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Applications of Physics to Astronomical Systems
Lecture 10
17)
Astronomical Detectors
Almost all the information we have about astronomical objects comes from observations of various
forms of electromagnetic radiation. (The main exceptions being “cosmic rays”, which are high
energy particles, and recently neutrinos. The detection of gravitational waves and dark matter are
hopes for the future.) The wavelengths used cover an enormous range and the techniques employed
to detect the photons vary accordingly.
(i)
Extremely High Energy Gamma Rays ~1012 to ~1015 eV
These are produced by collisions in natural particle accelerators such
as pulsars. They interact with atoms high in the earth’s atmosphere
and produce a shower of particles. These are travelling faster than the
speed of light in the air and this produces a cone of light (called
Cerenkov radiation) somewhat analogous to a shock-wave from a
supersonic aircraft. This is seen as a pulse on the ground and can be
collected by large relatively crude “light-buckets”.
(ii)
Gamma Rays ~106 to ~108 eV
There is a flux of these from the milky way produced by collisions
between cosmic rays and interstellar atoms. More interesting are the
gamma-ray bursts which last for just a few seconds but are now
known to be from such distant galaxies that they are the most
energetic events yet observed. Such photons are stopped high in the
atmosphere so space satellites are required. The detection is based on
production of electron-positron pairs when the gamma passes close to
an atomic nucleus. The ionisation these produce can then be observed
and the energy and direction of the gamma-ray reconstructed.
(iii)
X - rays ~103 to ~105 eV
These are produced by a wide range of very hot or energetic objects
ranging from white dwarf stars to the accretion disks around massive
black holes. Satellites still have to be used as the atmosphere is very
opaque to X-rays. At these wavelengths (1 keV = 1.2 nm) it becomes
possible to make images by using grazing incidence mirrors. These
are “nested” to get increased collecting area.
The detectors
themselves are generally based on the ionisation produced by the
photons. The amount of charge is proportional to the energy of the Xray photon so this acts as a crude spectrometer. Higher resolution
spectrometers can be made by using diffraction, e.g. Bragg diffraction
from crystal lattices.
(iv)
UV, visible and near IR ~1 to ~10 eV (1.2 microns to 120 nm)
This waveband remains the workhorse of astronomy. Obviously it enables us to see the stars but
also to observe spectra which contain a wealth of information including abundances of elements,
temperatures and (via the Doppler shift) velocities. The atmosphere is transparent from about 2
microns to 330 nm. Originally detection was by the human eye, and then photographic plates.
These have very large capacity (of order 10 9 “pixels”) and so were excellent for surveys, but the
efficiency (in terms of photons detected) and the linearity are poor.
Electronic detectors have now replaced photographic plates in almost
all applications. The first of these relied on the photo-electric effect
where electrons are expelled from a phospher into a vacuum.
Applying a voltage accelerates the electrons and if these are allowed
to fall on a further phospher or metal surface each will produce a
shower of electrons. Multiple stages provide the basis of the
photomultiplier tube, which produces a clear pulse for each photon.
Counting the photons provides linear and efficient detection, with
good time resolution (< 1 microsec) but no spatial information.
Almost all modern detectors are however based on solid-state devices.
Here the photons are made to generate electron-hole pairs in a semiconductor. Silicon is used for wavelengths shorter than about 1
micron. The electrons are held in potential wells, forming the pixels,
until the end of the exposure. The charges are then shuffled across
and down by applying voltages to sets of electrodes under the cells.
This gives rise to the name Charge Coupled Device (CCD). The
charges are then read out by a low noise amplifier and digitized. Pixel
sizes are in the range 10 to 30 microns and up to 10 7 can now be made
on a chip. For astronomy they normally operate at about 180K to
reduce “dark current”. Efficiency can be > 50% and linearity is very
good.
(v)
mid- IR ~(20 to 2 microns) and far-IR (200 to 20 microns)
Semi-conductors can also be used here but one has to go to more
exotic materials such as indium antimonide and gallium-doped
germanium is order to have the energy gap between the valence band
and the conduction band less than the photon energy. This means that
we have to go to progressively lower temperatures as the wavelength
increases to limit thermal excitation of electron-hole pairs. Liquid
helium cooling (4K) is typical. As yet the detectors have fewer pixels
than in the optical and for the longer wavelengths have to mounted on
top of silicon CCD’s for the readout.
(vi)
Sub-millimetre (1mm to 200 microns). 1 mm is 300 GHz.
Here the photons are too weak to detect individually. For broad-band
signals we use a “bolometer” to measure the heat absorbed by a small
substrate in a very cold ( ~ 0.1 K) cavity. Modulation is essential but
sensitivity of better than 10 –15 W can be achieved.
For spectral lines it is better to use – the coherent techniques. Here
the electric field is measured directly. For frequencies above 100 GHz
we cannot yet make amplifiers so a mixer is used to down-convert the
signal. The best mixers are Superconductor-Insulator-Superconductor
(SIS) devices which have a very sharp non-linearity and can be made
very small.
(vii)
Radio – below ~ 100 GHz.
The electric field is collected by an aerial (typically a horn) and
converted into a current which ca then be amplified, down-converted
and digitised. Because the signals have been gathered coherently
interferometry is possible. Very high angular resolution can then be
obtained, in the extreme case by tape-recording the voltages at
separate telescopes and then correlating them later.