VEGA: a Visible spectrograph and polarimeter for CHARA Science Cases Description

Transcription

1 VEGA: a Visible spectrograph and polarimeter for CHARA Science Cases Description Philippe Stee a, Denis Mourard a, Daniel Bonneau a Paul Berlioz-Arthaud b, Armando Domiciano de Souza a,c, Renaud Foy b, Petr Harmanec d, Slobodan Jankov c, Pierre Kervella e, Pavel Koubsky f, Stephane Lagarde a, Jean-Baptiste Le Bouquin g, Philippe Mathias a, Antoine Mérand e,h, Nicolas Nardetto a, Romain G. Petrov c, Karine Rousselet-Perraut g, Chantal Stehlé j, Gerd Weigelt i a Département GEMINI, Observatoire de la Côte d'azur, BP 4229, F Nice Cedex 04, France b Centre de Recherche Astronomique de Lyon, F Saint Genis-Laval Cedex, France c Université de Nice-Sophia Antipolis, Parc Valrose, F Nice Cedex 02, France d Astronomical Institute, Charles University Prague, V Holesovickach 2, CZ Praha, Czech Republic e Observatoire de Paris, LESIA, Paris, France f Astronomical Institute. Academy of Sciences of the Czech Republic g LAOG, 414 rue de la Piscine, DOmaine Universitaire, Saint Martin d Héres, France h CHARA (GSU), Mount Wilson Observatory, USA i Max Planck Institute für Radioastronomie, Auf den Hügel 69, D Bonn, Deutschland j Observatoire de Paris, LUTH, Paris, France 1. Introduction Interferometry has been intensively done at long wavelengths, starting with the radio interferometers in the years 50 th since it was easier to guide radio wavelengths in cable while keeping the phase information or using a local oscillator and a correlator to recombine "a posteriori" the beams over intercontinental distances. In the optical a lot of work as been done at IR and near-ir wavelengths since it was technically easier, or we must say, less difficult to recombine directly the optical beams since the coherence length is larger and the turbulence slower than at shorter wavelengths. Therefore, the visible domain of the electromagnetic spectrum is not covered at the same level than near or mid infrared. Some very nice and important results have been however obtained with the GI2T interferometer in south of France, the Mark III interferometer on the Mount Wilson, USA, the NPOI array in Flagstaff, USA or the SUSI interferometer in Australia. We will present in this paper the science cases of a new but already existing and tested instrument: the REGAIN focal instrument which was designed and built for the GI2T. This instrument, in his CHARA adaptation, called VEGA will open new fields in a wide range of Astrophysical topics only addressable in the visible domain. It will provide a spectral resolution up to within the spectral range µm and a spatial resolution of less than 1mas for up to 4 telescopes in its X- special configuration. A polarimetric device (SPIN) measuring simultaneously the polarization in 2 directions either circular or linear is also implemented in this instrument. Since VEGA was already tested on the sky on 1.5 m telescopes it is also very well suited for the 1m CHARA array and will only need minor adaptations for the injection of the CHARA beams. This paper will focus on some of the most promising science drivers only possible with this visible instrument. 2. Details of the VEGA instrument and expected performances VEGA is proposed with a dispersed fringe mode (called X- ): o Simultaneous combination of 4 telescopes, giving access simultaneously, for one measurement and per spectral channel, to 6 squared visibilities, 6 differential phases and 3 phase closures. o Spectral resolution: 1500, 5000 and and spectral range: from 0.45 to 0.87 m o Polarimetric capabilities allowing to simultaneously record interferograms in two polarizations, either linear or circular. Advances in Stellar Interferometry, edited by John D. Monnier, Markus Schöller, William C. Danchi, Proc. of SPIE Vol. 6268, 62683R, (2006) X/06/$15 doi: / Proc. of SPIE Vol R-1

2 2.1. Limiting magnitude and signal to noise considerations At short wavelengths, there is currently no adaptive optics devices available, so that images are not diffraction limited : this is the so-called multimode regime. In this regime we usually consider two domains for the signal to noise ratio (S), depending on the number of photons per speckle and per single exposure. If this number is larger than 1 (N>>1), then the signal to noise ratio increased as the square root of the number of photons (S N 2 ), whereas if it is smaller than 1 (N<<1), the signal to noise ratio increased as the number of photons (S N), (Roddier, 1988). In consequence, we first defined the limiting magnitude as the magnitude giving 1 photon per speckle and per single exposure (N=1). The signal to noise ratio is then equal to one (S=1) and should be multiplied by the square root of the number of speckle and by the square root of the number of images. We usually consider in speckle techniques that the limiting magnitude is defined as this limit. It is clear however that the signal to noise ratio could reach high values at this limit, since one can integrate a large number of single exposures containing a large number of speckles (Fig.1). In the differential regime, it has been shown by Petrov (1986), that the signal to noise ratio is the geometrical mean of the signal to noise ratio in the reference channel and in the science channel. Then a good signal to noise ratio could be achieved also in a science channel where the number of photons per speckle and per frame is much less than 1 (N<<1), thanks to the fact that we can have more than 1 photon per speckle and per frame in the reference channels (Fig 2). These considerations will be used for the following calculations. We also consider the following hypothesis for the calculations: 1. For a V=0 star, the number of photons received is equal to N 0 =1000 ph/s/cm 2 /A 2. Transmission in the visible Q Total =0.15% assuming a. Q CHARA =0.03 b. Q Instrument =0.15 ( grating the slit 0.3) c. Q Detector = Exposure time t 0 =20ms, integration time=1800s 4. r 0 estimations (for 650 nm) at Mt Wilson. These informations have been extracted from the Nils Turner presentation «A Hodge-Podge of CHARA Array Housekeeping functions» : a. Median conditions, r 0 =8.0*(650/500) 6/5 =11.0 cm (seeing=1.25 ) b. Excellent conditions, r 0 =15.0*(650/500) 6/5 =20.6 cm (seeing 0.7 ) I: Limiking magnifude a a) an C) S an I an! an Figure.1: Limiting magnitude for 30mn of integration, a signal to noise ratio of 10 on the differential visibility measurement with 1 = 50 nm (resp. 40, 6.7) and 2 = 0.4 nm (resp. 0.13, 0.02) for the Low LR (resp. Proc. of SPIE Vol R-2

3 medium MR, high HR) spectral resolution and a central wavelength equal to 650nm. The vertical lines correspond to the median and excellent seeing conditions I I I I I I I I I I I I I : LR HR N N. O1OT N 0.01 I I I I I I I I I I I I magnifude V Figure 2: SNR calculations for the three different spectral resolutions. The hypotheses on the spectral bands are the same than in figure 1. The calculations are made for a seeing of 1.25 (median conditions) and 30mn of integration time. As already explained, the principle of multimode operation relies on a coherence tracker and not on a fringe tracker. Therefore the previous calculations do not assume an external fringe tracker system but only a coherence tracking system able to stabilize the optical path difference to less than 2 microns over the duration of the integration of short exposures (typically 20 ms). A good model of baselines and an optical path difference measurement at a period of a few seconds are enough to a correct operation of VEGA. Real time processing of the data is used for the optical delays stabilisation as already demonstrated on the GI2T with the RAFT experience (Koechlin et al. 1996) 3. Science drivers for a visible instrument on CHARA Fundamental stellar parameters Physical processes working in the stellar interiors as well as the evolution of stars are based on some fundamental parameters, for instance the mass, the radius, the luminosity and abundances. The effective temperature, the surface gravity and the mean density are useful parameters defined from these fundamental parameters. Some other physical parameters like mass loss rate, pulsation period, rotation period or magnetic field must be interesting for the study of peculiar evolutionary stage. A classical way to test the stellar interior models is to compare the predicted and observed location of a star on theoretical evolutionary tracks in a H-R diagram. This can be done only for stars for which the mass, radius, luminosity and abundances are well known. To obtain significant results, accuracy of 1%-2% for the mass and the radius and an error on the effective temperature better than 50K is mandatory. Then a precision of about 5% - 10% on luminosity is sufficient. The direct determination of stellar mass and radius is only possible from high angular observations in two cases: Proc. of SPIE Vol R-3

4 - for a single star with accurate parallaxe (thanks to Hipparcos and Gaïa) - for stars in resolved spectroscopic binary Mass [M0] Figure 3: Radii and masses relation of very-low-mass stars from Ségransan et al. (2003) from the VLTI (blue points), PTI (red points) and from eclipsing binaries (black points). The stellar effective temperature can be computed combining the angular diameter (1% accuracy needed) with bolometric photometric data. An angular resolution of the order of 1 mas is required to resolve the visual orbit of near all the known spectroscopic binaries. In the near IR with an angular resolution of about 2 mas with a 100m baseline, the majority of the stars resolvable are red giants and only a very small number of main sequence stars (solar type or red dwarfs, PTI and VLTI 2003 see Fig. 3). A gain in angular resolution of about 4 will be obtain by observing in the visible with CHARA and it will be possible to measure the angular diameter of main sequence stars of spectral type ranging from O to M in the solar neighborhood. Moreover, the determination of the orientation of stellar rotation axes is also possible as well as the global alignment of binaries rotation axes through the measurement of the differential phase between the red and blue wings of photospheric absorption lines, numerous in the visible domain. Special care should be devoted to rapidly rotating hot O and B stars where also the effects of deviations from a spherical shape, limb and gravity darking play role. One can characterize the stellar flux either by the effective temperature which is defined locally and depends on the stellar latitude or to keep the original definition where the surface of the spherical star with radius R is replaced by the surface of the rotationally distorted star. A combination of spectrointerferometry with spectroscopy and photometry for suitably selected binaries will significantly help to improve the knowledge of these important physical properties of hot stars. The idea is to observe both, eclipsing and non-eclipsing spectroscopic binaries. Since even a partial resolution of such objects can permit determination of their spatial orbits (see, e.g. Taylor et al. 2003), one obtains the orbital inclinations and therefore accurate masses for non-eclipsing spectroscopic binaries. A comparison of radiative properties of eclipsing and non-eclipsing components in detached binaries with rapidly rotating hot stars should allow an empirical determination of gravity darkening for such objects. More generally, the study of little evolved stars in detached binaries can substantially improve the knowledge of all basic physical properties needed to test and further improve the theory of stellar structure and evolution. Although binaries in general are complex objects both from the observational and modeling point of view, their study can provide Proc. of SPIE Vol R-4

5 fundamental data to refine the theory of single-star formation and constrain possible scenarios of double stars formation. Reaching angular resolution close to 1 mas allows not only to separate the components of spectroscopic binaries but also to resolve the photospheric disks of the components. In such a case, the angular diameters can be combined with the distance to obtain the stellar radii. If the flux ratio can be determined as function of the wavelength, the observed spectro-photometric flux distribution of individual stars can be combined with the angular diameters to derive the effective temperatures and luminosity of the components. For more evolved and strongly interacting close binaries, additional physical processes have to be taken into account due to interaction between the two components, such as tidal distortion, mutual irradiation, mass transfer and/or mass loss or wind collision. The geometrical structure of such systems then involves not only the stellar disks but also various components of circumstellar matter: the accretion disk, scattering envelope and jet-like structures outside the orbital plane and also the gaseous stream between the components. As expected with GI2T-REGAIN (Thureau et al., 2000), spectro-interferometry in the visible with VEGA on CHARA, combined with classical techniques like photometry and spectroscopy may provide a deep insight into the morphology of such objects and important constraints for their future modeling Stellar activity: Interferometry can contribute to the study of stellar activity by detecting and mapping star spots (Fig. 4). A limited number of stars have features large enough to allow interferometric imaging. The gain in resolution allowed by the visible will substantially increase the number of targets. For instance, one of the main result of helioseismology is the location of the tachocline which is thought to be the engine of the solar magnetic field that gives rise to spots. In the Sun, spots are relatively small (1% of the solar radius), and, usually being formed in the equatorial region, tend to move to largest latitude. It seems that no solar spots have been observed close to the poles. Using stellar rotation, it is possible to reconstruct the stellar surface through inversions of high resolution spectroscopic observations such as Doppler Imaging technics, based on the line profile variation s(,t) as a function of time in spectrally resolved lines. The main result is that in most cases, very large spots (up to 30% of the visible surface) have been pointed out in the polar region. Problems remain about the application of this method, concerning the input physical parameters (stellar rotation, temperature differences...) and also the uncertainty in the latitude distribution. Figure 4: Numerical simulation of the red supergiant Betelgeuse surface from Freytag 2005 see URL: ( Proc. of SPIE Vol R-5

6 Although this technique has produced spectacular and numerous results, Doppler Imaging has a number of drawbacks (see for instance Fig. 8 in section asteroseismology ). First, the technique does not work for stars seen equator-on because there is no way to decide if a feature is in the northern or southern hemisphere. This leads to artifacts in images of stars which are viewed nearly equator-on. The areas near the poles are poorly reconstructed, with a tendency of the technique to produce polar spots and a not completely closed debate about the reality of the numerous polar spots detected. The technique is based on the assumption of rigid rotation. Differential rotation or meridian motions of spots in times comparable with the rotation period are very hard to take into account. It also assumes that the profile of the line without rotation is correctly modeled. As a result the uniqueness and the accuracy of many images is questionable. Interferometry allows the combination of s(,t) with the photocenter variation with wavelength and time (,t) deduced from the fringe phase (,t) even for unresolved objects. It has been demonstrated that this additional interferometric information eliminates the north-south ambiguity in equator-on stars and that it allows the separation of spatial and spectral features (Petrov 1989; Lagarde, 1994). It also allows the extension of the technique to much slower rotators, when the static line width is comparable or larger than the rotation broadening. Again, in some cases this program will be undertaken with the infrared Amber instrument. However, it needs clean, sharp and deep photospheric lines observed at a spectral resolution comparable with the static line width, i.e. of about or This requirement clearly point toward a visible instrument. In the IR, the reasonable size for the Amber spectrograph limited its highest spectral resolution to In the visible, with the same volume one can built a spectrograph with a resolution of Also, even if the technique is applied to non resolved features, the observation time needed to achieve a given result varies as the square of the interferometer resolution. An approximate rule of thumb allows to compare the SNR on s(,t) and on (,t), obtained with the same telescopes and instruments: SNR[s(,t)]/SNR[ (,t)] * /( /B) where * is the stellar angular diameter and ( /B) is the interferometer angular resolution. The angular diameter of most of the stars observed by Doppler Imaging is between 0.2 and 4 mas. This explains why the technique performs poorly with single aperture telescopes, where ( /D)>> *, and why it should perform fairly well with the CHARA array in the visible. It will be possible to adjust the baseline to have * 0.7( /B) for many systems (if * is too close to ( /B), the fringe contrast is destroyed) Exemple of stellar atmosphere studies: Abundance and magnetic field of Ap stars Chemically Peculiar A and B stars (CP stars) exhibit strong chemical abundance inhomogeneities of one or more chemical species (such as helium, iron-peak elements and/or rare earths) and a large-scale organized magnetic field that produces a typical signature in circularly-polarized spectra (classical S -shape of the V-Stokes spectral lines). As a major class of the known magnetic stars in the solar neighborhood, CP stars constitute an ideal lab for studying the magnetic field effects on physical processes occuring inside stellar atmospheres. Simultaneously mapping abundance distributions and magnetic fields is an important step in addressing the fundamental question of the origin of the magnetic field in CP stars: both the fossil and the core-dynamo theories have difficulties in explaining all the observed magnetic characteristics of CP stars (Moss, 2001). Secondly, the magnetic field and the abundance inhomogeneities are so closely related that maps have to be obtained simultaneously to understand well the key role of magnetism in atmosphere structuration (Leblanc et al., 1994), in ion migration across the stellar surface (Michaud, 1970), and in chemical stratification (Ryabchikova et al., 2002). However, very few abundance maps of CP stars are available today (e.g. Kochukhov et al. (2002) for 2CVn) and very few maps of magnetic fields have been reconstructed via Zeeman- Doppler Imaging (e.g. Kochukhov et al. (2002) for 2CVn) or by inversion of spectro-polarimetric data (Bagnulo et al. (2000) for CrB). Moreover, such inversion methods often lead to several magnetic field models that cannot be disentangled by classical spectro-polarimetric techniques. Within this context, the color-differential interferometry is very attractive to derive 2D abundance and magnetic maps since the differential fringe phase can be used to locate spots on stellar surfaces (Jankov et al., 2003). As it is equivalent to the first moment of the intensity distribution, this observable has a large sensitivity in the limb regions and gives access to absolute orientations via its sign, even for edge-on or pole-on geometries (Rousselet-Perraut et al., 2004). Proc. of SPIE Vol R-6

7 Coupling high spectral resolution with high angular ability is obviously of strong interest for adressing these stellar activity issues. A spectral resolution of is enough even for slow rotators like CrB and the medium one is convenient for faster rotators Differential rotation and stellar inclination Owing to phenomena such as meridional circulation, turbulence rotationally induced global and local hydrodynamical instabilities, magnetic stresses, non-radial pulsations, more or less steady horizontal shear velocity fields may exist. If a latitudinal external differential rotation actually exists, departures from symmetry in the rotationally broadened spectral line profiles should be expected. A powerful technique for measuring subtle features in stellar spectral lines is the Fourier transform (FT) method (e.g. Gray 1973, 1975). However, when it comes to differential rotation, several authors have shown that it cannot be disentangled from the stellar inclination angle if only spectroscopic data are used (Gray 1977, Bruning 1981, Reiners & Schmitt 2002 and 2003). However, differential rotation creates projected velocity fields (Fig.5 left), which give rise to non-symmetrical intensity maps (Fig. 5 center) inside spectral lines so that not only these lines are deformed but also there is a clear displacement of the stellar photometric barycenter, i.e., the photocenter across the line (Fig.5 right). Due to this photocenter signature we can combine differential interferometry (DI) with the Fourier transform method to study the stellar surface differential rotation. The advantages of combining FT of DI data to determine the stellar angular diameter were put forward by Chelli &Petrov (1995a,b) and Petrov et al. (1995). They showed that by this technique the influence of the local profile is removed in all Fourier frequencies and that consistent results can be obtained. Recently, Domiciano de Souza et al. (2004) developed an interferometric method to study differential rotation from DI combined to the Fourier transform of photocenter components and stellar spectra. Compared to spectroscopic techniques, their method has the double advantage of allowing the use of all available Fourier frequencies as well of disentangling the stellar differential rotation rate and the inclination angle. However, rather high instrumental performances are needed to determine consistently these parameters. In particular it is crucial to have a spectral resolution R high enough to result in as many as possible measurements across the spectral line profile (R > should be necessary) and spatial resolution able to partially resolve stars smaller than 1~mas to increase the number of available targets. Additionally, differential phases should be measured with precisions high enough so that the uncertainty in the photocenter is < 10-3 of the stellar angular radius, i.e. <10% of the photocenter amplitudes expected for differential rotators. i=ily:u U.5 j=75:{y={).6 :1 Figure 5: Left: Radial velocity maps for several combinations of i and, assuming a solar-like differential rotation law (l) = V eq R (1 sin 2 l), where l is the latitude. A positive velocity corresponds to displacements towards the Proc. of SPIE Vol R-7

8 observer (blue shifts). For differential rotation, regions of constant projected velocity (equal velocity strips) are not straight vertical ones as is the case for rigid rotation ( =0). Center: Intensity maps inside a photospheric line for i=45 and = 0.6. The curved dark patterns correspond to Doppler shifts of the local line profile caused by differential rotation. These non-symmetrical intensity maps result in a displacement of the stellar photometric barycenter, i.e., the photocenter (filled circles), relative to the geometrical center (opened circles). Right: Corresponding photocenter components, z (top; vertical direction) and y (middle; horizontal direction), and the normalized spectral flux (bottom). Photocenter components are given in units of angular stellar radius. The letters indicate selected wavelengths corresponding to a region in the blue line wing ( a ), the highest y value ( b ), the central wavelength ( c ), and the most positive value of z ( d ). Note that z =0 for uniform rotation. Figures from Domiciano de Souza et al. (2004).} Asteroseismology Asteroseismology studies the temporal variation of the photometric flux or of the spectrum s(,t) to identify the resonant modes of stars (Fig. 6). Since these modes propagate through the entire star (and the low order modes actually go through or close to the center), asteroseismology allows the exploration of the third dimension of the star. As in the case of Doppler Imaging, the combination of the spectral information of asteroseismology with differential interferometry can yield much improved results. As for the study of stellar spots the study of the spectrum alone yields a lot of problems, mainly to measure very small velocity variations, to identify the modes, to access higher orders and to solve windowing effect. For almost all pulsating stars, the angular diameter is below 1 mas, thus strongly reducing the contribution of resolved interferometry. However, in the case of non resolved stars, it has been shown that adding (,t) to s(,t) solves many problems. The demand for instrumental stability is relaxed, since for example it is possible to calibrate (,t) using as a reference s(,t), which has been recorded at the same time by the same pixels in the same instrument. This also allows a lower dependence on a perfect modelling of the non-oscillating spectrum. The element of spatial resolution introduced by (,t) permits to access higher modes. The fact that many modes affect very differently s(,t) and (,t) strongly helps identifying them. This last feature might also somehow ease the temporal coverage and window problems, but this remains to be studied. As for stellar spots, Doppler Imaging can be applied to the reconstruction of the pulsation modes. Jankov (2001) has recently demonstrated by fully simulating the image reconstruction process that the combination of s(,t) and (,t) if far more discriminating than the use of spectroscopic data alone (see Fig. 9). Figure 6: Spherical projection of the stellar surface brightness perturbation due to the non-radial pulsation m=4, l=5 mode on a star tilted at i=85 For stars close to the main sequence, many modes may be excited in different classes of variables, such as Cephei, SPB, Scuti, roap, Doradus and solar-like stars. Considering these latter, the Sun is the best example, and we have now a quite accurate picture of the inside of this star: extension of the nuclear core, radiative region, rotation structure (tachocline)... In particular, the sound velocity (i.e. thermodynamical conditions) is now modeled with an accuracy of Proc. of SPIE Vol R-8

9 about 0.1%. These results are only possible for the Sun, because it is a bright star (!) spatially resolved. Consequently, million modes can be measured, with high degrees (up to l=3000). However, because only flux is measured rather than intensity in all other stars, the highest degrees measurable for moderate rotators is l=4. The situation is less dramatic for faster rotators, where degrees up tu l=15 can be detected. However, other problems arise, such as the coupling between rotation and pulsation, diffusion... It is thus of prime importance to have access to a resolved stellar surface to detect numerous modes with high degrees, and visible interferometry is a promising tool for such a purpose. Scientific goals are very numerous: can the solar structure be generalized for all solar-like stars, how this structure evolves with time (avoided crossing, mixed-modes), what is the internal structure of more massive stars... Until now, modes have been identified through different photometric and/or spectroscopic methods, but not directly visualized. A first step would be to valid these different methods (flux amplitude ration, moments method, Fourier Doppler imaging) observing in the same conditions, i.e. visible spectrum (allowing thus a better spatial resolution than usual IR interferometers) and also with a quite high spectral resolution (30000 is enough for p-modes and for low-degree g- modes). The technique could be applied to solar type stars with a dedicated instrument optimized for this type of observations with CHARA: it would be some kind of Echelle spectrograph allowing to combine information from all lines in the visible spectrum. Such an instrument could only exist in a somehow later future. A VEGA visible instrument would yield access to Scuti and B stars and possibly to red giants. As for stellar spots, the need for photosphere lines, high spectral resolution and a spatial resolution allowing to approach resolving the stellar diameter point toward a visible instrument. For the observation of stellar spots and for asteroseismology they are significantly more interesting spectral features below 0.6 µm than between 0.6 and 0.8 µm. This is a reason to lower the shorter wavelength accessible to a CHARA visible instrument as much as possible. S 0.25 be PHASE PHASE ii gjg QT o.s a Figure 7: Pole-on (Top) and Mercator (Bottom) projections of the visible surface on a star tilted at i=85. Left: The input image of the surface temperature distribution due the non-radial pulsation m=4, l=5. Right: Maximum Entropy reconstructions from normalized flux spectra alone (Doppler Imaging) Cepheids: distance scale and pulsating atmospheres Long-baseline interferometers currently provide a new quasi-geometrical way to calibrate the Cepheid Period- Luminosity relation. Indeed it is now possible to determine the distance of galactic Cepheids up to 1kpc with the Proc. of SPIE Vol R-9

10 Interferometric Baade-Wesselink method (Kervella et al. 2004, see also Sasselov & Karovska 1994). Interferometric measurements lead to angular diameter estimations over the whole pulsation period, while the stellar radius variations can be deduced from the integration of the pulsation velocity curve deduced from line profiles by the so-called projection factor p (see Fig. 8). The visible band is particularly well adapted to extend the work on Cepheids already started with the VLTI/VINCI, FLUOR/CHARA and in a near future with AMBER/VLTI (Mourard 1996). Limb darkened angular diameter (mae) roe. LO Angolar dirn. (,a) Figure 8: left - VINCI observations of the pulsation of the Cepheid variable l Car (P = 35.5 days, red dots) and the adjusted radius curve (blue line), as deduced from the integration of the radial velocity measured on this star over its pulsation period (Kervella et al. 2004). right - FLUOR/CHARA observations of Cep (Mérand et al. 2005). First, the main limitation for the Period Luminosity calibration is the number of stars for which a direct distance determination can be determined by directly measuring with a good accuracy the amplitude of variation of the angular diameter. Figure 9a shows the gain in resolution permitted by the visible and subsequently the increase of Cepheids observable. Cepheids from Moskalik & Gorynya (2005) are represented accordingly to their mean angular diameter and the relative pulsation amplitude. Zones of detection of the pulsation at 20 (considering only two measurements at maximum and minimum radius) are represented considering a relative precision on the visibility of v 2 /V 2 =1.5% at V 2 =30%. The detection of the pulsation at 20 corresponds to a precision on the distance of 5% (with only two observations by cycle). A precision of 2% can be achieved for several points. Note that VEGA/CHARA is in the north hemisphere, thus we find that concerning Cepheids pulsating in the fundamental mode that : ~10 stars are observable for B<100m, ~15 stars are observable for 110m<B<220m, and ~5 stars are observable for 220<B<330m. FLUOR/CHARA (B=330m, Infrared) is comparable to VEGA/CHARA at B=100m in the visible. Figure 9b represents the number of Cepheids observed and the precision on the distance required to achieve a 0.01 magnitude precision on the zero point of the P-L relation. With a precision of 5% on the distance, the observation of 100 stars is required, while a precision of 2% need only 20 stars. Such objective seems to be possible with VEGA. Moreover, considering a minimum visibility detection of V 2 =5% it seems to be challenging but possible to detect the pulsation of first overstone Cepheids, which has never been done. It is clear that, due to the small angular size of Cepheids, increasing the angular resolution by more than a factor 4 is crucial for this programme. However, the most powerfull advantage of the instrument VEGA is its spectral resolution of With such a resolution it would be possible to constrain the projection factor. Proc. of SPIE Vol R-10

11 (a) (b) Relative precision on the Po relative radial pulsation amplitude Q P - P V ci DII (n.1..,. * 0) 33Om. 220m. 01 1Om IIpr. Cepheids data from Moskalik & Gorynya (2005) 0, ,02 0,00 Number of elate Relative precision on the distance 2% Number of elate Figure 9: left - Cepheids of fundamental mode (blue) observable with the VEGA/CHARA interferometer right Precision (in magnitude) obtained on the zero-point of the P-L relation as a function of Cepheids observed : with a precision of 5% on the distance, the observation of 100 stars is required, while a precision of 2% need only 20 stars. Second, a very important limiting aspect of the interferometric Baade-Wesselink method is the projection factor, which is actually related to the atmospheric velocity gradient present between the photosphere of the star and the line-forming region. Nardetto et al quantified this effect to 6% on the derived distance in the case of Cep. It corresponds to a 3km/s effect on the pulsation velocity amplitude. This theoretical result has been confirmed using the FLUOR/CHARA (low spectral resolution) instrument and the precise distance determination of the HST, by Mérand et al (Fig. 8). To determine the projection factor corresponding to each Cepheid of our sample (or at least for several Cepheids with period ranging from 3 days to 40 days), the visibility and phase differential measurements with the spectral and spatial resolution provided by VEGA are very promising. Nardetto et al. 2005a using a hydrodynamical model has calculated for each spectral channel in the line profile the corresponding visibility signature in the case of Cep. Such effect is clearly reinforced with the high spectral resolution of VEGA (Fig 10ab) and also for high spatial resolution (long baseline). Considering this visibility signature together with the differential phase measurements would permit to constrain Cepheid physical parameters, and in particular velocity gradients (and thus the projection factor). The wavelength- and phase-dependent visibility signatures are represented for different spectral lines in Fig. 10c. It is also possible by phase differential measurements to follow the Cepheid angular diameter curve, and thus to increase our sample of Cepheids. Distant Galactic Cepheids or angularly small Cepheids can constitute very interesting targets to test for example the metallicity effect on the Cepheids period luminosity relation (Mourard, 2005). The expected performances of differential phase measurements permit to reach a 1 as resolution. For a pulsating and slowly rotating Cepheid, the phase shift is of the order of 20% of the angular diameter. It appears then that VEGA could measure the angular diameter of many Galactic Cepheids by this way. These applications require spectral resolution between and 40000, which is much easier to achieve in the visible (Petrov, 1989). Proc. of SPIE Vol R-11

12 (a) Fs r e<por (c) (b) 0,51 0,50 0,49 0,49 0,47 0,46 0,45 00o0R=45000 R= R , ,0 6002,5 6003,0 6003,5 Wavelenpth Velocity (km/s) Figure 10: (a) A modelled 3D diagram which represents the spectral line profile (FeI A) for the maximum expansion velocity ( =0) with the corresponding intensity distributions. (b) Visibility as a function of wavelength for a projected baseline of 65 meters (GI2T) and for different spectroscopic resolutions. (c) The visibility is represented as a function of the velocity in the case of four spectral lines and for each pulsation phase. The pulsation phase is indicated on the right side of the diagram (Nardetto et al. 2005a). The high spectral and spatial resolution of VEGA are required to constrain Cepheids parameters and in particular the projection factor Conclusion : synopsis of the program on Cepheids. 1/ Observation of Cepheids (Moskalik & Gorynya 2005) in low resolution and high base to increase the sample of Cepheids observed by VINCI/VLTI, AMBER/VLTI and FLUOR/CHARA. Expecting a 2% precision on Cepheids distances, it will be possible to calibrate the zero-point of the Cepheids P-L relation at a level of 0.01 magnitude. 2/ If a wrong projection factor is used to determine the distance of the Cepheid, an error of 0.01 magnitude is possible on the zero-point. Using the high resolution (R=30000) and high baselines (B=300) of VEGA/CHARA it is possible to constrain the projection factor for several Cepheids. The priority targets are the most resolved Cepheids. This program is all the more important as there is currently a debate on a possible relation between the period of the Cepheid period and the projection factor (Gieren et al. 2005, Nardetto et al. 2005b). Moreover the phase differential measurements with high precision can be used to determine the Cepheid angular diameters and consequently their distances. This point is important to increase the sample of Cepheids observed: distant or angularly small Cepheids. Moreover we can mention that the very high precision of VEGA/CHARA would be useful for challenging programs as the determination of limb-darkening phase-dependence (effect of about 1%), and/or detection of Alpha Polaris amplitude of variation (less than 1%). In this last case the polarimetric mode could be very interesting. Proc. of SPIE Vol R-12

13 Mira stars and related objects At least three kind of programs can be addressed by VEGA/CHARA, and only by it. One is the measurement of the diameter of large sample of Miras, including Population II ones. The goal is to determine the pulsation mode, which is now undoubtely the fundamental one for 5 of them (Perrin et al, 2004). To what extent this result, critical for our understanding of the inner structure of these stars, can be generalized to the whole class, including Population II ones? The observable volume in the Galaxy, and in particular out of the disk, is much larger at visible than at IR wavelengths, and at CHARA than at the VLTI. These observations should better be undertaken at wavelengths below 440nm (Bonneau et al., 1982), where TiO bands are no longer present. The second program aims at studying the surface inhomogeneities of the closest Miras. Convection at least is a source of such in inhomogeneities (Schwarzschild, 1975; Freytag, 2005). Shock waves is another one. And it is an open question to know whether smock waves and convection interact. Observations at visible wavelengths with the spatial resolution of CHARA are able to bring new insight in this field. They require high spectral resolution (>~ 10000) to be able to select individual emission lines, and components of these lines. The critical feasibility issue of this programme is the accuracy in the visibility measurements at low visibilities. For instance, to detect a 1marcsec dark spot over a 20 mas uniform disk would require to measure 1% squared visibility difference at V 2 ~ 3%. The flux from the closest Miras is enough to measure such low visibilities if we are in the shot noise regime. The main limitation could come from calibrations. Differential phase measurementsmay help, depending on the spectral domain considered (within the visible range). Modeling of different spatial energy distributions is being done currently at CRAL. Since the properties of the velocity fields and of the convection pattern vary significantly, depending of the velocity on the phase of the light cycle, time coverage is definitely required. Finally, a third programme would highly useful to constrain model atmosphere of Miras : the measurement of the center-to-limb darkening (or possibly brightening, depending on the wavelength). It is much weaker at IR wavelengths than at visible ones, due to the relative continuum opacities. In addition, one knows since Labeyrie et al (1977) that the extension of Mira atmospheres is rapidly varying with the wavelength, due to molecular and atomic absorptions. The effect is markedly stronger at visible wavelengths than at IR, in particular because of the much higher number of lines in the visible. The center-to-limb darkening will vary accordingly (Jacob et al., 2000). Measurements of the spectrum of the center-to-limb darkening with VEGA/CHARA will provide the tool to setup empirical atmosphere models, as a function of the phase of the light cycle of the stars Active hot stars The envelopes of Be stars have been one of the favorite topics for visible interferometric observations. However, the mechanisms causing the mass loss and shaping the resulting circumstellar envelopes remain unclear. As well as the formation and the vanishing of the circumstellar disk as a function of time: is it the radiative pressure which excavate the disk or the gas which is slowly diluted in the interstellar vacuum? Be stars are also an excellent test bench for the interferometric study of many other astrophysical problems, such as the morphology and kinematics of envelopes or accretion disks around other types of stars or of the BLR of AGN. One important parameter is the typical size of the envelope for a given wavelength. Visible diameters are known but there is a debate about the IR ones. This illustrates that observations at different wavelengths actually allows exploring the 3 rd dimension of the structure. Furthermore one can try to understand the morphology and the kinematics of the envelope. Modeling by Stee & Bittar (2001) shows that images in narrow wavelength bins will allow extremely strong constraints on the envelope structure. However such images have not been observed so far, because no interferometer so far combined an imaging capability with the adequate spectroscopic resolution. Recently the VLTI as observed the Be star Ara with the MIDI instrument. Unfortunately the ~ 100m baseline was not large enough to resolve the circumstellar disk at 10 µm. Nevertheless, these measurements put an upper limit to the Proc. of SPIE Vol R-13

14 envelope size in the N band of max = 4 mas, corresponding to 14 R *, assuming R *=4.8R * and the Hipparcos distance of 74 pc. On the other hand the disk density must be large enough to produce the observed strong Balmer line emission the disk density must be large. In order to estimate the possible circumstellar and stellar parameters we have used the SIMECA code developed by Stee (1995) and Stee & Bittar (2001). Optical spectra taken with the échelle instrument Heros and the ESO-50cm telescope, as well as infrared ones from the 1.6m Brazilian telescope have been used together with the MIDI spectra and visibilities. These observations put complementary constraints on the density and geometry of Ara circumstellar disk. The possibility of larger distance that indicated by the Hipparcos parallax is discussed, as well as a potential truncation of the disk by a putative companion (see Chesneau et al. 2005). It is clear that with an increase by a factor of 16 in spatial resolution from MIDI at 10 µm to VEGA at 0.6 µm observations we will be able to study the connection between Ara central star and its deep disk structure. Moreover the physical processes between the visible and the IR are very different and more energetic phenomena will be accessible to visible observations. A lot of work has been focused on the information which can be extracted from the wavelength dependence of simple parameters such as the differential visibility V( )/V( ref ) or phase ( ) ( ref ). Stee has shown that the function that constrains most the kinematics is the differential phase. This parameter can be obtained with a single three-telescopes configuration, with two approximately perpendicular baselines, even if the source is unresolved by the interferometer. In this case, one obtains the variation of the source photocenter with wavelength. For envelopes with an axial symmetry, it has been possible to show that only three parameters extracted from the photocenter figure, which is the figure drawn by the photocenter through an emission line, provide very strong constraints on the kinematics (Stee, 1996). The overall diameter of the photocenter figure is connected to the size of the envelope. The angle between the main axis of the photocenter figure and the stellar rotation axis gives the ratio between rotation and expansion velocity. The angular extension of the curve is closely related to the velocity law (solid, Keplerian, etc ). Very recent results from the AMBER instrument with a spectral resolution of 1500 already seems to evidence what as been predicted by Stee (1996), i.e. the evidence of a Keplerian circumstellar rotating disk around Ara seen across the Br emission line (see Fig. 11). Figure 11: A schematic view of Ara seen from recent VLTI/MIDI observations (Chesneau et al. 2005) Proc. of SPIE Vol R-14

15 A careful fit of the high SNR photocenter curve by models might finally determine which velocity law is valid within the disk, but a final solution will be found only through full imaging of the envelope in a large number of narrow spectral channels selected in one or more emission line(s). For this image reconstruction, the resolution gain in H compared to Br will be decisive. The typical angular size for the disks of the closest Be star is about 3 milliarcseconds. This is too close to the resolution limit at 2.16 µm with a 200 m baseline (2.2 mas) to expect good images in the infrared. In the visible, one can expect images with up to 10x10 resolution elements, which should be enough to discriminate between several possible structures. The direct resolution of small structures within the disk would require kilometric baselines. Such small structures should evolve very rapidly in time, typically in a few hours. Following these variations might yield interesting information, given that the SNR is high enough and that we have a sufficient spectral resolution such the provided by the VEGA instrument. As discussed below, it is also possible to have information about stellar activity through differential interferometry, even if the star is not resolved. Shift of the photocenter in the North South direoton ID I I I 1,1! I! I! I! I! I! I,1!I! I! I! I! I 5 / o // IOL Ang t WAVELENGTH I I!I!1 Figure 12: Left: Predicted photocenter displacement as a function of wavelength with the SIMECA code from Stee (1996) for different rotational velocity laws and preliminary results from AMBER observations showing the same phase effect as a function of wavelength for the Be star Ara (Right). In summary, for a large number of stars, it will be possible to draw the photocenter curve as a function of wavelength. This curve will allow the derivation of many disk parameters and therefore allow statistical studies of the numerous variants of Be stars. If the curve can be obtained with a high SNR and a high spectral resolution, its fit will allow stronger model constraints. In some cases, it will be possible to resolve the central star (diameter, Vsini, position angle of the rotation axis) and even to follow its activity (spots or non radial oscillations). However, to be able to trust the interpretation of the photocenter measurements, it will be necessary to have full image reconstruction for some systems. All these points can, and will, be tackled by the infrared instrument (except probably the study of the central object). However, the gain provided by a visible extension will be decisive in many aspects. Only the angular resolution permitted in the red will permit images with a sufficient number of resolution elements. The highest spectral resolution will allow the detection and monitoring of small scale structures in the envelope. The photocenter and visibility variations with wavelength will be measured with much higher SNR. Typically the gain will be of the order of 15 because of the combined effect of stronger emission lines and better spatial resolution. This allow the observation of much more systems, which is decisive given the large number of variants of the Be phenomena. For instance Figure 13 presents our current view of the HAe/Be star MWC 297 which has been modeled by combining an accretion disk model Proc. of SPIE Vol R-15

16 and the SIMECA code in order to interpret AMBER MR-K measurements. It is obvious for such an object that the observation of the emitting region producing the observed H line profile (with an intensity 120 times the continuum!) will put very strong constrains on our modeling by, for instance, setting if H originates within the equatorial disk or in the spherical stellar wind. For a given star visible interferometry will also dramatically decrease the time necessary to detect small structures and therefore allow us to follow their evolution. Finally, the resolution of central star parameters and detection of stellar activity signature, which require very high SNR photocenter curves in photosphere lines, is very likely to be possible only in the visible. Figure 13: A schematic view of the HAe/Be star MWC 297 using an accretion disk model with the SIMECA code constrained from recent VLTI/AMBER observations (in press) Hot emission-line stars in binaries Spectrointerferometry also provides a unique chance to study what is the role of duplicity for the Be phenomenon. Many bright Be stars were found to be components of binary systems and the number of such discoveries is increasing (see Harmanec 2001 for a recent catalogue of hot emission-line stars in binaries). It is obvious that for Be stars which are components of binaries, the Roche lobe represents the maximum extent which their envelopes can occupy. Besides, variations of such envelopes must be somehow affected by the presence of the companion. Although there are cases when the companions to Be stars were directly observed, in many other cases the evidence of duplicity rests only on regular radial-velocity variations of Be stars in question. As explained by Taylor et al. (2003), it should be possible to detect companions up to 5 magnitude difference between the components. Spectrointerferometry can therefore not only Proc. of SPIE Vol R-16

17 allow to detect companions to many Be stars but also to find out their properties and nature. It will also be possible to see how the geometry of Be envelopes in binaries differs from that in single stars. As illustrated Figure 14 in the case of Lyrae (jet-like structures evidenced from interferometric observations with the GI2T and spectroscopic observations during the 1994 international campaign, Harmanec et al., 1996) high angular resolution observations combined with classical techniques like photometry and spectroscopy can provide deep insight into the morphology of the star and its circumstellar environment in order to constrain the modeling of such phenomena. For the distance of ±270 pc according to Hipparcos parallax and adopting the orbital radius of 58.5 R (Harmanec, P., 2002) the maximum angular separation of the stellar components of Lyrae is expected to be around 1 mas and should be resolved with CHARA array at wavelengths close to H. With VEGA on CHARA, a spectral resolution of about 5000 allows to use the Interferometric Differential Imaging Technique to measure the relative phase of the fringe visibility across H or HeI 6678 emission lines and to determine the relative position of the different emitting regions with a spatial resolution of a few µas. orb 0,10 1 mas orb 0,25 Figure 14: An artist's view of Lyrae system with various components of its circumstellar matter: gas stream, jet-like structure ans scattering halo/star wind above the hidden mass-accreting star (from Harmanec, 2002). An other interesting target for VEGA on CHARA will be the system Sgr (P = 138 d, d=513 pc) with a A type low mass supergiant, is the brighter member of the type of extremely hydrogen-deficient binary stars (HdB stars). The HdB stars are a rather rare class of evolved binary systems that are in a second phase of mass transfer were the primary has ended the core helium burning phase. In Sgr spectrum, the H line appears with a complex emission profile. The variations observed in this profile do not appear to have direct connection with the orbital motion of the system and a possible explanation is that it is associated with a non uniform shell surrounding the system (Frame et al., 1995). Circumbinary material is also evidenced by a strong infrared excess and in particular the prominent silicate dust signature at 9.7 µm (Fig. 15). Adopting the parameters from Dudley and Jeffery (1990) the orbital radius is found to be about 2 mas and the binary could be easily resolved with the CHARA array. Proc. of SPIE Vol R-17

18 A OAU 1AU Figure 15: Large scale structure of the circumbinary envelope of Sgr as proposed by Nariai (1967) WR stars CHARA with its unique combination of long baselines and observations in visible offers a unique opportunity to access to new stellar spectral types, among which Wolf-Rayet stars. The closest WRs are typically at 1 kpc, and their 'pseudophotosphere' radii ranges from 3 to about 25 solar radii, that scales their continuum. The closest system, 2 Velorum was observed in 1968 with the Narrabri Intensity Interferometer (Hanbury Brown et al. 1970) with baselines reaching 190 meters. Since then, no WR were observed until the observation of extended dust environment around WR stars with the Keck using aperture masking (Tuthilll et al. 1999, Monnier et al. 1999) and in mid-ir with MIDI/VLTI (Danchi and collaborators). The VLTI/AMBER instrument has observed again 2 Velorum almost 40 years after the Narrabri Interferometer but, as the instrument operates in K with baseline limited to 130m, the WR star can hardly be resolved. Hence, to date, the Narrabri observations of a resolved line forming region in a WR star are unique. Determining the extension of the lines formed in the wind of these stars would provide a direct constraint on the physical conditions in the wind. This can be done in an absolute manner (provided that the distance of the object is well determined), or in a relative one by comparison with the continuum. For lines in the visible, baselines ranging from 100m to 300m are mandatory to provide an effect detectable and measurable accurately for the closest WR stars. Hence this study is perfectly well suited for VEGA. Observing the line forming regions of these objects with long-baseline optical interferometry offers the opportunity to their wind in the following ways: 1. Set an upper limit for the angular size of the forming regions for different lines. 2. Perform a differential study of line extensions compared to continuum, 3. Examine the deviation from sphericity of these objects, 4. Examine the deviation from purely radial motion of the wind of these objects, Over the past years, the modeling of Wolf-Rayet winds has reached maturity. A code like CMFGEN (Hillier & Miller 1998) is now able to solve the radiative transfer equation in the comoving frame (assuming spherical symmetry and a steady state), and is capable of handling line and continuum formation, both in regions of small and high velocities. WR outflows are optically thick up to a few stellar radii above the hydrostatic Proc. of SPIE Vol R-18

19 surface. The denser the wind, the larger the radius of the effective photosphere where photons escape, and the more so at longer wavelengths due to the increase in free-free opacity. Photons falling in spectral regions where they experience line opacity will escape at still larger radii than photons experiencing exclusively continuum opacity. Such optical depth effects lead to larger ``effective photospheres'' in spectral regions occupied by lines. Despite their compact structure, this makes the angular size of WR stars in the corresponding spectral bands comparable or larger than that O stars. Hence, the CHARA observations can directly be compared to quantitative predictions. As an example, the observations from the Narrabri Interferometer can be compared to a good model of WC8, the spectral type of the WR component of WR (provided that the distance of the object is well constrained). Hanbury Brown and collaborators provided an angular diameter of CIII4647 line forming region of 2.05+/-0.19mas. In Fig. 16, the extension of this line, computed by the CMFGEN code is shown (courtesy of L. Dessart). In visible, the brightest and most extended lines are the carbon lines for WC stars (CIII 4647, 5696, CIV 5805) and mostly Helium lines for WN stars (HeII 4686, 5412, 6560, 6683, He I 5876, CIV 5805). 0 La P'oI Flux[arbftraiyunits] Figure 16: Model of line formation for the WR component of 2 Vel (Courtesy from L. Dessart). Bottom: grayscale image of the quantity pi(p) as a function of p, impact parameter in cylindrical coordinates and scaled wavelength x = ( / 0-1)c, where p is the impact parameter (in units of the hydrostatic radius R*), I(p) the specific intensity along p (at x). 0 corresponds to the rest wavelength of CIII4646 and c is the speed of light. The overplotted thick black curve gives the line-of-sight velocity location where the integrated continuum optical depth, along z and at a given p, equals 2/3. Top: line profile flux, directly obtained by summing pi(p) over the range of p. Thus, the line flux at x in the top panel corresponds to the cumulative sum of all contributions pi(p) at x shown in the bottom panel, giving a vivid illustration of the sites at the origin of the observed line profile and their extension. Proc. of SPIE Vol R-19