clearly visible in the corresponding amplitude spectra of the star
(Figure 7.12).
Figure 7.11: Phase plots with a period of 8.76 hours for IC 4996 106 in V (top) and
B (bottom) filters.
7.2.4
IC 4996 46
With (B − V ) = 0.72 mag and V = 15.3 mag (Delgado et al. 1998) IC 4996 46 is
also located in the region of the classical instability strip in the HR-diagram. The
frequency analysis resulted in the detection of different periods in V and B filters:
two periods of approximately 6.0 and 6.5 hours with amplitudes of 5.4 and 3.8 mmag
in V and two periods of 4.8 and 3.7 hours with amplitudes of 4.2 mmag each in B.
Hence, the star is considered as suspected PMS pulsator. Additional, longer time
series of better quality are needed to decide on the variability.
7.2. Pulsating PMS stars in IC 4996
7.2.5
65
Summary of PMS pulsators in IC 4996
Two bona fide, δ Scuti-like and one suspected γ Doradus-type pre-main sequence
stars have been found in IC 4996. Their frequencies and amplitudes are given in
Table 7.4. All errors were calculated using equations 6.1 and 6.2 (see Chapter 6).
star
no
37
40
106
f1
f1
f1
frequency
[d−1 ]
31.875(9)
33.569(7)
2.746(7)
V amp.
[mmag]
4.6(5)
7.6(5)
10.2(8)
B amp.
[mmag]
5.1(9)
8.5(7)
17.7(9)
Table 7.4: Frequencies and amplitudes determined for the three PMS pulsators in
IC 4996, where the errors in the last digits of the corresponding quantities are given
in parentheses.
66
7. Observational Results
Figure 7.12: Amplitude spectra of IC 4996 106 in V (top panel) and B (bottom
panel) filters; the identified frequencies are marked with arrows.
7.3. Pulsating PMS stars in NGC 6530
7.3
67
Pulsating PMS stars in NGC 6530
In NGC 6530 six PMS pulsators were discovered with classical δ Scuti type frequencies (NGC 6530 5, 82, 85, 263, 278 and 281), where for a seventh star pulsation
can only be suspected (NGC 6383 288). One star (NGC 6530 265) shows somewhat
longer periods indicating its possible PMS γ Doradus nature.
Table 7.5 shows an overview of all new pulsating PMS stars in NGC 6530, where
V , (B − V ) and (U − B) are taken from the WEBDA database and the spectral
types were derived by van den Ancker et al. (1997).
star
WEBDA
5
82
85
263
265
278
281
288
159
78
53
57
161
38
13
28
Sung No.
#
798
678
549
567
550
455
315
411
V
[mag]
13.59
13.97
13.07
13.67
13.75
12.17
13.35
13.23
(B − V )
[mag]
0.43
0.61
0.65
0.63
0.58
0.53
0.45
0.41
(U − B)
[mag]
0.26
0.30
0.35
0.35
0.14
0.36
0.27
0.35
sp
A1 III
A0/A5
-
Table 7.5: Astrophysical parameters of the bona fide and suspected PMS pulsators
in NGC 6530.
7.3.1
NGC 6530 5
NGC 6530 5 is situated in the region where the two observed fields overlapped each
other (see Chapter 6). Unfortunately it was the only of the 44 stars located in this
part of the CCD found to be a pulsating PMS star. For star 5 (WEBDA #159, Kilambi #159, Sung #798) no information about its spectral type or cluster membership is available in the literature, but U BV CCD photometry gives V = 13.59 mag,
(B − V ) = 0.43 mag and (U − B) = 0.26 mag (Sung et al. 2000). Hence, from its location in the cluster HR-diagram it is most likely to be a pre-main sequence member
situated in the region of the instability strip.
Two significant frequencies at 46.596 c/d (i.e. a period of 31 minutes) with
amplitudes of 1.4 mmag in V and 1.8 mmag in B and at 53.417 c/d (i.e. a period of
27 minutes) with amplitudes of around 1 mmag in both filters have been found and
are clearly visible in the amplitude spectra (see Figure 7.13). As the amplitudes are
rather low and the frequencies are rather high, the variation is not as prominently
visible from the star’s light curve. Figure 7.14 shows the V (top) and B (bottom)
light curves for the first seven nights, and Figure 7.15 for the second seven nights.
The gaps in the data sets - especially during the first week - are due to bad weather
conditions during the observations, when no data had been acquired.
68
7. Observational Results
Figure 7.13: Amplitude spectra of NGC 6530 5 in V (top panel) and B (bottom
panel) filters; ‘f1’ and ‘f2’ mark the identified frequencies.
7.3. Pulsating PMS stars in NGC 6530
69
Figure 7.14: Differential light curve of NGC 6530 5 obtained during nights 1 to 7
overplotted with a two frequency sine fit (solid line); top: V filter, bottom: B filter
(shifted for better visibility).
70
7. Observational Results
Figure 7.15: Differential light curve of NGC 6530 5 obtained during nights 8 to 14
overplotted with a two frequency sine fit (solid line); top: V filter, bottom: B filter
(shifted for better visibility).
7.3. Pulsating PMS stars in NGC 6530
7.3.2
71
NGC 6530 82
NGC 6530 82 (WEBDA #78, Kilambi #78, Sung #678) is located in field 1 of the
observations (see Chapter 6) and again no spectral type or membership information
is available in the literature. U BV CCD photometry (Sung et al. 2000) yields
V = 13.97 mag, (B − V ) = 0.61 mag and (U − B) = 0.30 mag, which places the star
inside the instability strip in the HR-diagram and makes it a likely cluster member.
In the frequency analysis three frequencies between 24.829 c/d and 38.531 c/d
with amplitudes between 1.7 mmag and 2.8 mmag have been found to be significant
(see Table 7.6). A simultaneous multi-sine fit with these three frequencies represents
the shape of the light curve well (Figure 7.16). The frequencies are also clearly visible
in the according amplitude spectra of the star in both filters (Figure 7.17). More
frequencies with even lower amplitudes might be buried in the noise, but to decide
on this longer and additional time-series observations would be needed.
Figure 7.16: Differential light curve of NGC 6530 82 overplotted with a three frequency sine fit (solid line); top: V filter, bottom: B filter (shifted for better visibility).
72
7. Observational Results
Figure 7.17: Amplitude spectra of NGC 6530 82 in B (bottom panel) and V (top
panel) filters; the identified frequencies are marked with arrows.
7.3. Pulsating PMS stars in NGC 6530
7.3.3
73
NGC 6530 85
NGC 6530 85 (WEBDA #53, Kilambi #53, Sung #549, van Altena #173) was also
situated in field 1 of the observations (see Chapter 6) and was studied frequently by
different authors. Kilambi (1977) obtained V = 12.96 mag, (B − V ) = 0.62 mag and
(U − B) = 0.30 mag for it and reported its cluster membership. These values match
the U BV CCD measurements of Sung et al. (2000), who found V = 13.07 mag,
(B − V ) = 0.65 mag and (U − B) = 0.35 mag, quite well.
The proper motion study by van Altena & Jones (1972) yielded a membership
probability of 78% for star 85. They also reported an E(B − V ) = 0.35 mag and
absolute magnitude MV = +0.81 mag. Chini & Neckel (1981) determined a spectral
type of A0 for NGC 6530 85. Together with the star’s position in the HR-diagram
placing it in the region, where instability can be expected, it is most probably a
pulsating, pre-main sequence member of the cluster.
The frequency analysis using Period98 and SigSpec allowed to identify simultaneous pulsation with five frequencies between 10.585 c/d and 31.148 c/d and amplitudes in the range of 39.1 mmag to 1.8 mmag (Table 7.6). The shape of the light
curve demonstrates the multi-periodic nature of the star beautifully (Figure 7.18).
The amplitude spectra (Figure 7.19) show the location of the frequencies and give
rise to the possibility that additional frequencies may exist and could not be resolved
with these observations.
Figure 7.18: Differential light curve of NGC 6530 85 overplotted with a five frequency
multi-sine fit (solid line); top: V filter, bottom: B filter (shifted for better visibility).
74
7. Observational Results
Figure 7.19: Amplitude spectra of NGC 6530 85 in B (bottom panel) and V (top
panel) filters; the identified frequencies are marked with arrows.
7.3. Pulsating PMS stars in NGC 6530
7.3.4
75
NGC 6530 263
NGC 6530 263 is one of the stars located in field 2 of the observations (see Chapter 6).
U BV CCD photometry (Sung et al. 2000) yielded V = 13.67 mag, (B−V ) = 0.63 mag
and (U − B) = 0.42 mag for star 263 (WEBDA #57, Kilambi #57, Sung #567,
van Altena #178), which coincides well with earlier reported measurements (e.g.
from Kilambi 1977). No information on the star’s spectral type is available in the
literature.
NGC 6530 263 was also contained in the proper motion study by van Altena &
Jones (1972), but it belonged to the 135 stars which have not been included into
their “primary” absolute parameter solution. They indicate that the star might
only be considered a probable cluster member. So, no final conclusion about the
membership of star 263 to NGC 6530 can be drawn from their investigation. From
the star’s position in the HR-diagram it was interpreted as a most likely member of
the cluster located in the region of the instability strip.
Indeed, the frequency analysis of both filters showed a single significant frequency
at 19.223 c/d, corresponding to a period of 1.25 hours, with an amplitude of 8.3 mmag
in B and 7.1 mmag in V . This typical δ Scuti like pulsation frequency clearly shows
up in the according amplitude spectra (Figure 7.21). The sine fit reproduces the
shape of the light curve of NGC 6530 263 reasonably well (Figure 7.20), but there
might be additional frequencies with lower amplitudes buried in the noise.
Figure 7.20: Differential light curve of NGC 6530 263 overplotted with a single
frequency fit (solid line); top: V filter, bottom: B filter (shifted for better visibility).
76
7. Observational Results
Figure 7.21: Amplitude spectra of NGC 6530 263 in B (bottom panel) and V (top
panel) filter; the identified frequencies are marked with arrows.
7.3. Pulsating PMS stars in NGC 6530
7.3.5
77
NGC 6530 265
NGC 6530 265 (WEBDA #161, Kilambi #161, Sung #550) is also located on field
2 of the observations (see Chapter 6). U BV CCD photometry (Sung et al. 2000)
gave V = 13.75 mag, (B − V ) = 0.58 mag and (U − B) = 0.14 mag. As no spectral
classification is available for this star, its probable cluster membership was estimated
from its location in the CMD and HR-diagram of the cluster.
The light curve in both filters (Figure 7.22) shows variability on a longer time
scale than typical for δ Scuti stars. The frequency analysis (the amplitude spectra
are shown in Figure 7.24), both with Period98 and SigSpec, showed two significant
peaks in the low frequency domain at 2.821 c/d (i.e. a period of 8.5 hours) and at
2.117 c/d (i.e. a period of 11.3 hours) with amplitudes around 4 mmag (see Table
7.6). Such long periods are rather difficult to explain with δ Scuti type pulsation;
they are also not correlated with instrumental or alias frequencies. Taking into account that the star is located near the area in the HR-diagram, where the γ Doradus
instability strip overlaps with the δ Scuti domain, that convection is believed to be
responsible for γ Doradus pulsations and that convection is very active in pre-main
sequence stars, NGC 6530 265 was considered as second suspected PMS γ Doradus
star.
Figure 7.22: Differential light curve of NGC 6530 265 overplotted with a twofrequency fit (solid line); top: V filter, bottom: B filter (shifted for better visibility).
78
7. Observational Results
The phase plots shown (Figure 7.23) show a sort of ‘lack of points’ around phases
of 0.5 in the plot of the residuals with f2 (lower panels). As the second period of
11.3 hours is close to half a day and the observations always have a gap due to the
day-night-cycle, each night the star was observed at nearly the same phases. At its
brightness minimum only few data points could be acquired, which can be also seen
regarding the light curve of star 265 (Figure 7.22).
Figure 7.23: Phase plots of NGC 6530 265; left: V filter, right: B filter; top panels
show the original data with frequency f1 of 2.821 c/d; bottom panels show residuals
to f1 with frequency f2 of 2.117 c/d.
7.3. Pulsating PMS stars in NGC 6530
79
Figure 7.24: Amplitude spectra of NGC 6530 265 in V (top panel) and B (bottom
panel) filter; the identified frequencies are marked with arrows
80
7.3.6
7. Observational Results
NGC 6530 278
NGC 6530 278 (WEBDA #38, Kilambi #38, Sung #455, van Altena #157) has also
been one of the stars observed on field 2 (see Chapter 6). Kilambi (1977) obtained
V = 12.16 mag, (B − V ) = 0.44 mag and (U − B) = 0.39 mag for this star, reported its
cluster membership and found indication for variability. The more recent values for
V = 12.17 mag, (B − V ) = 0.53 mag and (U − B) = 0.36 mag from CCD photometry
(Sung et al. 2000) confirm these values.
The proper motion study by van Altena & Jones (1972) gave a membership probability of 68% for NGC 6530 278. They also report a value for E(B − V ) = 0.35 mag
and absolute magnitude MV = -0.12 mag for star 278. Chini & Neckel (1981) used
their U BV and Hβ observations to derive a spectral type of A0 for this star .
Figure 7.25: Differential light curve of NGC 6530 278 overplotted with a nine frequency multi-sine fit (solid line); top: V filter, bottom B filter (shifted for better
visibility).
The light curve of NGC 6530 278 (Figure 7.25) clearly illustrates its multiperiodicity. The formal solution of the detailed frequency analysis yielded nine
significant frequencies in both filters between 4.0 c/d and 15.6 c/d with amplitudes
ranging from 13.8 mmag down to 2.6 mmag (see Table 7.6). Regarding the amplitude
spectra (Figures 7.26 and 7.27) it becomes evident that specific caution was necessary
to identify and prewhiten the frequencies one by one and working parallel with the
data of both filters. Due to several aliases related to 1 c/d an excess of peaks in the
frequency range between 0 and 15 c/d exists. At least three more frequencies lying
7.3. Pulsating PMS stars in NGC 6530
81
between 20 c/d and 30 c/d (marked with arrows in the lowest right panels in Figures
7.26 and 7.27) are showing up in both filters, but have not been significant in the
analysis. Whether all nine frequencies are indeed intrinsic has to be investigated
using additional, longer photometric time-series observations, but the pulsational
variability of NGC 6530 278 is evident.
82
7. Observational Results
Figure 7.26: Amplitude spectrum of NGC 6530 278 in the V filter; the identified
frequencies are marked with arrows.
7.3. Pulsating PMS stars in NGC 6530
83
Figure 7.27: Amplitude spectrum of NGC 6530 278 in the B filter; the identified
frequencies are marked with arrows.
84
7.3.7
7. Observational Results
NGC 6530 281
NGC 6530 281 (WEBDA #13, Kilambi #13, Sung # 315) lies in field 2 of the
observations (see Chapter 6). U BV CCD photometry (Sung et al. 2000) give
V = 13.35 mag, (B − V ) = 0.45 mag and (U − B) = 0.27 mag for this star, which coincides well with values derived earlier by different authors. No spectral classification
is available, but Kilambi (1977) reported that the star is most probably a cluster
member. Its location in the cluster CMD and HR-diagram confirms this suggestion.
Pulsation with seven periods simultaneously was found as formal solution in
the detailed frequency analysis. The light curve of the star demonstrates its multiperiodicity (Figure 7.28) and shows the beating effect caused by close frequencies.
The frequencies lie between 30.6 c/d and 43.4 c/d corresponding to periods of 47
to 33 minutes (Figures 7.30 and 7.29). The according amplitudes lie in the range
between 5.1 mmag and 1.4 mmag (see Table 7.6). Similar to NGC 6530 278 the
identification and prewhitening of intrinsic frequencies was done with special care
due to the numerous peaks caused by aliases in the frequency range between 35 c/d
and 45 c/d.
Figure 7.28: Differential light curve of NGC 6530 281 overplotted with a seven
frequency multi-sine fit; top: V filter, bottom: B filter (shifted for better visibility).
7.3. Pulsating PMS stars in NGC 6530
7.3.8
85
NGC 6530 288
NGC 6530 288 was also observed in field 2 of the observations (see Chapter 6). With
V = 13.23 mag and (B − V ) = 0.41 mag NGC 6530 288 (WEBDA #28, Kilambi #28,
Sung #411) falls into the region of the classical instability strip in the HR-diagram.
Hence, it was also one of the prime candidates to search for pulsation. In the V
filter data a frequency of 17.996 c/d corresponding to a period of 1.33 hours with
an amplitude of 0.8 mmag is significant. The same frequency can be found in the B
filter data, but it is not above 4·S/N and has a SigSpec significance 3 of only 5.04 in
the analysis using SigSpec. To decide on the star’s possible variability longer time
series are needed, especially as the amplitude of the suspected pulsation is very low.
3
For the definition of SigSpec significances see Chapter 6.
86
7. Observational Results
Figure 7.29: Amplitude spectrum of NGC 6530 281 in the V filter; the identified
frequencies are marked with arrows.
7.3. Pulsating PMS stars in NGC 6530
87
Figure 7.30: Amplitude spectrum of NGC 6530 281 in the B filter; the identified
frequencies are marked with arrows.
88
7.3.9
7. Observational Results
Summary of PMS pulsators in NGC 6530
Six bona fide, δ Scuti-like and one suspected γ Doradus-type pre-main sequence
stars have been found in NGC 6530. Their frequencies and amplitudes are given in
Table 7.6. All errors were calculated using equations 6.1 and 6.2 (see Chapter 6).
star
no
5
f1
f2
f1
f2
f3
f1
f2
f3
f4
f5
f1
f1
f2
f1
f2
f3
f4
f5
f6
f7
f8
f9
f1
f2
f3
f4
f5
f6
f7
82
85
263
265
278
281
frequency
[c/d]
46.596(9)
53.417(9)
38.531(6)
34.671(7)
24.829(9)
15.579(1)
12.700(1)
15.531(2)
10.585(4)
31.148(8)
19.223(2)
2.821(4)
2.117(5)
7.200(2)
12.121(2)
13.218(2)
4.178(2)
9.488(3)
6.013(4)
11.984(3)
15.684(5)
13.896(6)
43.418(4)
40.017(4)
37.457(6)
41.702(9)
40.367(8)
30.691(9)
38.245(7)
V amp.
[mmag]
1.4(3)
1.0(3)
2.4(3)
2.2(3)
1.8(3)
30.2(3)
17.1(3)
8.2(3)
3.5(3)
1.8(3)
7.1(2)
4.5(2)
3.5(2)
6.6(2)
9.4(2)
9.9(2)
6.2(2)
5.0(2)
3.6(2)
4.9(2)
2.9(2)
2.6(2)
4.2(2)
3.9(2)
2.4(2)
1.7(2)
1.9(2)
1.4(2)
2.1(2)
B amp.
[mmag]
1.8(3)
1.1(3)
2.8(3)
2.4(3)
1.7(3)
39.1(3)
23.0(3)
11.4(3)
4.7(3)
2.0(3)
8.3(3)
3.8(2)
3.9(2)
9.4(3)
12.4(3)
13.8(3)
8.0(3)
6.6(3)
5.0(3)
6.7(3)
3.7(3)
3.5(3)
4.7(2)
5.1(2)
2.7(2)
1.5(2)
2.2(2)
1.5(2)
2.1(2)
Table 7.6: Frequencies and amplitudes determined for all PMS pulsators in NGC
6530, where the errors in the last digits of the corresponding quantities are given in
parentheses.
7.4. Other variables
7.4
89
Other variables
Several stars located in the fields of the three clusters that do not fall within the
region of the classical instability strip display variability on very different time scales.
Some other objects are variable, but might not be members of the corresponding
clusters, or their variability remains inconclusive in the data. For completeness, the
most interesting different types of other variable stars are discussed below in detail,
while all definitive and suspected variables are listed in the corresponding tables for
each cluster (Tables 7.8, 7.9 and 7.10). The location of all detected variable and
suspected variable stars (including the pulsating PMS stars) in the sky is shown in
Figures 7.35, 7.40, 7.45 and 7.46.
7.4.1
Variable stars in NGC 6383
NGC 6383 15
V = 10.03 mag and (B − V ) = 0.34 mag together with its position in the HRdiagram indicate that NGC 6383 15 is not a member of the cluster. Two frequencies,
corresponding to periods of 1.645 and 1.414 hours, were found to be significant (see
Table 7.7). As it is most likely a foreground object, it seems to be a classical δ Scuti
star (see Fig. 7.31). Being as bright as 10th magnitude, the star was saturated on
the CCD for most of the time except for some hours during four nights. This is the
reason why its light curve is much shorter than for the pulsating PMS stars. The
errors were computed using equation 6.1 (see Chapter 6).
star
no
15
f1
f2
frequency
[d−1 ]
14.587
16.972
V amp.
[mmag]
8.5(3)
4.0(3)
B amp.
[mmag]
8.4(3)
6.2(2)
Table 7.7: Frequencies and amplitudes determined for NGC 6383 15, which most
probably is a foreground star; the errors in the last digits of the corresponding
quantities are given in parentheses.
NGC 6383 25
No colours and magnitudes were available from the literature for NGC 6383 25. Our
transformation yields V = 16.77 mag and (B − V ) = 1.58 mag. A frequency of
1.198 c/d, i.e. a period of 20.04 hours, leads to a phase plot shown in Figure 7.32.
Data without subtraction of nightly means were used for the analysis.
(B − V ) = 1.58 mag corresponds to a spectral type of M5 and is associated
to M/M¯ = 0.21, and R/R¯ = 0.27 according to Schmidt-Kaler (1965). Assuming
a rotation period of 20.04 hours, the equatorial rotational velocity is 16.37 km s−1 ,
90
7. Observational Results
Figure 7.31: Differential light curve of NGC 6383 15: top: V filter, bottom: B filter
(shifted for better visibility).
which can be calculated as given below:
vrot = 2 π
R R¯
2 π · 0.27 · 696000
=
= 16.37 kms−1
R¯ Prot
20.04 · 3600
(7.1)
This velocity would be in agreement with a rotating, active and weak-lined T Tauri
star.
NGC 6383 64
No information about this star was found in the literature. According to our observations, the star has only 16.72 mag in V and (B − V ) = 1.63 mag. A single
frequency of ∼ 2.499 c/d (corresponding to a period of 9.605 hours) with a peak-topeak amplitude of approximately 20 mmag is found to be significant in both B and
V light curves and leads to the phase plot shown in Figure 7.33.
NGC 6383 71
No astrophysical information was available from the literature for NGC 6383 71.
Our calculations give V = 15.37 mag and (B − V ) = 1.35 mag. If the star belongs to
the cluster, it is an early K type star according to Schmidt-Kaler (1965).
Two frequencies of 2.759 and 2.240 c/d, i.e. periods of 8.688 and 10.714 hours,
respectively, were detected (Figure 7.34). Variability on this time scale at such a
7.4. Other variables
91
Figure 7.32: Phase plot of NGC 6383 25: top: B filter, bottom: V filter.
high (B − V ) cannot be explained assuming cluster membership. Permitting NGC
6383 71 to be more distant than the cluster itself, interstellar reddening may shift
its position in the HR-diagram into the SPB, or β Cephei domain. A clear decision
can only be drawn from spectroscopy.
92
7. Observational Results
-0.2
B
mag
-0.1
0
0.1
0.2
0
0.2
0.4
0.6
0.8
1
0.6
0.8
1
phase
-0.2
V
mag
-0.1
0
0.1
0.2
0
0.2
0.4
phase
Figure 7.33: Phase plot of NGC 6383 64: top: B filter, bottom: V filter.
Figure 7.34: Differential light curve of NGC 6383 71: top: V filter, bottom: B filter
(shifted for better visibility). The solid line represents a multi-sine fit to the data.
7.4. Other variables
93
Summary of all other variable stars in NGC 6383
star
#
15
25
64
66
71
84
91
93
98
106
108
111
116
122
136
164
167
220
221
239
258
ref
T 47
T 28
F 10
F 20
F 11
F6
F8
T 52
T5
F3
T 77
-
V
mag
10.08
16.77
16.72
12.59
15.37
15.98
15.30
11.42
15.10
13.77
15.61
12.82
15.13
16.44
12.33
11.31
10.30
16.67
16.54
12.50
14.56
(B − V )
mag
0.34
1.58
1.63
0.33
1.35
1.31
0.90
0.17
1.06
0.52
1.45
0.32
1.31
1.42
0.72
0.01
0.29
2.03
1.19
0.85
0.69
sp
B8
A6
-
var. in filter
B/V /B&V
B&V
B&V
B&V
B
B&V
B
B
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B
B&V
B
V
B&V
B&V
B&V
remarks
foreground δ Scuti star
T Tauri star
probably not a cluster member
inconclusive in V
probably not a cluster member
inconclusive in V
inconclusive in V
different periods in B & V
T Tauri candidate
known IR excess
P∼7.26 hours, inconclusive
P∼5.65 hours, inconclusive
T Tauri candidate
unresolvable
inconclusive in V
P∼2 days, amplitude ∼ 30 mmag
saturated in V
P∼2.87 hours in V
irregular variable, T Tauri candidate
probably not a cluster member
different periods in B & V
Table 7.8: Other variables and suspected variables in the field of NGC 6383: ‘star’
denotes our star number and ‘ref’ the cross reference with the literature (according
to F ... Fitzgerald et al. 1978, T ... Thé 1965), the spectral types (sp) are taken
from the literature.
94
7. Observational Results
2000
258
1800
239
221
1600
220
1400
164
px
1200
198
167
170
152
111
106
136
1000
800
122
116
71
600
98 91
93
108
84
66
64
400
15
200
25
0
0
200
400
600
800 1000 1200 1400 1600 1800 2000
px
Figure 7.35: Schematic map of the observed field of NGC 6383 (fov ∼ 13.5’× 13.5’),
South is at the top and East is to the left) with all stars measured in Johnson B &
V in pixels, where 1 pixel corresponds to 0.396 arcseconds. The numbers correspond
to the identified new variable or suspected variable stars discussed in the text.
7.4. Other variables
7.4.2
95
Variable stars in IC 4996
IC 4996 67
IC 4996 67 has V = 12.82 mag and (B − V ) = 0.52 mag (Delgado et al. 1998), but
no spectral classification exists for this star. The light curves in both filters look
variable with relatively high amplitudes. The frequency analyses using Period98
and SigSpec yielded two significant frequencies, that seem to be intrinsic (Figure
7.36). The first frequency is at 1.928 c/d (i.e. a period of 12.4 hours) with amplitudes
of 10.3 mmag in V and 10.5 mmag in B, the second at 2.088 c/d (i.e. a period of 11.5
hours) with amplitudes of 9.7 mmag in V and 10.2 mmag in B. Permitting the star
to be less distant than the cluster, it could be a foreground γ Doradus type star.
Further studies are needed to decide on the nature of the star’s variability, where a
reliable spectral classification would be most important.
Figure 7.36: Differential light curve of IC 4996 67 in V (top) and B filter (bottom,
shifted for better visibility). The solid line corresponds to the multi-sine fit with the
two significant frequencies.
IC 4996 80
IC 4996 80 is located at the blue edge of the instability region in the HR-diagram.
V = 14.04 mag, (B − V ) = 0.58 mag and (U − B) = 0.15 mag was found by Delgado
et al. (1998), but again no spectral classification is available. In the data sets of
both filters one significant frequency at 3.578 c/d, i.e. a period of ∼6.7 hours, with
96
7. Observational Results
an amplitude of 2.0 mmag in V and 2.9 mmag in B was detected (Figure 7.37).
If star 80 is slightly hotter than indicated from its color, it could fall into the
region of the SPB instability domain in the HR-diagram.
Figure 7.37: Phase plots for IC 4996 80 in V (top) and B (bottom) filters with a
frequency of 3.578 c/d.
IC 4996 86
IC 4996 86 is again bluer than the blue edge of the classical instability region and
has V = 13.87 mag, B − V = 0.47 mag and U − B = 0.08 mag measured by Delgado
et al. (1998). Two frequencies at 3.294 c/d (i.e. a period of 7.3 hours) and 3.380 c/d
(i.e. a period of 7.1 hours) indicate a possible β Cephei- or SPB-type variable.
IC 4996 95
For IC 4996 95 no magnitudes and colors are available from the literature. Our
calibration yields V = 13.72 mag and B − V = 0.61 mag, which makes the star hotter
than the blue edge of the classical instability strip. The two detected significant
frequencies at 2.807 c/d (with 5.4 mmag amplitude in V and B) and 2.942 c/d (with
6.9 mmag amplitude in V and 5.1 mmag in B) lie very close to each other and produce
a beat in the light curve, Figure 7.38 shows the ‘fish-shaped’ light curve in the V
7.4. Other variables
97
filter with a simultaneous multi-sine fit (solid line). The corresponding periods are
8.6 and 8.2 hours, respectively. It also seems that the star is not a cluster member,
and the origin of its variation can only be suspected. It may be due to rotation or
caused by β Cephei- or SPB-type pulsations.
IC 4996 95
V filter
0.02
0.01
0
-0.01
-0.02
2520
2522
2524
2526
2528
2530
Figure 7.38: Differential light curve of IC 4996 95 in V with a multi-sine fit (solid
line) of the two significant frequencies.
98
7. Observational Results
IC 4996 104
With V = 14.03 mag, (B − V ) = 0.64 mag and (U − B) = 0.02 mag the star is also
lying near the blue edge of the classical instability region in the HR-diagram. Two
frequencies at 2.209 c/d (i.e. a period of 10.9 hours) with an amplitude of 7.5 mmag
in V and 10.4 mmag in B, and at 1.786 c/d (i.e. a period of 13.4 hours) with an
amplitude of 6.6 mmag in V and 7.7 mmag in B were found to be significant and a
simultaneous multi-sine fit to the data reproduces the shape of the light curve well
(Figure 7.39). As there is no spectral class available for this star, the nature of its
variation can only be suspected. The star could be a slowly pulsating B star or its
variation is caused by rotation.
Figure 7.39: Differential light curves for IC 4996 104 in V (top) and B (bottom,
shifted for better visibility) filters; the solid line corresponds to a multi-sine fit with
the two significant frequencies.
7.4. Other variables
99
Summary of all other variable stars in IC 4996
star
#
10
18
19
38
42
51
67
74
80
81
86
89
95
104
108
ref
D 73
D 54
D 56
D 91
D 13
D 14
D 55
D 83
D 18
D 61
D 17
W 110
W 50
W 115
D 96
V
mag
14.62
12.36
12.93
15.63
13.34
13.66
12.82
15.25
14.04
13.60
13.87
13.81
13.72
14.03
15.71
(B − V )
mag
0.84
0.66
0.39
1.86
0.52
0.59
0.52
1.13
0.58
0.59
0.47
0.55
0.61
0.64
0.89
var. in filter
B/V /B&V
V
B&V
B&V
B
B&V
B
B&V
V
B&V
B&V
V
V
B&V
B&V
V
remarks
inconclusive
different periods in B & V
different periods in B & V
inconclusive
different periods in B & V
inconclusive
2 periods
period is not significant in B
P ∼ 6.708 h
inconclusive
2 periods, unclear
period is not significant in B
2 close periods
2 periods
inconclusive
Table 7.9: Other variables and suspected variables in the field of IC 4996: ‘star’
denotes our star number and ‘ref’ the cross reference with the literature (according
to D ... Delgado et al. 1998, W ... number given in the WEBDA database), the
spectral types (sp) are taken from the literature. Values printed italic were derived
from our calibration because no values were available in the literature.
100
7. Observational Results
1000
104
95
900
800
81
80
86
89
74
67
700
600
px
38
51
42
46
37
40
500
400
300
108
18
19
200
106
100
0
0
10
100 200 300 400 500 600 700 800 900 1000
px
Figure 7.40: Schematic map of the observed field of IC 4996 (fov ∼ 6’ × 6’) with
all stars measured in Johnson B & V in pixels, where 1 px corresponds to 0.33
arcseconds. Identifiers refer to newly discovered variable and suspected variable
stars discussed in the text.
7.4. Other variables
7.4.3
101
Variable stars in NGC 6530
NGC 6530 13
NGC 6530 13 (V = 12.65 mag) was located in the overlapping region of the CCDs
(see Chapter 6). It was classified as B8 star by van den Ancker et al. (1997). Van
Altena & Jones (1972) estimate its cluster membership probability only to 25%. The
frequency analysis yielded a single significant frequency at 3.432 c/d, corresponding
to a period of 7.0 hours, with amplitudes of 3.6 mmag in V and 4.7 mmag in B (see
Figure 7.41). The star could be a SPB-type variable, but additional observations
are needed to decide on the origin of its variability.
Figure 7.41: Amplitude spectra of NGC 6530 13 in V (top) and B (bottom) filter,
where the significant frequency is indicated with arrows.
102
7. Observational Results
NGC 6530 21
NGC 6530 21 was observed in the overlapping field of the CCDs (see Chapter 6).
With V = 14.17 mag and (B − V ) = 0.97 mag it is located redwards of the instability
region in the HR-diagram. No spectral type or membership information is available
for this star, but from its (B−V )0 = 0.62 mag (dereddened with E(B−V ) = 0.35 mag
adopted from Sung et al. 2000) it should be an early G spectral type with log Teff =
3.77. Fourier analysis with Period98 yields two significant frequencies in the data
sets of both filters, which are confirmed by SigSpec. A simultaneous multi-sine
fit to the data with the two significant frequencies detected with both methods,
at 0.277 c/d (i.e. a period of 3.613 days) with amplitudes of 8.7 mmag in V and
9.7 mmag in B and at 0.166 c/d (i.e. a period of 6.037 days) with amplitudes
of 2.8 mmag in V and 2.3 mmag in B reproduces the shape of the light curve well
(Figure 7.42). The reason for such a variability can only be suspected to be rotational
modulation due to a spotty surface of a T Tauri type star.
Figure 7.42: Part of the light curve of NGC 6530 21 in V (top) and B filter (bottom).
NGC 6530 55
NGC 6530 is situated blueward of the instability region and was classified as spectral
type B6 by van den Ancker et al. (1997). The shape of the light curve indicates
variability (see Figure 7.43). The frequency analyses using Period98 and SigSpec
detect a single significant frequency of 2.285 c/d corresponding to a period of 10.5
hours. The star could be a slowly pulsating B star (SPB); rotation may also cause
such a variability.
7.4. Other variables
103
Figure 7.43: Differential light curve of NGC 6530 55 in V (top) and B filter (bottom,
shifted for better visibility).
NGC 6530 73
NGC 6530 73 (V = 11.94 mag) was classified as spectral type B5e by van den Ancker
et al. (1997). From the shape of its light curve it is most likely a binary star: as
the eclipse is deeper in the V than in the B filter, the redder star seems to occult
its bluer companion (see Figure 7.44).
104
7. Observational Results
Figure 7.44: Light curve of the eclipsing binary NGC 6530 73 in V (top, red) and
B (bottom, blue) filter.
7.4. Other variables
105
Summary of all other variable stars in NGC 6530
star
#
13
21
42
55
70
73
81
83
88
93
96
109
117
119
248
252
256
274
285
305
ref
WEBDA
W 94
W 75
W 41
W 84
W 1681
W 114
W 79
W 1641
W 1504
W 1458
W 40
W 112
W 141
W 116
W 19
W 165
W 20
W 144
W 230
W 151
V
mag
12.65
14.17
12.53
11.87
15.08
11.94
14.59
13.77
14.26
14.61
13.45
14.84
11.60
11.30
10.76
14.61
14.10
11.20
14.09
12.19
(B − V )
mag
0.27
0.97
0.34
0.22
1.30
0.39
1.18
1.09
1.22
1.08
1.20
1.02
0.45
0.41
0.24
1.06
0.70
0.21
0.71
0.72
var. in filter
B/V /B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
B&V
remarks
late B type star
2 periods
inconclusive
1 (4) periods ?
binary ?
binary
binary ?
1 period
2 periods
different periods in B
1 period, inconclusive
different periods in B
different periods in B
1 period, inconclusive
suspected
1 period, inconclusive
suspected
2 periods
inconclusive
different periods in B
&V
&V
& V , binary?
&V
Table 7.10: Other variables and suspected variables in the field of NGC 6530: star
denotes our star number and ref the cross reference with numbers given in the
WEBDA database, V and (B − V ) values were taken from the literature.
106
7. Observational Results
2000
55
1800
1600
70
1400
73
85
81
83
px
1200
88
93
82
1000
96
800
13
109
600
21
400
5
119
200
0
117
0
200
42
400
600
800 1000 1200 1400 1600 1800 2000
px
Figure 7.45: Schematic map of the observed field 1 of NGC 6530 (fov ∼ 13.5’ × 13.5’),
South is at the top and East is to the left) with all stars measured in Johnson B &
V in pixels, where 1 px corresponds to 0.396 arcseconds. Identifiers refer to newly
discovered variable and suspected variable stars discussed in the text. The area
marked in red corresponds to the overlapping region.
7.4. Other variables
2000
107
13
1800
1600
248
21
5
256
1400
42
252
1200
px
263
1000
281
265
274
800
278
600
285
288
400
305
200
0
0
200
400
600
800 1000 1200 1400 1600 1800 2000
px
Figure 7.46: Schematic map of the observed field 2 of NGC 6530 (fov ∼ 13.5’ × 13.5’),
South is at the top and East is to the left) with all stars measured in Johnson B &
V in pixels, where 1 px corresponds to 0.396 arcseconds. Identifiers refer to newly
discovered variable and suspected variable stars discussed in the text. The area
marked in red corresponds to the overlapping region.
108
7. Observational Results
7.5
Summary of cluster properties
In total ten new bona fide pulsating pre-main sequence stars have been discovered in
the three young clusters NGC 6383, NGC 6530 and IC 4996, supplemented by two
potential PMS γ Doradus type stars and three candidate PMS pulsators (details see
Table 7.11). Their significance in connection with young clusters, their importance
for the study of young stellar objects in general and their pulsational properties are
discussed below.
cluster
NGC 6383
IC 4996
NGC 6530
Total
total
number
286
113
194
593
pre-main sequence
bona-fide puls. cand. γ Dor
2
1
2
1
1
6
1
1
10
3
2
other
variables
21
15
20
56
Table 7.11: Numbers of variable and suspected variable stars in NGC 6383, NGC
6530 and IC 4996.
7.5.1
NGC 6383
Using our own calibrated data for NGC 6383, it was possible to draw a dereddened,
observed HR-diagram - in the MV − (B − V )0 plane - of the cluster (Figure 7.47).
The values for the distance modulus of (V − MV ) = 11.9 ± 0.25 mag and reddening
of E(B − V ) = 0.33 ± 0.02 mag taken from Fitzgerald et al. (1978) were used to
deredden the V and (B − V ) measurements.
Of 15 cluster members that fall in the region of the classical instability strip (see
Figure 7.47), for only two, NGC 6383 170 and NGC 6383 198, pulsation could be
clearly detected, whereas for NGC 6383 152 variability can only be suspected. This
corresponds to ≤20% variable stars within the region of the classical instability strip
in NGC 6383. This amount is somewhat below the corresponding percentages for
other main sequence variables. For example, for λ Bootis stars in the instability
region ≥50% are claimed to show pulsation (Paunzen et al. 1997). Breger (2000)
reported that between 1/3 to 1/2 of the stars situated in the lower instability strip
show photometrically detectable light variations due to pulsation.
The panels in Figure 7.48 show again HR-diagrams for NGC 6383, where each
star is color coded with respect to the amplitude noise level (ampnoise ) based on its
time-domain point-to-point scatter (σpt2pt ). The amplitude noise level in the Fourier
domain is given as:
√
ampnoise = 2 · σpt2pt / N
where N denotes the number of data points per star.
(7.2)
7.5. Summary of cluster properties
109
Figure 7.47: Observational HR-diagram of all stars in the field of NGC 6383: bonafide PMS pulsators (filled blue circles), candidate PMS pulsator (open blue circle),
other variables (red triangles) and other suspected variables (black triangles); the
ZAMS values are taken from Schmidt-Kaler (solid line) and the borders of the classical instability strip (dashed lines) have been transformed into the MV − (B − V )0
plane.
It can be clearly seen that, relying on the 4 · S/N criterion (Breger et al. 1993,
Kuschnig et al. 1997), the lowest detectable amplitudes for the brightest of the
observed stars lie between 1.0 and 2.0 mmag in the Fourier domain. The expected
continuous increase of the measured amplitude noise towards fainter stars is visible
in both filters: in V the bins of the same amplitude noise are parallel to the X
axis (top panel in Figure 7.48), while in B the ‘lines’ of same amplitude noise are
somewhat inclined (bottom panel in Figure 7.48). The reason for this is that in the
MV − (B − V )0 plane, the amplitude noise for the B filter is plotted and a star
with MB = 2.0 mag can either have a MV = 1.5 mag and (B − V )0 = 0.5 mag or its
MV = 1.0 mag and (B − V )0 = 1.0 mag.
110
7. Observational Results
Figure 7.48: HR-diagram of all measured stars in the field of NGC 6383 with the
expected amplitude noise in each star used for color coding; top: V filter, bottom:
B filter; the values for the ZAMS (solid line) are taken from Schmidt-Kaler (1965).
7.5. Summary of cluster properties
7.5.2
111
IC 4996
38 stars lie within the region of the classical instability strip in IC 4996, hence have
been prime candidates to search for PMS pulsation. Three stars, IC 4996 37, IC
4996 40 and IC 4996 106, had large enough amplitudes to detect pulsation unambiguously, while pulsation can only be suspected for IC 4996 46. From the location
of the PMS pulsators in the HR-diagram they are most likely to be members of the
cluster. A definite decision on their cluster membership can only be drawn after a
detailed proper motion or radial velocity study. However, the number of discovered
pulsating pre-main sequence stars corresponds to ∼10%. As the observations have
been performed in service mode and the integration time per frame was not long
enough for stars in the magnitude range between 14 < V < 16 mag, additional PMS
pulsators may be discovered in a follow-up observing run, because their amplitudes
might be too small to be detectable within these data.
Of the total 113 stars analysed in the field of the cluster, 16 stars have been found
to be variable including the PMS pulsators, which means only ∼15%. The reason
for such a low percentage lies in the not satisfactory quality of the measurements:
Especially towards the fainter, and hence redder stars, the noise level increased
dramatically and made it impossible to detect variability. Figure 7.49 shows the
location of all detected variables and suspected variable stars in IC 4996 in the
HR-diagram, where the lack of variable, red stars is clearly illustrated.
Figure 7.50 shows again HR-diagrams for IC 4996, where each star is color coded
with respect to the expected amplitude noise level (for detailed explanation see
above) in both filters (top: V and bottom: B filter). The brightest star, IC 4996 34,
has a V = 10.45 mag and was partly saturating the CCD chip, which is the reason
why it has one of the highest values for the amplitude noise.
7.5.3
NGC 6530
In NGC 6530 seven pulsating pre-main sequence stars could be discovered in total,
where six show δ Scuti like pulsation and one is a suspected PMS γ Doradus type
star. Of those pulsators one star is located in the overlapping region (see Chapter 6),
two in field 1 and four in field 2, which shows that the decision to split the observing
time for two different fields was successful. Taking into account that 25 stars have
been primary candidates to search for pulsation, as they are located in the region
of the classical instability strip, the seven detected stars correspond to 28%. This
fraction of close to 1/3 of the observed stars in the region of interest is higher than
for the other analysed clusters. It is also similar to the values of 1/3 to 1/2 pulsating
stars in the lower instability strip reported by Breger (2000) for the classical, postand main sequence δ Scuti stars.
Among the 194 observed stars in the fields of NGC 6530, at least 12 other stars
were found to be variable.
Figure 7.51 shows the location of all detected variable and suspected variable
stars in NGC 6530 in the HR-diagram.
112
7. Observational Results
Figure 7.49: Observational HR-diagram of all stars in the field of IC 4996: bona-fide
PMS pulsators (filled blue circles), candidate PMS pulsator (open blue circle), other
variables (red triangles) and other suspected variables (black triangles); the ZAMS
values are taken from Schmidt-Kaler (solid line) and the borders of the classical
instability strip (dashed lines) have been transformed into the MV − (B − V )0 plane.
Like for the other two clusters, the expected amplitude noise level based on the
point-to-point scatter in the time domain has been calculated and Figure 7.52 shows
the HR-diagrams, where each star is color coded according to its amplitude noise
bin (top: V , bottom: B filter). Please note that again some blue stars are too
bright and were partly saturated on the CCD chip. Hence, their number of data
points is somewhat lower and their accuracy is worse than for their slightly fainter
counterparts.
From their location in the HR-diagrams of NGC 6530 it becomes evident that
some stars in the field of the cluster might not be members, but are fore- or background field stars.
7.5. Summary of cluster properties
113
Figure 7.50: HR-diagram of all measured stars in the field of IC 4996 with the
expected amplitude noise in each star used for color coding; top: V filter, bottom:
B filter; the values for the ZAMS (solid line) are taken from Schmidt-Kaler (1965).
114
7. Observational Results
Figure 7.51: Observational HR-diagram of all stars in the field of NGC 6530: bonafide PMS pulsators (filled blue circles), candidate PMS pulsator (open blue circle),
other variables (red triangles) and other suspected variables (black triangles); the
ZAMS values are taken from Schmidt-Kaler (solid line) and the borders of the classical instability strip (dashed lines) have been transformed into the MV − (B − V )0
plane.
7.5. Summary of cluster properties
115
Figure 7.52: HR-diagram of all measured stars in the field of NGC 6530 with the
expected amplitude noise in each star used for color coding; top: V filter, bottom:
B filter; the values for the ZAMS (solid line) are taken from Schmidt-Kaler (1965).
Chapter 8
Modelling pulsation
As the main objective of this work was to discover new members of the group of
pulsating pre-main sequence stars in young open clusters, an asteroseismic investigation can only be based on limited information and considered as a first estimate.
Additional, longer, photometric time-series observations with multi-site campaigns
or from space are needed to characterize the stars asteroseismologically in greater
detail.
8.1
Pulsation constants
Using the empirical calibration by Reed (1998), which is described below, it is possible to derive Mbol , log Teff and log L/L¯ from (B − V )0 and Mv for all the newly
discovered pulsating pre-main sequence stars. For stars with log Teff ≤ 3.961 (i.e.
Teff ≤ 9140 K) the relation
(B − V )0 = −3.648 (log Teff ) + 14.551
(8.1)
can be used to determine log Teff . The bolometric magnitude, Mbol , is further given
as:
Mbol = Mv + BC(T ),
(8.2)
where the bolometric correction, BC(T ) can be determined using
BC(T ) = −8.499 (log Teff − 4)4 + 13.421 (log Teff − 4)3 −
−8.131 (log Teff − 4)2 − 3.901 (log Teff − 4) − 0.438
(8.3)
To compute the pulsation constant, Q, additional calculations of the mass ratios
and log g values have been necessary. This was done using the following relations
(e.g. Voigt 1991):
log(M/M¯ ) = 0.59 − 0.13 Mbol
(8.4)
4
4
g/g¯ = (M/M¯ ) · (Teff
/Teff,¯
)/(L/L¯ ),
(8.5)
where the following solar values have been adopted: Teff,¯ = 5780 K and g¯ =
2.7 · 104 cms−1 . Q was calculated using equation 3.3 (see Chapter 3). The radial
116
8.2. Pulsation models
117
Q values are 0.033 d for the fundamental mode, 0.025 d for the first, 0.020 d for the
second and 0.017 d for the third overtones; values smaller than 0.017 d are identified
as higher overtones. Regarding the resulting Q values given in Table 8.1 it becomes
evident that most of the periods detected correspond to higher overtone modes.
Only the first period of NGC 6530 278 (P = 0.139 d) and possibly the fourth period
of NGC 6530 85 (P = 0.094 d) seem to be radial fundamental modes. The Q values for
IC 4996 106 and NGC 6530 265 (marked with asterisks) are much larger indicating
that these two objects are most likely no PMS δ Scuti type stars. Indeed, they are
suspected to be the young counterparts of the classical γ Doradus type stars, as
mentioned before.
8.2
Pulsation models
The majority of the classical δ Scuti stars pulsate with a large number of nonradial acoustic modes simultaneously. Similar behaviour is expected to occur in
the PMS δ Scuti-type pulsators. As non-radial pulsation models for PMS stars are
currently under development, it was only possible to calculate linear, non adiabatic
pulsation for radial modes of PMS models for all stars found in this study. Of
course, for most of the stars the radial models cannot reproduce all frequencies
simultaneously, but indicate coexistence of radial and nonradial modes. The results
of the model calculations given below can therefore only be used as a first estimate of
the pulsational properties of the newly detected stars. Additional observations from
multi-site campaigns or using space telescopes as well as non-radial pulsation models
for PMS stars are needed for detailed asteroseismic analyses. With the results of
this work, the number of known PMS pulsators showing more than a single period
increased substantially, hence building an important incentive for the development
of the PMS pulsation theory combining radial and non-radial modes.
8.2.1
Discussion of observed frequencies
The calculations of the radial pulsation models for NGC 6383 170 resulted in
three possible solutions, but none reproduces all five frequencies simultaneously.
The radial model fitting the observed frequencies best, has a stellar mass of 2.5 M¯ ,
log L/L¯ = 1.68, Teff = 8100 K, and pulsation in third (f1) and fifth overtones (f2).
The solution seems to be optimal, because it is closest to the parameters derived
spectroscopically by van den Ancker et al. (2000): log L/L¯ = 1.69 and Teff = 8090 K
(filled symbol in Figure 8.1). According to the calculated Q values the first frequency
is probably the third and the second an even higher radial overtone mode, because
its Q is even smaller.
As no spectral classification is available for NGC 6383 198, the V and (B − V )
values were used to derive empirical ranges of luminosity and effective temperature
based on the transformations given by Kenyon & Hartmann (1995). These ranges
are indicated by the dashed box in Figure 8.1. Only for pulsation in the third radial
overtone (TO) the theoretical models have temperatures and luminosities close to
118
8. Modelling pulsation
star
M/M¯
log Teff
log L/L¯
log g
NGC 6383 170
2.96
3.849
1.53
3.716
Mbol
[mag]
0.915
NGC 6383 198
IC 4996 37
IC 4996 40
IC 4996 106
NGC 6530 5
3.04
2.26
2.50
1.99
2.31
3.869
3.892
3.904
3.865
3.923
1.57
1.17
1.31
1.00
1.20
3.774
4.136
4.088
4.136
4.239
0.828
1.820
1.474
2.238
1.747
NGC 6530 82
2.00
3.875
1.01
4.170
2.215
NGC 6530 85
2.62
3.864
1.37
3.887
1.326
NGC 6530 263
NGC 6530 265
2.19
2.15
3.869
3.883
1.13
1.11
4.069
4.140
1.921
1.984
NGC 6530 278
3.47
3.896
1.75
3.762
0.383
NGC 6530 281
2.47
3.918
1.29
4.156
1.519
period
[d]
0.070
0.051
0.073
0.121
0.057
0.053
0.031
0.030
0.364
0.022
0.019
0.026
0.029
0.040
0.064
0.079
0.064
0.095
0.032
0.052
0.354
0.472
0.139
0.083
0.076
0.239
0.105
0.166
0.083
0.064
0.072
0.023
0.025
0.027
0.024
0.025
0.033
0.026
Q
[d]
0.015
0.011
0.016
0.027
0.013
0.013
0.015
0.013
0.183
0.012
0.011
0.014
0.015
0.021
0.020
0.024
0.020
0.029
0.010
0.023
0.177
0.235
0.032
0.019
0.017
0.055
0.024
0.038
0.019
0.015
0.017
0.011
0.012
0.013
0.012
0.012
0.016
0.013
Table 8.1: Pulsation constant Q for the newly discovered PMS pulsators (see text
for detailed explanation).
8.2. Pulsation models
119
Figure
8.1:
Linear,
non-adiabatic
radial
pulsation
models
for
NGC 6383 170 and 198: solid lines are PMS evolutionary tracks for 1.5, 2.0,
2.5 and 3.0 M¯ (Palla & Stahler 1993); the open circle denotes the model for star
170 reproducing the observed frequencies best as a third (TO) and fifth radial
overtone (FiO), the filled triangle marks the position of the observed values for star
170 derived by van den Ancker et al. (2000); the asterisk (∗) shows the optimum
model, third radial overtone (TO) pulsation, for NGC 6383 198; the dashed box
indicates the empirical ranges of log Teff and log L/L¯ for star 198; the shaded area
is the theoretical PMS instability strip for the first three radial modes (Marconi &
Palla 1998).
the observations. Relying on the cluster membership of NGC 6383 198, it seems
reasonable that the star pulsates with a single frequency in the third radial overtone
having 2.0 M¯ , log L/L¯ = 1.3 and Teff = 7345 K. The Q value of 0.013 d indicates
an even higher overtone mode for this star.
The solutions for the linear, non adiabatic radial pulsation models for IC 4996
37 are shown in Figure 8.2 (green symbols). They correspond to monoperiodic
pulsation with the parameters listed in Table 8.2. The model with the best solution
in comparison with the observations is pulsation in the fourth radial overtone mode
with Teff = 8200 K and log L/L¯ = 1.262. The Q value of 0.016 d indicates at least
pulsation in the third overtone mode, maybe even higher.
The Q value of 0.013 d computed for IC 4996 40 suggests high overtone pulsation. This agrees well to the solutions of the model calculations shown in Figure
8.2 (magenta symbols). They correspond to monoperiodic radial pulsation with the
parameters listed in Table 8.2. The model with the best solution in comparison with
the observations is pulsation in the fifth radial overtone mode with Teff = 8460 K and
log L/L¯ = 1.39.
120
8. Modelling pulsation
star
IC 4996 37
IC 4996 40
overtone
mode
3rd
4th
5th
6th
3rd
4th
5th
6th
log L/L¯
1.090
1.262
1.377
1.481
1.125
1.240
1.390
1.485
Teff
[K]
8000
8200
8300
8400
8260
8360
8460
8560
M/M¯
1.75
1.90
2.00
2.10
1.80
1.90
2.00
2.10
Table 8.2: Parameters for monoperiodic radial pulsation models of IC 4996 37 and
40. The best fit model is marked in bold letters.
The best-fit radial pulsation model for NGC 6530 5 reproduces the observed
frequencies as the sixth (f1) and seventh (f2) radial overtones with a mass of 1.95 M¯ ,
log Teff = 3.923 and log L/L¯ = 1.20 (red square in Figure 8.3). As the Q values
evaluate to 0.013 d and 0.011 d respectively, they also suggest high overtone mode
pulsation.
According to the radial pulsation models, NGC 6530 82 seems to oscillate
simultaneously in the fourth (f2) and fifth (f1) radial overtone modes, while the
third frequency (f3) is consistent with the frequency of the second radial overtone.
These calculations are also supported by the computed Q values that indicate higher
than third overtones for f1 and f2, but second overtone oscillation for f3. Only the
position of the star in the HR-diagram contradicts second overtone pulsation with f3,
because the star is hotter than the second overtone blue edge. The radial pulsation
models also give 1.7 M¯ , log Teff = 3.875 and log L/L¯ = 1.01 for this star. Most
likely, not all three frequencies correspond to purely radial pulsation and permitting
non-radial oscillations for star 82 could change the picture.
NGC 6530 85 is reproduced as a PMS pulsator with 2.1 M¯ definitely oscillating in a mixture of radial and nonradial modes, because not all observed periods
can be explained by radial pulsation only. The linear non-adiabatic radial pulsation
models identify the fourth frequency (f4) as the first radial overtone, the second frequency (f2) as the second and the first frequency (f1) as the third overtone modes.
Furthermore the models also yielded log Teff = 3.862 and log L/L¯ = 1.4 for NGC
6530 85. The computed Q values differ slightly from the model predictions and
would suggest the fourth frequency (f4) to be the fundamental mode, the second
frequency (f2) the first harmonic, first (f1) and third frequency (f3) as second harmonic and the fifth (f5) as higher overtone. This slight discrepancy can be also
interpreted as a hint towards non-radial pulsation existing in star 85. The close pair
of the frequencies f1 and f3 could possibly be explained by fast rotation of the star
causing rotational splitting of the modes. Further studies, e.g. the determination of
v · sin i , are needed to draw a final conclusion about the origin of these two similar
8.2. Pulsation models
121
Figure 8.2: Linear, non-adiabatic radial pulsation models for IC 4996 37 and 40:
solid lines are PMS evolutionary tracks for 1.5, 2.0, 2.5, 3.0, 3.5 and 4.0 M¯ (Palla
& Stahler, 1993); the blue shaded area marks the region of the PMS instability strip
for the first three radial modes (Marconi & Palla 1998).
frequencies.
The Q value for the pre-main sequence star NGC 6530 263 pulsating monoperiodically is 0.023 d. Using the linear, non-adiabatic radial pulsation models its
single period was identified as the radial second overtone mode with the star having
1.8 M¯ , log Teff = 3.865 and log L/L¯ = 1.13.
NGC 6530 278 shows nine significant frequencies in the data and the radial
pulsation models derive a mass of 2.75 M¯ , log Teff = 3.896 and log L/L¯ = 1.75 and
oscillation in the second (f5), third (f2), fourth (f9) and fifth (f8) radial overtones.
Again, only a combination of radial and nonradial modes can explain the star’s
pulsational behaviour and the close pairs of frequencies that cannot be both radial.
Determination of the Q values yielded f1 as the fundamental, f5 as the first harmonic,
the close pair f2 and f7 as the second and f9 as third overtones. The other frequencies
are most likely nonradial modes. Clearly, also the Q values indicate a mixture of
radial and nonradial modes in NGC 6530 278.
For NGC 6530 281 only three frequencies could be identified as radial modes,
while the others should be of nonradial nature. With 1.9 M¯ , log Teff = 3.918 and
log L/L¯ = 1.29 from the models, the star’s frequencies could be reproduced as the
fourth (f6), the sixth (f2) and the seventh (f1) radial overtone modes. These findings
are supported by the Q values (see Table 8.1) being smaller than 0.017 d for all seven
observed periods.
It is quite evident that the higher overtone modes seem to have higher amplitudes
in pre-main sequence stars and hence are easier to detect. The reason for that may
be that the instability region for the fundamental modes is quite narrow with respect
122
8. Modelling pulsation
Figure 8.3: Linear, non-adiabatic radial pulsation models for the six δ Scuti like
PMS pulsators in NGC 6530: solid lines are PMS evolutionary tracks for 1.5, 2.0,
2.5, 3.0, 3.5 and 4.0 M¯ (Palla & Stahler, 1993); the blue lines mark the region of
the PMS instability strip for the first three radial modes (Marconi & Palla 1998).
to the first and second overtone ones according to the nonlinear convective models
(Marconi M., private communication). Fundamental pulsation is restricted to a
150 – 200 K wide zone close to the red boundary of the whole instability strip for
PMS stars. In this region the amplitudes are expected to decrease drastically due
to the increasing efficiency of convection which suppresses pulsation.
Our results confirm the suggestion that PMS δ Scuti-type stars pulsate with
a mixture of radial and non-radial modes, because for most of the stars it was
impossible to find purely radial pulsation models. Once non-radial pulsation models
for pre-main sequence stars are successfully developed, the pulsational properties of
the stars found in this work have to be reanalysed accordingly.
8.2.2
PMS γ Doradus type pulsators
For IC 4996 106 and NGC 6530 265 the linear non adiabatic, radial pulsation
models could not reproduce the observed frequencies at all. This evidence suggests
our speculation that pre-main sequence γ Doradus pulsation can be assumed for
the two stars, where the driving mechanism is different to the κ-mechanism in the
δ Scuti-like pulsating pre-main sequence stars. Convection may be responsible for
γ Doradus-type pulsation (see before), and hence γ Doradus-type pulsation is most
likely to occur in PMS stars as well. The computations of the Q values for these
8.2. Pulsation models
123
stars are quite interesting. Handler & Shobbrook (2002) have shown that the ‘classical’ γ Doradus-type stars show Q values larger than 0.23 d. However, in their
publication they also show a histogram where the γ Doradus stars can also pulsate
with frequencies yielding log Q values as small as -0.7 corresponding to Q = 0.186 d
(Figure 9 in their paper).
For NGC 6530 265, the second frequency has Q = 0.242 d, hence being an
indicator for γ Doradus type pulsation. The first frequency of this star results in a
Q value of 0.182 d. This may also be attributed to γ Doradus type pulsation taking
into account slight errors in the computation of the fundamental parameters for the
star and the maybe not accurate enough determination of the pulsation frequency.
For IC 4996 106 the Q value of its single period is 0.186 d which also can be
explained as γ Doradus type pulsation given the errors mentioned before.
The discovery of this new class of PMS γ Doradus type stars would be of special
importance for the study of stellar structure and evolution. As they are believed
to pulsate with g-modes originating in the stellar interior, a dense enough pulsation
frequency spectrum would enable a detailed asteroseismic analysis and allow to
probe the inner structure of stars still contracting towards the ZAMS. On the other
hand the existence of PMS γ Doradus stars could also help to resolve the driving
mechanism for their evolved counterparts, as it is currently only suggested to be
connected to convection. The only two stars showing γ Doradus like characteristics
have to be reobserved with a better time resolution for periods on the order of several
hours to a day and additional stars with similar pulsational properties have to be
found, before a final conclusion can be drawn.
Chapter 9
The empirical PMS instability
strip
9.1
All known pulsating PMS stars
Compared to the situation at the beginning of this work, the number of pulsating PMS stars has increased significantly and allows to probe the PMS instability
strip observationally. In the year 2000 eight pulsating pre-main sequence stars were
known, where four were members of the clusters NGC 2264 and NGC 6823 and four
were HAEBE type field stars. At that time, PMS stars have been thought to pulsate
purely radially with only one or two periods.
Until the end of this work (as of September 2005), the total number of pulsating
pre-main sequence stars has increased to 37, of which 25 stars are bona fide PMS
δ Scuti-like pulsators, two are potential PMS γ Doradus stars and ten remain pulsating PMS candidates. Within this work, eight bona fide δ Scuti PMS pulsators, the
two potential PMS γ Doradus stars, which define a new type of variability for PMS
stars, and three candidates for PMS pulsation could be discovered in three young
clusters. Very soon during this work it became clear that they are not pulsating
purely monoperiodically (e.g. NGC 6383 170 was one of the first stars with five
detected frequencies) and that the frequencies cannot be explained by radial pulsations only. A complete catalog of the known pulsating PMS stars and candidates
including an overview of their parameters is given in Table 9.1. It is now possible for
the first time to probe the instability strip for pre-main sequence stars empirically.
9.2
The new PMS instability strip
Regarding the HR-diagram (Figure 9.1) with the location of all known PMS pulsators it becomes evident that there seems to be a lack of stars toward the region
of the red edge of the classical instability strip. Whether this is only a selection
effect caused by poor number statistics or has some astrophysical reason, can only
be speculated.
124
V 589 Mon
V 588 Mon
NGC 6823 HP57
NGC 6823 BL50
NGC 6383 152
NGC 6383 170
NGC 6383 198
IC 4996 37
IC 4996 40
IC 4996 46
IC 4996 106
IC 348 H 254
NGC 6530 5
NGC 6530 82
NGC 6530 85
NGC 6530 263
NGC 6530 265
NGC 6530 278
NGC 6530 281
NGC 6530 288
V 351 Ori
V 346 Ori
UX Ori
IP Per
HR 5999
HD 35929
HD 142666
HD 104237
CQ Tau
BF Ori
HD 34282
V 1247Ori
β Pic
VV Ser
V 375 Lac
WW Vul
PX Vul
Name
DEC (2000.0)
[dd:mm:ss]
+09:42:04.1
+09:41:03.4
+23:16:37.8
+23:17:49.6
-32:35:30.9
-32:36:17.9
-32:37:24.0
+37:39:31.0
+37:39:32.8
+37:42:26.5
+37:42:26.5
+32:06:22.1
-24:18:03.5
-24:23:42.1
-24:24:55.7
-24:15:46.9
-24:14:05.2
-24:13:28.0
-24:15:02.6
-24:12:21.1
+00:08:40.4
+01:43:48.3
-03:47:14.3
+32:31:53.7
-39:06:18.3
-08:19:38.4
-22:01:40.0
-78:11:34.6
+24:44:54.0
-06:35:00.6
-09:48:35.4
-01:15:21.7
-51:03:59.5
+00:08:39.0
+40:40:05.0
+21:12:31.0
+23:53:49.0
F2 III
A7 III/IV
F0 Ve
A5 IIIp
A5
A4
F0/A8 III-IV
A1 III
A0/A5
A7 IIIe
A5 III
A3e
A7 V
A7 III/IVe
F0 IIIe
A8 Ve
A4 V
F2 IVe
A5 II-IIIevar
A0e
A5 III
A5 V
A2e
A7e
A3e
F0 Ve
sp
V
[mag]
10.32
9.73
14.60
14.50
12.34
12.60
12.83
15.21
15.03
15.30
15.71
10.60
13.59
13.97
13.07
13.67
13.75
12.17
13.35
13.23
8.90
10.10
9.60
10.40
6.98
8.20
8.81
6.60
10.70
10.30
9.85
9.82
3.86
11.50
12.94
10.51
11.67
3.850
3.900
3.857
3.860
3.884
3.908
3.870
3.914
3.927
3.900
3.865
3.854
3.923
3.875
3.864
3.869
3.883
3.896
3.918
3.929
3.870
3.886
3.940
3.889
3.845
3.857
3.880
3.930
3.830
3.940
3.857
3.910
3.850
3.860
-
log Teff
1.51
2.05
1.25
1.60
1.77
1.70
1.30
1.26
1.26
1.41
1.00
1.62
1.20
1.01
1.37
1.13
1.11
1.75
1.29
1.35
1.15
0.98
1.49
0.97
2.12
1.92
1.03
1.50
1.48
1.15
1.20
1.05
2.13
2.08
-
log L/L¯
Type
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
cluster member
HAEBE
HAEBE
HAEBE
HAEBE
HAEBE
PMS ?
HAEBE
HAEBE
HAEBE
HAEBE
HAEBE
PMS ?
PMS ?
HAEBE
HAEBE
HAEBE
HAEBE
Table 9.1: All known PMS pulsators and pulsating candidate stars.
RA (2000.0)
[hh:mm:ss]
06:39:28.5
06:39:05.9
19:43:06.8
19:43:09.1
17:34:55.1
17:34:37.0
17:34:48.0
20:16:22.0
20:16:30.0
20:16:43.9
20:16:29.4
03:44:31.2
18:04:42.3
18:04:30.8
18:04:20.7
18:04:21.8
18:04:20.8
18:04:13.9
18:04:00.2
18:04:09.9
05:44:18.8
05:24:42.8
05:04:30.0
03:40:47.0
16:08:34.3
05:27:42.8
15:56:40.2
12:00:05.1
05:35:58.0
05:37:13.3
05:16:00.5
05:38:05.3
05:47:17.1
18:28:49.0
22:34:40.9
19:25:58.7
19:26:40.3
puls. freq.
#
19
12
2
2
suspected
5
1
1
1
suspected
1
4
2
3
5
1
2
9
7
suspected
5
4
suspected
9
1
1
1
2 (3?)
1
1 (?)
2
1
2 (3?)
2 (3?)
2
1 (?)
1 (?)
9.2. The new PMS instability strip
125
126
9. The empirical PMS instability strip
One possible explanation would be that the amplitudes of pulsation towards
the red side of the PMS instability strip are not high enough to be detected. In
this region of the HR-diagram the influence of convection is quite high and is supposed to damp the pulsation amplitudes. This is also supported by the detection of
mostly higher overtones and fewer fundamental modes in pre-main sequence stars
compared to their post- and main sequence counterparts. Another possibility is that
the phenomena of pulsation among the PMS F-type stars mix with the sometimes
strong acitivity of T Tauri stars, hence being even more difficult to detect. A final
explanation can only be given after observing especially pre-main sequence stars
located near the red edge of the classical instability region and after studying their
properties in more detail.
Figure 9.1: HR-diagram with PMS evolutionary tracks (D’Antona & Mazzitelli
1994), the borders of the classical δ Scuti (black lines; taken from Breger & Pamyatnykh 1998) and the red edge of the PMS instability strip (blue dashed lines; taken
from Marconi & Palla 1998), all PMS pulsators: cluster (colored stars) and field stars
(black dots), possible PMS γ Doradus stars (black open circles) and PMS pulsator
candidates (open stars).
9.2. The new PMS instability strip
127
Figure 9.2 shows the newly discovered pulsating PMS stars and candidates (red
symbols) together with the PMS pulsators detected by other authors (black filled
circles) in the observational HR-diagram. A slightly different inclination of the
borders of the empirical PMS instability region compared to the classical δ Scuti
instability strip is suggested (blue dashed lines). The instability region for PMS
stars is drawn under the plausible assumption that it coincides with the classical
instability strip on the ZAMS.
Figure 9.2: Observational HR-diagram showing the borders of the classical δ Scuti
instability strip (black lines) taken from Breger & Pamyatnykh (1998), the ZAMS
(thick black line) taken from Schmidt-Kaler (1965), the bona fide, pulsating PMS
cluster members (red filled stars), the PMS candidate pulsators (red open stars) and
the potential PMS γ Doradus stars (red open circles) discovered in this work as well
as the other known PMS pulsators as of September 2005 (black dots). The blue
dashed line marks the observational borders of the instability region for pre-main
sequence stars.
It will be also of special interest to observe stars towards the cooler side of
the instability regions and search for more γ Doradus-type young stars. As the
128
9. The empirical PMS instability strip
mechanism driving the pulsation in the post- and main sequence counterparts is
still not clear yet, asteroseismic studies of confirmed PMS γ Doradus stars in the
future would help to understand the physical processes working in these stars. As
the pulsation periods - like for the classical γ Doradus stars - are close to one day,
such objects would be perfect candidates for observations with satellites. Except
for multi-site campaigns, ground based observations would always show significant
aliases at a frequency of 1 c/d due to the day-night cycle. The detection of this
potential new group of PMS pulsators was only possible because the photometric
investigation of the cluster was not limited to the temperature and period range
expected for δ Scuti-type pulsation only.
The comparison of the post- and main sequence γ Doradus stars with their
suspected PMS counterparts could help to better understand the driving mechanism
in γ Doradus stars in general. The location of the two potential PMS γ Doradus
type stars, NGC 6530 265 and IC 4996 106, in the observational HR-diagram in the
MV – (B − V )0 plane is shown in Figure 9.3. Obviously, IC 4996 106 falls directly
into the γ Doradus instability region, while NGC 6530 265 is slightly hotter than
the observed blue edge. As it is still not clear up to which temperatures γ Doradus
pulsations can be excited for the evolved stars, this is not surprising. Some classical
γ Doradus objects are hotter than the observed blue edge as well, because e.g. they
are members of a binary system (Handler 2005).
Figure 9.3: Location of the two potential PMS γ Doradus stars, NGC 6530 265
(marked as ‘265’) and IC 4996 (marked as ‘106’), in the HR-diagram with the region
of the classical (thin black lines) and the observational (dashed red line) γ Doradus
instability strips; the ZAMS (solid line) was taken from Schmidt-Kaler (1965).
Chapter 10
Conclusions
This work was dedicated to discover new members of the group of pulsating PMS
stars and to define the borders of the PMS instability strip observationally. More
than one third of the stars known until now (status 2005) have been found within
this study. The location of the theoretical red edge of the PMS instability strip is in
reasonable agreement with the observations, even the slightly different inclination
than for the empirical red edge of the classical instability strip is reproduced. Although the observations have been quite short for asteroseismic purposes, estimates
on the excited modes could also be given illustrating the need for longer ground- or
space-based observations for more detailed investigations.
Often not very much additional information is available for the stars, e.g. for
most of the objects spectral classification would urgently be needed. For the couple of
pulsating stars with known spectral types, mostly emission lines have been reported
in the literature and also very often an excess in the IR was observed. From stellar
evolution theory it is suspected that such young stars are fast rotators, but for nearly
all known PMS pulsators no v · sin i determination has been performed yet, making
it difficult to confirm this concept. All discovered pulsating young stars seem to lie
in the intermediate mass range between 1.5 and 3.5 M¯ . Additionally the detected
periodicities due to pulsation always lie on top of longer, sometimes irregular light
variations caused by the circumstellar material. All these features are characteristics
of the HAEBE stars indicating their pre-main sequence nature.
From the pulsational analysis of the 13 bona fide and suspected pulsating premain sequence stars discovered in this work, it is evident that the higher overtone
modes seem to have higher amplitudes in pre-main sequence stars and hence are easier to detect. This is confirmed by the asteroseismic investigations of other detected
PMS pulsators (e.g. Ripepi et al. 2005). At the moment it can only be speculated
about the reason for this behaviour. The instability region for the fundamental
modes is quite narrow with respect to the first and second overtone ones according
to nonlinear convective models (Marconi M., private communication). Pulsation in
the fundamental modes is restricted to a 150 – 200 K narrow zone close to the red
boundary of the theoretical PMS instability strip. In this region the amplitudes are
129
130
10. Conclusions
expected to decrease drastically due to the increasing efficiency of convection which
suppresses pulsation.
The discovery of a new class of PMS pulsators, suggested by the present study,
namely pre-main sequence γ Doradus stars, would be of special importance for the
study of stellar structure and evolution. As the classical γ Doradus stars are believed
to pulsate with g-modes originating in the stellar interior, a dense enough pulsation
frequency spectrum of a PMS counterpart would enable a detailed asteroseismic
analysis and allow to probe the inner structure of stars still contracting towards
the ZAMS. On the other hand the existence of PMS γ Doradus stars could also
help to resolve the driving mechanism for the post- and main sequence γ Doradus
stars, as it is currently only suggested to be connected to convection. The only two
PMS cluster members showing γ Doradus like characteristics have to be reobserved
with a better time resolution for periods on the order of several hours to a day and
additional stars with similar pulsational properties have to be found, before a final
conclusion can be drawn.
Appendix A
Photometric data
A.1
Stars in the field of NGC 6383
For all stars in the field of NGC 6383, V and (B −V ) values were computed using our
transformation (see in Chapter 6). Assuming cluster membership the according MV
and (B − V )0 values were obtained using E(B − V ) = 0.33 mag and (V − MV ) =
11.9 mag (Fitzgerald et al. 1978). Table A.1 lists the respective values together with
the star numbers assigned in this work (no), cross references with the literature
(ref ), if available, and the coordinates in pixels on the CCD (X, Y ).
Table A.1: Photometric data for all observed stars in the field of NGC 6383.
no
ref
X
Y
V
(B − V )
MV
#
[px]
[px]
[mag]
[mag]
[mag]
1
T 32
1199.72
18.23 13.431
0.713
1.531
2
664.68
20.66 16.263
1.290
4.363
3
T 13
1924.37
68.24 12.056
0.210
0.156
4
T 38
466.83
69.62 13.627
0.758
1.727
5
T 19
1573.26
74.82 11.645
1.303
-0.255
6
375.99
82.95 15.563
0.878
3.663
7
1164.79
86.15 16.818
1.327
4.918
8
817.81
87.60 16.852
1.209
4.952
9
1248.96
91.15 16.866
1.144
4.966
10
397.82 102.68 16.438
1.060
4.538
11
T 39
651.10 114.00 13.374
0.672
1.474
12 EV 281 1589.78 122.32 14.596
1.706
2.696
13
644.77 140.70 15.357
1.046
3.457
14
T 12
2019.93 143.29 13.038
0.792
1.138
15
T 47
164.02 178.78 10.030
0.339
-1.870
16
1887.08 181.27 17.003
1.250
5.103
cross references according to: T ... Thé (1965),
EV ... Lloyd Evans (1978), F ... Fitzgerald et al. (1978)
131
(B − V )0
[mag]
0.383
0.960
-0.120
0.428
0.973
0.548
0.997
0.879
0.814
0.730
0.342
1.376
0.716
0.462
0.009
0.920
132
A. Photometric data
Table A.1: Photometric data for NGC 6383 – continued
no
ref
X
Y
V
(B − V )
MV
#
[px]
[px]
[mag]
[mag]
[mag]
17
2039.42 184.29 16.689
1.225
4.789
18
T 31
1220.05 189.32 14.006
0.736
2.106
19
1812.81 190.70 16.166
1.133
4.266
20
T 40
628.59 193.29 13.264
0.605
1.364
21
1605.13 199.85 16.077
1.518
4.177
22
31.87 201.90 16.383
0.993
4.483
23
619.57 202.99 15.158
1.178
3.258
24
1994.76 208.01 16.794
1.219
4.894
25
1729.50 217.61 16.767
1.578
4.867
26
658.03 223.70 17.094
1.273
5.194
27
1049.01 225.60 16.861
1.186
4.961
28
517.64 236.29 17.093
2.194
5.193
29
T 41
727.62 259.08 10.913
1.192
-0.987
30
1664.79 261.29 15.050
0.843
3.150
31
T 30
1202.61 268.73 13.379
0.914
1.479
32
1097.97 275.27 16.731
1.346
4.831
33
364.16 280.00 16.741
1.164
4.841
34
1961.79 281.93 14.550
1.162
2.650
35
T 29
1548.39 286.17 13.655
0.818
1.755
36
115.06 288.75 16.779
1.049
4.879
37
T 46
304.32 316.07 12.675
0.759
0.775
38
1439.01 325.23 14.780
0.778
2.880
39
1764.32 338.96 16.081
1.105
4.181
40
836.25 351.06 16.542
1.428
4.642
42
788.92 381.86 16.293
1.240
4.393
43
102.83 390.66 16.382
1.196
4.482
44
1941.02 405.92 14.340
0.825
2.440
45
975.65 408.16 15.494
1.073
3.594
46
1231.15 439.34 17.018
1.304
5.118
47 EV 109 1335.71 439.39 16.672
1.114
4.772
48
1823.24 440.78 14.399
0.857
2.499
49
1559.90 441.50 16.519
1.588
4.619
50
593.61 454.80 16.441
1.152
4.541
51
1819.27 483.77 16.148
1.062
4.248
52
992.62 488.38 16.745
1.425
4.845
53
1796.40 489.68 16.211
1.112
4.311
54
668.14 520.89 16.908
1.558
5.008
55 EV 108 1357.83 522.08 15.690
1.204
3.790
56
954.27 530.73 15.653
1.136
3.753
57
T 45
307.36 535.32 13.276
2.410
1.376
58 EV 107 1492.03 543.14 14.841
1.002
2.941
59
T 96
1721.10 553.00 11.403
0.704
-0.497
60 EV 111 1155.64 570.06 15.681
0.915
3.781
61
1709.27 570.10 14.552
0.920
2.652
62 EV 113
323.40 591.39 15.427
1.248
3.527
cross references according to: T ... Thé (1965),
EV ... Lloyd Evans (1978), F ... Fitzgerald et al. (1978)
(B − V )0
[mag]
0.895
0.406
0.803
0.275
1.188
0.663
0.848
0.889
1.248
0.943
0.856
1.864
0.862
0.513
0.584
1.016
0.834
0.832
0.488
0.719
0.429
0.448
0.775
1.098
0.910
0.866
0.495
0.743
0.974
0.784
0.527
1.258
0.822
0.732
1.095
0.782
1.228
0.874
0.806
2.080
0.672
0.374
0.585
0.590
0.918
A.1. Stars in the field of NGC 6383
133
Table A.1: Photometric data for NGC 6383 – continued
no
ref
X
Y
V
(B − V )
MV
#
[px]
[px]
[mag]
[mag]
[mag]
63 EV 380
925.06 595.27 14.190
0.667
2.290
64
870.31 595.94 16.721
1.634
4.821
65
855.11 602.99 15.481
1.335
3.581
66
T 28
1468.18 605.70 12.532
0.163
0.632
67
923.89 611.24 15.374
1.115
3.474
68
233.11 621.07 15.116
0.982
3.216
69 EV 114
492.68 621.79 15.942
1.034
4.042
70 EV 110 1366.40 637.78 16.758
1.186
4.858
71
908.28 646.88 15.373
1.348
3.473
72
45.21 656.10 16.907
1.204
5.007
73
1686.12 662.89 15.174
0.926
3.274
74 EV 112 1051.58 670.23 17.131
1.262
5.231
75 EV 115
595.97 710.18 16.084
1.091
4.184
76
1595.67 714.97 14.665
0.891
2.765
77
T 44
234.60 736.71 13.560
0.706
1.660
78
T 43
412.39 762.81 13.103
0.445
1.203
79 EV 158 1603.79 765.28 13.797
0.691
1.897
80
1903.06 788.90 16.290
1.571
4.390
81
755.16 792.04 16.050
1.335
4.150
82
T 50
162.81 806.29 13.527
0.943
1.627
83
1380.35 830.01 15.077
1.358
3.177
84
1246.26 841.33 15.982
1.315
4.082
85
1549.76 852.24 16.491
1.283
4.591
86
T 17
1389.39 853.14 12.649
0.659
0.749
87
383.82 854.74 16.256
1.213
4.356
88
511.36 860.75 15.543
1.114
3.643
90
F 21
1003.22 867.06 12.000
0.773
0.100
91
F 10
1125.23 881.18 15.312
0.952
3.412
92
F2
1277.07 885.23 10.403
-0.017
-1.497
93
F 20
1173.62 888.08 11.473
0.094
-0.427
94
1325.99 890.34 15.870
1.013
3.970
95
983.26 892.09 16.770
1.742
4.870
96
348.81 893.05 16.332
1.133
4.432
97
996.66 894.20 15.576
0.955
3.676
98
F 11
1051.19 896.15 15.155
1.143
3.255
99
462.23 898.94 16.814
1.409
4.914
100
756.70 899.89 15.735
1.348
3.835
101
891.73 902.12 17.578
1.650
5.678
102
T7
1695.73 904.24 12.752
0.339
0.852
103
T 42
673.15 910.60 13.660
0.798
1.760
104
1826.27 911.37 16.469
1.250
4.569
105
707.38 911.91 14.043
1.142
2.143
106
F6
1144.36 914.89 13.833
0.598
1.933
107
752.96 920.26 14.324
0.852
2.424
108
1247.24 926.33 15.608
1.452
3.708
cross references according to: T ... Thé (1965),
EV ... Lloyd Evans (1978), F ... Fitzgerald et al. (1978)
(B − V )0
[mag]
0.337
1.304
1.005
-0.167
0.785
0.652
0.704
0.856
1.018
0.874
0.596
0.932
0.761
0.561
0.376
0.115
0.361
1.241
1.005
0.613
1.028
0.985
0.953
0.329
0.883
0.784
0.443
0.622
-0.347
-0.236
0.683
1.412
0.803
0.625
0.813
1.079
1.018
1.320
0.009
0.468
0.920
0.812
0.268
0.522
1.122
134
A. Photometric data
Table A.1: Photometric data for NGC 6383 – continued
no
ref
X
Y
V
(B − V )
MV
#
[px]
[px]
[mag]
[mag]
[mag]
109
F7
1122.63
931.87 12.662
0.255
0.762
110
970.63
933.44 16.568
1.441
4.668
111
F8
1130.81
944.62 12.900
0.333
1.000
112
1462.80
963.98 15.573
1.100
3.673
113
482.24
964.66 16.406
1.044
4.506
114 EV 132 1920.00
964.68 16.575
2.313
4.675
115
594.66
967.04 14.272
0.987
2.372
116
1055.88
969.96 15.128
1.306
3.228
117
559.77
970.03 14.200
0.848
2.300
118
F9
1112.65
971.08 10.885
0.051
-1.015
119
F 23
867.23
971.34 13.824
1.039
1.924
120
295.87
977.09 17.157
1.659
5.257
121
427.70
980.10 16.458
1.318
4.558
122
1054.97
984.20 16.442
1.416
4.542
123
313.07
984.63 14.516
1.048
2.616
124 EV 101
629.94
985.65 14.616
0.993
2.716
125
947.79
986.76 15.805
1.239
3.905
126
1110.64
994.16 15.072
1.298
3.172
127
938.11 1000.83 16.425
1.493
4.525
128
F5
1313.96 1007.28 12.915
0.474
1.015
129
49.92 1013.16 14.420
1.379
2.520
130
1170.83 1015.02 17.255
1.962
5.355
131
1377.98 1018.00 17.168
1.582
5.268
132
347.13 1028.38 16.958
1.574
5.058
133 EV 118
798.36 1041.11 15.004
0.921
3.104
134
1777.01 1046.14 16.344
1.157
4.444
136
T 52
168.39 1057.81 12.333
0.719
0.433
137
764.65 1060.37 15.063
1.368
3.163
138
T 53
373.26 1062.19 12.364
0.845
0.464
139
F 14
1163.20 1064.25
9.908
0.011
-1.992
140 EV 117
600.77 1068.02 16.079
3.436
4.179
141 EV 341
292.97 1072.18 14.870
0.673
2.970
142
755.79 1072.22 16.298
1.206
4.398
143
1142.33 1085.32 16.181
1.561
4.281
144
F 25
957.10 1104.88 12.546
0.193
0.646
145
755.94 1107.58 14.528
0.920
2.628
146
768.08 1109.97 14.881
1.161
2.981
147
F 24
867.63 1117.11 11.401
0.099
-0.499
148
214.36 1119.42 15.477
1.107
3.577
149
F 22
822.58 1123.93 12.344
0.586
0.444
150
1137.82 1125.75 15.797
1.118
3.897
151
F 18
1210.31 1127.16 13.414
0.895
1.514
152
T 54
642.20 1143.42 12.343
0.574
0.443
153
153
813.09 1160.87 15.900
1.043
4.000
154 EV 131
460.58 1172.03 17.089
1.410
5.189
cross references according to: T ... Thé (1965),
EV ... Lloyd Evans (1978), F ... Fitzgerald et al. (1978)
(B − V )0
[mag]
-0.075
1.111
0.003
0.770
0.714
1.983
0.657
0.976
0.518
-0.279
0.709
1.329
0.988
1.086
0.718
0.663
0.909
0.968
1.163
0.144
1.049
1.632
1.252
1.244
0.591
0.827
0.389
1.038
0.515
-0.319
3.106
0.343
0.876
1.231
-0.137
0.590
0.831
-0.231
0.777
0.256
0.788
0.565
0.244
0.713
1.080
A.1. Stars in the field of NGC 6383
135
Table A.1: Photometric data for NGC 6383 – continued
no
ref
X
Y
V
(B − V )
MV
#
[px]
[px]
[mag]
[mag]
[mag]
155
1894.82 1181.04 15.942
1.326
4.042
156
1723.72 1191.24 16.841
1.260
4.941
157
1829.80 1196.28 16.631
1.169
4.731
158
T 84
1921.94 1202.35 12.173
0.177
0.273
159
897.28 1208.26 16.722
1.459
4.822
160
1733.69 1211.13 15.075
1.104
3.175
161
1671.01 1211.58 14.758
0.802
2.858
162
972.84 1214.98 17.186
1.603
5.286
163
1067.84 1218.17 17.137
1.521
5.237
164
T5
499.27 1228.13 11.307
0.011
-0.593
165
766.29 1231.46 17.090
1.361
5.190
166 EV 105 1428.35 1232.13 15.391
1.282
3.491
167
F3
1300.22 1238.68 10.329
0.281
-1.571
168 EV 119
746.33 1248.59 14.920
0.952
3.020
169
454.11 1262.42 17.140
2.916
5.240
170
F4
1185.09 1262.77 12.900
0.702
1.000
171
811.75 1266.25 17.083
1.403
5.183
172
387.18 1267.98 16.583
1.102
4.683
173 EV 120
650.21 1270.25 15.923
1.085
4.023
174
354.41 1271.08 14.232
1.321
2.332
175 EV 106 1578.70 1285.38 15.113
0.979
3.213
176
T 58
312.02 1289.90 12.425
0.339
0.525
177 EV 121
578.36 1291.90 15.956
1.126
4.056
178
213.07 1291.97 16.642
1.322
4.742
179
T 83
1845.63 1296.59
9.581
0.020
-2.319
180
418.82 1306.86 16.830
1.286
4.930
181
316.03 1310.79 17.028
1.935
5.128
182 EV 104 1303.33 1319.82 15.217
1.000
3.317
183
414.92 1321.43 14.659
0.837
2.759
184
153.17 1340.29 15.929
1.199
4.029
185
1935.35 1355.92 15.803
1.011
3.903
186
1868.70 1357.27 15.862
1.005
3.962
187
992.20 1359.94 16.896
1.131
4.996
188
182.27 1360.00 16.524
1.882
4.624
189
136.86 1361.14 14.973
0.814
3.073
190 EV 103 1308.64 1367.22 16.275
1.087
4.375
191
1678.35 1369.76 15.107
0.876
3.207
192
986.65 1388.20 15.407
1.158
3.507
193
1634.02 1391.19 15.577
0.950
3.677
194
T4
2011.94 1391.19 13.425
1.053
1.525
195
T 82
1707.91 1404.01 13.222
0.963
1.322
196 EV 102 1360.30 1409.28 14.729
0.802
2.829
197 EV 514 1548.60 1412.10 14.088
0.633
2.188
198
T 55
849.95 1418.74 12.827
0.627
0.927
199
1237.07 1419.28 16.019
1.093
4.119
200
1797.97 1435.41 15.211
0.929
3.311
cross references according to: T ... Thé (1965),
EV ... Lloyd Evans (1978), F ... Fitzgerald et al. (1978)
(B − V )0
[mag]
0.996
0.930
0.839
-0.153
1.129
0.774
0.472
1.273
1.191
-0.319
1.031
0.952
-0.049
0.622
2.586
0.372
1.073
0.772
0.755
0.991
0.649
0.009
0.796
0.992
-0.310
0.956
1.605
0.670
0.507
0.869
0.681
0.675
0.801
1.552
0.484
0.757
0.546
0.828
0.620
0.723
0.633
0.472
0.303
0.297
0.763
0.599
136
A. Photometric data
Table A.1: Photometric data for NGC 6383 – continued
no
ref
X
Y
V
(B − V )
MV
#
[px]
[px]
[mag]
[mag]
[mag]
201
T 56
863.37 1439.33 13.828
0.893
1.928
202
492.36 1443.38 16.411
1.590
4.511
203
T 57
418.35 1454.55 10.690
0.190
-1.210
204 EV 127
904.69 1458.22 14.628
0.915
2.728
205
42.27 1462.05 14.396
2.072
2.496
206
663.53 1463.84 15.912
0.966
4.012
207
1765.15 1465.00 16.533
1.117
4.633
208
84.05 1474.30 16.125
1.286
4.225
209
223.36 1476.16 16.487
1.421
4.587
210
1455.78 1489.89 16.119
1.355
4.219
211 EV 126
794.32 1490.92 15.320
0.881
3.420
212
T 81
1630.73 1491.70 12.691
1.900
0.791
213 EV 128 1552.35 1516.75 15.825
1.331
3.925
214
1679.66 1526.65 15.995
0.996
4.095
215
906.80 1527.87 16.612
1.157
4.712
216
1479.14 1528.07 16.472
1.038
4.572
217
374.19 1529.61 16.395
1.522
4.495
218
495.27 1543.18 16.943
1.102
5.043
219
156.18 1550.33 14.952
1.051
3.052
220
1047.89 1558.93 16.671
2.027
4.771
221
154.93 1558.95 16.539
1.190
4.639
222
811.08 1559.72 16.820
1.531
4.920
223
48.92 1570.62 15.104
1.995
3.204
224
1102.75 1582.05 16.358
1.123
4.458
225
1865.21 1586.25 15.578
0.886
3.678
226
T 80
1494.10 1600.21 13.342
0.780
1.442
227
1853.22 1609.97 16.519
0.584
4.619
228
540.81 1629.21 15.321
0.923
3.421
229
1003.84 1630.04 16.503
1.555
4.603
230
485.73 1643.99 15.383
1.320
3.483
231
1356.33 1681.92 16.414
1.181
4.514
232
1167.04 1682.17 16.032
0.995
4.132
233
1149.45 1684.80 15.258
0.929
3.358
234
1855.72 1696.06 16.588
1.236
4.688
235
T 79
1558.33 1700.24 13.480
0.783
1.580
236
1437.65 1701.57 14.253
0.779
2.353
237
248.32 1702.52 17.885
1.871
5.985
238
T 62
82.90 1703.12 12.942
0.828
1.042
239
T 77
1876.77 1711.86 12.507
0.916
0.607
240
1652.13 1714.77 16.039
1.002
4.139
241
368.10 1720.81 15.411
1.304
3.511
242 EV 125
787.44 1721.21 15.718
1.054
3.818
243 EV 140
755.81 1732.73 14.914
1.529
3.014
244
345.34 1734.22 14.181
2.003
2.281
245
365.95 1750.00 15.843
1.045
3.943
246
319.94 1752.38 16.910
2.713
5.010
cross references according to: T ... Thé (1965),
EV ... Lloyd Evans (1978), F ... Fitzgerald et al. (1978)
(B − V )0
[mag]
0.563
1.260
-0.140
0.585
1.742
0.636
0.787
0.956
1.091
1.025
0.551
1.570
1.001
0.666
0.827
0.708
1.192
0.772
0.721
1.697
0.860
1.201
1.665
0.793
0.556
0.450
0.254
0.593
1.225
0.990
0.851
0.665
0.599
0.906
0.453
0.449
1.541
0.498
0.586
0.672
0.974
0.724
1.199
1.673
0.715
2.383
A.1. Stars in the field of NGC 6383
137
Table A.1: Photometric data for NGC 6383 – continued
no
ref
X
Y
V
(B − V )
MV
#
[px]
[px]
[mag]
[mag]
[mag]
247
901.91 1752.96 16.854
1.268
4.954
248
1639.58 1772.67 16.486
1.071
4.586
249 EV 122
558.13 1775.02 14.052
0.754
2.152
250
1255.14 1777.12 15.439
1.443
3.539
251
1080.34 1784.27 15.613
0.895
3.713
252 EV 123
665.71 1792.36 14.396
1.174
2.496
253 EV 124
711.27 1793.36 15.343
1.170
3.443
254
72.90 1797.44 15.741
1.221
3.841
255
409.24 1799.02 16.521
1.188
4.621
256
1186.89 1799.37 15.286
0.919
3.386
257
T 71
1253.82 1808.40 12.794
0.867
0.894
258
2016.97 1814.63 14.563
1.018
2.663
259
125.07 1833.07 16.374
1.533
4.474
260
846.68 1838.83 16.057
1.238
4.157
261
T 78
1713.86 1844.37 12.375
0.779
0.475
262
T 63
193.20 1848.21 12.487
0.671
0.587
263
915.88 1852.10 17.024
2.334
5.124
264
106.20 1859.70 15.449
0.997
3.549
265
1820.86 1863.28 15.305
1.167
3.405
266
T 64
290.91 1864.59 11.234
0.769
-0.666
267
T 10
574.28 1865.66 10.048
-0.016
-1.852
268
1769.05 1868.72 15.043
1.169
3.143
269
1606.88 1870.61 16.628
1.209
4.728
270 EV 133 1302.32 1875.69 16.390
2.501
4.490
271
783.40 1880.84 13.844
1.246
1.944
272
1376.96 1894.95 15.457
2.191
3.557
273
T 70
985.36 1912.30 13.074
0.502
1.174
274
2016.04 1921.81 14.643
0.912
2.743
275
T 65
443.72 1927.54 10.629
1.296
-1.271
276
1021.39 1928.14 15.421
0.823
3.521
277
807.00 1932.98 15.108
1.265
3.208
278
995.31 1948.07 16.408
2.347
4.508
279
1194.13 1957.44 15.576
1.264
3.676
280
1201.84 1975.30 14.674
0.858
2.774
281
T 69
876.93 1980.01 12.777
0.896
0.877
282
1948.21 1983.77 15.008
0.865
3.108
283
176.35 1987.67 15.584
1.110
3.684
284
716.69 1991.23 14.702
2.120
2.802
285 EV 316 1754.04 1992.89 14.539
0.739
2.639
286
35.14 2016.29 16.176
1.735
4.276
287
1224.00 2018.87 15.302
1.108
3.402
288
1162.16 2022.13 14.807
0.997
2.907
289
424.21 2022.17 16.411
1.878
4.511
cross references according to: T ... Thé (1965),
EV ... Lloyd Evans (1978), F ... Fitzgerald et al. (1978)
(B − V )0
[mag]
0.938
0.741
0.424
1.113
0.565
0.844
0.840
0.891
0.858
0.589
0.537
0.688
1.203
0.908
0.449
0.341
2.004
0.667
0.837
0.439
-0.346
0.839
0.879
2.171
0.916
1.861
0.172
0.582
0.966
0.493
0.935
2.017
0.934
0.528
0.566
0.535
0.780
1.790
0.409
1.405
0.778
0.667
1.548
138
A.2
A. Photometric data
Stars in the field of IC 4996
For all stars in the field of IC 4996, V and (B − V ) values were computed. Assuming cluster membership the according MV and (B − V )0 values were obtained
using E(B − V ) = 0.70 mag and (m − M ) = 13.28 mag (taken from the WEBDA
database). Table A.2 lists the respective values together with the star numbers assigned in this work (no), cross references with the literature (ref ), if available, and
the coordinates in pixels on the CCD (X, Y ).
Table A.2: Photometric data for all observed stars in IC 4996.
no
ref
X
Y
#
[px] [px]
1
D9
335 624
2
68 106
3
231
85
4
D 74
434 116
5
521
29
6
614
13
7
654
12
8
627
73
9
659
77
10
D 73
898 123
11
D 78
827 157
12
D 64
669 200
13
D 69
556 252
14
D 89
479 237
15
D 68
375 258
16
D 82
208 223
17
94 266
18
D 54
222 273
19
D 56
400 306
20
D 29
558 364
21
D 62
906 360
22
D 84
872 400
23 D 31p 841 444
24
D 88
743 372
25 D 39p 675 440
26
D 60
620 455
27
D 50
551 480
28
117 380
29
59 392
30
D 77
360 494
31
D 65
395 520
32 D 36p 409 491
33
D 57
574 519
cross references according
P ... Purgathofer (1964)
V
(B − V )
MV
(B − V )0
[mag]
[mag]
[mag]
[mag]
12.29
0.54
-0.16
-0.99
14.64
0.74
0.04
1.36
14.90
1.69
0.99
1.62
14.71
1.78
1.08
1.43
14.29
2.05
1.35
1.01
14.69
1.15
0.45
1.41
14.90
1.79
1.09
1.62
14.66
0.80
0.10
1.38
14.91
0.87
0.17
1.63
14.51
0.95
0.25
1.23
14.79
0.89
0.19
1.51
14.04
1.25
0.55
0.76
14.25
1.55
0.85
0.97
15.37
1.05
0.35
2.09
14.26
0.91
0.21
0.98
14.93
1.03
0.33
1.65
14.62
0.86
0.16
1.34
12.26
0.78
0.08
-1.02
12.78
0.53
-0.17
-0.50
14.83
0.82
0.12
1.55
13.91
0.79
0.09
0.63
15.21
1.01
0.31
1.93
15.10
0.88
0.18
1.82
15.32
0.97
0.27
2.04
15.68
1.16
0.46
2.40
13.31
0.69
-0.01
0.03
11.31
1.31
0.61
-1.97
15.04
1.00
0.30
1.76
15.24
1.02
0.32
1.96
14.67
1.07
0.37
1.39
14.01
0.65
-0.05
0.73
15.44
1.02
0.32
2.16
12.99
0.63
-0.07
-0.29
to: D ... Delgado et al. (1998) &
A.2. Stars in the field of IC 4996
139
Table A.2: Photometric data for IC 4996 – continued
no
ref
X
Y
V
(B − V )
#
[px]
[px] [mag]
[mag]
34
D 47
669 500 10.42
0.40
35
D 51
924 507 11.55
1.37
36
D 79
871 516 14.83
0.97
37
D 32
941 567 15.23
0.91
38
D 91
844 618 15.56
1.88
39
D 75
800 612 14.73
0.82
40
D 30
709 579 14.91
0.87
41
D 25
688 589 14.74
0.79
42
D 13
612 558 13.18
0.64
43
D 80
512 565 14.85
0.87
44
D 23
424 583 14.37
0.70
45
D 24
346 565 14.11
0.76
46
D 85
297 569 14.66
0.89
47
P 69
90 555 12.69
0.51
48
D 86
180 600 15.28
0.89
49
P 68
129 669 14.13
0.62
50
D 12
509 624 13.11
0.61
51
D 14
601 602 13.52
0.71
52
D 27
647 623 14.47
0.88
53
D 16
670 630 13.72
0.69
54
D 28
724 629 14.45
0.92
55
D7
732 673 11.78
0.61
56
D 70
534 673 14.33
0.69
57
D 10
663 676 12.57
0.64
58
D6
604 708 11.66
0.54
59
D 49
588 720 11.17
0.53
60
D 11
541 721 12.91
0.56
61
D 71
490 714 14.42
1.97
62 P 68a
82 722 13.00
1.60
63
D 59
203 745 13.19
0.56
64
151 763 15.16
1.07
65
D 66
397 779 14.27
0.80
66
D 67
486 762 14.25
0.99
67
D 55
581 784 12.71
0.66
68
D 63
647 792 14.09
0.68
69
646 791 14.09
0.68
70
D 15
616 804 13.63
0.66
71
D 93
714 792 15.68
1.03
72
D 19
739 795 14.13
0.72
73
D5
701 764 16.58
1.20
74
D 83
793 793 15.14
1.21
75
1017 765 14.97
0.98
cross references according to: D ... Delgado et
P ... Purgathofer (1964)
MV
(B − V )0
[mag]
[mag]
-0.30
-2.86
0.67
-1.73
0.27
1.55
0.21
1.95
1.18
2.28
0.12
1.45
0.17
1.63
0.09
1.46
-0.06
-0.1
0.17
1.57
0.00
1.09
0.06
0.83
0.19
1.38
-0.19
-0.59
0.19
2.00
-0.08
0.85
-0.09
-0.17
0.01
0.24
0.18
1.19
-0.01
0.44
0.22
1.17
-0.09
-1.50
-0.01
1.05
-0.06
-0.71
-0.16
-1.62
-0.17
-2.11
-0.14
-0.37
1.27
1.14
0.90
-0.28
-0.14
-0.09
0.37
1.88
0.10
0.99
0.29
0.97
-0.04
-0.57
-0.02
0.81
-0.02
0.81
-0.04
0.35
0.33
2.40
0.02
0.85
0.50
3.30
0.51
1.86
0.28
1.69
al. (1998) &
140
A. Photometric data
Table A.2: Photometric data for IC 4996 – continued
no
ref
X
Y
#
[px] [px]
76
D 81
895 838
77 D 34p 959 904
78 D 102 845 873
79
D 22
789 862
80
D 18
759 844
81
D 61
713 885
82
D 52
679 892
83
D 53
563 846
84 D 35p 538 850
85
D 21
492 819
86
D 17
446 811
87 D 100 468 847
88
D 76
179 832
89 P 66b 103 804
90
P 66
120 846
91
90 867
92 P 66a
41 810
93
197 952
94
241 988
95
P 63
287 946
96
P 60
362 953
97
408 981
98
436 989
99
P 58
552 944
100
P 57
575 963
101
691 968
102
P 38
723 952
103
749 959
104
P 40
787 990
105 D 37p 970 207
106
D 95
719 147
107 D 40p 435 601
108
D 96
177 261
109
44 263
110
D 72
399 373
111 D 42p 367 416
112 D 38p 269 210
113 D 41p 363 161
cross references according
P ... Purgathofer (1964)
V
(B − V )
MV
(B − V )0
[mag]
[mag]
[mag]
[mag]
14.93
0.85
0.15
1.65
15.46
0.98
0.28
2.18
15.97
1.22
0.52
2.69
14.45
0.76
0.06
1.17
14.00
0.71
0.01
0.72
13.49
0.75
0.05
0.21
12.13
0.62
-0.08
-1.15
12.17
1.83
1.13
-1.11
15.52
1.12
0.42
2.24
14.37
0.68
-0.02
1.09
13.78
0.62
-0.08
0.50
15.89
0.82
0.12
2.61
14.86
1.06
0.36
1.58
14.01
0.92
0.22
0.73
13.91
0.62
-0.08
0.63
15.23
1.26
0.56
1.95
13.32
0.55
-0.15
0.04
14.84
0.87
0.17
1.56
14.93
1.68
0.98
1.65
13.72
0.61
-0.09
0.44
12.40
0.70
0.00
-0.88
15.70
1.28
0.58
2.42
14.92
1.24
0.54
1.64
14.50
1.68
0.98
1.22
14.18
1.59
0.89
0.90
15.42
1.33
0.63
2.14
13.52
0.81
0.11
0.24
15.69
1.31
0.61
2.41
13.92
0.74
0.04
0.64
15.55
1.18
0.48
2.27
15.61
1.01
0.31
2.33
15.84
1.03
0.33
2.56
15.66
0.99
0.29
2.38
15.64
1.68
0.98
2.36
14.40
1.78
1.08
1.12
16.07
1.09
0.39
2.79
15.59
1.09
0.39
2.31
15.78
1.15
0.45
2.50
to: D ... Delgado et al. (1998) &
A.3. Stars in the fields of NGC 6530
A.3
141
Stars in the fields of NGC 6530
For all stars in the field of NGC 6530, V and (B − V ) values were computed.
Assuming cluster membership the according MV and (B − V )0 values were obtained
using E(B − V ) = 0.33 mag and (m − M ) = 11.65 mag (taken from the WEBDA
database). Tables A.3 and A.4 list the respective values together with the star
numbers assigned in this work (no), cross references with the WEBDA database
(ref ), if available, and the coordinates in pixels on the CCD (field 1: X1 , Y1 ; field
2: X2 , Y2 ).
Table A.3: Photometric data for all observed stars in field 1 of NGC 6530.
no
ref
X1
# [WEBDA]
[px]
1
67
1342
2
1774
867
3
1827
699
4
1785
806
5
159
780
6
105
760
7
101
806
8
98
891
9
97
890
10
108
689
11
1800
767
12
1750
951
13
94
934
14
1748
955
15
82
1090
16
89
1008
17
87
1044
18
83
1096
19
91
1003
20
90
1016
21
75
1243
22
155
1356
23
156
1370
24
62
1392
25
71
1288
26
69
1311
27
1557
1499
28
1580
1447
29
262
1472
30
1609
1370
31
1593
1415
32
1587
1435
33
1472
1719
34
27
1869
cross reference according
Y1
[px]
560
15
191
273
295
372
424
430
500
666
857
802
765
733
774
621
502
398
317
244
428
73
104
365
551
517
570
667
669
768
778
797
726
717
to the
V
(B − V )
MV
[mag]
[mag]
[mag]
11.94
0.40
0.29
14.92
0.80
3.27
14.57
3.40
2.92
14.47
0.96
2.82
13.57
0.46
1.92
10.55
0.18
-1.10
14.57
1.15
2.92
13.99
0.64
2.34
11.35
0.16
-0.30
14.06
0.94
2.41
14.87
2.51
3.22
18.20
-0.06
6.55
12.64
0.30
0.99
17.25
-0.73
5.60
12.35
0.33
0.70
13.52
0.79
1.87
12.85
0.97
1.20
10.56
0.15
-1.09
13.50
0.76
1.85
14.69
1.22
3.04
14.08
1.11
2.43
13.70
0.73
2.05
13.94
0.69
2.29
14.07
1.15
2.42
13.30
0.94
1.65
12.23
0.90
0.58
14.59
1.26
2.94
14.44
1.53
2.79
13.89
0.92
2.24
14.96
1.10
3.31
12.57
1.23
0.92
13.95
1.14
2.30
15.05
0.99
3.40
12.81
0.75
1.16
WEBDA database.
(B − V )0
[mag]
0.06
0.47
3.07
0.62
0.13
-0.15
0.82
0.31
-0.17
0.61
2.18
-0.39
-0.04
-1.06
0.00
0.46
0.63
-0.18
0.43
0.89
0.78
0.40
0.36
0.82
0.60
0.56
0.93
1.19
0.59
0.76
0.89
0.81
0.66
0.41
142
A. Photometric data
Table A.3: Photometric data for field 1 of NGC 6530 – continued
no
ref
X1
# [WEBDA]
[px]
35
22
1937
36
26
1884
37
39
1726
38
36
1770
39
34
1821
40
49
1654
41
48
1670
42
41
1732
43
33
1835
44
1368
2017
45
129
1106
46
128
1600
47
127
1531
48
1577
1450
49
1578
1449
50
46
1645
51
1551
1511
52
1618
1345
53
263
1365
54
264
1332
55
84
1045
56
136
873
57
135
819
58
110
589
59
1861
566
60
117
206
61
59
62
301
63
321
64
134
653
65
1848
635
66
314
561
67
1871
534
68
1867
543
69
1718
1070
70
1681
1165
71
72
1267
72
74
1230
73
114
493
74
139
30
75
313
240
76
137
301
77
99
78
119
cross reference according
Y1
[px]
684
609
529
350
251
179
77
89
47
213
2004
2018
1973
1897
1885
1876
1768
1743
1746
1703
1797
1843
1817
1843
1762
1692
1593
1548
1551
1681
1646
1572
1552
1534
1658
1569
1481
1399
1439
1358
1236
1193
1130
1065
to the
V
(B − V )
[mag]
[mag]
12.78
1.21
11.59
0.27
14.22
1.32
14.29
1.11
13.81
0.90
11.09
0.16
13.44
2.34
12.49
0.37
13.32
0.87
14.68
0.74
12.97
0.96
14.09
0.75
14.10
1.24
14.26
0.78
14.51
0.87
11.55
0.25
14.31
0.89
14.13
0.93
11.43
0.23
14.03
0.86
11.88
0.25
14.14
0.86
14.08
0.68
11.20
0.18
14.50
0.93
10.79
0.26
14.08
0.90
13.96
0.82
14.29
0.96
12.90
0.90
15.06
1.04
13.54
1.36
14.95
0.89
15.06
1.03
14.54
1.42
14.59
1.54
14.00
0.68
10.64
0.14
11.94
0.42
12.33
1.54
12.88
2.03
11.67
0.27
15.40
2.14
14.41
1.01
WEBDA database.
MV
[mag]
1.13
-0.06
2.57
2.64
2.16
-0.56
1.79
0.84
1.67
3.03
1.32
2.44
2.45
2.61
2.86
-0.10
2.66
2.48
-0.22
2.38
0.23
2.49
2.43
-0.45
2.85
-0.86
2.43
2.31
2.64
1.25
3.41
1.89
3.30
3.41
2.89
2.94
2.35
-1.01
0.29
0.68
1.23
0.02
3.75
2.76
(B − V )0
[mag]
0.88
-0.06
0.99
0.77
0.57
-0.17
2.01
0.04
0.54
0.40
0.63
0.42
0.90
0.44
0.54
-0.09
0.55
0.59
-0.11
0.53
-0.08
0.53
0.34
-0.15
0.59
-0.07
0.56
0.49
0.63
0.57
0.71
1.03
0.56
0.69
1.09
1.21
0.35
-0.19
0.09
1.21
1.70
-0.07
1.81
0.67
A.3. Stars in the fields of NGC 6530
143
Table A.3: Photometric data for field 1 of NGC 6530 – continued
no
ref
X1
# [WEBDA]
[px]
79
102
756
80
80
745
81
79
1118
82
78
1171
83
1641
1284
84
160
1307
85
53
1517
86
1566
1484
87
54
1490
88
1504
1641
89
37
1719
90
31
1823
91
29
1836
92
24
1874
93
1458
1739
94
64
1356
95
44
1677
96
40
1719
97
1658
1233
98
1672
1211
99
315
1131
100
81
1091
101
96
888
102
99
858
103
602
104
1884
464
105
447
106
34
107
114
108
319
263
109
112
539
110
113
524
111
158
387
112
409
113
71
114
140
99
115
164
116
88
117
141
94
118
372
119
116
421
120
512
121
111
601
122
109
655
cross reference according
Y1
[px]
1188
1112
1210
1139
1175
1164
1323
1173
1124
1221
1323
1409
1304
1214
1132
990
901
875
931
962
1058
1017
1029
997
859
893
829
874
811
707
623
527
510
506
493
391
358
219
142
249
268
245
290
491
to the
V
(B − V )
[mag]
[mag]
13.69
0.95
14.68
1.03
14.62
1.31
13.94
0.63
13.72
1.19
12.67
0.39
13.00
0.70
15.18
0.94
11.53
0.32
14.28
1.36
14.59
0.83
11.69
0.26
13.07
0.91
11.84
0.21
14.30
1.18
11.63
0.24
14.70
1.50
13.18
1.32
15.08
1.43
14.50
1.44
14.11
0.93
11.70
0.19
12.15
0.23
10.77
0.14
15.60
2.19
14.90
2.16
15.16
1.36
13.54
1.45
14.58
3.69
11.84
3.11
14.79
1.13
13.86
2.04
12.69
1.96
11.49
0.22
15.27
2.16
13.41
-0.01
30.24
-8.31
15.13
1.02
11.53
0.43
15.18
1.00
11.26
0.31
15.49
0.91
12.61
0.31
14.55
0.88
WEBDA database.
MV
[mag]
2.04
3.03
2.97
2.29
2.07
1.02
1.35
3.53
-0.12
2.63
2.94
0.04
1.42
0.19
2.65
-0.02
3.05
1.53
3.43
2.85
2.46
0.05
0.50
-0.88
3.95
3.25
3.51
1.89
2.93
0.19
3.14
2.21
1.04
-0.16
3.62
1.76
18.59
3.48
-0.12
3.53
-0.39
3.84
0.96
2.90
(B − V )0
[mag]
0.62
0.70
0.97
0.30
0.85
0.05
0.37
0.60
-0.01
1.03
0.50
-0.07
0.57
-0.12
0.85
-0.09
1.17
0.99
1.10
1.10
0.59
-0.14
-0.10
-0.19
1.85
1.83
1.02
1.12
3.36
2.77
0.80
1.71
1.63
-0.11
1.82
-0.35
-8.64
0.69
0.10
0.66
-0.02
0.57
-0.02
0.55
144
A. Photometric data
Table A.4: Photometric data for all observed stars in field 2 of NGC 6530.
no
ref
X2
# [WEBDA]
[px]
1
67
666
2
1774
192
3
1827
23
4
1785
130
5
159
104
6
105
84
7
101
130
8
98
215
9
97
214
10
108
13
11
1800
91
12
1750
272
13
94
259
14
1748
273
15
82
414
16
89
332
17
87
368
18
83
421
19
91
327
20
90
340
21
75
568
22
155
681
23
156
694
24
62
717
25
71
613
26
69
635
27
1557
823
28
1580
772
29
262
797
30
1609
695
31
1593
740
32
1587
760
33
1472
1043
34
27
1193
35
22
1262
36
26
1209
37
39
1050
38
36
1095
39
34
1145
40
49
978
41
48
994
42
41
1056
43
33
1158
44
1368
1341
245
17
1396
cross reference according
Y2
[px]
1726
1181
1357
1440
1461
1538
1590
1597
1667
1832
2024
1879
1931
1894
1941
1787
1668
1564
1483
1410
1594
1239
1270
1531
1717
1683
1736
1833
1835
1934
1944
1964
1892
1883
1850
1775
1695
1516
1417
1345
1243
1255
1213
1379
1978
to the
V
(B − V )
[mag]
[mag]
11.94
0.40
14.92
0.80
14.57
3.40
14.47
0.96
13.57
0.46
10.55
0.18
14.57
1.15
13.99
0.64
11.35
0.16
14.06
0.94
14.87
2.51
18.20
-0.06
12.64
0.30
17.25
-0.73
12.35
0.33
13.52
0.79
12.85
0.97
10.56
0.15
13.50
0.76
14.69
1.22
14.08
1.11
13.70
0.73
13.94
0.69
14.07
1.15
13.30
0.94
12.23
0.90
14.59
1.26
14.44
1.53
13.89
0.92
14.96
1.10
12.57
1.23
13.95
1.14
15.05
0.99
12.81
0.75
12.78
1.21
11.59
0.27
14.22
1.32
14.29
1.11
13.81
0.90
11.09
0.16
13.44
2.34
12.49
0.37
13.32
0.87
14.68
0.74
11.90
0.19
WEBDA database.
MV
[mag]
0.29
3.27
2.92
2.82
1.92
-1.10
2.92
2.34
-0.30
2.41
3.22
6.55
0.99
5.60
0.70
1.87
1.20
-1.09
1.85
3.04
2.43
2.05
2.29
2.42
1.65
0.58
2.94
2.79
2.24
3.31
0.92
2.30
3.40
1.16
1.13
-0.06
2.57
2.64
2.16
-0.56
1.79
0.84
1.67
3.03
0.25
(B − V )0
[mag]
0.06
0.47
3.07
0.62
0.13
-0.15
0.82
0.31
-0.17
0.61
2.18
-0.39
-0.04
-1.06
0.00
0.46
0.63
-0.18
0.43
0.89
0.78
0.40
0.36
0.82
0.60
0.56
0.93
1.19
0.59
0.76
0.89
0.81
0.66
0.41
0.88
-0.06
0.99
0.77
0.57
-0.17
2.01
0.04
0.54
0.40
-0.14
A.3. Stars in the fields of NGC 6530
145
Table A.4: Photometric data for field 2 of NGC 6530 – continued
no
ref
X2
# [WEBDA]
[px]
246
1340
1445
247
167
1521
248
19
1293
249
1359
1372
250
164
1644
251
166
1759
252
165
1705
253
1278
1685
254
290
1549
255
6
1992
256
20
1280
257
21
1285
258
1374
1313
259
1335
1464
260
1377
1298
261
1393
1270
262
25
1229
263
57
804
264
63
714
265
161
836
266
52
877
267
51
937
268
47
1007
269
1541
619
270
77
559
271
1759
252
272
1788
122
273
1811
56
274
144
26
275
1693
463
276
294
800
277
295
927
278
38
1071
279
23
1241
280
1404
1231
281
13
1538
282
11
1622
283
231
1875
284
232
1838
285
230
1695
286
162
1271
287
163
1276
288
28
1210
289
1355
1395
290
1345
1427
cross reference according
Y2
[px]
2002
1816
1668
1529
1396
1355
1304
1207
1208
1269
1373
1289
1278
1333
1188
1175
1112
1121
1046
985
943
932
993
971
861
1079
1045
910
811
917
865
843
775
935
970
1012
1046
912
743
840
693
665
608
580
575
to the
V
(B − V )
[mag]
[mag]
14.67
1.46
14.36
1.03
10.82
0.25
14.73
0.85
13.96
0.86
13.60
0.84
14.70
1.19
15.42
0.86
13.92
1.74
10.79
0.11
14.06
0.75
14.55
0.90
15.50
0.95
15.54
0.94
14.80
1.09
15.26
0.92
11.40
0.16
13.63
0.68
12.26
0.90
13.75
0.61
12.11
0.37
14.22
2.22
12.28
0.30
14.31
0.71
11.99
0.54
15.22
0.88
14.69
0.85
13.84
2.48
11.20
0.23
15.00
1.17
14.19
1.18
13.94
0.81
12.16
0.57
12.41
1.21
15.12
1.04
13.32
0.49
14.77
0.76
14.19
0.67
14.21
0.80
14.05
0.76
12.00
0.79
12.69
0.31
13.20
0.45
14.76
2.21
15.15
1.00
WEBDA database.
MV
[mag]
3.02
2.71
-0.83
3.08
2.31
1.95
3.05
3.77
2.27
-0.86
2.41
2.90
3.85
3.89
3.15
3.61
-0.25
1.98
0.61
2.10
0.46
2.57
0.63
2.66
0.34
3.57
3.04
2.19
-0.45
3.35
2.54
2.29
0.51
0.76
3.47
1.67
3.12
2.54
2.56
2.40
0.35
1.04
1.55
3.11
3.50
(B − V )0
[mag]
1.13
0.70
-0.09
0.52
0.52
0.51
0.86
0.52
1.40
-0.22
0.41
0.56
0.62
0.61
0.75
0.59
-0.18
0.35
0.56
0.27
0.03
1.89
-0.03
0.38
0.21
0.55
0.52
2.15
-0.10
0.84
0.85
0.47
0.24
0.88
0.71
0.16
0.43
0.34
0.47
0.43
0.45
-0.03
0.12
1.88
0.67
146
A. Photometric data
Table A.4: Photometric data for field 2 of NGC 6530 – continued
no
ref
X2
# [WEBDA]
[px]
291
1306
1580
292
229
1620
293
1456
294
293
958
295
1516
929
296
1526
901
297
1492
990
298
157
953
299
1503
969
300
798
301
152
622
302
311
375
303
150
370
304
292
498
305
151
541
306
291
698
307
688
308
153
925
309
283
1307
310
16
1453
311
284
1556
312
1446
313
1698
314
1836
315
285
1796
cross reference according
Y2
[px]
624
625
517
744
721
693
717
682
626
499
468
692
170
128
193
151
108
211
252
394
307
111
178
61
339
to the
V
(B − V )
MV
[mag]
[mag]
[mag]
13.38
0.77
1.73
13.79
2.29
2.14
14.76
0.77
3.11
12.98
2.41
1.33
14.36
2.37
2.71
15.02
1.52
3.37
14.49
1.94
2.84
14.18
0.80
2.53
14.80
1.01
3.15
14.35
1.00
2.70
12.98
0.70
1.33
13.94
1.05
2.29
13.13
0.72
1.48
13.90
0.76
2.25
12.16
0.85
0.51
13.10
1.60
1.45
14.66
0.88
3.01
12.54
0.88
0.89
13.19
1.40
1.54
12.14
0.30
0.49
13.83
0.99
2.18
14.13
0.80
2.48
14.60
0.80
2.95
14.30
0.98
2.65
14.20
0.93
2.55
WEBDA database.
(B − V )0
[mag]
0.44
1.96
0.44
2.07
2.03
1.19
1.61
0.47
0.68
0.67
0.37
0.71
0.38
0.42
0.52
1.27
0.55
0.55
1.06
-0.03
0.65
0.47
0.47
0.64
0.60
Abbreviations
The following abbreviations are used in the text:
AA – Astronomy & Astrophysics
A&AS – Astronomy & Astrophysics Supplement
AJ – Astronomical Journal
ApJ – Astrophysical Journal
ApJS – Astrophysical Journal Supplement
Ap&SS – Astrophysics & Space Science
CCD – charge coupled device
CMD – colour-magnitude diagram
Comm. Ast. – Communications in Asteroseismology
CTIO – Cerro Tololo Interamerican Observatory
DSN – Delta Scuti Network
HAEBE – Herbig Ae/Be
HR-diagram – Hertzsprung-Russell diagram
IMF – initial mass function IR – infrared
IRAF – Image Reduction and Analysis Facility
KPNO – Kitt Peak National Observatory
MNRAS - Monthly Notices of the Royal Astronomical Society
MOMF – Multi Object Multi Frame
NOAO – National Optical Astronomy Observatory
OSN – Observatory Sierra Nevada
PASP – Publications of the Astronomical Society of the Pacific
PMS – pre-main sequence
PSF – point-spread function
SED – spectral energy distribution
UV – ultraviolet
ZAMS – zero-age main sequence
147
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ASP Conf. Ser. 173, 181
Use was made of the First Digitized Sky Survey (http://archive.eso.org/dss/dss),
which is copyrighted by the Space Telescope Science Institute (STScI Digitized Sky
Survey, 1993, 1994, AURA, Inc. all rights reserved), the WEBDA database, operated by E. Paunzen (http://www.univie.ac.at/webda/) and the SIMBAD database.
Curriculum Vitae
Mag.a Konstanze Zwintz
born May 4th , 1974
in Vienna, Austria
June 1992
October 1992
July 1996
December 1996
December 1998
March 1999
In 2000
March - April 1999
October 1999
May - June 2000
Graduation from high school in Vienna, Austria
Commencement of studies Astronomy and Astrophysics
at the University Vienna, Austria
Summer student at the Fermi National Accelerator Laboratory (FERMILAB) with Prof. Alberto Santoro in
Batavia, Illinois, USA
Begin of masters’ thesis Hubble Deep Field Guide Star
Photometry in the group of Prof. Werner W. Weiss
Scholarship at the University of Vienna, Austria
Graduation with the Master of Natural Sciences at the
University of Vienna
PI of the project Microvariability with the Hubble Space
Telescope Fine Guidance Sensors sponsored by the Jubiläumsfonds der Österreichischen Nationalbank
PI of the ASTROVIRTEL Project Asteroseismology With
The Hubble Space Telescope Fine Guidance Sensors with
the ST-ECF in Garching, Germany
Research at the University of British Columbia, Vancouver, Canada, in preparation of the satellite project MOST
with Prof. Jaymie Matthews and Dr. Rainer Kuschnig
Begin of PhD thesis in the group of Prof. Werner W.
Weiss at the University Vienna, Austria
Research at the University of British Columbia, Vancouver, Canada, in preparation of the satellite project MOST
with Prof. Jaymie Matthews and Dr. Rainer Kuschnig
Start of the development of the VISAT database
(http://ams.astro.univie.ac.at/visat/)
151
2000 - 2005
July 2003
July 2004
Since October 2004
Tutor for the lectures Astronomische Instrumente I and Astronomische Instrumente II at the Institute of Astronomy, University Vienna, Austria
Participation as student in the Space Summer School Alpbach
2003
Tutor during the Space Summer School Alpbach 2004
Coordinator and organiser of the COROT PMS Thematic
Theme
Conferences
May
December
October
August
July
April
July
May
September
September
May
January
2005
2004
2004
2004
2004
2003
2002
2002
2001
2000
2000
2000
April 1998
June 1997
COROT Week 8, Toulouse, France
COROT Week 7, Granada, Spain
First Brazilian COROT Workshop, Natal, Brazil
MOST Science Team Meeting, Vienna, Austria
IAU Symposium 224, The A-Star Puzzle, Poprad, Slovakia
The Second Eddington Workshop, Palermo, Italy
Asteroseismology Across the HR-diagram, Porto, Portugal
Gesamtösterreichische Astronomentagung, Graz, Austria
COROT Science Week, Vienna, Austria
COROT Science Week, Paris, France
MOST Science Workshop, Vancouver, Canada
The Third MONS Workshop: Science Preparation and Target
Selection, Aarhus, Denmark
COROT Kick Off Meeting, Nice, France
A Half Century of Stellar Pulsation Interpretations - A Tribute
To Arthur N. Cox, Los Alamos, New Mexico, USA
Observing Runs
September 2004
June 2004
January 2004
November 2002
September 2002
August 2002
August 2001
August 1999
0.9m Telecope at the Cerro Tololo Interamerican Observatory (CTIO), La Serena, Chile (14 nights)
3.6m Telescope, ESO La Silla, Chile (3 nights)
1.8 m Telescope, Osservatorio Asiago, Italy (3 nights)
Organisation of the multi-site campaign for the pre-main
sequence pulsators V 588 Mon und V 589 Mon:
0.65 m Telescope, Mauna Kea, Hawaii (14 nights)
0.75 m APT, Fairborn Observatory, Arizona (14 nights)
0.9 m Telescope, Observatorio Sierra Nevada (OSN), Spain
(14 nights)
1.0 m Telescope, Loiano Observatory, Italy (6 nights)
0.75 m APT, SAAO, Sutherland, South Africa (14+ nights)
1.5m Telescope at the Observatorio Sierra Nevada (OSN),
Spain (14 nights)
0.9m Telescope at the Cerro Tololo Interamerican Observatory (CTIO), La Serena, Chile (14 nights)
0.9m Telescope at the Cerro Tololo Interamerican Observatory (CTIO), La Serena, Chile (14 nights)
1.9m Telescope at the South African Astronomy Observatory (SAAO), Sutherland, South Africa (14 nights)
Publications
Ripepi V., Bernabei S., Marconi M., Palla F., Arellano Ferro A., Bonanno A.,
Ferrara P., Frasca A., Jiang X.J., Kim S.-L., Marinoni G., Mignemi G., Monteiro M.J.P.F.G., Oswalt T.D., Reegen P., Rimas J., Rodriguez E., Rolland A.,
Ruoppo A., Terranegra L., Zwintz K., 2005, A multisite photometric campaign
on the Pre-Main-Sequence δ Scuti pulsator IP Per, AA, submitted
Kallinger T., Zwintz K., Pamyatnykh A.A., Guenther D.B., Weiss W.W., 2005,
Pulsation of the K 2.5 giant star GSC 09137-03505?, AA 433, 267
Zwintz K., Marconi M., Reegen P., Weiss W.W., 2005, Search for pulsating premain sequence stars in NGC 6383, MNRAS 357, 345
Taraba M., Zwintz K., Bombardelli C., Lasue J., Rogler P., Ruelle V., Schlutz J.,
Schüßler, O’Sullivan S., Sinzig B., Treffer M., Valavanoglou A., Van Quickelberghe M., Walpole M., Wessels L., 2005, Project M3 - A Study for a Manned
Mars Mission in 2031, Acta Astronautica, in press
Zwintz K., Marconi M., Kallinger T., Weiss W., 2004, Pulsating pre-Main sequence
stars in young open clusters, in proceedings of the IAU Symposium No. 224.
‘The A Star Puzzle’, J. Zverko, J. Ziznovsky, S.J. Adelman, and W.W. Weiss
eds., Cambridge University Press, p. 353
Ripepi V., Bernabei S., Marconi M., Palla F., Arellano Ferro A., Bonanno A., Ferrara P., Frasca A., Jiang X. J., Kim S. L., Marinoni S., Mignemi G., Oswalt
T. D., Reegen P., Rimas J., Rodriguez E., Rolland A., Ruoppo A., Terranegra
L., Zwintz K., 2004, Multisite observations of the pre-Main-Sequence δ Scuti
star IP Per, in proceedings of the IAU Symposium No. 224. ‘The A Star Puzzle’, J. Zverko, J. Ziznovsky, S.J. Adelman, and W.W. Weiss eds., Cambridge
University Press, p.799
Zwintz K., Weiss W.W., 2004, Pulsating pre-main sequence stars as possible Eddington targets, in proceedings of ‘Second Eddington Workshop: Stellar structure and habitable planet finding’, F. Favata, S. Aigrain and A. Wilson eds.,
ESA SP-538, p.105
Weiss W. W., Aerts C., Aigrain S., Alecian G., Antonello E., Baglin A., Bazot
M., Collier-Cameron A., Charpinet S., Gamarova A., Handler G., Hatzes A.,
154
Hubert A.-M., Lammer H., Lebzelter T., Maceroni C., Marconi M., de Martino
D., Janot-Pacheco E., Pagano I., Paunzen E., Pinheiro F. J. G., Poretti E.,
Ribas I., Ripepi V., Roques F., Silvotti R., Surdej J., Vauclair G., Vauclair S.,
Zwintz K., 2004, Additional Science Potential of COROT, in proceedings of
‘Second Eddington Workshop: Stellar structure and habitable planet finding’,
F. Favata, S. Aigrain and A. Wilson eds., ESA SP-538, p.435
Paunzen E., Zwintz K., Maitzen H.M., Pintado O.I., Rode-Paunzen M., 2003, New
Variable Stars in Open Clusters I: Methods and Results for 20 Open Clusters,
AA 418, 99
Kallinger T., Kaiser A., Stuetz Ch., Weiss W.W., Zwintz K., Bigot L., 2003,
MOST and COROT high precision photometry simulations of the roAp star
10 Aquilae, in proceedings of ‘Asteroseismology across the HR Diagram’, Astrophys. & Space Science, Vol. 284, No.1
Kallinger T., Zwintz K., Kaiser A., Mittermayer P., Weiss W.W., 2003, VISAT VIenna Selection of Astronomical Targets, Comm. Ast., No. 143, 43
Oehlinger J., Kaiser A., Kallinger T., Mittermayer P., Weiss W.W., Zwintz K.,
2003, The MOST and COROT prime target fields: A target inventory, Comm.
Ast., No. 143, 36
Zwintz, K., Weiss W.W., 2003, Pulsating Pre-Main Sequence Stars in NGC 6383?,
in proceedings of ‘Asteroseismology Across the HR Diagram’, Astrophys. &
Space Science, Vol. 284, No. 1
Kuschnig R., Matthews J., Lanting T., Walker G.A.H., Zwintz K., 2000, Ultrapresice photometry from space: Simulations of the MOST space telescope
performance, in proceedings of the ‘MOST Science Workshop’, Vancouver,
Canada
Weiss W.W., Kuschnig R., Zwintz K., 2000, Variability Survey with the HST in
proceedings of ‘The Impact of Large Scale Surveys on Pulsating Star Research’,
L. Szabados & D.W. Kurtz eds., ASP Conf. Ser. Vol. 203, p. 38
Zwintz K., 2000, Stellar Photometry with the HST Fine Guidance Sensors, in proceedings of the ‘MOST Science Workshop’, Vancouver, Canada
Zwintz K., Kuschnig R., Weiss W.W., Witeschnik A., 2000, Photometric Properties
of the Hubble Space Telescope Fine Guidance Sensors, in proceedings of ‘The
Impact of Large Scale Surveys on Pulsating Star Research’, L. Szabados &
D.W. Kurtz eds., ASP Conf. Ser. Vol. 203, p. 82
Zwintz K., Weiss W.W., 2000, Asteroseismology with the HST Fine Guidance Sensors: The Microvariability Survey, in proceedings of ‘The Third MONS Workshop: Science Preparation and Target Selection’, Aarhus Universitet, Denmark, 24-26 January 2000
Zwintz K., Weiss W.W., Kuschnig R., Frandsen S., Gray R., Jenkner H., 2000,
Variable HST Guide Stars (I), AAS 145, 481
Weiss W.W., Zwintz K., Kuschnig R., Witeschnik A., 1999, Deadtime Correction,
Color term and Sensitivity of the Hubble Space Telescope Fine Guidance Sensors, Comm. Ast., No. 129
Zwintz K., Kuschnig R., Weiss W.W., Gray R.O., Jenkner H., 1999, Hubble Deep
Field Guide Star Photometry AA 343, 899
Zwintz K., 1999, Hubble Deed Field Guide Star Photometry Master Thesis, University Vienna
Kuschnig R., Weiss W.W., Zwintz K., 1998, Microvariability survey based on photometry with the HST Fine Guidance Sensors, in proceedings of ‘A Half Century of Stellar Pulsation Interpretation: A Tribute to Arthur N. Cox’, P. A.
Bradley & Joyce A. Guzik eds., ASP Conf. Ser. 135, 362
Zwintz K. , 1998, Investigation of Systematic Effects in the HST Fine Guidance
Sensors Photometry, COROT (Seismology of Stars: Convection and Rotation): Minutes of the Kick-Off Meeting, Nizza, 27. - 29. April 1998
Weiss, W. W., Kuschnig R., Mkrtichian D. E., Kusakin A. V., Kreidl T. J., Bus S.
J., Osip D. J., Guo Z., Hao J., Huang L., Sareyan J.-P., Alvarez M., Bedolla
S. G., Zverko J., Ziznovsky J. V., Mittermayer P., Zwintz K., Polosukhina N.,
Mironov A. V., Dorokhov N. I., Goranskij V. P., Dorokhova T. N., Schneider
H., Hiesberger F., 1998, Photometry of ET Andromedae and pulsation of HD
219891, AA 338, 919
Kuschnig R., Weiss W.W., Zwintz K., 1997, Stability of FGS Photometry HST
Calibration Workshop, STScI, S. Casertano et al., eds
Danksagungen
Zu allererst bedanke ich mich bei Werner W. Weiss, der mich auf das spannende
Thema meiner Dissertation aufmerksam gemacht und mich in all den Jahren tatkräftig bei allen Arbeiten unterstützt und gefördert hat. Ohne seinen unermüdlichen
Einsatz um verschiedene Forschungsprojekte, hätte ich meine Zeit nicht voll für die
Wissenschaft aufwenden können. Ich danke ihm auch dafür, dass ich mit Problemen
immer zu ihm kommen konnte und er sich stets Zeit für mich genommen hat.
Michel Breger sei auch für seine wertvollen Ratschläge für meine Arbeit im Lauf
der Jahre gedankt. In seinen Privatissima habe ich einiges über photometrische
Problemstellungen lernen können.
Gerald Handler danke ich für das Korrekturlesen dieser Arbeit und seine wichtigen Kommentare.
Special thanks goes to Marcella Marconi and Alosha Pamyatnykh, who showed
a lot of patience in explaining the basics of stellar pulsation theory to an observer
like me. I am also grateful for the several model calculations they have performed
for some of my stars.
Meinen Kolleginnen und Kollegen auf der Sternwarte und speziell in meiner
Arbeitsgruppe ein herzliches Dankeschön für die gute Zusammenarbeit, das nette
Umfeld und die vielen gemeinsam getrunkenen Liter Kaffee.
Spezieller Dank gebührt meinem Großvater Ing. Alfred Beyrl. Er hat mir die
Ruhe und Möglichkeit zur Entspannung gegeben, wenn ich viel um die Ohren hatte.
Mein wärmster Dank ergeht an die besten Freunde, die man haben kann: Petra,
Sigrid, Marion, Monika, Rainer, Verena, Victor, Clarissa, Christa, Peter, Michi,
Werner, Silvia – ihr habt mir mit eurer Geduld, euren Schultern zum Anlehnen und
eurer Unterstützung in jeder Form den für mich so wichtigen Rückhalt in guten und
schlechten Phasen gegeben.
Zuletzt möchte ich meinen Eltern posthum dafür danken, dass sie meine Interessen von klein auf gefördert, meine Ausbildung ermöglicht und immer an mich
geglaubt haben.
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