Pulsating Pre-Main Sequence Stars In Young Open Clusters

Transcription

1 Pulsating Pre-Main Sequence Stars In Young Open Clusters Dissertation eingereicht von Mag a. 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

2 Meinen Eltern, Dr. Edgar und Mag. Sigrid Zwintz, gewidmet.

3 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 10 6 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

4 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.

5 Contents Abstract 1 1 Early Stellar Evolution Star formation Evolution of pre-main sequence stars The birthline for low-mass stars The birthline for intermediate-mass stars Evolutionary tracks Pre-main sequence stars T Tauri stars Classes of T Tauri stars Variability Herbig Ae/Be stars Variability Evolutionary stage Asteroseismology Introduction Pulsation δ Scuti stars γ Doradus stars The classical instability strip Pulsation in PMS stars Historical background The PMS instability strip Theoretical investigations Comparison with observations (status 2000) Seismology of PMS stars Young open clusters Basic definitions Embedded and exposed clusters

6 4 CONTENTS 5.2 Sequential star formation PMS stars in young open clusters NGC Historical background IC Historical background NGC Historical background Cloud collapse and star formation Proper motion studies Observations and data reduction NGC Bias level variations Color dependent extinction IC SigSpec NGC Observational Results Pulsating PMS stars in NGC NGC NGC NGC Summary of PMS pulsators in NGC Pulsating PMS stars in IC IC IC IC IC Summary of PMS pulsators in IC Pulsating PMS stars in NGC NGC NGC NGC NGC NGC NGC NGC NGC Summary of PMS pulsators in NGC Other variables Variable stars in NGC Variable stars in IC

7 CONTENTS Variable stars in NGC Summary of cluster properties NGC IC NGC Modelling pulsation Pulsation constants Pulsation models Discussion of observed frequencies PMS γ Doradus type pulsators The empirical PMS instability strip All known pulsating PMS stars The new PMS instability strip Conclusions 129 A Photometric data 131 A.1 Stars in the field of NGC A.2 Stars in the field of IC A.3 Stars in the fields of NGC Abbreviations 147 Bibliography 148 Curriculum Vitae 151 Publications 154 Danksagungen 157

8 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, follow- 6

9 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 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 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,

10 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

11 1.3. Evolutionary tracks 9 barrier. If it reaches 10 6 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 T eff 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.

12 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.

13 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 Taurus- Auriga 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 K to 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, H α 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, anal- 11

14 12 2. Pre-main sequence stars ogous to the solar chromosphere. Features which cannot be explained yet include the strong H α, 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 Classes of T Tauri stars Classical T Tauri stars (CTTSs) were discovered from H α 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 H α 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) 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 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.

15 2.2. Herbig Ae/Be stars 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 H α 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.

16 14 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 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.

17 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 RE obs denotes the empirical red edge, BE the blue edge for the radial overtones and BE F the blue edge for the fundamental mode (Breger & Pamyatnykh 1998).

18 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 10 7 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): ξ nlm (r, θ, φ, t) = ξ nl (r)y l m (θ, φ)e iω nlmt (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

19 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 ω nlm 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 ν nlm ω nlm /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 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 K.

20 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 = ) 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.

21 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 can be also written as (Breger 1979): Q = P (ρ/ρ ) 1/2 (3.2) log Q = log P log g M bol + log T eff (3.3) For the fundamental mode in δ Scuti stars Q evaluates to d, for the first overtone mode Q is d, for the second harmonic it is d and for the third d. The smaller the Q values the higher is the corresponding radial overtone pulsation mode for δ Scuti stars 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

22 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 , a γ Doradus star with most likely tidally excited pulsation, the cross relates to HD , 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 (p 1 ) and seven overtones (p 2 - p 8 ) 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

23 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 (RE obs 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 p 1, the blue edges for the according overtones are marked with p 2 - p 8, respectively. RE obs 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.

24 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.

25 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 The position of V588 Mon (HD , NGC ) and V589 Mon (HD , NGC ) 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 , 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 T eff is 7800 K and the surface gravity log g 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

26 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 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 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 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 must be a high overtone PMS pulsator.

27 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 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 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 UBV (RI) C CCD time-series photometry. Both stars show two periods simultaneously: BL 50 with 1.7 hours (I C amplitude = 18 mmag) and with 2.4 hours (I C amplitude = 6 mmag); HP 57 with 1.9 hours (I C amplitude = 27 mmag) and with 1.5 hours (I C 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 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 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

28 26 4. Pulsation in PMS stars Name RA (2000.0) DEC (2000.0) sp V log T eff log L/L f [hh:mm:ss] [dd:mm:ss] [mag] # V589 Mon 06:39: :42:4.1 F2 III V588 Mon 06:39: :41:3.4 A7 III/IV NGC 6823 HP57 19:43: :16: NGC 6823 BL50 19:43: :17: V351 Ori 05:44: :08:40.4 A7 IIIe HR :08: :06:18.3 A7 III/IVe HD :27: :19:38.4 F0 IIIe HD :00: :11:34.6 A4 V Table 4.1: Parameters of the eight known PMS pulsators in the year 2000: name, right ascension (RA) and declination (DEC) at the epoch , spectral type (sp) if available, V magnitude, effective temperature (log T eff ) and luminosity (log L/L ). The last column specifies the number of frequencies found in each star as of 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 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 T eff 7100 K and the blue edge between 7100 T eff 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 Kelvin- Helmholtz time scale. For a 1.5 M star this accounts to 10 6 yr and for a 4.0 M star it is only 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 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 known eight PMS pulsators (coloured symbols). Also, the borders of the classical δ Scuti instability strip for the theoretical, fundamental blue edge (BE F, 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

29 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

30 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 T eff = 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).

31 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 time 1 is greater than 10 8 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

32 30 5. Young open clusters 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 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

33 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.

34 32 5. Young open clusters 5.4 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 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 (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. α h 34.8 m δ age 1.7 ± 0.4 Myr diameter 20 distance 1.5 ± 0.2 kpc Table 5.1: Main cluster properties for NGC 6383.

35 5.4. NGC 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 He also determined the distance modulus for the cluster to be 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 ) 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 H β 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 , 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 UBV 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 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 ULBV 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 , is found to be a double-lined binary with an effective temperature of 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 , has a mass loss rate smaller than

36 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. R v = 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, H α 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.

37 5.5. IC 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 10 7 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 (taken from the First Digitized Sky Survey). α h 16 m 30 s δ age 8.87 Myr diameter 6 distance kpc Table 5.2: Properties of IC 4996 taken from the WEBDA database Historical background Alfaro et al. (1985) obtained uvby and H β observations for 15 stars brighter than V = 12 mag with spectral types earlier than A0 and report a mean distance modulus of ± The age of IC 4996 was estimated by them to be years.

38 36 5. Young open clusters Vansevi cius 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 ± years. A total of 1120 stars was measured by Delgado et al. (1998) in a field of 7 7 of IC They found the average values of the colour excess and true distance modulus to be E(B V ) = 0.71 ± 0.08 mag and (V 0 M V ) = 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 ± 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.

39 5.6. NGC 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; Author V 0 M V E(B V ) age distance [mag] [mag] [Myr] [kpc] Walker (1957) 3.0 Kilambi (1977) Chini & Neckel (1981) 0.39±0.09 Mc Call (1990) 11.35± ±0.07 Sung (2000) 11.25± Table 5.3: Astrophysical properties of the PMS pulsators in NGC 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

40 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 M8 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 UBV 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 10 7 years ago and probably is continuing up to now. 37 pre-main sequence stars with H α emission were detected in NGC 6530 by Sung et al. (2000) using UBV RI and H α 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 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

41 5.6. NGC 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 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 H α 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 M V (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 Proper motion studies For 363 stars brighter than 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).

42 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 (see Figure 6.1) with a scale of /px. In total, 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 , South is at the top and East is to the left). 40

43 6.1. NGC 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.quad 1 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.

44 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 px px Figure 6.3: Schematic map of the observed field of NGC 6383 (fov , 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): σ(a) = 2 N σ(m) (6.1) σ(f) = 6 N 1 π T σ(m) A, (6.2)

45 6.1. NGC 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 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!).

46 44 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 (Sterken & Manfroid 1992): m = m 0 (k 0 + k CI) X, (6.3) where m 0 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 ) (B-V) literature 2 slope k (B-V) instrumental (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 8 th 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

47 6.2. IC a 0 = ± 0.052, a 1 = ± and a 2 = ± (B V ) instr values for all stars could then be transformed to the standard system according to: (B V ) trans = a a ((B V ) instr a 0 ) a 2 a 1 2 a 2 (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, 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 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

48 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 = 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 For all 113 identified stars a detailed frequency analysis was performed using primarily again Period98, but also the new software SigSpec developed by Reegen

49 6.2. IC px 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 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[φ FA (A)], (6.5) where Φ FA is the false alarm probability. A significance of 8, for example, means that in one out of 10 8 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.

50 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 NGC 6530 Figure 6.9: Composed false-color image of the two observed fields of NGC The overlapping region is clearly visible. Each frame has a fov of , 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 with a scale of /px. In total, 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).