VEGA: a Visible spectrograph and polarimeter for the VLTI Science Cases Description

Transcription

1 VEGA: a Visible spectrograph and polarimeter for the VLTI Science Cases Description Philippe Stee a, Denis Mourard a, Daniel Bonneau a A. Domiciano de Souza c, Renaud Foy d, Stephane Lagarde a, Jean-Baptiste Le Bouquin e, Philippe Mathias a, Romain G. Petrov b, Karine Rousselet-Perraut e, Gerd Weigelt c a Département GEMINI, Observatoire de la Côte d'azur, BP 4229, F Nice Cedex 04, France; b Université de Nice-Sophia Antipolis, Parc Valrose, F Nice Cedex 02, France; c Max Planck Institute für Radioastronomie, Auf den Hügel 69, D Bonn, Deutschland; d Centre de Recherche Astronomique de Lyon, F Saint Genis-Laval Cedex, France; e LAOG, 414, rue de la Piscine, Domaine Universitaire Saint-Martin d Hères, France. 1. Introduction Interferometry has been intensively done at large 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 said, less difficult to recombine directly the optical beams since the coherence length is larger and the turbulence slower. Nevertheless, 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 and also with 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 Grand Interféromètre à 2 Telescopes (GI2T) in southern France. This instrument, in his VLTI adaptation, called VEGA (Visible spectrograph and polarimeter) will open new fields in a wide range of Astrophysical topics only possible in the visible domain. It will provide a spectral resolution up to at 0.6 µm and a spatial resolution of less than 1mas for up to 4 telescopes in his X-λ special configuration. A polarimetric device (SPIN) measuring simultaneously the polarization in 2 directions either circular or linear is also implemented in this instrument. A multiple band-passes mode (so called COURTES mode) is also implemented for wide field observations in 4 simultaneous wavelengths for 3 telescopes which will open the extragalactic domain up to magnitudes 13 in the visible. Since VEGA was already tested on the sky on 1.5 m telescopes it is also very well suited for the 1.8 m VISA array on the VLTI and will only need minor adaptations for the injection of the VLTI 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 two major modes of observations (see Vega Concept Study Report for more details): A high spectral resolution operation with a dispersed fringe mode (called X-λ): o Simultaneous combination of 4 telescopes, giving access simultaneously, for one measurement, to 6 squared visibilities, 6 differential phases and 3 closure phases. o Spectral resolution: 1500, 5000 and and spectral range: from 0.5 to 0.9 µm (if possible from 0.4 to 1.0 µm) o Polarimetric capabilities allowing to simultaneously recording the interferograms in two polarizations, either linear or circular. o 3 in the slit direction with the ATs. A high sensitivity operation with a multiple band passes mode (called COURTES) o Simultaneous combination of 3 telescopes o Spectral resolution of 100 to 2500 and spectral range: from 0.5 to 0.9 µm (if possible from 0.4 to 1.0 µm) o 4 simultaneous field of view of 1.5 x1 (at four different wavelengths for the ATs).

2 VEGA: a visible spectrograph and polarimeter for the VLTI 2 In multimode operation, we usually consider two domains for the signal to noise ratio, depending on the number of photons per speckle and per single exposure. If this number is larger than 1, then the signal to noise ratio increased as the square root of the number of photons, whereas if it is smaller than 1, the signal to noise ratio increased as the number of photons (Roddier, 1988). In consequence, we first defined the limiting magnitude as the magnitude giving 1 photon per speckle and per single exposure. The signal to noise ratio is then equal to one 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 be very high at this limit, since one can integrate a large number of single exposures containing a large number of speckles. 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, thanks to the fact that we can have more than 1 photon per speckle and per frame in the reference channels. 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 a. Q VLTI =0.1 (see figure of ICD Issue 3.0 where we have also supposed that the M9 is not a mirror but a dichroic with T=R=50% in order to equally aliment STRAP and VEGA). b. Q Instrument =0.4 ( grating for the chromatic OPD correctors) c. Q Detector =0.3 d. Q Total =1.2% 3. Exposure time t 0 =20ms 4. r 0 estimations (for 650 nm) at Paranal a. Median seeing 0.8, r 0 =12.6*(650/500) 6/5 =17.3cm b. Best conditions (20% of time) seeing 0.5, r 0 =27.7cm c. Excellent conditions, seeing 0.3, r 0 =46.1cm Figure 1: limiting magnitude for 1 photon per speckle and per frame and for a 50 nm spectral band. This limit is independent of the kind of telescopes and is only defined by the area of a coherence cell. The calculations are made for a wavelength of 650nm. This magnitude corresponds to the value where the signal to noise ratio is equal to 1 for a single exposure and per speckle.

3 VEGA: a visible spectrograph and polarimeter for the VLTI 3 Median Good Excellent Seeing r 0 (650nm) 17.3 cm 27.7 cm 46.1 cm V (1ph/sp/fr) for a 50 nm spectral band Time for 1% in V (COURTES) for the above magnitude 400s 1000s 2850s Time for 5% in V (COURTES) for the above magnitude 16s 40s 110s Time for 1% in V (X-λ) for the above magnitude 6320s 10000s 16900s Time for 5% in V (X-λ) for the above magnitude 253s 400s 680s Table 1: Exposure time calculations for 1% and 5% accuracy in both, COURTES and X-λ, modes for the different typical seeing conditions in Paranal. The calculations are made for a differential measurement with a reference channel of 50 nm and a science channel of 0.1 nm. This duration are calculated for the limiting magnitude where the number of photons per speckle and per frame in the reference channel is equal to 1. These results show that, under good seeing conditions, a V=11.2 source, could be measured with a 1% accuracy in a spectral channel of 0.1 nm in the COURTES mode in 20 mn of integration. Under excellent conditions, the same measurement could be achieved on a V=12.3 star in 45mn. In order to estimate the ultimate capabilities of the instrument, we have also consider a measurement with a wide spectral channel of 50 nm and a science channel of 1 nm, quite well adapted for the large emission lines of active galactic nuclei in the visible domain. Median Good Excellent Seeing r 0 (650nm) 17.3 cm 27.7 cm 46.1 cm Time for 1% in V for a V=15 object 76h 30h 11h Time for 5% in V for a V=15 object 3h 1.2h 26mn Table 2: Same hypothesis as table 1 but for a science channel of 1 nm instead of 0.1 nm. We calculate the exposure time required for a V=15 source, well fainter than the limit of 1 photon per speckle and per frame. The regime of signal to noise ratio is different and the numbers are quite larger. However, under good conditions and if one accept 5% of visibility accuracy, such sources are reachable. This last table shows that sources as faint as V=15 could be observed with VEGA in the COURTES mode, with a reference channel of 50 nm and a science channel of 1 nm, if one accept an accuracy of 5% in visibility measurements. At this level of magnitude a large number of extra galactic sources are reachable and unique science, even with 5% accuracy on the visibility, is possible. 3. Science Drivers for a Visible VLTI Instrument 1. Stellar activity: Interferometry can contribute to the study of stellar activity by detecting and mapping star spots. 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(λ,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

4 VEGA: a visible spectrograph and polarimeter for the VLTI 4 method, concerning the input physical parameters (stellar rotation, temperature differences...) and also the uncertainty in the latitude distribution. Figure 2: Numerical simulation of the red supergiant Betelgeuse surface from Freytag 2005 see URL: ( 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 questioned. Interferometry allows the combination of s(λ,t) with the photocenter variation with wavelength and time ε(λ,t) deduced from the fringe phase φ(λ,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(λ,t) and on ε(λ,t), obtained with the same telescopes and instruments: SNR[s(λ,t)]/SNR[ε(λ,t)] Φ * /(λ/b) where Φ * is the stellar angular diameter and (λ/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 (λ/d)>> Φ *, and why it should perform fairly with the

5 VEGA: a visible spectrograph and polarimeter for the VLTI 5 ATs in the visible. It will be possible to adjust the baseline to have Φ * 0.7(λ/B) for many systems (if Φ * is too close to (λ/b), the fringe contrast is destroyed). Exemple: Spots on magnetically active stars Scientific objectives Energy of late type stars is mainly carried outwards by convection whose study is a key in modelling the inner structure as well as the atmosphere of these objects. Mapping the turbulent eddies of such stars and deriving their size and luminosity (or temperature) is of strong interest to better understand stellar convection. Late type stars might exhibit asymmetric and even highly fragmented mass-loss events, which are believed to be linked to stellar surface parameters such as limb-darkening, effective temperature or surface features. Observing Strategy Very High Angular Resolution is mandatory: the longest baselines with the ATs and the shortest wavelengths are used. Accurate visibility measurements in the small visibility range (i.e. in the second lobe of the visibility function) is required. Closure phase, and better, phase measurements are powerful to detect asymmetries and surface spots. Precursor/supporting observations First results on limb darkening determination of late-type giants have already been obtained with NPOI (Wittkowski etal., 2001) and with VLTI/VINCI (Wittkowski etal., 2004). AMBER observations are expected in the next years. Requirements - Spectral Resolution : no requirement - Spectral coverage : shorter wavelengths are required to improve angular resolutions. Note also that stellar chromospheric activity of late type stars is studied via the calcium triplet at λ8542 A. - Polarization : no requirement but probably useful - Visibility accuracy : Accuracies as high as AMBER ones are required. Target statistics Wittkowski et al. (2002) produces a list of more than 70 targets within VLTI observation context. Their list mainly includes RS CVn type, BY Dra type, W UMa type, FK Com type, T Tau type and single giants. Estimated angular diameters range from 0.1 mas up to 3 mas. The authors only select 4 targets with the selection criterium of angular diameter larger than 2 mas to be able to put 2 2 pixels in the images in the nearinfrared range. Observing in the visible range allows to increase the target number of the Wittkowski s sample from 4 up to 17. Exemple of stellar atmosphere studies : Abundance and magnetic field of Ap stars Scientific objectives 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 difficulty 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 α2cvn) and very few maps of magnetic fields have been reconstructed via Zeeman-Doppler Imaging (e.g.

6 VEGA: a visible spectrograph and polarimeter for the VLTI 6 Kochukhov et al. (2002) for α2cvn) or by inversion of spectro-polarimetric data (Bagnulo et al. (2000) for β CrB). Moreover, such inversion methods often lead to several magnetic field models that cannot be disentangled by classical spectro-polarimetric techniques. Observing Strategy 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). To determine magnetic field topologies, color-differential interferometry has to be applied in magnetically sensitive lines, through the circularly-polarized components of the Zeeman pattern. As an illustration, Figure 3 displays the SPIN signals for two Bagnulo s models of magnetic topology of the prototype βcrb (Bagnulo et al., 2000) plotted in Figure 4. The SPIN signals correspond to the fringe phase difference between the blue and red wings of the spectral line observed in the left circularly polarized light. The two models can be disentangled from the SPIN signals, even with the two telescopes and the 50-m baseline of the GI2T Interferometer. Obviously the SPIN signals will be larger with the larger baselines of the VLTI and using 4 telescopes simultaneously within this context of stellar surface imaging will be more efficient. Figure 3: Top: SPIN signals (see text for explanation) computed for two magnetic field models of β CrB (Bagnulo et al., 2000) versus stellar phase (left) and versus baseline angle (right). Bottom: SPIN signals versus baseline angle computed for Model #1 (left) and for Model #2 (right). The dashed lines correspond to the top cuts. Computations are performed for their online Fe λ5018a with the GI2T/REGAIN parameters: spectral resolution of 30000, baseline of 50 m and observability of ±2h around the transit (dashed are as on the bottom maps).

7 VEGA: a visible spectrograph and polarimeter for the VLTI 7 Figure 4: Magnetic topologies of βcrb for different stellar phases reconstructed by Bagnulo et al. (2000) from spectro polarimetric data. The first model is mainly dominated by a dipolar component while the second one is close to a quadrupolar configuration. Local magnetic field strength is displayed in gray scale and magnetic field direction is displayed by a vector pointing toward the surface (blue) or outwards (red). Note that both models have a similar inclination (star seen by the southern pole) whereas their azimuths significantly differ. Requirements: Spectral Resolution: Model atmospheres, abundance maps, magnetic topologies are mainly constrained by spectro-(polari)metric observations at high spectral resolution and via inversion techniques. Coupling high spectral resolution with high angular ability is obviously of strong interest for adressing stellar activity issues. A spectral resolution of is enough even for slow rotators like βcrb. For faster rotators, the medium spectral resolution can be convenient. Spectral coverage: CP stars exhibit numerous and deep metallic lines in the visible range. Moreover the shorter the wavelength the better the angular resolution. Measurement accuracy: An accuracy on the differential phase between two close spectral channels better than 1 is required. Target statistics We focus our target search on the Renson catalog General catalog of Ap and Am stars (Renson et al., 1991). We select: peculiar stars only V limiting magnitude smaller than 8.0 (to consider realistic conditions of medium spectral resolution mode) a declination ranges of [-84 ; +36 ] We compute the histogram of the distribution (Figure 5) without taking into account for: - angular diameter (generally unknown) - magnetic field strength (which varies from a few hundred Gauss up to a few thousand Gauss over our target selection) - abundance inhomogeneity knowledge

8 VEGA: a visible spectrograph and polarimeter for the VLTI 8 Figure 5: Histograms of number of CP stars vs. V magnitude for a declination range of [-84 ; +36 ]. There are 133 targets. Note that such histograms need to be detailed interms of magnetic field strength and/or abundance inhomogeneity scale since interferometric signals would be directly proportional to these quantities. As an example we plot in figure 6 the histograms of Hmin and Hmax for targets whose magnetic field strength has been measured. We clearly show a large spread of magnetic field values, and that targets for which high magnetic field values are known is obviously less numerous. Figure 6: Histograms of number of CP stars vs. magnetic field strength, Hmin (left) and Hmax (right). There are 67 targets for the first case and 37 for the second one. 2. 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

9 VEGA: a visible spectrograph and polarimeter for the VLTI 9 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.7 left), which give rise to non-symmetrical intensity maps (Fig. 7 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.7 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 as 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 ρ, i.e. <10% of the photocenter amplitudes expected for differential rotators. Figure 7: Left: Radial velocity maps for several combinations of i and α, assuming a solar-like differential Veq rotation law Ω ( l) = (1 α sin 2 l), where l is the latitude. A positive velocity corresponds to R displacements towards the 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 (α=0). Center: Intensity maps inside a photospheric line for i=45 and α = 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, ε z (top; vertical direction) and ε y (middle; horizontal direction), and the normalized spectral flux (bottom). Photocenter components are given in units of angular stellar radius ρ. The letters indicate selected wavelengths corresponding to a region in the blue line wing (λ a ), the highest ε y value (λ b ), the central wavelength (λ c ), and the most positive value of ε z (λ d ). Note that ε z =0 for uniform rotation. Figures from Domiciano de Souza et al. (2004).}

10 VEGA: a visible spectrograph and polarimeter for the VLTI Asteroseismology Asteroseismology studies the temporal variation of the photometric flux or of the spectrum s(λ,t) to identify the resonant modes of stars (Fig. 8). 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 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 ε(λ,t) to s(λ,t) solves many problems. The demand for instrumental stability is relaxed, since for example it is possible to calibrate ε(λ,t) using as a reference s(λ,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 ε(λ,t) permits to access higher modes. The fact that many modes affect very differently s(λ,t) and ε(λ,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(λ,t) and ε(λ,t) if far more discriminating than the use of spectroscopic data alone (see Fig. 9). Figure 8: 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 β Cephei, SPB, δ Scuti, roap, γ 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 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).

11 VEGA: a visible spectrograph and polarimeter for the VLTI 11 Figure 9: 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). The technique could be applied to solar type stars with a dedicated instrument optimized for this type of observations with the ATs: 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 relatively simple VEGA extension to the visible would yield access to δ 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 visible VLTI instrument as much as possible. 4. AGNs Studies of the morphology and kinematics of AGN is a key objective of infrared and visible VLTI observations. Classical interferometry will allow studies of structures in the inner NLR (e.g., possibly torus or jet structures), whereas differential interferometry (e.g., measurements of the wavelength dependence of the visibility function with high spectral resolution) will provide results on the size, structure, and kinematics of the BLR. The observational technique is similar as described in the section of Be star envelopes. The important advantages of interferometry in the visible are the higher resolution, the higher emission line signal and the photon-counting detectors. For the BLR studies by differential interferometry, a very preliminary SNR estimate: SNR flux. (Baseline/λ). Line/Continuum. It indicates that the ATs in H a would do about eight times better than the UTs in Brγ. These SNR values are given for single mode operation. The Strehl ratio in the visible with the ATs is assumed to be one half of the value in the K band with the UTs. This allows much more precise ε(λ,t) measures and much stronger constraints on the BLR kinematics and therefore on the central mass or fainter limiting magnitude for a given accuracy.

12 VEGA: a visible spectrograph and polarimeter for the VLTI 12 A visible instrument working with the current installed telescopes at Paranal (telescopes+ao) must be operated in the multimode regime. Under these conditions, UTs will be only 4 times more efficient than the ATs and the limiting magnitude, defined as the limit where one get 1 photon per speckle and per single short exposure, does not depend on the telescope but only on the seeing conditions. Therefore at this limit, differential measurements will permits high signal to noise ratio by simply integrating a large number of single exposures containing a large number of individual speckles. Calculations have been made showing that under good seeing conditions (0.5 ), with a wide reference channel of 50 nm and a science channel of 0.1 nm, a 1% accuracy could be achieved on a V=11.2 source in 20mn. Under excellent seeing conditions (0.3 ), 45mn are necessary on a V=12.3 sources. If one consider a wide reference channel of 50nm and a science channel of 1nm, a V=15 source could be measured with a 5% accuracy in 1.2 h under good seeing conditions (0.5 ). This time go down to 26 mn under excellent seeing conditions (0.3 ). 5. 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 and the luminosity. 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 and luminosity 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 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: - for a single star with accurate parallaxes (thanks to Hipparcos and Gaïa) - for stars component of resolved spectroscopic binary. Figure 10: Radii and masses of the four very-low-mass stars now observed with the VLTI, GJ 205, GJ 887, GJ 191 (also known as "Kapteyn's star") and Proxima Centauri (red filled circles; with error bars). For comparison, planet Jupiter's mass and radius are also plotted (blue triangle). The two curves represent theoretical models for stars of two different ages (400 million years - red dashed curve; 5 billion years - black fully drawn curve. The stellar effective temperature can be compute 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 2002). A gain in angular resolution of about 4 is

13 VEGA: a visible spectrograph and polarimeter for the VLTI 13 expected by observing in the visible (R band) with the VLTI and it will be possible to measure the angular diameter of main sequence stars of spectral type ranging from O to M in the stellar neighbor. 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. Although double stars are complex objects both from the observational and modelling point of view, their study provide fundamental data to refine the theory of single stars formation and evolution 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 disk 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 close binaries, additional physical processes have to be taken into account due to the interaction between the two components. Tidal interaction, mutual reflection, mass transfer, or mass loss are the kind of processes which characterize such systems. The structure of this system presents several components: mass stream, accretion disc, scattering envelope and jet-like structures. Spectro-interferometry in the visible with VEGA on the VLTI, combined with classical techniques like photometry and spectroscopy will provide deep insight on the morphology of the star and its circumstellar environment in order to constrain the modeling of such phenomena. 6. Cepheids: distance scale and pulsating atmospheres Interferometry can improve the accuracy of the 0 point of the period-luminosity in Cepheids, which is a fundamental parameter for the distance scale in the Universe. This is based on simultaneous measurements of cyclic radial velocity curves and angular diameter variations. It has been shown (Mourard, 1996) that, the gain in resolution permitted by the visible substantially increases the precision of the diameter measurement and therefore we obtain access to a much larger sample of Cepheids. Figure 11: VINCI observations of the pulsation of the Cepheid variable L Car (P = 35.5 days, red dots) and the adjusted radius curve (green line), as deduced from the integration of the radial velocity measured on this star over its pulsation period. The visible band is particulary well adapted to continue the work on Cepheids, already started with the VLTI/VINCI and in a near future with AMBER. The main limitation for the Period Luminosity calibration is the number of stars for which a direct distance determination could be determined by directly measuring with a good accuracy the amplitude of variation of the angular diameter. 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. The next figure shows a histogram of Cepheids as a function of the amplitude of the variation of the visibility as a

14 VEGA: a visible spectrograph and polarimeter for the VLTI 14 function of the pulsational phase. This figure clearly shows that the V band is the best spectral window for this program: Figure 12: number of Cepheids as a function of the amplitude of variation of visibility during the pulsation cycle. The different colors correspond to different spectral bands. In the following figure, we have plotted the angular diameter of a sample of LMC Cepheids (in blue, by using a J,J-H relation and in pink by using a J,J-K relation). This figure shows that a few Cepheids are brighter than V=11, within the reach of VEGA. The expected performances of differential phase measurements permit to reach a 1µas 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 a few LMC Cepheids, giving for the first time access to a second level of the distance scale problem. More details could be found in Mourard, Figure 13: Angular diameters of a sample of LMC Cepheids as a function of their K magnitude. (in blue, by using a J,J-H relation and in pink by using a J,J-K relation)

15 VEGA: a visible spectrograph and polarimeter for the VLTI 15 In addition to the period-luminosity relation, many other important Cepheid projects performed, such as the study of surface brightness and velocity structure variations through the pulsation cycle and the subsequent improvement of the models of pulsating atmospheres. For the brightest Cepheids, the resolution with 200 m baseline in the visible is large enough to obtain detailed information on the surface structure. Measurements of the temporal variation of the limb-darkening profile allow studies of the stellar photosphere. Adding the spectral dimension permits to track the pulsation in the different layer of the atmosphere and study in greater detail its mechanism (Breitfellner & Gillet, 1993). Furthermore, studying in a line profile the variation of the spectrum s(λ,t) together with this of the photocenter displacement ε(λ,t) measured by differential interferometry allows the extension of the Cepheid study to unresolved objects. These two last applications require spectral resolution between and 40000, which is much easier to achieve in the visible (Petrov, 1989). 7. Mira stars and related objects High-resolution optical observations of Mira stars allow, for example, detailed studies of the stellar disk, of surface inhomogeneities, and of the wavelength and phase dependence of Mira star diameters (see, e.g., Bonneau & Labeyrie 1973, Labeyrie et al. 1977, Bonneau et al. 1982, Haniff et al. 1992, 1995, Quirrenbach et al. 1992, Weigelt et al. 1996, Hofmann et al. 2000). In strong TiO absorption bands (at 671 nm or 714 nm) the diameter of Mira variables is much larger than in the continuum (Bonneau et al. 1982). For example, the first diffraction-limited images of the Mira star R Cas obtained with the bispectrum speckle interferometry method and the Russian 6 m telescope have shown that the average uniform-disk diameters of R Cas are 49 mas at 714 nm (strong TiO absorption band), 37 mas at 700 nm (moderate TiO absorption), and 30 mas at 1045 nm (continuum). Furthermore, the 671 nm, 714 nm and 700 nm images of R Cas are non-uniform and elongated. For example, at 714 nm the size of the elongated R Cas disk is 42 mas x 56 mas (Weigelt et al. 1996, Hofmann et al. 2000). Adopting the HIPPARCOS parallax, a photospheric R Cas radius of 377 solar radii was obtained. The described observations were compared with the theoretical wavelength dependence of the diameter (and shape of the visibility) predicted by Mira star models (e.g. Watanabe & Kodaira 1979, Scholz 1985, Bessell et al. 1989, Bessell et al. 1996) to test the models. The asymmetric stellar disk of R Cas discussed above and the unusually high K-band visibility values of R Leo at high spatial frequencies (Perrin et al. 1999) observed with the IOTA interferometer and of R Cas observed with the GI2T interferometer (Weigelt et al. 2000) show that the stellar surface of Mira stars is more complex than previously thought. Huge convection cells (predicted by Schwarzschild) or bright hot spots are most likely responsible for the the unexpected visibilities. In future, both infrared JHK interferometry and optical interferometry in TiO bands are required to study the nature of the surface structures. 8. 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 α Ara with the MIDI instrument. Unfortunately the 200m baseline was not large enough to resolve the circumstellar disk at 10 µm. Nevertheless, these measurements put an upper limit to the envelope size in the N band of φmax = 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

16 VEGA: a visible spectrograph and polarimeter for the VLTI 16 stellar parameters we have used the SIMECA code developed by Stee (1995) and Stee & Bittar (2001). Optical spectra taken with the echelle 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 α Ara circumstellar disk. The possibility of larger distance that indicated by the Hipparcos parralax is discussed, as well as a potential truncation of the disk by a putative companion (see Chesneau et al. 2005, A&A, in press). It is clear that with an increase by a factor 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 α 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(λ)/V(λ ref ) or phase Φ(λ) Φ(λ 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. Figure 14: A schematic view of α Ara seen from recent VLTI/MIDI observations (Chesneau et al. 2005, A&A, in press) 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 and preliminary 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 circumstellar rotating disk around α Ara seen across the Brγ emission line (see Fig. 14). 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 α compared to Br γ will be decisive. The typical angular size for the disks of the closest Be star is about

17 VEGA: a visible spectrograph and polarimeter for the VLTI 17 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 4x4 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. Figure 15: 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 α 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 16 presents our current view of the HAe/Be star MWC 297 which has been modeled by combining an accretion disk model 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 Hα line profile (with an intensity 120 times the continuum!) will put very strong constrains on our modeling by, for instance, setting if Hα originates within the equatorial disk or in the spherical stellar wind.

18 VEGA: a visible spectrograph and polarimeter for the VLTI 18 Figure 16: 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 preparation). 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. 9. Other hot star envelopes Many aspects that can be studied for Be stars are interesting for Wolf Rayet stars as well as for Large Blue Variables. Quite some time ago the Narabry intensity interferometer has resolved the gas envelop of λ 2 Vel, but without being able to find where in this binary system the circumstellar material is located (Davis et al. 1996). Image information though the measure of a small number of phase closures as well as spectrally resolved differential interferometry would resolve this type of systems. The largest and brightest LBV star η Carinae has been extensively studied by speckle interferometry, adaptive optics and the Hubble Space Telescope to investigate its small-scale structure and try to understand better its rapidly varying huge mass loss. Observations show structures at all scales with important differences between the images at different wavelengths. Reconstructing interferometric diffraction-limited images at different wavelengths is a key to the understanding of this class of objects. REFERENCES Bagnulo, S., Landolfi, M., Mathys, G., Landi Degl Innocenti, M. 2000, A&A, 358, 929 Bonneau D., Foy R., Blazit A., Labeyrie A. 1982, A&A 106, 235