Lecture notes - Desy

16. Photon detectors
16.1 Ground-based detectors
To actually measure cosmogenic gamma rays we can use one of two methods: first, we may
try to actually observe the electromagnetic shower produced by the photon in the Earth’s
atmosphere. This can be done by measuring the shower particles on the ground, either using
scintillator arrays in an airshower array like TIBET or in a water Cherenkov detector like
MILAGRO, in which photomultiplier tubes measure the Cherenkov light emitted by charged
particles in the water tank. These ground-based detectors can only work with the particles
that actually reach the ground, so even if they are operated at high altitude like 4000 m above
sea level, the threshold energy for measuring and identifying gamma rays is usually in excess
of 1 TeV. Detectors of this type have the benefit of working at all times and having a large
field-of-view, whereas the photon angular localization is moderate.
Fig.16.1: Schematic view of an array of telescopes that observes atmospheric Cherenkov radiation
from a gamma-ray induced airshower.
A more direct way of measuring electromagnetic showers in the earth’s atmosphere involves
atmospheric Cherenkov radiation, which is emitted by charges, most notably electrons and
positrons, that move faster than the local phase velocity of light. Each electromagnetic shower
would therefore emit a brief flash (∼ 10 ns) of blueish light that can be observed with appropriate
telescopes. The technique was shown to actually work in the late Eighties with the 10-m
1

telescope at Whipple observatory, and about 70 sources of high-energy gamma-rays have been
detected to date (see http://tevcat.uchicago.edu/). Currently operational observatories that
use this technique involve the H.E.S.S. array in Namibia, the MAGIC telescope of the Canary
Islands, and the VERITAS array in southern Arizona. It should be noted that these telescopes
can only detect gamma rays at energies higher than about 100 GeV, for sufficiently many of
the radiating electrons have to be faster than the phase velocity of light in the atmosphere to
produce a detectable amount of Cherenkov light.
With all techniques a very large effective area can be realized: in the case of an airshower
array one observes all showers that have some shower particles hit the detectors, and therefore
the detection probability is not related to the actual size of the scintillation detectors. In the
case of atmospheric Cherenkov telecopes we only need the telescope located within the light
pool on the ground, which has about 200 m radius. This makes for a geometric area of 10 5 m,
which is also roughly the effective area at about 1 TeV, where the light yield is high enough
to in all cases overcome telescope inefficiencies related to, e.g., the mirror reflectivity or the
PMT quantum efficiency. At lower energies around 200 GeV the effective area is a bit smaller.
The technique works only in clear nights, so only about 10% of the time may be used for
observations. Also the field-of-view is usually small with about 4 degrees diameter. On the
other hand, the single-photon localization is good with about 0.1 ◦ .
The real problem for all ground-based gamma-ray detectors is to distinguish between gammaray
showers and cosmic-ray showers, which can be about 10 5 times as numerous. Cosmic-ray
showers usually involve a large number of pions and are therefore muon-rich, whereas gammaray
showers are muon-poor, but do contain some. Cosmic-ray showers also have particles with
larger transverse momenta and therefore appear wider and more fluffy than gamma-ray showers.
Air shower arrays mainly discriminate hadronic showers via the muon content, which provides
about a factor of 10 in gain, so there are only 10 4 hadronic showers per gamma-ray shower. This
relatively poor background rejection largely offsets the gains arising from the large field-of-view
and the high duty cycle. Atmospheric Cherenkov telecopes reject background by imaging the
shower in Cherenkov light, preferentially by three or more telescopes so a full three-dimensional
image can be derived. Cuts are then applied to select gamma-ray-like events based on the shape
of the showers. The background rate after cuts of contemporaneous atmospheric Cherenkov
telecopes is only about 1 Hz/deg 2 . The Crab nebula, a pulsar-wind nebula and the standard
candle of high-energy gamma-ray astrophysics, produces a rate after cuts of about 0.15 Hz in
the same detector, so it is obvious that a very good background rejection is achieved.
The excellent photon localization allows a further background suppression for point sources,
for which one searches for a peak in rate at the location of the source compared with the quasiuniform
background. Essentially all events associated with the Crab nebula should be recorded
within 0.2 ◦ of its true position, and the background rate in that solid angle element is also only
2

0.1 Hz, so the observed rate would actually more than double and the Crab is detectable with
high statistical confidence of 5σ within 4 minutes.
Fig.16.2: World map with the currently operating or constructed ground-based detectors of
high-energy gamma rays.
16.2 Satellite-based detectors
Lower-energy gamma rays can be measured by placing detectors outside the Earth’s atmosphere
by using balloons or satellites. A common characteristic is that the effective area of these
detectors is small, for interaction must occur within the detector material that has to fit on a
satellite. On the other hand we can control to a certain extent the environment in which those
interactions take place.
In the energy range between about 20 MeV and 100 GeV one uses pair creation on the Coulomb
field of high-Z nuclei such as Tungsten. By measuring the path taken by the electron-positron
pair one can deduce the direction of the incoming gamma ray. In actual experiments one
interleaves thin foils of, say, Tungsten with positionally sensitive charged-particle detectors, for
example silicon-strip detectors, to minimize the amount small-angle scattering on the Coulomb
3

field of the Tungsten nuclei that would change the particle track. Below this detector array
one would place a calorimeter to measure the energy of the electron-positron pair, ideally by
stopping it in the detector. For that purpose the calorimeter has to be thicker than the total
range of the electrons in the detector material. In the Nineties the EGRET experiment used
this technique to observe more than 270 sources of GeV-ish gamma rays. Its successor, Fermi,
has been launched in 2008.
charged-particle
anticoincidence
shield
γ ray
pairconversion
foils
particletracking
detectors
calorimeter
e+ e-
Fig.16.3: Schematic view of a pair creation telescope like GLAST. The detector is embedded in
a shield of scintillators that allow to veto the signals caused by cosmic rays. Position-sensitive
detectors measure the track of an electron-positron pair that has been creation in one of the thin
foils of high-Z material, before the energy is measured in a calorimeter.
To suppress background from cosmic rays entering the detector, the satellite experiments use
an anti-coincidence shield, which is nothing more than an envelope of scintillators that give
a signal when a charged particles enters the instrument. By requiring that any event in the
tracking detector and the calorimeter be not coincident with an signal in the scintillators one
effectively vetoes all background-induced events, so that very weak emission components like
extragalactic diffuse gamma-ray emission can be observed.
The photon localization is compromised by small-angle scattering of the electron or positron
in the tracker material, so essentially by Coulomb scattering and ionization or excitation. The
efficacy of these processes in changing the momentum of an electron falls off with energy, so
the single-photon localization is bad at low energies and good at high energies, where, however,
4

the photons are few.
At energies below 20 MeV one can use Compton scattering to measure gamma-rays. The
detector would consist of two planes and one would require a Compton scattering to happen
in each plane. Generally one measures the energy of the electron on which the photon has
scattered, and the location where the scattering took place. If the gamma-ray is absorped in
the second detector plane, we know the gamma-ray energy after the first scattering in the upper
detector plane. The energy deposited during the first scattering then gives the scattering angle,
so one can reconstruct the arrival direction to a angular ring around the axis connecting the
locations of the two scattering. The photon energy follows from the energy deposited in both
detectors, if one can make the usually correct assumption that the photon doesn’t leave the
lower detector plane. Background arising from activation of detector or satellite material is a
major concern.
Fig.16.4: Schematic view of a Compton telescope like COMPTEL.
At even lower energies below a few MeV one uses solid-state detectors that are equipped with
collimators or coded masks to give some imaging capability. An example is INTEGRAL.
5