Ground-based Observations

In this work we have made use of data collected at a number of ground-based opti- cal telescopes, and as such we will provide a description of how these instruments function in general. In a ground-based spectrometer, light is collected by the main telescope before being sent into the spectrograph, and the amount of light col- lected by the telescope — hence the spatial resolution and faintness detectable — is dependent on the size of the primary mirror. The main ground-based telescope we have used is the Plaskett telescope at DAO (Plaskett, 1927), which was planned to be, but never was, the largest telescope in the world. This telescope has a 78” primary mirror with a spectrograph mounted at the Cassegrain focus. We have also made use of data collected at the James Lick Telescope (Holden, 2008), a 38” refracting telescope (the largest refracting telescope in the world up to 1897) outside San Jose, California. Finally, we also made use of data collected at Calar Alto Observatory (Els¨asser, 1981; Graser & Hopp, 1991), Granada, using the 87” telescope. Most data was collected by R.E.M. Griffin, or in some cases taken from the DAO archive. All observations were reduced, and data provided, by R.E.M. Griffin. An observation log is provided in Chapter 5.

Light from the main telescope, once collected at the primary mirror, is sent into the spectrograph, which is placed at the focus of the instrument. In Fig. 3.1

Mirror

Slit

Telescope

Light from

Collimating

Diffraction

grating

Camera

Mirror

Detector

Figure 3.1: Diagram of a spectrograph with a reflecting diffraction grating. we see a general schematic of a spectrograph. Light enters through a slit, and a collimating mirror transforms the incident beam into parallel rays. At this point, the rays are directed onto a diffractor. A diffraction grating is an optical element with narrow, closely spaced lines at approximately the wavelength of the light, and the grating may transmit the light (a transmission grating), or reflect the light (a reflection grating — as seen in Fig 3.1). An echelle grating may also be used, which is a variant of the diffraction grating with a relatively low groove density, but a groove shape which is optimized for use at a high incidence angle. Echelle gratings are highly blazed, and hence the light is dispersed into high orders, with some overlap. A prism is then required to split the overlapping orders.

Once the light has been split into its spectral components it is focussed onto the detector. Historically, and for some of the observations used in this thesis, the detector was comprised of photographic emulsion (in our case Kodak IIa- 0). This photographic exposure has the disadvantage of reacting non-linearly to

incident radiation, making it difficult to calibrate spectra to an absolute flux scale. However the emulsion has the advantage of being physically robust, once the image is chemically fixed, and it is well stored, it will remain unchanged for decades. Hence the archival DAO data used in this thesis, some of which was collected by K.O. Wright in his early observations ofζAurigae in the 1950’s, has not declined in quality. These data have been digitized by R.E.M. Griffin as part of a long running campaign, and made available for our use. In the modern era photographic plates have been replaced with Charge-Coupled Devices (CCDs). CCDs react linearly to incident radiation, and have a very high quantum efficiency (_∼95%), allowing high signal-to-noise observations on an absolute flux scale. The CCD is comprised of a photo-reactive layer of silicon divided into columns by insulating boundaries called channel stops, and rows by strips of aluminium, creating the individual pixels. The (p-doped, MOS) silicon layer is biased above the threshold for inversion and hence incoming photons will free electrons as a result of the photoelectric effect. This charge is held in place by the potential well created by positively charged electrodes (the electrode positioning is what defines the position of the pixels). The pixels are then read out row by row, with each pixel passing its charge to its neighbour, and the final pixel in each row of the device passing its charge to a charge amplifier. The resultant voltages are then stored as a digital image.

Traditionally in ground-based spectroscopy line-profiles were very difficult to measure accurately as a result of the limitations of photographic plates, seeing, and the faintness of the object (requiring that the spectrograph’s slit be open fully) reducing the spectral resolution. In order to overcome these limitations and perform accurate quantitative spectroscopy astronomers measured the Emission Equivalent Width (EEW) of spectral lines, and the EEW will be the subject of our study. It is computed from the observed spectrum as

W = Z λmax λmin 1₋Fcont Fλ dλ (3.1)

whereFλ is the measured flux, and Fcont is the local continuum flux. The EEW is the width of a feature (hence it has units of wavelength) whose total intensity is zero, and whose total flux deficit is equal to that of the line. This value is unaffected by the spectral resolution of the instrument as flux in a given wavelength band is

Figure 3.2: Diagram of the Hubble Space Telescope. This diagram is post- installation of the STIS instrument, which took the place of GHRS. Image Credit: NASA.

conserved. The EEW is often used in conjunction with curve-of-growth analysis, and this has been a fruitful area in the study of ζ Aurigae binaries. In particular, the analysis of Wilson (1957) of EEW’s measured at DAO.