Kennicutt (1992) Galaxies are shown in order of increasing Hubble type from top to bottom. Continuum: sharp break and rise after H and K lines (hence brighter in red than blue); dominated by red star light, strong absorption lines; E and S0 have K-star type spectrum  Flattening of continuum (reduction of red light compared to blue light) Spirals: F-K stars dominate spectrum  Decrease in strength of absorption lines  Increase in emission lines and strength of emission lines; A-star dominate spectrum but also - lots of SF, many O, B stars: - abs. lines of He (typical of O, B stars) - em. lines of ionized gas H,K G Mg Na OII OIII H ,NII Hβ NeII Hγ He SII Spectra of galaxies from E  Im:

AST 3003 MEASURING GALAXY REDSHIFTS Notes from Anthony Fairall – based on chapter 2 on Measuring Galaxy redshifts of his book “Large-Scale Structures in the Universe”, 1998, Wiley- Praxis Series in Astornomy and Astrophysics Most galaxies are distant, small angular size, only slightly larger than the ‘seeing’, with light centrally concentrated – that light is captured and fed to the slit of a spectrograph.

The telescope optics direct the light of the distant galaxy onto the slit. The width of the slit is a compromise between the seeing and the resolution of the spectrograph. A very narrow slit gives better resolution, but may only capture a fraction of the galaxy’s light. Usually the slit width is matched to the general seeing. The Collimator then converts light passing through the slit into a parallel beam. A miniature Cassegrain telescope, used in reverse mode, is a favoured design for a collimator. The focal ratio of the collimator must match that of the telescope.

Many spectrographs use reflection gratings, which work by the constructive interference of light: nλ = a(sinα + - sinβ) where n is the order number Angle θ is the blaze angle, which concentrates the light in a general direction, favouring one or more of the orders. The beams reflected off the grating are intercepted by a camera, which reverses the role of the collimator by bringing the parallel beams back to focus. The image of the slit is smeared sideways to form a spectrum. Slit width and dispersion determine the resolution Δλ, which can also be expressed as Resolution: R = λ/Δλ

The final image is focussed on a CCD detector –> a high quality version of those used in everyday cameras. CCDs have a matrix of potential wells within a silicon chip. These wells – each a pixel - trap photo-electrons released by incoming light. By moving the wells like a conveyor belt, the pixels are led to the corner of the CCD where the content of each is assessed. Special elongated CCDs accommodate extended (effectively one-dimensional) spectra. CCD’s are very efficient: 75% efficiency of light compared to 5% of photographic plates BUT; low sensitivity in blue

Multi-fibre spectrographs enable the spectra of many neighbouring galaxies to be recorded simultaneously. (SALT uses mulitple slits to allow for MOS-Mulit-object spectroscopy) They use optical fibres to isolate light from different galaxies in the focal plane of the telescope. One technique is to locate the image in the focal plane in register with that of a metal plate, prepared with drilled holes. Fibres are plugged into the holes. The 2dF/6dF systems (British-Australian) have the fibres connected to magnetic buttons (with miniature prisms). A robot sets up the field beforehand by drawing out each fibre from its storage and placing the button on a metal plate exactly where the light of a particular galaxy will fall. One field can be set up whilst another is being observed.

Galaxies are relatively faint objects and integrations have to be extended so that enough photons are gathered. At low photon counts, photon noise is significant – a case of signal versus noise. Determine Dark Counts (or bias) at beginning or end of night In reality, the situation is also complicated by cosmic rays, which dapple the image. Fortunately their damage is confined to single pixels, and can be tidied up by software routines.  easily removed by having 2 (or more) exposures Response of pixels in CCD: not uniform  measure flats

The pixel elements of a CCD do not have identical efficiencies in recording photons. Flat field exposures (e.g. normal white light) enable the relative sensitivities to be determined.

Calibration arcs – from low-density argon, neon or helium lamps – bracket the galaxy observations. The wavelengths of the arc lines are known precisely. Wavelength and pixel number are reconciled in a computer reduction, normally by an iterative fit to a polynomial relation. After that the galaxy spectrum can be re-binned in wavelength intervals.

Radial velocity standard stars, often observed in the twilight, also serve as a check. Spectra have good signal to noise and reveal the principle absorption features.

Sky subtraction is very necessary for faint objects. the night sky – even without moonlight – has significant continuum towards the blue and prominent emission features (forbidden Oxygen lines) at 5577, 6309 and 6364 Angstroem The galaxy spectra following also exhibit modest redshifts, spectral features have been shifted to slightly longer wavelengths (Doppler shift).

While stellar content may be assessed, the main purpose in obtaining such spectra is to extract a redshift: z = (λ / λ 0 ) – 1 when this is small, it can be converted into a velocity of recession via V = cz However, if the velocity is relativistic, the correct relationship is

The convention however is to publish cz, but confusion exists as some software packages apply the relativistic correction.

The other convention is to publish heliocentic velocities (i.e. correcting for the Earth’s orbital motion of 30 km/s). Further corrections may be made. In his ‘Reference Catalogues’, de Vaucouleurs advocated correcting to the local frame of rest within the Local Group of galaxies via V 0 = V sin l cos b where l and b are the Galactic longitude and latitude, though there is some debate as to this correction. By contrast, radio astonomers are fond of correcting to the ‘Local standard of rest’ after removing the Sun’s motion relative to neighbouring stars, which is believed to be 19 km.s towards RA 18 h, Decl. = +30 0

The modern day alternative to making manual measurements of the wavelengths of redshifted spectral features is to conduct a mathematical ‘cross-correlation’. To do this, the spectrum is first re-binned into channels of equal width in ‘log wavelength’, such that the Doppler shift is the same for all channels. Large- scale variations are then filtered out (i.e. only spectral lines left). The spectrum is then ‘slid’ against a template prepared from a known bright galaxy or star, and a correlation function derived. The peaks in the function show where the spectrum best fits the template. The highest peak is almost certainly the correct redshift.

Emission-line galaxies have wonderfully sharp spectral features that give precise redshifts. Again the correlation technique can be used, but with an emission-line template.

Except for very nearby galaxies, the redshifts normally reveal the cosmological expansion of the Universe, as reflected in Hubble’s law. (Licence has been taken in labelling it a ‘Doppler’ shift as that is not true, except it is equivalent to a Doppler shift, as the galaxies are moving away from our Galaxy.) However, peculiar velocities may add to or subtract from the cosmological velocity. The spectra also reveal the nature of the dominant optical emission from the galaxy. Usually it is stellar, i.e. absorption lines. Average spectral types are usually late G-early K, though a few ‘early-type’ spectra do appear. But traces of weak emission may show up in <50% of galaxies. About 10% of galaxies show strong emission lines, that light originating from nebular, not stellar, sources. But these usually have strong blue continua from hot young stars – the source of excitation. About 1-2% of galaxies reveal broad emission lines – of active galactic nuclei, the topic covered in a following section.