A stellar long-baseline interferometer consists of an array of several separate telescopes, which redirect starlight to a central location where interference fringes are formed. The angular resolution that the interferometer can achieve depends on the baseline (the telescope separation), while for a single telescope the resolution is defined by its aperture diameter. The angular resolution is proportional to the ratio of the wavelength to the baseline of the interferometer or the diameter of the telescope. The largest existing telescope is 10 m in diameter, whereas the largest interferometer has a baseline of 600 m, giving the latter a factor of 60 increased resolving power.
If the interferometer is coupled with a spectrograph, fringes can be analysed along spectral domains. Such a spectro-interferometer allows for eliminating disturbances by the Earth atmosphere and measure differential phases between two wavelengths, i.e., carry out colour differential interferometry (Petrov, 1988). For magnetically active stars such observations can provide orientation of the stellar rotation axis with respect to the line of sight (Lagarde et al., 1995), which is important for interpretation of spectroscopic and photometric observations of spotted stars. Also, differential fringe phases can be used to locate spots on the stellar surface (Wittkowski et al., 2002; Jankov et al., 2003; Rousselet-Perraut et al., 2004), even in such difficult cases as pole-on and equator-on orientations of the star which cannot be resolved by traditional Doppler spectroscopy. In addition, the stellar differential rotation rate and the inclination angle can be disentangled (de Souza et al., 2004).
A spectro-interferometer equipped with a polarimeter allows for recording the visibility function and fringe phase for different polarisation states and, thus, can provide measurements of local magnetic fields on the stellar surfaces (Rousselet-Perraut et al., 2004), which is the key information for studying stellar magnetic activity.
A list of the available and planned for the near future instruments is given in Table 3.
|GI2T/REGAIN||Nice||R = 28 500 – 35 000*||400 – 850 nm|
|AMBER||VLTI/Paranal||200||R = 10 000**||600 – 2500 nm|
|VINCI||VLTI/Paranal||200||8 – 13/17 – 26 µm|
|Keck I/II||Hawaii||85||5 – 24 mas||1 – 10 µm|
|OHANA||Hawaii||800||0.5 mas||K band|
|0.3 mas||J band|
|SI||Space based||0.06 & 0.12 mas||155 & 280 nm|
|Darwin||Space based||20 – 200||6 – 18 µm|
© Max Planck Society and the author(s)