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Astro Stats and Measurement
Measurement Part D : Detection of Light
Andy Lawrence, University of Edinburgh
Sept-Dec 2013
To actually make a measurement, we need to detect the incoming light - i.e. the light must cause
an effect of some kind that we can quantify or record. In this chapter we will concentrate on the
physical principles of detection, followed by a survey of the main kinds of detector in use. We will
not discuss the details of all the ancillary structure needed - readout electronics, calibration units,
digital processors, re-imaging optics etc.
First however, a key point concerning nearly all detection systems : the need for amplification. The
effect caused when a photon interacts with matter is typically very small. To end up with a measurable
signal therefore always requires either accumulation or amplification. Two old-fashioned examples
can make the idea clear. (i) On a photographic plate, a single photon can cause a chemical change in
a tiny amount of material. This would not be visible to the eye, but because the change is irreversible,
a prolonged exposure will eventually accumulate a noticeable effect. If the chemical change were
to “relax back”, this would not work. (ii) In a Geiger counter, an incoming gamma-ray or particle
will ionise an atom in the gas contained, liberating one or more electrons. Left to themselves, these
electrons would soon recombine and we would be none the wiser. But if a large voltage is applied
across the gas, the liberated electrons are accelerated towards the anode, causing more ionisations
and forming a cascade. The result of this amplification is an easily measurable pulse of current, so
that we can detect individual particles.
We start by reviewing the main ways that light interacts with matter. Section 4.1 considers the main
physical processes, and section 4.2 focuses on the processes in man-made structures such as capacitors, p-n junctions, and superconductors. We then look at how these processes are used in actual
detection devices - first in devices that are based on individual photon detection (section 4.3), then
on devices that rely on the net heating effect (section 4.4), and finally on coherent wave detection
devices, i.e. antennas and receivers (section 4.5).
1
Interaction of light with matter
In this section we look briefly at the principle ways in which light interacts with matter. Interaction
can be either coherent - phase-sensitive detection of the wave itself - or incoherent - interactions based
on the energy of photons. The main incoherent interactions are graphically summarised in Fig. 1.
1.1
Coherent interaction
Light is an electromagnetic wave, which is an oscillation in the electric field (and in the accompanying
magnetic field) that permeates space. Free electrons, or electrons in a conducting material, respond to
the electric field via the Coulomb force. The oscillating field caused by a wave can therefore produce
an alternating current/voltage in a conducting material. Loosely speaking, any arrangement of such
conducting material can be seen as an antenna. The oscillating voltage will switch between +V and
−V so that the time averaged voltage is zero. However the average of V 2 is always positive, and the
power produced is P ∝ V 2 . The electrical power produced in the antenna is therefore proportional
to the power in the incoming wave.
1
Astro Measurement D : Detection
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Figure 1: Three basic categories of photon interaction. Adapted from a figure originally made by Dr
A.Woodcraft.
1.2
External photo-electric effect
In this case light is detected as a particle rather than a wave; the photon is absorbed, and if its energy
is sufficient, it ionises an atom. We refer to the “external photo-electric effect” when an electron
is completely removed from the surface of a material. Substances that show this effect are known
as photo-emitters. This is the effect for which Einstein got the Nobel prize. The minimum photon
energy needed to remove an electron is called the work function W for the material concerned. For
photons above this energy, the rate of photons determines the emitted rate of electrons. Note that we
therefore measure
Z
ν=∞
Z
(Sν /hν)dν
∞
rather than
ν=W/h
Sν dν.
W/h
The measured quantity is a current. However, this is typically tiny and therefore needs some kind of
amplification to produce a measurable effect.
1.3
Internal photo-electric effect : ionisation
If the liberated electron remains inside the volume of the material, we refer to the internal photoelectric effect. However, such an electron would quickly recombine. In order to produce a measurable
effect, the usual technique is to apply a voltage that will accelerate the electron, which will then
cause further local ionisations, producing a cascade which can be measured at the anode as a pulse of
arriving charge. In this way we can therefore count photons individually.
Ionisation in gas-filled chambers is a standard technique for laboratory detection of particles and highenergy photons, i.e. X-rays and gamma-rays, and for some domestic uses such as smoke detectors. It
was originally the main technique used in X-ray astronomy, but now has been superseded by solid-
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Figure 2: Energy level structure in solids. In insulators, the valence and conduction bands are well separated,
whereas in conductors they overlap. In semi-conductors the “band gap” is small, so that an absorbed photon
is capable of raising a valence electron into the conduction band.
state devices. Note that for a given element the cross-section for photo-ionisation is largest at its
ionisation edge ν0 , but declines as ∼ (ν/ν0 )−3 at higher frequencies. Hydrogen, with its photoionisation edge at an energy of 13.6eV is therefore very inefficient at detecting X-rays; by an energy
of 1keV the cross section is 400,000 times smaller. Heavier elements have ionisation edges at higher
energy. The normal choice for gas-based photon-counting detectors is to use an inert gas such as
Xenon or Argon.
1.4
Internal photo-electric effect : band gaps
In solids, close-by atoms distort and split the energy levels of their neighbours, both in energy and
spatially. The result is a set of broad energy “bands” across the whole structure. (See Fig. 2). At the
lowest energies we have strongly bound (inner shell) electrons. Next there is the valence band, then
the conduction band, and above this the ionisation limit. Electrons in the conduction band can easily
migrate spatially. In a conductor the valence band and conduction band overlap. In an insulator, there
is a band gap of size Eg , making it hard for electrons to move. In a semiconductor the gap is present
but is not very big. Electrons can then be boosted into the conduction band, either thermally, or by the
absorption of a photon with E > Eg . For example in Silicon, Eg = 1.2eV; any photon with E > Eg
can be absorbed. In wavelength terms, this corresponds to photons with λ < hc/Eg = 1.1µm. When
this happens, an electron-hole pair is formed, with a negative electron in the conduction band, and a
positive hole in the valence band.
It can help to think of the physical structure. An element like Silicon has four valence electrons,
half the number needed to make a filled shell. The orbitals for these electrons are elongated fingers
orientated in a tetrahedral structure around the centre of the atom. Neighbouring atoms orient themselves in a diamond lattice so that every valence electron bonds tightly with a valence electron in a
neighbouring atom. Every atom then “sees” eight electrons, making a very stable structure, and a
filled set of states. Fig. 4 shows a flattened representation of the Silicon lattice. An electron raised to
a higher energy state is also moved to a physically different location. This leaves an unfilled location
in the lattice, which is the hole. Another electron (not necessarily the original one) can move into this
gap, but this leaves another gap, and so on. This shifting pattern can be seen as either a movement of
(negative) electrons or (positive) holes.
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Figure 3: Using a semi-conductor as a photo-conductor. An absorbed photon creates an electron-hole pair,
and an applied voltage then creates a current.
Fairly soon, the excited electron would decay back to fill the gap in the lattice. However, if a voltage
is applied across the photo-absorbing material, the electron-hole pair will drift through it, creating
a current. The value of the current will be proportional to the rate of photons that have E > Eg
impinging on the material. This is using the material as a photo-conductor. More complicated and
interesting things can be done with semi-conductor materials, but we will discuss these in Section 2,
after completing our survey of types of light interaction.
1.5
Compton scattering
At high energies, E > 50 keV, photo-ionisation becomes very inefficient, because of the ν −3 fall
off of the photo-ionisation cross section for all elements. However, a photon can deposit energy
in a medium by Compton scattering. (See Fig. 1.) In the classical view, the electric field of an
incoming EM wave excites an oscillation in free electrons. (For high energy photons such as X-rays,
even a bound electron behaves as if it were a free electron). The oscillating electrons then radiate
as dipoles. The result is outgoing radiation that is the same frequency as the incoming radiation,
but scattered into a range of directions. (This is “Thomson scattering”.) In the quantum view, each
outgoing photon can be seen as having a certain probability of being scattered into various different
directions, with the probability distribution following the classical dipole radiation pattern. However,
in quantum physics, as well as having energy E = hν the photon also has momentum p = E/c. In
order to conserve both energy and momentum during a collision between a photon and an electron,
this requires the electron to carry off some KE and the photon to lose some energy, such that
E1 − E2
1
=
(1 − cos θ),
E1 E2
me c2
which for E1 ∼ E2 can be expressed as
E
∆E
∼
(1 − cos θ),
E
me c2
where E is the original photon energy, me is the mass of the electron, and θ is the scattering angle,
between the original and new photon directions. As me c2 = 8.19 × 10−14 J = 511 keV , you can
see that visible light photons suffer negligible energy change, but it becomes a significant effect for
X-rays.
Just like with the internal ionisation effect, the result of Compton scattering is a free electron given
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Figure 4: Flattened schematic view of the Silicon lattice, and the process of photo-absorption. In 3D reality,
Silicon has a tetrahedral lattice. Every valence electron is shared by two atoms. The absorption of a photon
raises the energy of an electron so that it can drift between atoms, and also leaves behind a hole in the lattice
structure.
excess energy, which can be accelerated by an electrical voltage to produce a cascade and so a measurable pulse. However, unlike the internal ionisation effect, the original photon is not destroyed, and
so can carry on to produce further events.
1.6
Pair production
It is in principle possible for a photon to convert spontaneously into matter. To conserve charge, it
has to make an electron-positron pair. To conserve energy, the photon must have at least twice the
rest mass energy of an electron E = 2me c2 = 1.02 MeV, i.e. it must be a gamma-ray. However, the
event also has to conserve momentum. This is impossible at the energy threshold, when the created
electron and positron are at rest, and very unlikely at all energies. In practice therefore, pair creation
always needs a third body to be involved, in order to absorb the momentum. This can be another
gamma-ray : photon-photon collisions are important in some exotic types of astrophysical object.
More important for our purposes here, it can be a massive body such as an atomic nucleus, which
can absorb the momentum without changing much in velocity. High density materials, such as lead,
are more effective pair-creators. The pairs created, just like electrons from ionisation or Compton
scattering, can be cascaded and so produce a measurable pulse. Pair production is therefore not
normally a spontaneous process, but one where photons interact with matter and deposit energy. At
the highest photon energies, this is the dominant way that light interacts with matter.
Astro Measurement D : Detection
(a) p-type
6
(b) n-type
Figure 5: Lattice doping. On the left, the inclusion of Boron atoms leaves a (positive) hole in the structure,
and so is known as a p-type material. On the right, the inclusion of a Phosphorus atom adds an extra (negative)
electron, creating an n-type material.
1.7
Secondary light production : Scintillation and Cerenkov radiation
This is not a distinct interaction of light with matter, but an important follow on effect. Whenever
a high-energy photon causes photo-ionisation, Compton interaction, or pair production, the affected
electron can leave behind a hole in the atom in an inner shell. When outer electrons drop back down
to fill the hole, visible-wavelength light is emitted. The secondary electrons in the “local cascade”
can also cause this effect, and so we see a tiny but detectable flash of light from the location of the
interaction. This is known as scintillation.
When the very highest energy gamma-rays (>300GeV) interact with matter, the particles produced
in the cascade will have extremely large energies, such that their velocities can be greater than the
local light speed in the material concerned. They are forced to decelerate. This deceleration causes
radiation known as Cerenkov radiation. This radiation is seen from the walls of nuclear reactors, and
is also used in particle physics experiments to detect high energy particles. The effect can also take
place when high-energy gamma-rays (or cosmic rays) enter the Earth’s atmosphere, producing faint
but detectable flashes of light. In other words, we can use the Earth’s atmosphere as a detector. We
shall see later how this is put into operation.
1.8
Heating
This is likewise not a distinct type of interaction but rather a net effect. Following some initial
interaction, such as an ionisation, or a Compton scattering, there may be a chain of effects, such
as recombination, further ionisations, excitation of lattice vibrations, and so on. Energy will be
conserved in this chain of events, but gradually dissipated, ending up as heat. The final effect of the
absorption of light may then simply be to heat the material. This heating effect can itself be used to
measure the amount of light absorbed.
2
Interaction with man-made structures
The previous section surveyed the ways that light interacts with naturally occuring materials. Those
effects are all used in practical light detection systems of various kinds. However, we can go further,
constructing artificial structures with desirable properties. Structures made from semi-conductors are
of course the basis of the whole modern electronics industry; but they are also the basis of the most
important classes of light detector in astronomy. Here we look briefly at several key types of artificial
structure.
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Figure 6: Structure of a Metal-Oxide-Semiconductor (MOS) capacitor, which is the basis of a single CCD
pixel.
2.1
Doping
The first trick comes from doping, that is introducing a small number of alien atoms into the Silicon
lattice. This process makes what is known as an extrinsic semi-conductor, as opposed to pure Si
which is an intrinsic semi-conductor. (See Fig. 5.) If a lattice position is filled with a Boron atom,
which has only three valence electrons, one of the eight surrounding valence electron slots is left
empty. This is known as a “p-type” or acceptor material. It doesn’t take much energy for a neighbouring electron to hop into this spot - about 0.045 eV, compared to the band gap of 1.2eV. However,
this leaves a hole in the valence band, which can in turn be filled by a neighbouring electron, and so
on. The result is that the valence band becomes conducting by the motion of holes. Alternatively, if
a lattice position is filled with a Phosphorus atom, which has five valence electrons, there is a spare
electron. This is known as an “n-type” or donor material. As the spare electron cannot occupy any
of the filled states, it must be in a state of higher energy (typically only ∼ 0.05 eV less than the conduction band) and can therefore be relatively easily excited into the conduction band. Such extrinsic
semi-conductors can be used directly as photo-conductors, sensitive to longer wavelength light than
intrinsic Si.
2.2
Trapping : MOS capacitors
The simplest light detecting method using a semi-conducting material is to deploy it as a photoconductor. However, we can also trap charge by combining Silicon with other components. The
usual arrangement is shown in Fig. 6. The surface of the silicon is covered with an insulating layer
such as SiO, with a conducting electrode over this, and a voltage applied across the region. The
photon passes through the insulator and is absorbed in the silicon, creating an electron-hole pair. The
applied voltage separates the electron and the hole, which cannot then recombine. The electrons are
stopped at the insulator, creating a tiny capacitor storing the charge. This arrangement is known as a
Metal-Oxide-Semiconductor (MOS) capacitor. The importance for light detection is that one can
integrate, with the accumulated charge being proportional to the total number of photons detected.
This is the basis of the CCD, which we shall discuss shortly.
2.3
p-n junctions : photodiodes
The third trick is to join p-type and n-type materials to make a p-n junction. Excess electrons move
from the n-type material to fill the holes in the p-type material; conversely holes move from the n-type
to neutralise the electrons on the p-type side. The result is a non-conductive layer depleted of charge
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Figure 7: Behaviour of a p-n junction, which is the basis of the photodiode. Adapted from a figure available
on Wikimedia under the GNU Free Documentation license, originally uploaded by user “The Noise”.
carriers (known as the “depletion zone”). However, this region is charged, because the migrating
electrons/holes leave ions behind, making a “built-in voltage” across the junction. The diffusion
stops when the built-in voltage reaches the equilibrium level. (See Fig. 7). The main importance for
electronics is that a p-n junction is a diode; electrons can easily flow from n to p but not the other way.
However, a p-n junction can also be used as a photo-diode. If a photon is absorbed in the depletion
layer, the electron and hole are swept apart by the built-in voltage. The junction therefore acts as a
capacitor. This can then be run either as a photo-conductor (e.g. for a solar power cell) or in photovoltaic mode. In the latter case, the junction is held at zero current and the voltage drop depends on
how much charge has been collected.
2.4
Superconductors
The phenomenon of superconductivity at very low temperature involves pairs of electrons, known as
”Cooper pairs” which travel locked together through the lattice structure of the conductor. The pair
of electrons also produce a distortion of the lattice, which travels with them, i.e. a kind of phonon
excitation. In a Superconducting Tunnel Junction (STJ), made of two superconductors separated by
a thin insulating layer, Cooper pairs can spontaneously tunnel across the junction. If a magnetic field
is applied across the junction, this suppresses the spontaneous current. However photons can break
up the Cooper pairs. This process also produces quasi-particles - electron-like or hole-like travelling
lattice vibrations - that tunnel across, making a pulse of current. Only a very small photon energy is
required to produce this effect - of the order milli-eV, so devices based on this effect can detect low
frequency photons (submm) or alternatively can produce large pulses from single optical or X-ray
photons.
A property of superconductors which is of growing importance is that of kinetic inductance. Impedance
to current flow is in general made of resistance plus reactance. For static currents, there is only resistance, which relates to the difficulty of electrons having to difffuse through the lattice, with a very
short time between collisions. For time varying currents (for example in AC circuits) additional effects can impede the current, and also produce lags. The two well known such reactance effects
are inductance, where the changing magnetic field induces a voltage which resists the change, and
capacitance, where accumulated charge likewise acts against the current. These can be seen as “inertial” effects. A similar inertial effect arises from the fact that when a voltage is applied, electrons
do not immediately establish a constant drift velocity, but have to be accelerated. This is generally
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Figure 8: Amplification in a photomultplier tube. (Taken a wikimedia image first created by Colin Eberhardt
and released into the public domain.)
called “kinetic inductance” although it has nothing to do with (magnetic) inductance; sometimes the
term “surface reactance” is used. At high temperatures it is negligible compared to other forms of
impedance, but for alternating currents in superconductors it is an important effect. Photon absorption
alters the value of kinetic inductance because the created quasi-particles occupy some of the available
quantum states.
3
Photon detection devices
Having completed our survey of how light interacts with matter, we will now look at how these
principles are put into action to make working detectors. We start in this section with devices based
around detecting individual photons. In the following two sections, we will first look at devices using
the heating effect (bolometers) and then look more closely at coherent wave detection systems (radio
receivers).
3.1
Photo-multiplier tubes
The current produced by a photo-electric material is extremely small even for the brightest stars.
Measurement therefore requires amplification. In a photo-multiplier tube (PMT) this is achieved by
accelerating the liberated electrons, as illustrated in Fig. 8. The first light sensitive surface is also the
cathode. The accelerated electrons can produce further ionisations at a series of secondary “dynodes”,
thus producing a final current that is several orders of magnitude larger. Historically, as the first
electronic device used in astronomy, the photo-multiplier tube was important for three reasons; it was
many times more sensitive than previous visual or photographic methods; it was linear, producing a
signal that was proportional to the impinging flux; and it produced an electrical signal which could
be directly fed into computers, thus starting the digital age. PMTs also have very rapid response, and
so can adjust to changes in photon rate very quickly, which also means that they can act as photoncounting devices. PMTs however have two big disadvantages : they do not integrate, and they are not
two-dimensional devices, so cannot be used to record an image. For these reasons, in optical and Xray astronomy, CCDs are now almost always the detector of choice. However, where rapid response
is needed, such as measuring very fast variables, or short flashes of light such as those caused by
the atmospheric Cerenkov effect, PMTs still have an advantage. Furthemore, the same principles are
now used in the micro-channel plate, which we will meet later, where an image is recorded. Finally,
if individual photon counting can be achieved, integration can be digital.
Astro Measurement D : Detection
3.2
10
Charge Coupled Devices (CCDs)
As we have seen above, a MOS capacitor can accumulate charge, which makes it able to detect
extremely faint sources of light. A single MOS capacitor is around 20µm across, which can be
matched well to the plate scale at large telescopes. A two dimensional array of such light detecting
pixels can therefore record an image. At the end of the exposure, we need to measure the charge
on each pixel. The scheme for doing this is illustrated in Fig. 9. Charge is transferred in parallel
from one column to the next, with the final column being transferred to an extra readout column. The
charges in the pixels in the readout column are then transferred in the orthogonal direction. As the
final pixel is discharged, it produces a voltage which is then passed to circuitry which converts the
analogue voltage to a binary number proportional to the number of electrons in the capacitor, i.e. the
number of detected photons. The cycle is then repeated until the whole chip is read out.
How is the charge transferred from one pixel to the next? This is where the “charge coupling” part
comes in. The metallic contacts or gates at the top of the MOS capacitor are constructed so that
they overlap neighbouring pixels. The centre of a pixel can be held at a different voltage from the
parts overlapping its neighbours to left and right. By cycling the three voltages, the charge can be
encouraged to drift across the capacitor and then into the next capacitor. A detailed example of how
this works in practice can be seen in Chapter 8 of Tomey (2010). In a single pixel-to-pixel transfer
the efficiency of charge transfer is very high - 99.999 to 99.9999% - but millions of transfers have to
be made for a complete readout of the CCD. The streaks you sometimes see around bright stars in
CCD images is due to this not quite perfect charge transfer efficiency.
Another practical problem is that purely by thermal effects, electrons are being lifted up into the conduction band all the time. In astronomical terms, this makes what is known as dark current. This will
depend on how many thermal electrons have energies comparable to the band gap. This depends on
their energy distribution, given by the Boltzman distribution, for which the relative number of electrons at energy level E goes as exp −E/kT . Integrating above Eg and making some other detailed
corrections it is found that the expected dark current goes as
q ∝ T 1.5 exp
−Eg
.
2kT
For domestic digital cameras this effect is not important, but in the very low light level conditions
of astronomy it can be crucial. It is therefore normal to cool CCD detectors, typically using a liquid
nitrogen dewar, to temperatures of the order 150K.
A further practical issue is that the quantum efficiency (i.e. the probability of a photon being absorbed) can vary from one pixel to another. This is yet another effect that has to be empirically
calibrated, by observing something that is uniformly illuminated (a “flat field”). This can be an out
of focus region on the inside of the telescope dome, or the twilight sky.
Astronomy was one of the first areas to adopt CCD technology, followed by digital photography.
Modern domestic digital cameras mostly do not use CCDs, but instead use the closely related CMOS
technology. Each CMOS pixel is a little more complicated, and has its own readout circuitry, so that
there is no need for charge transfer. This gives a big advantage in readout speed, and each pixel can be
read non-destructively. However, it is hard to make large sensor arrays to the quality and consistency
needed for astronomical imaging, so CCDs remain the detector of choice for optical astronomy - but
not for IR astronomy, as we discuss next.
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Figure 9: The readout scheme for a CCD. The columns are sometimes known as the parallel registers and the
readout column as the serial registe. “ADC” stands for Analogue to Digital Converter.
Astro Measurement D : Detection
Name
Silicon
Germanium
Mercury cadmium telluride
Mercury cadmium telluride
Indium antimonide
Germanium gold
Silicon arsenic
Germanium copper
Germanium gallium
Stressed germanium gallium
Formula
Si
Ge
Hgx Cd1−x Te
Hgx Cd1−x Te
InSb
Ge:Au
Si:As
Ge:Cu
Ge:Ga
Ge:Ga
12
Type
Intrinsic
Intrinsic
Intrinsic
Intrinsic
Intrinsic
Extrinsic
Extrinsic
Extrinsic
Extrinsic
Extrinsic
λmax (µm)
1.11
1.85
2.5 (x = 0.55)
5 (x = 0.7)
5.6
8.3
23.1
30.2
115
∼ 200
Table 1: Maximum useful wavelengths for some extrinsic and intrinsic semiconductors.
Figure 10: Bump bonded hybrid detector. (Figure kindly provided by Dr A.Woodcraft).
3.3
IR arrays
Because of the size of the band-gap, the longest wavelength light which a Silicon CCD can detect is
at 1.1µm. Other semi-conductors, especially doped substances, can have smaller band-gaps and so
detect light at longer wavelengths (see Table 1). However, it is not practical to make IR CCDs from
these substances, as it is much harder to construct the necessary circuitry. Because of its importance
to the electronics industry, nearly all the necessary research and development has been done with
Silicon. In practice this means that while the detection material may not be silicon, the readout
circuitry must be.
This drives us to using hybrid devices as illustrated in Fig. 10. These have a sandwich structure
: a detector wafer and a readout wafer, joined by many individual electrical connections, generally
using indium solder, a process known as bump bonding. Small dots of indium are patterned on the two
surfaces, which are then forced together to make contact. The silicon wafer is typically not operated
as a CCD, in that charge is not moved from pixel to pixel. Rather, each pixel is a CMOS circuit, and
a network of transistors is used to switch readout circuitry to each pixel in turn. This technique is
known as multiplexing, so that the readout wafer is also referred to as a multiplexer wafer . Being
able to address each pixel separately has a big advantage in that the array can be read out without
destroying the charge, so that for example one can perform multiple readouts to improve the accuracy
even in the presence of readout noise.
At the very longest wavelengths (above around 100 µm) it is necessary to apply high pressure, of the
order of 0.5 GPa, to the semiconductor in order to achieve sufficiently small band-gaps. These are
called stressed detectors. However, currently it is very hard to make such devices larger than a few
tens of pixels.
Like CCDs, IR arrays can suffer from thermal noise, dark current, and readout noise. The good news
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is that readout noise can be gotten around by using multiple readouts. The bad news is that at a given
temperature the dark current is much worse because the band-gap is smaller, so IR arrays have to be
cooled to much lower temperatures, especially when working in the far-IR, where temperatures of
only a few degrees above absolute zero are used. This is achieved by placing the detector inside a
de-pressurised vacuum vessel containing liquid Helium.
There are other practical issues, such as the sensitivity to mechanical stresses between the two layers
of the sandwich, and the fact that because of the CMOS and multiplexing circuitry, it is much harder
to make mosaics of arrays. However by far the biggest drawback of IR arrays is that they are very
expensive - a science grade CCD array may cost you $100,000, but a science grade IR array will cost
$1M.
3.4
Microchannel plates
In principle, silicon CCDs work at any wavelength with λ < hc/Eg . However, for λ < 450nm,
CCDs become less efficient, especially as we move into the UV. This is because photons are easily
absorbed in the electrodes and other structures on the top of the MOS capacitors, before they reach
the silicon. A technique used sometimes is to turn the CCD upside down, making it a so called back
illuminated device. The main problem with this is that the charge tends to spread out before reaching
the depletion region, unless the chip is thinned to around 10µm. However this leaves the chip fragile,
and susceptible to fringing effects - Newton’s rings-like patterns appearing in the image.
One approach used in the UV is employ a converter - a surface with a coating which absorbs UV
light and emits optical light - a phosphor. The optical light can then be detected by a normal CCD.
However, another serious problem is visible light contamination. Many objects for which we wish to
measure the UV flux are much much brighter (by a factor of a million or more) in the visible. It is
hard to perfectly filter out this visible wavelength light. A better technique is therefore to use a photoemissive device, like a PMT, rather than a photoconducting device like a CCD. Plenty of substances
have work functions such that they are only sensitive to UV light.
Standard PMTs are single pixel devices. To produce an image, a solution is to make a microchannel
plate (MCP), which is a slab of glass with many parallel microscopic pores of width 10–100µm,
coated with a good emitter of secondary electrons. A high potential (kV) is applied across the faces,
so that each channel acts like a PMT and produces a cascade of electrons, as shown in Fig. 11. Gain
factors are of the order 105 , and using stacks of such plates, can be up to 108 . One advantage of
MCPs, like PMTs, is that they can be used as photon counting devices - each splash of electrons is
individually detectable.
The electrons emerging from the MCP still have to be read out somehow. One technique is to use a
phosphor screen to convert the electrons to visible light, and then to detect this with a normal silicon
CCD. An arriving pulse of electrons can of course be easily measured at the anode, but somehow
one needs a two dimensional position-sensitive anode. This can be achieved by using two grids of
perpendicular anodes, and simply to see which anode the pulse is detected at in both x and y. In
practice the pulse produces a splash of detections at several anodes, but the signal is centroided, and
can therefore be located to better than the separation of the anodes. This kind of device is known as
a multi-anode microchannel array (MAMA), and is for example used in the Hubble Space Telescope
ACS camera.
3.5
Proportional counters
Traditional laboratory particle and gammma-ray detectors use a gas filled volume in which internal
ionisations take place, and an applied voltage accelerates the electrons. In the early years of X-ray
Astro Measurement D : Detection
14
Figure 11: The operation of a microchannel plate. Light is incident on a surface with a grid of fine channels.
The inside of each channel acts as a photo-emitter. and an applied voltage produces a cascade of secondary
electrons. (Adapted from a figure found on Wikimedia which has been released into the public domain. Originally created by user “Andreas06”)
astronomy, this was also the method used. The liberated electron moves off with excess KE, and can
cause further local ionisations even before the acceleration by the applied voltage takes effect. (In
fact the initial photo-ionisation can liberate a second electron from the first atom, known as an Auger
electron.) The size of this local mini-cascade depends on the energy of the incoming photon; and so
the size of the final pulse at the anode also depends on the photon energy. If the conditions in the
detector are carefully set up, the relation between photon energy and pulse size will be linear. This
is known as a proportional counter. Such a device will then count individual photons, and allow us
to estimate the energy of each photon, as discussed further in Chapter 5. It is possible to make a
two dimensional detector by using crossed anodes, as in the case of microchannel plates discussed
above. However the spatial discrimination and the spectral discrimination are relatively poor. In soft
X-rays (< 0.5keV), where such proportional counters are not very efficient, MCPs have sometimes
been used, producing excellent images but losing the ability to measure energy.
3.6
X-ray CCDs
Between 0.1keV and 10keV, silicon CCDs have become the detector of choice. The absorption in the
surface electrodes that makes UV detection inefficient decreases in the X-ray regime. Furthermore,
back-illuminated CCDs are also better than in the UV, because the X-rays travel further through the
silicon before being absorbed, so there is less pixel-to-pixel smearing of the detected charge. The big
advantage of X-ray detection, just as with gas counters, is that Eph >> Eg . The electron released into
the conduction band has large energy and causes multiple secondary impact ionisations, the number
of which will be proportional to the photon energy. The energy resolution is significantly better than
for gas-filled proportional counters. We will consider the issue of energy resolution more closely in
Chapter 5.
Astronomical X-ray detectors are normally used at low count rates. Rather than accumulating charge
over time, the CCD is therefore read out repeatedly looking for single events. However, it can take
several seconds to read out a complete frame, and some sources of interest do have count rates of this
order, leading to potential “pile up”. The effect on the measured flux of a source can be corrected,
but this is still a significant drawback of CCDs. A new generation of CCDs have frame readouts that
are several times faster, but the future probably lies with CMOS and related technologies.
Astro Measurement D : Detection
3.7
15
Hard X-ray detectors
For harder X-rays (>10 keV) standard silicon CCDs become less efficient, because the absorption
path length is too long, and most photons pass straight through the silicon layer. In principle you
could make the detector thicker, but then standard CCD construction techniques don’t work. Semiconductors with higher atomic number are more efficient but just as with IR-arrays, it is too difficult/expensive to make CCD-like circuitry with materials other than Silicon, so detectors are small
and expensive arrays built in university labs. A popular choice used all the way through hard X-ray
to lower energy gamma-rays (∼ 10 MeV) is Germanium. At these large energies, the absorption
length is long, so pixels can be several cm across. At intermediate ranges (10-500 keV) there are new
materials (eg CdTe and CdZnTe) which are reasonably efficient and don’t have to be too large, so that
as in the IR, arrays of pixels can be bump-bonded to a standard silicon CCD for read out.
An alternative to direct detection is to use scintillators, which convert X-ray photons to photons
at visible wavelengths, which are then detected with semiconducting or photoemissive detectors.
These possess energy resolution, but the resolution is worse than for semiconductor detectors, and
they are less sensitive (the fraction of X-ray photons converted to light is typically 10%). Their main
advantage is that they can be made into a much larger volume than semiconductor detectors. Common
scintillator materials are alkali halides such as NaI or CsI. The term scintillator is usually reserved
for bulk detectors; thin films operating in the same manner are called phosphors.
3.8
STJs and KIDS
Superconducting structures can be configured to make devices which detect individual photons. In an
STJ, recording an absorbed photon is relatively simple - a pulse of current is detected. In principle,
such devices could replace semiconductor devices such as CCDs over a very large wavelength range.
However, although such devices have been experimented with for some years, they have proved difficult to multiplex, i.e. to make large arrays of pixels. At the time of writing, Kinetic Inductance
Detector Systems (KIDS) seem more likely to be the long term replacement for CCDs. As explained
in section 2.4, absorption of light alters the kinetic inductance, but how is this measured? Kinetic
inductance involves a characteristic lag time, which means that for an applied AC voltage, the resulting current amplitude depends on the AC frequency; there will be a resonant frequency, which
depends on the value of the kinetic inductance. Experimental KIDS devices make a micro-resonator
by combining a capacitor with a superconducting strip. Absorption of a photon causes a temporary
change in the resonant frequency. Recent experiments show that pixels of this kind can be more easily
multiplexed than STJs.
3.9
Compton telescopes
For photons in the range 1MeV – 30MeV most detectors operate via Compton scattering. The simplest design consists of two levels, as shown in Fig. 3.1. In the top level, a gamma-ray is Compton
scattered, giving up some of its extreme energy to an electron. The scattered photon then travels
into a second level where it is absorbed completely. The levels consist of scintillators or semiconductors which determine the approximate interaction points. The two levels are optimised for the
required behaviour; the first level should have a low atomic number (e.g. silicon) for high Compton
and low photo absorption efficiency, and the second level should have high atomic number for good
absorption efficiency. A more sophisticated system will have several levels of each type of material.
Compton telescopes can be used to determine the energy of the photon (with a resolution of 5 to
10%), and also give some spatial information. As explained in Chapter 3, the track between the
two detection layers, together with the change in energy caused by the Compton scattering, gives a
measurement of the arrival direction of the incoming photon.
Astro Measurement D : Detection
3.10
16
Pair telescopes
At still higher energies (around 20 MeV to 300 GeV), pair telescopes are used. These consist of many
converter layers and detectors. The converter layers consist of a high atomic number material (e.g.
lead) which forms a target where pair production can take place. As with a Compton telescope, the
detectors may be scintillators or semiconducting detectors. By constructing the tracks of the electrons
and positrons as they pass through the device, it is possible to calculate the direction of the original
gamma ray photon, and thus its position on the sky. The initial energy can also be determined to a
resolution of around 20% by analysing the tracks and possibly using a further detector at the end of
the instrument which absorbs the electron and positron to determine the remaining energy.
3.11
Cerenkov telescopes
At the very highest energies (>300 GeV), we use Cerenkov telescopes. Such a high energy particle
actually causes a cascade of particles, known as an air shower. These particles can be formed with
such a high energy that they are moving faster than the speed of light in air. They rapidly decelerate
and so radiate, a form of radiation known as Cerenkov radiation. Atmospheric modelling shows
that the number of particles reaches a maximum at a characteristic height of ∼ 10km; the light then
produces a pool of light on the ground of radius ∼ 130m. A 1 TeV photon will produce around
100 photons/m2 in a flash lasting a few nanoseconds, peaking at a wavelength of 300-350nm. This is
reasonably bright but only for a very short space of time. This suggests using a PMT rather than CCD,
as these have extremely fast response, and can be read out continuously at MHz rates. Furthermore,
one wants to search for flashes over a wide angle, but only crude imaging is needed. The standard
technique is to use large (∼10m) but simple reflecting telescopes, with an array of PMTs at the focal
plane covering a field ∼ 5◦ across.
The arriving light is only symmetrical for an overhead shower. For other angles, the result is an
elliptical smear at the focal plane with long axis pointing towards the origin of the shower. If an
array of such telescopes is used, spread over the ground-pool, the origin of the shower is indicated
by the intersection of these ellipses, giving a final angular resolution of around 0.1◦ . High energy
cosmic rays as well as gamma-rays can cause such Cerenkov showers, and heavily outnumber the
gamma-rays from astronomical sources, but give a characteristically different shape in the image
plane.
4
Heat measurement systems : bolometers
At wavelengths longer than 200µm but shorter than radio wavelengths - i.e. in the submm - the main
technique available is to use bolometers. These are devices where we measure the temperature change
caused by the absorbed radiation. Some years back, bolometers were used in the mid and far-IR as
well, but now photo-conducting arrays are much better in the IR.
4.1
General principles of a bolometer system
The basic idea is illustrated in Fig. 12. Radiation is absorbed by a small thermal mass attached to a
heat reservoir kept at a fixed temperature, which leads to a change in temperature of the thermal mass,
which is sensed by some kind of thermometer, usually based on measuring resistance. Consider the
balance of heat in and heat out for the thermal mass. The heat absorbed per second is P , the radiation
power falling on the detector. Heat is meanwhile being conducted out via the thermal link. If this link
has thermal conductance G (heat conducted per per unit time per degree temperature difference) then
Astro Measurement D : Detection
17
Figure 12: Schematic of a typical bolometer system. Figure found at Wikimedia and used under the Creative
Commons Attribution license 3.0. Originally created by D.F.Santavicca (User “Tls60”)
the heat per second being lost is G∆T where ∆T is the temperature difference between the thermal
mass and the reservoir. The thermal mass will heat up until heat in = heat out, i.e. P = G∆T . The
equilibrium temperature difference is therefore
∆T =
P
.
G
So for a given radiation power, we need the conductance G of the link to be as low as possible, to give
a bigger effect. However this will also tend to make the system respond slowly to any change in P , as
we shall see below. This is a problem for standard observing techniques. Because the background is
very bright and variable at IR and submm wavelengths, the usual method is to switch rapidly (many
times a second) between the patch of sky containing the object being measured, and a neighbouring
patch of sky, repeatedly measuring the difference in radiation power from these two patches. (This is
known as “chopping”.)
So we need a bolometer system to respond on a short timescale. How long does it take to reach the
equilibrium ∆T derived above? The heat absorbed in time τ is Q = P τ . If the thermal mass has heat
capacity C = cm where c is the specific heat of the substance concerned, and m is its mass, then this
should result in ∆T = Q/C. We therefore find that the response time scale is
τ=
C
.
G
Given that we need G to be small for sensitivity, to achieve responsivity, we also need the heat
capacity C to be as small as possible. At high temperatures, most substances have roughly constant
specific heat (the Dulong and Petit law) but at very low temperatures c ∝ T 3 (the Einstein-Debye
law), so cooling the thermal mass makes a dramatic difference.
4.2
Measurement in bolometer systems
How do we measure a small change in temperature? The traditional method is to measure the change
in resistance. A device where resistance is dependent on temperature is known as a thermistor. Semi-
Astro Measurement D : Detection
18
conductors make excellent thermistors as well as photo-conductors; at higher temperatures, there is a
greater probability of thermal electrons making it into the conduction band.
New devices which are particularly sensitive rely on the phenomenon of superconductivity. A Transition Edge Sensor (TES) has a bilayer of a thin superconducting metal and a normal metal, and becomes superconducting below some transition temperature. If the operating temperature is near the
transition temperature, then resistance is extremely sensitive to tiny temperature changes. Originally
these devices were very noisy, but this problem has been solved, and so TES devices are becoming
very important.
Another new class of detector is the Kinetic Inductance Detection System (KIDS). As discussed in
section 3.8, these can be used in single photon-counting mode, but they can also be used in bolometric
mode, i.e. measuring the rate of overall incident light energy.
4.3
Operational issues for bolometers
Bolometers used for far-IR and submm detection are operated at very low (milli-kelvin) temperatures,
for three reasons. The first reason is to minimise the thermal background. The second reason is to
minimise the heat capacity of the thermal mass, so that it responds quickly to changes in heating
rate. The third reason is to maximise the sensitivity of the thermometer component to small changes
in temperature, and for TES sensors, to bring them near to the transition temperature. To get to
such very low temperatures, a simple bath of liquid helium won’t suffice, so one has to use other
(expensive) methods, like dilution refrigerators, or adiabatic demagnetisation refrigerators, which
use thermodynamic tricks to drive the temperature further down.
Even more than for IR arrays, the main problem with bolometers is money. CCDs are cheap because
the the huge electronics industry has driven down the price of components. For IR arrays, the devices
are expensive to construct, but at least the basic R&D has been done because of the civil and military
interest in IR detection. For submm bolometers, there are essentially no commercial or military
applications, so that all the research and development, as well as the construction, has to be done by
university and government labs.
5
Coherent wave detection systems
The basis of coherent detection of electromagnetic waves is simple, but a considerable amount of
complication is needed to arrive at a measurable signal. Typically, we need a concentration system,
an antenna, and a receiver. The receiver is itself a mult-stage device, as we shall discuss below.
The whole end-to-end system is sometimes known as a radiometer. For interferometer systems we
additionally require a correlator. Lets take a look at some of the key components.
5.1
Concentration and pointing system
At high radio frequencies, it is normal to collect and concentrate light with a parabolic dish, as
described in Chapter 3 and illustrated in Fig. 13. The angular response of the system will then be
determined by the diffraction pattern of the dish, and the sensitivity determined by the size of the dish.
Sometimes radio dishes are used in a prime focus arrangement, but more often they will be arranged
as Cassegrain systems, with a secondary dish feeding the waves to a feedhorn and a waveguide, with
a simple dipole at the base of this concentration system. One advantage of a dish is that it is steerable,
so that it can be pointed to the desired direction in the sky, and track a fixed position in celestial
co-ordinates as the sky rotates.
Astro Measurement D : Detection
(a) Simple dipole antenna
19
(b) Parabolic dish feed system
Figure 13: Two types of radio telescope. Left : At low frequency, radio waves are typically detected directly
by a dipole antenna, and the signal fed to the receiver system. Right : At high frequency, radio waves are first
concentrated by a parabolic dish and fed to the antenna via a feedhorn and waveguide. (Figure kindly provided
by Dr A.Woodcraft).
At low frequencies it is more normal to use antennas directly. A single antenna has a very broad
angular response, as discussed in Chapter 3. A dipole array can have a rather better angular response,
but still not good enough for most astronomical purposes. Usually therefore, astronomers use large
collections of antennas as interferometer arrays, so that the final angular resolution is determined by
the diffraction size of the whole array. As described in Chapter 3, phase delays can also be used to
effectively point such an array in the desired direction.
5.2
Antennas
The term “antenna” can mean anything that acts as a transducer between the waves in free space
and the receiver system. Usually in astronomy the antenna itself is a relatively simple device - at its
simplest, a conducting wire in which the electric field of the arriving wave drives a forced oscillation
in the conduction electrons, which corresponds to an AC current. The most common antenna is the
half wave dipole, illustrated in Fig. 13. Like any driven oscillator, the dipole has a resonant frequency,
given by ν = 4c/l where l is the length of one of the dipole sections. Sensitivity to other frequencies
falls off either side. The bandwidth can be defined in the usual way by the points at which sensitivity
falls to half (3dB down). It can be shown that this depends on the ratio l/d where d is the diameter of
the wire. Typical fractional frequency bandwidths are in the range 2–10%.
5.3
Receivers
A typical receiver can be seen as a chain consisting of an amplifier, a mixer, a detector, and a spectrometer, as illustrated in Fig. 14. The first task of the receiver is to amplify the signal, because the
voltages produced directly by the antenna are tiny.
Given the broad range of frequencies to which the antenna responds, a second task of the receiver
is frequency selection or tuning. In principle the amplification, filtering, and detection circuits could
all be tuned to a specific frequency, but this is very hard to do, especially at GHz frequencies. Furthermore electronic circuits are in general more stable and have better gain when operating at much
lower (MHz or kHz) frequencies. The solution is to include a second internal signal provided by
a local oscillator, and to mix this with the incoming radio frequency (RF) signal, producing a beat
signal at much lower frequency known for historical reasons as the intermediate frequency (IF). The
Astro Measurement D : Detection
20
Figure 14: Logical layout of the components of a heterodyne radio receiver system.
detection circuits can then work at a fixed frequency, and you tune the local oscillator so that the
desired component of the RF signal produces a beat signal at the fixed IF. This type of system, using
the mixing of two frequencies, is known as a heterodyne receiver. Historically, radio receivers were
designed to work with an IF at kHz frequencies, and this is still true for domestic radio receivers.
Modern electronics can work at much higher frequencies, but the principle of designing to use a fixed
frequency by mixing with a local oscillator remains. Real-world radio astronomy receivers in fact
typically use several local oscillators, making several different IFs for different parts of the system,
with frequencies ranging from kHz to GHz.
The third task is “detection” which in radio parlance means producing an output signal that is proportional to the received power. The antenna responds to the electric field amplitude of the incoming
wave by producing an oscillating voltage V . This switches direction during the cycle however, so
that the mean voltage is zero. Normal practice is to feed this voltage to a square law circuit which
multiplies the signal by itself and averages, producing an output proportional to the mean value of
V 2 . As well as being a well determined quantity, V 2 is proportional to the power of the incoming
wave (as opposed to the the electric field amplitude), which is the quantity we wish to measure.
The final stage is spectrometry. This could be achieved by a frequency sweep, or by autocorrelating
the signal, as we will discuss in Chapter 5.
5.4
Noise and calibration in receivers
A problem with measuring faint signals this way is that any circuit produces a signal simply due to
thermal noise, also known as Johnson noise or Nyquist noise. Thermal agitation means that even
when no voltage is applied across a resistor, there is always a voltage V , randomly fluctuating about
zero. Putting this fluctuating voltage through a square law detector produces a net signal. On short
timescales, the value of voltage from one moment of time to the next is correlated, but on long enough
timescales it is random, i.e. white noise. The power spectral density is per Hz of V 2 is
Φ(f ) =
2Rhf
hf
e kT −1
.
Here R is resistance in ohms. This is a flat spectrum up to a knee at frequency of f = kT /h. For
T = 293, f = 6 × 1012 Hz, corresponding to a timescale 0.16 picoseconds, and a light wavelength
of 49µm. So for radio frequencies, white noise is a very good aproximation. At low frequencies this
gives
Φ(f ) ' 2RkT.
Astro Measurement D : Detection
21
The total noise from zero to f = kT /h is therefore 2Rk 2 T 2 /h. The noise power is therefore sensitive
√
to temperature. Radio receivers normally work over some bandwidth ∆f so the RMS noise is V¯2 =
√
4kT R∆f .
The dependence of noise on temperature actually provides a method to calibrate a radiometer, by
switching between the external signal and a resistor of known temperature. It is common to characterise the strength of a signal by its effective temperature, i.e. the temperature that would given the
same signal strength just by noise alone. Note that because of square law detection, noise adds a bias
to the V 2 signal, rather than simply making a fluctuation about the expected signal - remember that
even a voltage with zero mean has a positive value of V¯2 . The relevant temperatures simply add. Suppose that, in the absence of any astronomical signal, the system produces power due to thermal noise
characterised by a “system temperature” Tsys . The astronomical signal produces adds an amount of
power that, if it had been due to thermal noise, would have looked like a temperature Tsig . Then the
net effect is observed power the same as would have been produced by a system with temperature
Tobs = Tsys + Tobs .