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High Energy Observational Astrophysics
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Introduction
In the 1st year you were introduced to astrophysics, stars and planets. You learned about galaxies, star
formation, coordinate systems, the difficulties of observing the Universe from our position on a rotating
sphere, and many other aspects of modern astronomy. The majority of the objects discussed, for example
stars and planets, are observable in the visible range of the electromagnetic spectrum.
Since almost all of the observations we can take of the Universe make use of photons or particles from
distant objects we would be massively restricting ourselves by only making use of electromagnetic radiation
between 400 and 700 nm!!!!! We address this with HEOA in year 2 and Radio and Infrared Astronomy in
year 3.
In HEOA we will be introduced to the big power players in the Universe – the X-ray and gamma ray
emitters. We will learn about the physical processes which drive the development of ever more sensitive
detectors and will find out about the various space missions and observatories which search the cosmos for
gamma ray bursts or exploding supernovae.
Although the majority of the course focuses on X-ray and gamma ray sources and their detection we will
also look at neutrinos, gravity waves and cosmic rays.
(* = need to memorise for exam)
Contents
1. Mechanisms for generating x-rays and gamma rays
1.1 Black body emission
1.2 Thermal Bremsstralung
1.3 Synchrotron radiation
2. Who are the usual suspects as sources for emission in the Universe?
2.1 Type I and II Supernovae
2.2 Neutron stars
2.3 X-ray pulsars
2.4 X-ray bursters
2.5 Gamma ray bursts
2.6 Quasars
2.7 Active Galactic Nucleii
2.8 Other X-ray sources
2.9 Note on X-ray source names
3. Problems associated with Earth based observation
3.1 Main sources of X-ray emission in the known Universe
3.2 Early attempts to observe high energy radiation
4. Techniques for detecting X-rays and gamma-rays
4.1 Interactions of photons with matter
4.1.1. Photoelectric effect:
4.1.2 Compton effect:
4.1.3. Pair production:
4.2 Bringing this all together and the mass absorption coefficient
4.3 Inelastic and elastic scattering
4.3.1 Rayleigh scattering
4.4 Detectors types used to detect x-rays and gamma rays
5 Main detector types
5.1 Proportional Counters
5.2 Scintillation Detectors
5.3 Influence of detector size and target material on energy spectrum produced
5.4 What is a semiconductor?
5.5 Semiconductor detectors
5.6 Microchannel plate detector
5.7 Multi Pixel Photon Counter (MPPC) and Charge Coupled Devices (CCDs)
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6. Imaging systems
6.1 Spark chamber
6.2 Compton telescope
6.3 Grazing incidence telescope
6.4 Rotational modulation collimator
6.5 Coded mask aperture collimator
6.6 A honeycomb collimator
6.7 A wire chamber telescope
7. Great Observatories
7.1 ROSAT (1990-1999)
7.2 Hubble Space Telescope (1990)
7.3 The Compton Gamma Ray Observatory (1991)
7.4 Chandra X-ray Observatory (1999)
7.5 Spitzer Space Telescope (2003)
7.6 Fermi Gamma ray Space Telescope (2008)
7.7 INTEGRAL (2002)
7.8 Constellation-X (2016)
8 Neutrino sources and detection methods
8.1 Stellar neutrinos
8.2 Supernova neutrinos
8.3 Detecting neutrinos
8.3.1 Neutrino capture reactions:
8.3.2 Neutrino scattering:
8.4 The solar neutrino problem
8.5 SNO and how it solved the solar neutrino problem
8.5.1 Charged current interaction
8.5.2 Neutral current interaction
8.5.3 Electron elastic scattering
8.5.4 Future possibilities for neutrino observations
9. Gravitational waves
10. Cosmic Ray Particles
10.1 So where do cosmic ray particles come from?
10.2 How old are cosmic ray particles?
10.3 What is the total flux of cosmic ray particles?
10.4 How do we detect cosmic rays?
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1 Mechanisms for generating x-rays and gamma rays
Before looking at the objects themselves, we first need to understand the processes responsible for photon
production so that we can interpret the physical conditions in the object. The three main sources are:
1.1 Black body emission
E photon  h f where h is known as Plank’s constant (6.626×10-34Js) and f is the frequency of the light in Hz.
The change in energy E of a charged particle q as it moves through a potential difference V is given by
E  qV . 1 electron-volt =1 eV =1.610-191 = 1.610-19 joules.
E photon  h f 
ch

(*) So a 1 nm light source has an energy of E photon 
ch

 1.99  10 13 J  1.24 MeV
The Maxwell-Boltzmann distribution defines the distribution of an amount of energy between identical but
distinguishable particles and gives the probability of a
particle in a gas at temperature T having a particular speed
(v) assuming thermal equilibrium.
The mean energy of a particle in a gas in thermal
equilibrium at temperature T is given by:
3
k B T (*) where kB is known as the Boltzmann
2
constant (1.38×10-23 JK-1) and T is the mean temperature of
the object in kelvins.
E avge 
When an object is heated, it radiates energy in the form of photons. The energy distribution of these photons
is determined by the temperature of the object whereas the intensity of the radiated heat is determined by the
number of photons emitted per unit time. As the temperature of the object increases the mean photon energy
shifts to higher energies. This is seen as a change in colour as an object is heated from red (long wavelength
/ low energy) to blue hot (shorter wavelength / higher energy). But this shift in mean photon energy is not
restricted only to the visible part of the electromagnetic spectrum. Very low temperature objects emit
principally in the infrared or microwave regions. At the other end of the spectrum very hot celestial bodies
emit photons peaked in the ultraviolet or even X-ray region of the electromagnetic spectrum.
This continuous
spectrum produced by
hot bodies is known as
the black body
spectrum and is used to
describe the behaviour
of any object that is a
perfect absorber or
emitter of radiation.
Black body emission
spectra have the same
characteristic shape
without absorption or
emission lines, the peak
photon energy
dependent only on the
temperature of the object as shown below.
The spectrum of wavelengths emitted by a body at a temperature, T, has a characteristic shape that is
strongly dependent on the wavelength (to the inverse fifth power). The figure shows the black-body
radiation curves, showing the wavelength distribution of emitted photons at different temperatures. Planck’s
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law describes the spectral distribution of radiation emitted by a black body. Very hot bodies (3000 - 20,000
K) like our Sun emit light at visible wavelengths. The Sun acts like a black body near 6000K, whereas the
Earth acts like a black body near 300 K.
The monochromatic flux from the surface of a black body emitter can be described by Planck’s Law for
which F(λ, T) is the amount of radiant energy emitted as a function of wavelength and temperature of the
object, and h and c are Plank’s constant and the speed of light.
F ( , T ) 
2 h c 2
5 e hc kT  1


in Wm-2nm-1.
The dependence of peak wavelength on temperature is found by differentiating the equation above with
respect to wavelength. This is known as Wien’s displacement law and is given by:
(*)
max T  2.898  10 3 in mK.
This is roughly the same as E photon  kT .
The Sun has a peak wavelength of 500nm. What is its effective temperature?
X-rays from about 0.12 keV to 12 keV are classified as "soft" x-rays. Using Wien’s displacement law this
corresponds to peak emission from objects between 280 kilokelvin and 28 megakelvin. All main sequence
stars reside in this range or below it.
X-rays from about 12 keV to 120 keV are classified as "hard" X-rays and correspond to temperatures
between 28 megakelvin and 280 megakelvin. Neutron stars at 107 kelvin reside in this range.
Gamma rays range from about 120 keV to 30 MeV corresponding to temperatures up to 7×1010 kelvin.
Research is ongoing into possible sources observed.
1.2 Thermal Bremsstralung
A charged particle undergoing an acceleration radiates photons. An example
of this is when electrons moving back and forth in antennae produce
electromagnetic radiation, such as transmitted by radio stations.
The photons are emitted in a characteristic double lobed form.
Bremsstrahlung, or braking radiation, is emitted by a charged particle
accelerating under the influence of another charged particle. The particle
emits energy in the form of electromagnetic radiation, at the expense of its
kinetic energy.
This is the reason Rutherford’s
model of the atom was flawed.
When an electron with velocity v,
passes a charge consisting of Z
protons, with total charge Ze+.
The impact parameter b of the
interaction represents the distance
of closest approach.
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The electron is accelerated during its interaction, and since the acceleration is not uniform it emits photons
with a range of wavelengths, i.e. a spectrum. The Bremsstrahlung spectrum becomes more intense and shifts
towards higher frequencies when the energy of the accelerated particles is increased.
If we now generalise the results to a population of electrons
with a certain velocity spread given by the MaxwellBoltzmann distribution at some uniform temperature, the
total emission by all particles in this population is called
thermal Bremsstrahlung.
Consider radiation from hot ionised gas at temperature T in
the vicinity of a star. In plasma, the free electrons are
constantly producing Bremsstrahlung radiation during
interactions with the ions.
It can be shown that the intensity
1
 E 

I  kT  2 ne ni Z 2 g B exp  
 where, gB is the gaunt
 kT 
factor and E is the photon energy.
1.3 Synchrotron radiation
Synchrotron radiation is very similar to Bremsstrahlung radiation, the difference being that in the former
charged particles are accelerated (either in a straight line or around a circle) whereas in the latter they deflect
from other charged particles (usually ions). In astrophysics the most common example of synchrotron
radiation occurs when fast moving electrons pass through a region in which there is a magnetic field causing
them to spiral around the field lines emitting radiation. The energy (or frequency) of the radiation depends
on the strength of the field and the component of the electrons motion perpendicular to it. If again we
assume a Maxwell-Boltzmann distribution of velocities then the intensity of the radiation of energy E
emitted can be shown to be:
I  AE 
where α and A are constants.
The intensity therefore follows a power law distribution as
shown below.
Please note that it is very possible to confuse
Bremsstrahlung and Synchrotron radiation signals at high
energy. To avoid this we must also measure the flux at low
energies. Furthermore synchrotron radiation is usually polarised unlike Bremsstrahlung radiation.
In the diagram the blue light from the jet emerging from M87, towards the lower right, is due to synchrotron
radiation.
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2. Who are the usual suspects as sources for X-ray and gamma ray emission in the Universe?
2.1 Type I and II Supernovae
Supernovae are some of the most spectacular astrophysical objects. For a period of days to weeks, a
supernova may outshine an entire galaxy. Emitting around 1046 joules (more than the Sun will emit in its
entire lifetime) in the space of a few weeks their optical luminosity at its peak can be 109 times the Sun’s.
Astronomers recognise two types based on their characteristic light curves; type I and type II supernovae.
Let’s have a brief recap…If a main sequence star has a mass of
over 8 times the mass of the Sun it is destined to be a type II
supernova, such as SN1987A shown in the figure. The pressure
produced by nuclear fusion in the core prevents gravitational
collapse. When all the hydrogen is fused the star collapses and
the core heats until it is possible for heavier elements to fuse.
This continues until iron is formed after which point fusion
actually absorbs energy rather than produces it. The star
immediately pre-supernova stage has an onion-like layer
structure. Gravity now takes over and now there is no energy
generation to prevent catastrophic collapse. In just a matter of
seconds the core shrinks from roughly 8000 km diameter to just
20 km, and the temperature rises 100 billion degrees or more. As the density rises protons and electrons
collide to form neutrons and a vast number of neutrinos.
(*)
p  e   n  e
Eventually a neutron star is formed if the core is less than three solar masses. Neutron stars are incredibly
dense - similar to the density of an atomic nucleus. They are typically 20 km across but more massive than
the Sun. Because they contain so much mass within such a small volume, the gravitational field at the
surface of a neutron star is immense. Neutron stars also have powerful magnetic fields which can accelerate
atomic particles around their magnetic poles producing powerful beams of radiation. These sweep around
like massive searchlight beams as the star rotates. If such a beam is oriented so that it periodically points
toward the Earth, we observe it as regular pulses of radiation that occur whenever the magnetic pole sweeps
past our line of sight. In this case, the neutron star is known as a pulsar.
For now let’s forget about the neutron
star at the centre and let’s consider the
outer layer of the star. Initially these
layers begin to collapse along with the
core. However the shockwave
rebounding from the core hits these outer
layers with enormous energy and they
are ejected violently outward creating
explosive nucleosynthesis. This shockfront continues to expand into interstellar
space. The majority of heavy elements
are created in supernovae, either in the
fusion stage as the supermassive star was
compressing down, or in the explosive
stage via nucleosynthesis as the
fragments were blown out into
interstellar space.
The figure shows an example of light emission curves from type I and type II supernovae. The key
differences are that whereas type I supernovae curves all have similar shape and maximum luminosity, type
II are much more variable and are dominated by emission lines from hydrogen. Hydrogen emission lines are
absent in type I supernovae. Astronomers have created models to try to explain these curves. They explain
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the variability in the light curves as being dependent on the variable size and composition of the star, and
explain the hydrogen emission lines as being created by the advancing shockwave. The similar type I light
curves are explained as being representative of a common origin.
On average, a supernova explosion occurs about once every hundred years in a typical galaxy. About 25 to
50 supernovae are discovered each year in other galaxies, but most are too far away to be seen without a
telescope. The brightest supernova in recent history is SN1987A which exploded in the Large Magellanic
Cloud in 1987.
So what is a type I supernova? Whatever mechanism we choose
has to explain the consistent light curve and peak luminosity.
The favoured model is the detonation of a white dwarf star in a
binary system.
The material in a white dwarf no longer undergoes fusion
reactions, so the star has no source of energy and is not
supported by the heat generated by fusion against gravitational
collapse. It is supported only by electron degeneracy pressure,
causing it to be extremely dense. The physics of degeneracy sets
a maximum mass for a white dwarf, the Chandrasekhar limit, at
approximately 1.4 solar masses. Beyond this it cannot be
supported by degeneracy pressure alone.
A type Ia supernova occurs when a white dwarf star acquires
additional mass by siphoning matter away from a companion star,
usually a red giant. This occurs in a binary system where one star
has evolved to a white dwarf whilst the other has become a red
giant. If close enough hydrogen rich material is pulled off the
giant and spirals across to the dwarf. When it reaches a critical
mass about 1.4 times the mass of the Sun, packed into a radius
equivalent to that of the Earth, the heat and pressure in the centre
of the star are sufficient to initiate runaway nuclear fusion, and
the white dwarf explodes. An example of this is SN 1006. Its
supernova remnant is shown in the figure right. But why do we
get runaway fusion in a white dwarf but not in a main sequence
star? In a normal star comprised of ordinary matter fusion in the centre would result in an increase of
pressure, the star would expand and cool. A white dwarf however is comprised of degenerate matter which
doesn’t follow gas laws. Increased temperature just makes the reaction occur faster resulting in runaway.
Since the initial conditions are about the
same in all cases, these supernovae tend to
have similar luminosity, and their ‘light
curves’ (how the luminosity changes over
time) are predictably similar as shown in
the figure for 6 separate type I supernovae.
They can therefore be recognised from the
characteristic shape of their light curves
and fixed peak brightness.
The similarity in the absolute luminosity
profiles of nearly all known Type I
supernovae has led to their use as standard
candles in extragalactic astronomy.
X-rays and gamma rays are produced during nucleosynthesis and also during radioactive decay of the
various products. They also originate from the shockwave which heats interstellar gas creating thermal X-ray
emission.
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2.2 Neutron stars
A neutron star is a supernova remnant following gravitational collapse of a massive star. Composed almost
entirely of neutrons, they are very hot. A typical neutron star has a mass roughly equivalent to 1 solar mass,
with a corresponding radius of about 15 to 30 km. Neutron stars have overall densities of 1017 kg/m3 which
are very similar to the density of an atomic nucleus of 3×1017 kg/m3. This density is approximately
equivalent to the mass of the entire human population compressed to the size of a sugar cube.
As the core of a massive star is compressed during early stages of a supernova, angular momentum and
magnetic flux are conserved. Since both the angular velocity and magnetic flux are inversely proportional to
the square of the radius, this contraction results in rotation periods between about 1.4 ms to 30 seconds and
magnetic fields of up to 108 T.
The temperature inside a newly formed neutron star is around 1012 kelvin. However, the huge number of
neutrinos it emits carries away so much energy that the temperature falls within a few years to around 106
kelvin.
At 1 million kelvin, what is the peak wavelength emission and what kind of radiation is this?
The outer crust of a neutron star is composed of ordinary
atomic nuclei crushed into a solid lattice with a sea of
electrons flowing through the gaps between them. Inside
this the inner crust is made up of neutron rich isotopes of
ordinary nuclei in a sea of neutrons and electrons. Inside
this the outer core is comprised principally of neutrons
with a small number of protons and neutrons. Below this
the inner core is filled with neutrons.
Let’s now look at two examples of neutron star behaviour
which can be studied via X-ray and gamma ray
observations.
2.3 X-ray pulsars
Although neutron stars do emit blackbody radiation, they
are not just spherical blackbody emitters like other stars.
Neutron stars have additional ways of emitting
electromagnetic radiation. The strong magnetic field, which
is bonded to the star’s ionised gas, coupled with the rapid
rotation of a neutron star make it a very potent electrical
generator. (Here on Earth, commercial electrical generators
work by rotating a series of magnets inside a coil of wires.
The essential point is that you need to have a magnetic field
in motion). The magnetic field generated by the rotating
magnetized neutron star is strong enough to rip charged
particles (such as electrons) away from the surface of the neutron star.
The charged particles follow the magnetic field lines to the north and south magnetic poles of the neutron
star.
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The accelerated particles produce intense but narrow beams of synchrotron radiation, pointing away from the
two magnetic poles. We can see one of these beams of light ONLY if it is pointing toward us, just as we see
the light from a lighthouse only when it is pointing toward us. Remember that synchrotron radiation is
emitted whenever high speed electrons move along curved paths through a magnetic field. Synchrotron
radiation has a continuous spectrum but one that is distinctly different from a black body spectrum. This
helped astronomers first identify for example that the light from the Crab Nebula was primarily synchrotron
radiation and not black body emission.
In many cases pulsars are found within the
remnant of a supernova, the exception
being when the supernova explosion is
asymmetric and the neutron star is fired
out of the supernova remnant. An example
is the Crab pulsar which is at the centre of
the Crab nebula which was formed
following the massive supernova of 1054.
The incredible images above show
emission from the Crab nebula in different
spectral ranges. X-ray emission (Chandra
X-Ray Telescope) shows how the pulsar’s
magnetic field funnels outgoing charged
material into two oppositely directed jets.
It also shows how fast moving material
from the pulsar creates additional
shockwaves forming X-ray emitting rings.
The Crab pulsar is one of the most rapidly
rotating pulsars with a period of 33
milliseconds. NB. Pulsars emit across the
entire EM spectrum and not just in X-rays.
A complicating factor is that on a neutron star, just as on Earth, the magnetic poles don't coincide with the
rotational poles. Thus, the beams of radiation pointing away from the magnetic poles are at an angle to the
rotational axis of the neutron star; as the neutron star rotates, the beams swing around in a cone.
If a beam happens to sweep across our location in space, we
see a brief flash of light. Neutron stars whose beam of
electromagnetic radiation happens to sweep across us are
called pulsars. The pulses come at the same rate as the
rotation of the neutron star, and thus, appear periodic. The
most rapidly rotating neutron star observed in 2004 in
Sagittarius, PSR J1748-2446ad, rotates at 716 revolutions per
second.
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2.4 X-ray bursters
Not all neutron stars are destined to lead a life of isolation. Some of them are born in binaries that survive
the supernova explosion that created the neutron star, and in dense stellar regions such as globular clusters
some neutron stars may be able to capture companions. In either case, mass, mostly hydrogen, may be
transferred from the companion to the neutron star (or black hole).
X-ray bursters differ observationally from X-ray pulsars in that they show a sharp rise time (1 - 10 seconds)
followed by peak emission which can last around a minute, followed by a gradual diminishment in emitted
light. The ratio of the burst flux to the persistent flux can be up to 103 but is typically about 100. Total output
is around 1031 J in the form of X-rays.
It is believed that accretion of hydrogen onto the
surface of a neutron star produces a layer where
hydrogen fusion takes place, creating a build up
of helium. Eventually the temperature in the layer
reaches the critical point at which helium fusion
occurs to form carbon and oxygen with the
release of more energy. Normally this increased
energy production would raise the temperature
and pressure and cause the material to expand and
cool. However ideal gas laws don’t apply on the
surface of a neutron star and a temperature
increase doesn’t lead to a pressure increase which
would otherwise have caused the layer to expand
and cool stopping the reaction.
This ‘failure’ in the feedback loop leads to thermal runaway leading to a thermonuclear flash. Energy is
released as high energy photons via black body radiation. This release of energy may be observed as an
increase in the star's luminosity with a space telescope, and is called an X-ray burst. The figure shows an Xray burst in Terzan 2 which is in Sagittarius.
Eventually so much energy is
produced that the pressure
begins to rise and material is
blown off the neutron star
surface and cools. Bursts are not
powerful enough to disrupt the
stability or orbit of either star,
and the whole process may
begin again. Most X-ray
bursters have irregular periods,
which can be on the order of a
few hours to many months,
depending on factors such as the
masses of the stars, the distance
between the two stars, the rate
of accretion, and the exact
composition of the accreted
material.
Confirmation of this model
comes from observations that
the longer the gaps between
bursts the more energy is
released. This is attributed to the
greater opportunity for fuel to accumulate. A few dozen X-ray bursters have been discovered in our galaxy.
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2.5 Gamma ray bursts
Gamma ray bursts (GRBs) are flashes of gamma rays associated with extremely energetic processes in
distant galaxies. The sources of most GRBs are billions of light years away from Earth, implying that their
origin is both extremely energetic (a typical burst releases 1044 joules, as much energy in a few seconds as
the Sun will in its entire 10 billion year lifetime) and extremely rare (a few per galaxy per million years).
They are the most luminous electromagnetic events known to occur in the Universe.
Although a typical burst lasts a few
seconds (short GRB), they can last as long
as a few minutes (long GRB) and the initial
burst is usually followed by a longer-lived
"afterglow" emitted at longer wavelengths
(anything between X-ray and radio
emission). Unlike X-ray bursters, sources of
both short and long GRBs only appear to
emit once in their history. Observed GRBs
are distributed isotropically across night sky
suggesting an extragalactic origin. New
GRB sources are being discovered at a rate
of about one per day.
GRBs were first detected in 1967 by the Vela satellites, a series of satellites designed to detect covert nuclear
weapons tests. Little information was available to verify their source until direct measurement of their
redshifts was made using optical spectroscopy clarifying their distance and luminosity and connecting them
to the deaths of massive stars. The dimmest gamma ray burst recorded so far has a redshift of z = 6, implying
that it was created only 750 million years after the big bang. If we can understand these objects, they could
show us the earliest states of the Universe.
Initial questions were concerned with
the sheer amount of power emitted,
no known process in the Universe
able to produce as much energy.
However, GRBs are now thought to
be highly focused explosions, with
most of the energy collimated into a
narrow jet from between 2 and 20
degrees in angular spread. Because
their energy is strongly beamed, the
gamma rays emitted by most bursts
are expected to miss the Earth and
never be detected. When a gammaray burst is pointed towards Earth,
the focusing of its energy along a
relatively narrow beam causes the
burst to appear much brighter than it
would have been were its energy
emitted spherically. When this effect
is taken into account, typical gammaray bursts are observed to have a true
energy release of about 1044 J. This is
comparable to the energy released in
a bright type I supernova. Indeed very bright supernovae have been observed to accompany several of the
nearest GRBs.
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The most widely-accepted mechanism for the origin of long lasting GRBs is that they are created in a
hypernova. Hypernovae are believed to result from the collapse of stars of greater than 30 solar masses
which are spinning rapidly. The black hole at the centre forms before the outer layers of the star have a
chance to contract very much. As a result the black hole quickly forms an accretion disc from the
surrounding stellar material. Some of the infalling material does not fall into the black hole but is ejected in
powerful back to back jets along the axis of rotation of the accretion disc. These jets of particles blow the
star apart. If one of these jets is directed towards Earth we would see a powerful GRB formed as the
relativistic particles slow down and convert their kinetic energy to gamma rays. The jets only have a short
existence as it doesn’t take long for the black hole to eat the accretion disc!
Short duration gamma ray bursts occur in systems with no star formation and where no massive stars are
present. The favoured theory for the origin of most short GRBs is the merger of a binary system consisting
of two neutron stars. According to this model, the two stars in a binary slowly spiral towards each other due
to the release of energy via gravitational radiation until the neutron stars suddenly rip each other apart due to
tidal forces and collapse into a single black hole. The infall of matter into the new black hole in an accretion
disk then powers the GRB.
Currently orbiting satellites detect an average of about one GRB per day. But because GRBs are visible to
distances encompassing most of the observable universe, a volume encompassing many billions of galaxies,
this suggests that they must be exceedingly rare events per galaxy. For a galaxy of approximately the same
size as the Milky Way, the expected rate (for long GRBs) is about one burst every 1,000,000 years.
2.6 Quasars
At first glance quasars and GRBs sound identical. However whereas quasars are thought to be powered by
accretion of material onto supermassive black holes in the centres of distant galaxies, GRBs are connected
with the death throes of massive stars going hypernova.
The nature of Quasars (quasi-stellar radio
sources) has only recently been
understood and prior to the 1960s their
origin was one of the biggest mysteries in
astrophysics. The typical spectra from
stars or galaxies, as already seen, is that of
any black body emitter with absorption
lines superimposed dependent on the type
of gas through which the light passes. But
light from these mysterious objects is very
different.
The spectra are not peaked in any one area
of the electromagnetic spectrum but
instead extended over a vast range from
X-ray to radio wave. The spectral lines
were also a mystery. Whereas the spectra
from stars contain absorption lines, the spectra from quasars consisted of emission lines as shown in the
figure right produced by the quasar 3C 273. Furthermore these emission lines do not align themselves with
the standard emission lines produced by any known chemical tested on Earth.
A breakthrough came in the mid 60’s when Schmidt realised that 4 of the absorption lines from a quasar
called 3C 273 located in the constellation of Virgo, were spaced relative to each other in exactly the same
ratio as 4 of the lines of the hydrogen Balmer series, the difference being that the lines produced by the
quasar had all been shifted to much longer wavelengths as shown right.
This provided astrophysicists with the means to deduce the Doppler shift parameter z for the light,
calculating a value of z = 0.158 for 3C 273. This is a massive value corresponding to a recessional velocity
of 4.4×107 m/s. Moving at this velocity it was impossible that the source of light could be produced from
within the Milky Way. In fact, according to Hubble’s Law, the quasar 3C 273 was actually 630 Mpc away
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from Earth. This provided a means of estimating the luminosity based on the brightness measured on Earth.
The result was astounding. Although many had previously believed quasars to be a new type of star, the
reality was that 3C 273 and others like it were each emitting over 100 times as much energy as our own
galaxy with its 200 billion stars!!!!! The mass of 3C 273 has since been measured to be approximately 900
million solar masses. Quasars were therefore something new!!!!!
Quasars are ultra luminous objects located at the centre of very distance galaxies. At the centre of every
quasar is a supermassive black hole. As matter is pulled into the black hole it releases vast amounts of
energy which explains both the emission lines
and the lack of a traditional black body
emission spectrum. If there is a strong magnetic
field in the accretion disc, the magnetic field
lines are aligned running toward the poles
perpendicular to the accretion disc. Gas moving
along this magnetic field then produces beams
of electrons and plasma seen in the jets to make
the radio lobes of radio emission. The figure
shows the false colour image of radio emission
from the quasar Cygnus A. A strong red shift
indicates that it is 230 Mpc from Earth, the
emission jets extending 70 kpcs in each
direction. More than 100,000 quasars have been discovered with luminosities ranging from 1038 to 1042
watts. The Milky Way in comparison has a luminosity of 1037 watts.
2.7 Active Galaxies
Active galaxies are galaxies which have a small core of emission embedded in an otherwise typical galaxy.
This core is called the active galactic nucleus (AGN). This core may be highly variable and very bright
compared to the rest of the galaxy. Models of active galaxies concentrate on the possibility of a
supermassive black hole which lies at the centre of the galaxy and the AGN is therefore typically a quasar.
For "normal" galaxies, we can think of the total energy
they emit as the sum of the emission from each of the stars
found in the galaxy. For the "active" galaxies, this is not
true. There is a great deal more emitted energy than there
should be... and this excess energy is found in the infrared,
radio, UV, and X-ray regions of the electromagnetic
spectrum.
There are several types of AGN although most scientists
believe that, even though these types look very different to
us, they are really all the same thing viewed from different
directions!
The photo to the right shows the active galaxy known as
Circinus.
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2.8 Other X-ray sources
There are of course other types of astrophysical object which emit X-rays via black body emission.
Galactic clusters are formed by the merger of individual galaxies
or groups of galaxies. The infalling material (which contains
galaxies, gas and dark matter) gains kinetic energy as it is
gravitationally attracted colliding with gas already in the cluster
causing it to be heated to between 107 and 108 K depending on
the size of the cluster. This very hot gas emits X-rays. The figure
to the right shows the Bullet cluster imaged by the Chandra Xray Observatory. The colour represents X-ray emission.
X-ray photons from the Sun were first detected using a pin-hole camera attached to a rocket in 1949. It had
not been predicted that the solar corona would be 1000 times hotter than the effective temperature of the Sun
and so the detection of X-rays was a surprise (remember atmospheric absorption of X-rays from Intro to
Astro course and the lack of any satellites in 1949). While neutron stars and black holes are the
quintessential point sources of X-rays, all main sequence stars are likely to have hot enough coronae to emit
X-rays. The planetary nebula around a white dwarf is also a weak source of X-rays.
(left) X-ray image of the SN 1572 Type Ia remnant as seen by Chandra Space Telescope. (right) X-ray
image of solar corona taken by the soft X-ray telescope on board the Yohkoh solar observatory spacecraft.
2.9 Note on X-ray source names
Initial approach was to name new X-ray sources according to the constellation in which they were found.
E.g. Cyg X-3 is the 3rd brightest X-ray source in Cygnus. Problems arose however when transient sources
were observed. For example Cen X-3 is now the brightest source in Centaurus. Current convention is to
name the source according to its position in the sky with a prefix to indicate which satellite or instrument
made the observation. E.g. HO324+28 tells us that this source was first observed by the High Energy
Astrophysics Observatory 1 satellite at a location of 3hr24min R.A. and +28deg dec.
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3. Problems associated with Earth based observation
The most luminous objects in the discovered universe are Quasars (1039-1042 W). There are however several
difficulties associated with measuring brightness including attenuation of light by dust and gas in the
interstellar medium, and by the Earth’s atmosphere. Only the visual, infrared, and radio regions of the
electromagnetic spectrum are directly observable from the Earth’s surface.
This is problematic since many very hot celestial objects emit X-rays. Optical frequencies are used to image
stars or patches of dust, whereas infrared frequencies are used to image the same objects to provide more
detailed information of any heat sources. Since radio waves penetrate dust, scientists use radio astronomy
techniques to study regions that cannot be seen in visible light, such as the dust-shrouded environments
where stars and planets are born, and the centre of our Galaxy, the Milky Way. The other frequencies are
attenuated by the Earth’s atmosphere or ionosphere, and special effort must be made to detect them.
Most visible and radio
wavelengths do reach the
ground (see "optical window"
and "radio window") and can be
observed by ground-based
telescopes. A limited amount of
infrared (IR) and ultraviolet
(UV) light also reaches the
ground. Key space telescopes
are: Compton Gamma Ray
Observatory (CGRO), Chandra
X-ray Observatory, Hubble
Space Telescope (HST), and
Spitzer Space Telescope.
High Energy Observational Astrophysics 2010
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Copyright: Chandra mission website and Space Telescope Science Institute
To observe other wavelengths,
instruments are placed in:





Satellites
Rockets
Balloons
Aircraft
Mountain Top Observatories
The figure shows the altitude by which
half of the incoming radiation has been
absorbed as a function of wavelength and
also the type of observation technique
which would be appropriate. A major
problem for X-ray and gamma ray
observation is that they are absorbed on
passing through our atmosphere. This
attenuation is exponential rather than
linear, the effect being that the flux is
most quickly attenuated on encountering
the outer layers.
Hence we have to get above the Earth’s atmosphere. Our options are…
1. A balloon. (only useful to observe very hard gamma rays)
2. A rocket.
3. A satellite.
But cost is a factor…
Experiment
Mountain observatory
Aircraft
Balloon
Rocket
Satellite
Total cost
£2,000,000
£240,000
£300,000
£500,000
£200,000,000
High Energy Observational Astrophysics 2010
Duration
10 years
1 day
1 day
10 minutes
5 years
Cost per hour
£50 per hour
£10,000 per hour
£12,500 per hour
£3,000,000 per hour
£10,000 per hour
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Already we see that high energy observational astronomy is a costly business compared to a few hundred
pounds per night for optical observing. So why bother?? Added to this most astronomers until the 1960s did
not believe that there was any demand for high energy photon observation. Why?? Well, if you remember
Wien’s displacement law then you know that peak wavelength and black body temperature are inversely
proportional. If you know that the Sun has an effective temperature of 5800 K and a peak wavelength in the
optical region of 500 nm then the effective temperature of a gamma ray emitter must be approximately
10,000 times the temperature of the Sun!! Many astronomers doubted that objects at such temperatures (107
K) could exist. In the 1960s the first discrete powerful astronomical sources were observed and all this
changed.
3.1 Main sources of X-ray emission in the known Universe
The figure below shows the strength of signal from low and high energy photon sources as a function of
galactic longitude.
3.2 Early attempts to observe high energy radiation
Early rocket flights:
 1946-1952 Naval Research Laboratory (NRL) in USA using captured
V2 rockets detects UV and X-ray emission from the Sun.
 1962 American Science and Engineering group (ASE) launches
Aerobee rocket and detects first X-ray source (Sco X-1).
 1964-1970 Numerous balloon and rocket flights detect other sources.
The figure below shows the variation in counting rate as a function of galactic
longitude from a rocket borne proportional counter flown in 1967. The hard
line represents the expected distribution based on known sources whilst the
circles represent the data obtained in that flight.
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But rockets had short duration so could not perform detailed scans and position resolution was very poor so
it was difficult to know where exactly any signal originated.
Early satellite flights:
 The first artificial satellite, Sputnik 1, was launched by
the Soviet Union in 1957.
 The space race initiated by Sputnik consigned research
satellites to low priority. Nonetheless the technology
associated with both the satellite and the detectors
improved quickly. Uhuru, also known as the Small
Astronomical Satellite 1 (SAS-1) was the first earthorbiting mission dedicated entirely to celestial X-ray
astronomy. It was launched in 1970 and operated for 3
years. It consisted of two proportional counters with
honeycomb collimators giving both wide (5º by 5º) and
narrow (0.5º by 5º) field of view. With these it made the
first comprehensive and uniform all sky survey.
The figure shows Uhuru (Copyright R. Giacconi and colleagues). Its detectors surveyed in the energy band
from 2 to 20 keV and could detect sources as weak as 1/10,000th the strength of Sco X-1. Sun and star
sensors onboard allowed its position to be accurately determined. Uhuru spun slowly, usually making one
revolution every 12 minutes whilst mapping out a scan of space either 0.5º or 5º wide.
The most significant discovery was probably the observation of a time variation in the strength of the signal
from Her X-1. This led to the realisation that this was a binary system composed of a neutron star accreting
matter from a normal star. As can be seen the source exhibited complex time variability, pulsing with a
period of 1.24 s due to the rotation of the neutron star, eclipsing every 1.70 days with the period of the
binary orbit, and also varying with a 35-day period believed associated with the precession of the accretion
disk. Binary star systems contain two stars that orbit around their common centre of mass. X-ray binaries are
made up of a normal star and a collapsed star (a white dwarf, neutron star, or black hole). These pairs of
stars produce X-rays if the stars are close enough together that material is pulled off the normal (companion)
star by the gravity of the dense, collapsed star. The X-rays come from the area around the collapsed star
where the in-falling material is heated to very high temperatures. This exchange of mass may eventually
bring their evolution to stages that single stars cannot attain and binary stars are the progenitors of both
novae and type Ia supernovae.
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Copyright : E. Schreier
Copyright : NASA

The second NASA Small Astronomy Satellite (SAS-2) was dedicated to gamma-ray astronomy in the
energy range above 35 MeV. SAS-2 was launched in 1972 and ran until1973 when a failure of the
low-voltage power supply ended the collection of data. Its detector was a 32-level wire sparkchamber aligned with satellite spin axis with an effective area of 540 cm2. SAS-2 provided the first
detailed look at the gamma-ray sky and established the high energy component of diffuse celestial
radiation. It also correlated the gamma-ray background with galactic structural features.
The figure below shows the distribution of gamma-rays above 100 MeV along the galactic plane.
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
COS-B, launched by the European Space Agency, provided more detailed mapping of gamma-ray
sources (~30 MeV-5 GeV) throughout the Universe. COS-B carried a single large experiment, the
Gamma-Ray Telescope comprising a spark chamber and a proportional counter. Launched in 1975,
COS-B operated until 1982. It’s highly elliptical orbit enabled long observation times enabling more
detailed mapping.
The figure below shows the gamma-ray intensity contours produced by COS-B (top 300 MeV to 5 Gev;
middle 150 MeV to 300 MeV; bottom 70 MeV to 150 MeV)
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
Vela satellites were
operated jointly by the
U.S. Department of
Defense and the U.S.
Atomic Energy
Commission, and
managed by the U.S. Air
Force. They were not
intended primarily for
astronomical studies but
rather to search for
clandestine nuclear
bomb tests. They did
however provide much
useful astronomical data.
Operated between 1967
and 1979 they detected
gamma-ray bursts (0.2 to
1.5 MeV) of 1 second
duration. Triangulation
showed these were not
confined to the galactic
plane and so must be
extra-galactic in origin.
Gamma-ray bursts are narrow
beams of gamma rays associated with extremely cataclysmic events in distant galaxies. They are the most
luminous electromagnetic events known to occur in the universe and can last from milliseconds to several
minutes, although a typical burst lasts a few seconds. Most observed GRBs are believed to be released
during a hypernova event, as a rapidly rotating, high-mass star collapses to form a black hole. All observed
GRBs have originated from outside the Milky Way. The plots above show the range of GRB profiles
recorded.
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4. Techniques for detecting X-rays and gamma-rays
4.1 Interactions of photons with matter
In order to make detectors we need to understand the way in which photons interact with matter. We can
split this interaction into three separate processes…
4.1.1. Photoelectric effect: A photon interacts with an atom in the detector material. The photon is absorbed
and energy is given to an electron which is emitted. This is called a photoelectron. The binding energy is the
minimum energy E0 required to release a photoelectron from a metal. When a photon is absorbed within a
metal, some of the photon’s energy will be used up in freeing the photoelectron from the metal (E0), and if
there is any energy remaining, then this will appear as kinetic energy of the ejected photoelectron.
So we can write Ekinetic energy  h f  E0 (*)
Key points of the photoelectric effect:a) The fact that the number of photoelectrons increases with the intensity of the radiation is explained by
each photon liberating exactly one photoelectron. A higher intensity of light implies that more photons are
present and so more photoelectrons are ejected.
b) The fact that the maximum kinetic energy of the photoelectrons depends on the frequency of light is
explained because photons corresponding to light of a higher frequency carry more energy. So after E0 has
been used up, there is more energy left over to appear as kinetic energy of the photoelectron.
c) The fact that there is a lower limit for the frequency of light, below which no photoelectrons are emitted is
due to the fact that since the minimum energy required to eject an electron is E0, then the minimum
frequency of light needed to do this is E0  h f .
The likelihood or probability that this interaction will occur is called the cross section (σ) and depends on the
energy of the photon and the Z (atomic number) of the detector atom.
For a fixed Z
  E3.5
For a fixed Eγ
  Z5
So the photoelectric effect is heavily favoured (much more likely) for heavy atoms such as Xenon.
Let’s look at an example.
Imagine a ray of green light of wavelength λ = 530 nm incident on a detector with a work function of 1.1eV.
What is the kinetic energy given to a photoelectron ejected from this target?
What is the lowest wavelength of light that can release an electron from this target?
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Einstein had proposed that despite all the evidence that light is a wave, it also has particle-like properties.
There was previously no connection between the energy of a light wave and its frequency. Later in 1923 de
Broglie, influenced by Einstein’s work, suggested that particles such as electrons could exhibit wave-like
qualities, and waves such as light could exhibit particle-like behaviour. He called this wave-particle duality.
4.1.2 Compton effect: The energy of the photon is linked to the frequency of the wave by the equation
E photon  h f . But E photon  m c 2 (*) where m is the mass of the photon (if it is assumed to be corpuscular).
The momentum p of any particle is equal to its mass multiplied by its velocity i.e. p  m c and since
E photon  h f  m c 2 , we can rewrite this as E photon  h f  p c (*) and so momentum of a wave can be
written as
p
h

. (*) This is the de Broglie equation.
In 1923 Arthur Compton set up a collision between X-ray
photons and electrons. The experiment showed that the Xray photons and electrons behaved exactly like ball
bearings colliding on a table top. Because the electron was
scattered, the photon must have transferred both
momentum and kinetic energy to it. This can only be
explained by assuming that photons have momentum. But
he observed something else. Before the collision the photon
had one wavelength and after the collision its wavelength
had increased. Clearly the electron had been given energy,
conservation of energy indicating that the scattered photon must therefore have lower energy than prior to
the collision. The increase in wavelength, corresponding to a drop in frequency, could then only be
explained by assuming the E photon  h f relationship.
The more general equation for Compton scattering in which a photon interacts with an electron at rest,
causing the photon to scatter at an angle θ with an altered wavelength and giving recoil energy to the
h
electron is:  f  i 
(1  cos  ) (*)
me c
where me is the rest mass of the electron
in kg.
The maximum energy loss occurs when
the photon is scattered by 180 degrees
i.e. directly backwards. The figure shows
the energies of a 500 keV photon and an
electron after Compton scattering.
The cross section for Compton scattering
varies slowly with energy of the incident
photon. It also depends on the electron
density and so varies with the Z and
density of the material.
In the context of high energy observational detectors, part of the energy of the X-ray or gamma ray is
transferred to the scattering electron, which recoils and is ejected from its atom. The rest of the energy is
taken by the scattered, "degraded" photon. The fast moving scattered electron then rips through the detector
material creating a path of ionisation.
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Incidentally the figure shows Compton scattered photon energy as a function of scattering angle for different
incident photon energies. Note how irrespective of the incident energy, a 180 degree reflection of the photon
reduces its energy to between 170 and 220 keV. Remember this for later!!
Let’s look at an example.
Let’s imagine that we collide a gamma ray photon (λ = 3×10-14 m) with an electron. What is the momentum
of the photon before the collision? What is the energy lost by the photon if following the collision its
direction changes by 60 degrees?
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4.1.3. Pair production: This refers to the creation of an elementary particle and its antiparticle, usually from
a photon. This is allowed, provided there is enough energy available to create the pair – at least the total rest
mass energy of the two particles.
In pair production, if the incoming photon has energy equal or greater than 1.022 MeV then it can nudge the
nucleus, transferring a small amount of momentum and the nucleus can then create an electron and a
positron (each of at least 0.511 MeV). The positron after creation produces a trail of ionisation much like the
electron until eventually it has expended all its energy. At this point it is likely that it will come to rest near
an electron at which point it will annihilate to create a pair of 0.511 keV gamma rays. This is a good
calibration source for any detector.
The cross section for pair production depends weakly on the energy of the photon for a constant Z detector.
The cross section
  Z2
for constant photon energy.
Ee  Ee  E  2me c 2
2
(*) ( 2me c  2  511 keV ) and Ee is the K.E. of the particle.
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4.2 Bringing this all together and the mass absorption coefficient
Putting together all this information we can
plot a graph of Z against photon energy
showing the various processes which are
favoured.
We can therefore state that for a fixed Z the
photoelectric effect is dominant at low
photon energies, whilst pair production is
dominant at high energies. Mid energy
interactions favour Compton scattering.
The mass absorption coefficient is a measurement of how strongly a substance absorbs photons at a given
energy. The graph below plots the absorption coefficient (basically the cross section) as a function of energy.
It clearly shows the various processes which are favoured for various photon energy ranges.

 

A narrow beam of monoenergetic photons with an incident intensity Io, penetrating a layer of material with
mass thickness x, mass absorption coefficient (μ/ρ), and density ρ, emerges with intensity I given by the
exponential attenuation law:
  
I ( x)  I o exp     x 
  
(*)
Using graph above what thickness of material is required to reduce the flux of 50 MeV gammas by 50 % ?
Unfortunately for our detectors it is a little more complicated and we must add a few things to the plot before
we can obtain a realistic value for I(x).
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4.3 Inelastic and elastic scattering
The graph below displays the absorption (and attenuation) coefficient as a function of photon energy for air.
This graph is a little more complicated in that it separates absorption and scattering processes. Scattering is
defined as a process in which a portion of the photons coming from a source scatter (bounce) off molecules
and other small particles in the atmosphere or target in the case of a detector.
The scattering process can be either elastic or inelastic but be careful as this definition is not the same as the
more familiar kinematic definition of collisions…
Collision definition = An elastic collision is an encounter between two bodies in which the total kinetic
energy of the two bodies after the encounter is equal to their total kinetic energy before the encounter. So
elastic collisions occur only if there is no net conversion of kinetic energy into other forms. An inelastic
collision is a collision in which kinetic energy is not conserved. In the case of macroscopic bodies energy is
typically lost to heating or deformation.
Scattering definition = In particle physics inelastic scattering is a process in which the kinetic energy of the
incident particle or photon is not conserved. In almost all scattering processes of interest to us, the energy
of the incident photon is lost. Scattering due to inelastic collisions will almost always be inelastic, but, since
elastic collisions often transfer kinetic energy between particles, scattering due to elastic collisions can also
be inelastic, as in the case of Compton scattering.
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4.3.1 Rayleigh scattering
Rayleigh scattering is the elastic scattering of light or other
electromagnetic radiation by particles much smaller than the
wavelength of the photons. It can occur when photons pass
through solids or liquids, but is most prominently seen in gases.
(In the case of our blue sky it is these scattered photons that
give the sky its brightness and colour.
Remember the Sun is a point source in the sky so why wouldn’t
you expect our sky to look any different than looking at a torch
on a dark night? Rayleigh scattering is inversely proportional to
the fourth power of wavelength, so the shorter wavelength blue
light will scatter more than the longer wavelengths.
Conversely, glancing toward the sun, the colours that were not scattered away—the longer wavelengths such
as red and yellow light—are visible, giving the sun itself a slightly yellowish hue. Viewed from outer space
the space is black and the sun is white). In the photo the beam of a laser pointer is visible at night due to
Rayleigh scattering on airborne dust.
So why are there two Compton lines on the graph? Well in the Compton effect the interacting photon passes
on in a new direction (Compton scatters) having given up part of its energy to smack into an electron
(Compton absorption). This means that the photon will have deviated from the direct line of sight between
source and detector. The line in the graph indicating the total absorption coefficient only receives a
contribution therefore from the Compton absorption line and disregards any contribution due to Compton
scattering.
Since elastic scattering does not diminish the kinetic energy of the incident photon but does alter its direction
a question exists concerning how best to display the data. If you have a detector with which you want to
measure the photon rate from a distant source then you would want to know the TOTAL attenuation
coefficient as you are interested in the absolute intensity arriving at your target. This is by far the most
common value to take (the highest line on the plot). If you (for some crazy reason) would just like to know
the intensity as a consequence only of absorption and wish to ignore completely the effect of scattering, then
you would want to know the total absorption coefficient at the energy of interest.
See http://physics.nist.gov/PhysRefData/XrayMassCoef/tab3.html for a comprehensive listing of the total
mass absorption and mass attenuation coefficients for all elements.
See http://pml.nist.gov/PhysRefData/XrayMassCoef/tab4.html for a listing of the total mass absorption and
mass attenuation coefficients for some popular/relevant materials.
The total mass attenuation coefficient is defined as μ/ρ whereas the total attenuation coefficient is just μ.
Equally the total mass absorption coefficient is μen/ρ, whereas the total absorption coefficient is just μen. The
huge number of different names we give to coefficients in this area of particle astrophysics is a bit irritating
but so long as you understand the terms you can just for example divide by density (or not) and press on.
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4.4 Detectors types used to detect x-rays and gamma rays
Remember our perfect detector has low ionisation energy (giving high charge yield and better resolution),
fast signal speed (so we can observe higher rates) and high density (so we can observe a greater energy
range). It is desirable for it also to have a high Z as for the same energy photon, a higher Z material will be
more likely to interact via photoelectric absorption or pair production rather than Compton scattering.
Remember in Compton scattering the photon gives a varying proportion of its energy to electrons releasing
its energy over a far longer path length in the target giving rise to a less resolved signal. (Incidentally the
signal speed is faster in solids because the speed of the charge carriers is greater).
The basic idea of all detectors is the same. They are designed to allow photons to interact with some sort of
target material causing ionisation as photons interact via the photoelectric effect or Compton effect or pair
production.
In the case of the photoelectric effect the photon is completely absorbed by a target atom and a photoelectron
is created. This then shoots off in a specific direction and bounces into other atoms creating a trail of
ionisation along its path until it loses all its energy.
In the case of the Compton effect the photon gradually loses its energy by momentum transfer to orbital
electrons until all the original photon energy is lost.
In the case of pair production the photon is absorbed and all its energy (minus the energy required to create
them i.e. the rest mass) is given in the form of kinetic energy to an electron and a positron which then create
ionisation trails.
Detectors then either measure this charge directly or measure it indirectly by for example recording the light
produced as the ionised atoms and electrons recombine. We will look at the following types…
Proportional counter
Semiconductor detector
MPPC
measure charge produced
Microchannel plate
Scintillation counter
-
measure light from scintillation
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5 Main detector types
Some detectors just measure the overall intensity of the radiation (i.e. the number of photons).
Some can also record the energy of each photon and so create a spectrum of the number of photons as a
function of their energy.
Some are position sensitive meaning that they can record the direction of the ionisation trail and therefore
obtain information on the path of the initial photon.
5.1 Proportional Counters
Basically this is just a gas filled
container with a central electrode to
attract the charge (ionisation)
created by the photon. The container
is filled with a suitable gas (e.g.
argon + 10% methane). The
incoming photon interacts to create
ionisation. Due to the cylindergeometry arrangement, the primary
electrons created in this way "see"
an increasing electrical field on
route to the wire. This strong electric
field pulls these electrons toward the
central anode wire. In so doing these
electrons gain energy and create additional ionisation as they collide with atoms in their path. These ‘new’
(called secondary) electrons go on to collide with other atoms creating even more charge, this cascade of
ionisation acting to effectively amplify the original signal. Eventually the cloud of electrons reaches the
central wire where a current pulse is recorded by the electronics. The gas gain (amplification purely by
collisional excitation in the gas) can be up to 10,000 and so the electronics can be pretty basic and therefore
rugged. Without this process of gas amplification, weak signals would not be able to be measured, as they
would not be sufficiently discernible from the electronic "noise."
If all the photon energy is given up inside the chamber then the size of the pulse is proportional to the energy
of the original photon. However proportional counters have limitations. The energy resolution is poor
because the average energy to create an electron-ion pair is about 30 eV and so a 10 keV photon or
photoelectron will produce only about 300 electrons. The uncertainty in the signal is…
Uncertainty 
N
300

 6%
N
300
This is far too high an uncertainty considering the money required to put such a detector in space! Bear in
mind also that this uncertainty is a function of the photon energy i.e. one possible photon from boron (0.185
keV) produces approx. 6 pairs whilst a photon from molybdenum (17.5 keV) produces approx. 583 pairs.
So the more electron-ion pairs we can generate the better. Surely we can just look at sources producing
higher energy photons? No. Although they are rugged and simple there are two massive limitations to gas
filled proportional detectors. The first is that the window must be thin enough to avoid absorbing photons
but it must be thick enough to hold in gas at high pressure and to survive launch stress. Beryllium is a
popular choice. It is strong, low Z and therefore has low absorption characteristics but it is very toxic.
Incidentally gas leakage through the thin window is the most common reason why proportional counters fail.
The second problem is that of capture. If the photon is to penetrate the solid high density window then it
must be sufficiently energetic. However these energetic photons must then be totally absorbed by a low
density gas such as argon in order to determine their energy. This places a threshold on the energy range. If
the photon energy is too low it will be absorbed in the glass. If the energy is too high it will not be stopped
within the target. Proportional counters are therefore only effective in the 0.1 keV to 20 keV photon energy
range.
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Let’s prove this with an example…
A gas proportional counter at room temperature consists of a 50 cm long 0.5 bar argon gas target at 1 bar
within a metal tube with a 0.1 cm thick beryllium window at one end. By calculating the proportion of
original photons of energy 1, 10, 100, 1000 keV which go on to interact with the argon target, confirm that
maximum sensitivity is around 10 keV. (Beryllium has a density of 1.85gcm-3 and 1 bar argon gas has a
density of 5.85×10-3gcm-3 at room temperature.
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Over the years proportional counters have been improved and some are now capable of discerning the path
of ionisation through the gas. Typically these are rectangular boxes containing grids of wires arranged
orthogonally to one another.
A track of ionisation, once created, is drifted within an
electric field toward these grids. Upon arrival it creates a
signal on both sets of wires and triangulation provides the x
and y coordinates. The z coordinate is determined by
measuring the drift time from the ionization event to the
wires. This allows the path of ionisation to be reconstructed.
These wire grids are called Multi-wire Proportional
Chambers (MWPCs) and any MWPC which uses the arrival
time of a pulse of charge to obtain information on the
direction of the original photon or particle is also called a
Time Projection Chamber (TPC). An example is DRIFT II
shown below. (see http://drift.group.shef.ac.uk/ for more
information).
The ionisation trail has been
determined using x ,y ,z, coordinates.
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5.2 Scintillation Detectors
Scintillation detectors are made up of materials which produce light when either ionising particles pass
through them or photons interact with target atoms. They can be gases or liquids or solids (but they must
always be transparent to the light they produce!) An incident photon interacts via the photoelectric effect or
Compton scattering or pair production exactly as before and the kinetic energy of the electrons goes on to
create a path of ionisation through the target.
This ionisation (with no electric field present of course) then recombines. This generates photons within the
scintillator, usually in the visible range of the electromagnetic spectrum. This light is guided to a
photomultiplier tube where is interacts with a photocathode releasing electrons. A photocathode is a
negatively charged electrode which is coated with a photosensitive compound with a very low work function
such as caesium iodide. When this is struck by light, the absorbed energy causes electron emission due to the
photoelectric effect.
These electrons are then guided under an electric field
to the 1st dynode which is coated with a material such
as beryllium oxide that emits multiple numbers of
secondary electrons via collisional excitation when
the incident electron impacts with the surface. These
electrons are then guided by strong electric fields to
the 2nd, 3rd, 4th , etc. dynode and at each step the
number of electrons increases until the electrons are
finally collected at the anode. One electron generated
by the photocathode can give rise to over 1,000,000
electrons at the anode. Examples of scintillation
target materials include sodium iodide (a), xenon (b),
and bismuth germanium oxide (BGO) (c).
(a)
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(c)
Hopefully the original output of visible photons from the scintillator will have been sufficiently high to
enable a respectable output signal from the PMT. However the reason that we need PMTs is that the output
signal from scintillators is low and typical energy resolutions are around 6%. The main problem is the sheer
number of steps involved. Remember a source photon creates electrons and ions by various processes which
then recombine to produce photons which then pass to the photocathode which generates electrons which are
then amplified using the dynode stages inside the PMT. The table below shows where the signal is
diminished.
The scintillation efficiency is the fraction of the original photon energy which is transferred to target
electrons (which then go on to ionise other target atoms). Light loss is the fraction of the total light produced
by the scintillator which does not reach the photocathode. (Just look at the crystal in (a)). The photocathode
quantum efficiency is the efficiency with which a photon incident on the photocathode is converted back to
an electron.
Putting these values together we can calculate that sodium iodide scintillation detectors need around 230 eV
of incident photon energy to create an electron which will then pass to the 1st dynode. This figure is even
greater for other inorganics such as BGO and terrible for cheaper organic liquid scintillators such as
diisopropylnaphthalene. Bear in mind also that a 10 kg inorganic crystal costs approximately £5000 whilst
the equivalent funding could buy over 2000 litres of organic liquid scintillator.
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The type of image you see from an interaction of multiple photons from a source with a characteristic energy
given by E0 in shown below.
A = the photopeak produced by photoelectric absorption. This is termed the full energy peak because it
represents complete capture of the photons energy by the target. There are many different ways that this
could happen. It could occur via simple photoelectric absorption. It is also possible that a photon interacts
via Compton scattering and that the recoiling photon is then absorbed via the photoelectric effect. These two
processes would occur so quickly in succession that the readout electronics would see only a single pulse.
B = the Compton background continuum. Remember that Compton scattering deposits a variable amount of
energy in the target per scatter which is a function of the recoil angle.
C = the Compton edge: this represents the maximum energy that can be given to the recoiling electron in the
first Compton scattering involving an incident photon. NB. If the detector was very very large the Compton
edge would not exist as the photon would Compton scatter until all its energy had been deposited. However
the range between scatters is usually of the order of metres and typically we only see a single scatter.
D = really there should be no events here but sometimes a couple of unrelated Compton scattering events
occur at the same time and the analysis software integrates them as one higher energy signal. It is also
possible that a photon may undergo Compton scattering twice and then leave the target without depositing
all its energy.
E = this is a backscatter peak and is formed when the source photon does not interact at all with the target
but instead scatters with general stuff in the lab via the Compton effect, its recoiling photon making an angle
of pretty much 120 to 240 degrees to its original travel. This causes the photon to lose a significant fraction
of its energy. The recoiling photon then enters the target and interacts to yield a peak. If you look again at
the graph of photon recoil energy as a function of angle for incident photons of 5 MeV or less you will see
that between these angles all recoiling photons have energy about 170 to 220 keV. This is responsible for the
peak.
F = contribution from background photons coming from the rest of the universe.
G = the electrical noise remembering that the PMT amplifies greatly any stray charge landing on any dynode
stage.
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5.3 Influence of detector size and target material on energy spectrum produced
The size of the detector and target material really influences the energy spectrum. For example the range of
the electron produced by the photoelectric effect (and all subsequent electrons created along its ionisation
path) is of the order of a few millimetres in a solid or liquid scintillator.
It is therefore likely that a peak due to
photoelectric absorption will always be
observed in a target so long as the
photoelectric effect is favoured in that
energy range. In the case for example of
sodium iodide crystals the high density
and high Z give rise to a clear peak due
to photoelectric absorption whereas in
diisopropylnaphthalene, a low density
liquid scintillator with a low Z,
Compton scattering is favoured and a
peak due to photoelectric absorption
may not be seen.
Compton scattering takes place over
much longer ranges and unless the
target is massive a photon will most
likely only undergo a single Compton
scatter. (This explains the Compton
continuum and edge in B and C
respectively in the figure above.
Equally for a small target we must assume in pair production that the 511 keV gammas produced when the
positron annihilates with an electron simply pass out of the detector and are lost. But remember that the
electron and positron which were originally created were given the energy of the original photon minus
1.022 MeV. Because the ionisation path length for charged particles such as electrons and positrons is so
small this will be completely deposited in the target giving rise to an energy peak 1.022 MeV less than the
photoelectric absorption peak. This is called the double escape peak as shown in the figure above. Of course
if one of the 511 eV gammas is contained in the target then there will also be a single escape peak in the
energy spectrum at 511 keV less than the photoelectric absorption peak as shown in the figure for an
intermediate detector below.
If the detector target is massive then we can absorb everything! Photoelectric absorption still occurs of
course. In the case of
Compton scattering
photons will now
undergo repeated
Compton scatters until
the photon energy is so
low that photoelectric
absorption occurs and
the final low energy
photon is absorbed. The
energy of all recoiling
electrons produced is
contained within the
target of course. In pair
production the two 511
keV gammas produced are captured within the target. The result should be a single peak in the energy
spectrum equivalent to the energy of the source of radiation.
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However in reality such large detectors are impractical and the plot below is more common.
The main advantage of scintillation detectors is that their high density and large Z allow them to detect high
energy gammas and they are most useful for photon energies between 20 keV and many MeV (and in some
cases GeV). A main disadvantage is that the PMTs tend to be fragile and must survive the launch.
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5.4 What is a semiconductor?
In a material the valence band contains electrons which form bonds between atoms (and are therefore bound
to individual atoms) whilst the conduction band contains electrons which are free to move from atom to
atom. This is represented by the figure below. In conductors the valence and the conduction bands overlap
meaning that there are always available electrons to move freely between atoms (e.g. copper) conducting
electricity.
If the band gap at room temperature is greater than 6 eV thermal energy is insufficient to promote an
electron from the valence band to the conduction band. An insulator is therefore any material with a band
gap greater than this e.g. wood.
A semiconductor, such as a silicon crystal, is a solid material that has electrical conductivity between a
conductor and an insulator. In semiconductors the band gap is around 1 eV and so at room temperature
thermal excitation can promote only a very small number of electrons from the valence band into the
conduction band. However they don’t stay there for long and usually drop back down. However the small
band gap allows other methods to be used to promote electrons to the conduction band enabling them to be
used for conduction.
We can exploit this to make a photon or particle detector. If we cool the semiconductor down to liquid
nitrogen temperatures (-197 degC) then the conduction band will be empty. However photons interacting in
the semiconductor can promote electrons into the conduction band and if a voltage is applied this charge
(both the electrons in the conduction band and the holes in the valence band) will be removed from the
detector.
Whilst such a target based around an intrinsic silicon crystal would function adequately its design can be
significantly improved. The term intrinsic here distinguishes between the properties of pure "intrinsic"
semiconductors and the dramatically different properties of doped n-type or p-type semiconductors. In an ntype semiconductor, a dopant e.g. phosphorus is added to the silicon melt as it is grown. This dopant
contributes extra electrons, dramatically increasing the conductivity. In a p-type semiconductor, the dopant
e.g. boron produces extra vacancies or holes, which likewise increase the conductivity.
When p-type and n-type materials are placed in contact with each other a diode is produced. In so doing,
some of the free electrons in the n-region which have reached the conduction band are free to diffuse across
the junction and combine with holes to form negative ions. In so doing they leave behind positive ions at the
donor impurity sites. This forms what is called a "depletion region".
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A space charge builds up across the depletion region which inhibits any further electron transfer unless it is
helped by putting a forward bias on the junction meaning that a positive voltage would be applied to the P
region in the figure above with a negative voltage applied to the N region. This assists electrons in
overcoming the Coulomb barrier of the space charge in the depletion region. Electrons will then flow with
very small resistance in the forward direction. Conversely reverse biasing the junction by applying a
negative voltage to the P region and a positive voltage to the N region further impedes the flow of electrons
across the junction.
Although the application of a reverse voltage to the pn junction will cause a brief current to flow as both
electrons (black dots) and holes (white dots) are pulled away from the junction, this current will cease when
the potential formed by the widened depletion layer equals the applied voltage. At this point nearly all the
potential difference or voltage drop is across the depletion region due to the lack of charge carriers and the
resulting large resistance.
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5.5 Semiconductor detectors
Semiconductor devices enable the principle of charge collection used in gas proportional counters to be
extended to a solid target. They use a semiconductor (usually silicon or germanium based) to detect
traversing charged particles or the absorption of photons. They are called semiconductor diode detectors if
they are based on a single diode. If they contain many diodes they are referred to as semiconductor detectors.
Semiconductor detectors are particularly effective at gamma and x-ray spectrometry and as particle detectors
because their far greater density and Z provide greater stopping power and a wider detection range of photon
energies.
Semiconductor detectors are made up of many pn junctions (diodes). To work as a detector, the pn diode is
reverse-biased by applying a positive voltage to the electrode on the N side and grounding the electrode on
the P side. Reverse biasing extends the depletion region through the detector forming the target volume. As
shown in the figure below the n type region is kept as small as possible whilst the depletion region is as large
as possible.
When a photon strikes a semiconductor, it can promote an electron from the valence band to the conduction
band creating an electron(-) - hole(+) pair. (This is analogous to electron-ion pairs generated in proportional
gas detectors). If a photon produces electron hole pairs outside the depletion region, they will recombine
very quickly, thus producing no net current. However if the photon produces electron hole pairs in the
depletion region, the strong internal field will rapidly separate the pairs before they recombine, electrons
drifting towards the anode, and holes to the cathode, resulting in a net current across the diode.
The concentration of these electron-hole pairs is dependent on the amount of photons striking the
semiconductor. The collected charge produces a current pulse on the electrodes, whose integral equals the
total charge generated by the incident particle, i.e. is a measure of the deposited energy. The charge
generated is passed through a charge-sensitive preamplifier and is then measured in an outer circuit. As the
amount of energy required to create an electron-hole pair is known, and is independent of the energy of the
incident photon, measuring the number of electron-hole pairs allows the energy of the incident radiation to
be found.
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Although Silicon based semiconductors used to be most common Germanium detectors (HPGe-high purity
germanium) are now growing in popularity. Germanium has a smaller band gap than silicon, it has a higher
density and a higher Z and lower impurities which allow larger depleted regions and easier manufacture.
The energy resolution is related to the low energy threshold: only 3.6 eV is necessary to produce an electronhole pair, a low value compared to the ionization energy in a gas (30 eV) or the approximately 300 eV
necessary to observe an electron from a sodium iodide crystal scintillator.
Because of this the energy resolution is approximately 1%
(compared with 6% for a gas detector). They also have good spatial
resolution which comes from the high density of silicon or
germanium, which reduces the range of the secondary electrons
(electrons excited by the primary ionising electron). The difference
in energy resolution between scintillator detectors (a) and
semiconductor detectors (b) is shown in the figure.
High purity germanium (HPGe) semiconductor
Semiconductor detectors have high stopping power and are
able to measure photon energies up to several MeV. They
have high resolution and are rugged. However there is no
charge multiplication in the semiconductor and so the
signal-to-noise ratio is a critical issue requiring low-noise
electronics. Space missions require liquid nitrogen cooling
to avoid thermal noise which places a limit on the lifetime
of the mission.
Sodium iodide crystal
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5.6 Microchannel plate detector
A microchannel plate is a slab made from highly resistive material of typically 2 mm thickness with a
regular array of tiny tubes or slots (microchannels) leading from one face to the opposite, densely distributed
over the whole surface. The microchannels are typically 10 micrometers in diameter and spaced apart by
approximately 15 micrometers; they are parallel to each other and often enter the plate at a small angle to the
surface (~8° from normal). Both the slab is sandwiched top and bottom by electrodes so that a high voltage
can be applied through the length of the channels.
An incident x-ray photon enters a channel and frees (via photoelectric emission) an electron from the
channel wall. An electron accelerating potential difference (approx. -1500 Volts) is applied across the length
of the channel. The initial electron strikes the adjacent wall, freeing several electrons (via "secondary
emission"). These electrons accelerate along the channel until they in turn strike the channel surface, giving
rise to more electrons.
Eventually this cascade process yields a cloud of several thousand electrons which emerge from the rear of
the plate. The electrons exit the channels on the opposite side where they are themselves detected by
additional means, often simply a single metal anode measuring total current. In some applications each
channel is monitored independently to produce an image.
A single x-ray photon interacting in a channel of the MCP produces a charge pulse of about 1000 electrons
that emerge from the rear of the plate. Since the individual tubes confine the pulse, the spatial pattern of
electron pulses at the rear of the plate preserve the pattern (image) of x-rays incident on the front surface.
This can provide spatial information.
Advantages of MCPs are high efficiency and high gain. However they have very poor energy resolution and
can only be used for soft x-rays (as the photon energy must be high enough to interact with the target via the
photoelectric effect but must not be so high that it passes straight through the target.
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5.7 Multi Pixel Photon Counter (MPPC) and Charge Coupled Devices (CCDs)
The Multi Pixel Photon Counter (MPPC) is a novel type of
semiconducting photon sensor developed in the last few years. It has
superb photon detection ability, excellent cost versus performance and is
very compact. It will most likely supersede expensive large PMTs in
near future. The MPPC consists of an array of APD (Avalanche PhotoDiode) pixels arranged on a substrate of area of around 1cm by 1cm.
Each pixel consists of p and n type semiconductor surrounded by a guard ring for isolation from other pixels.
Applying about 70 V of reverse bias voltage to the p-n junction creates a strong electric field in the depletion
region. If a photon is incident on a pixel, the photon produces an electron hole pair.
This electron then induces an avalanche in the depletion
region because the diode is so close to its reverse breakdown
voltage. Thus the pixel "fires" and produces a signal. Since
the diodes are not connected and it is very likely that more
than one photon arrives at the MPPC, the number of fired
diodes gives a measure of the total number of photons.
So the number of pixels on a sensor defines the dynamic
range, namely how strong the light is that can be measured
by the MPPC. If the light yield is low enough compared with
number of pixels on the sensor, one can have response linear
to the light yield. But in the case of measuring stronger light,
sometimes two or more photons go into a same pixel. In such
a case the second photon is not recorded and response is no
longer linear. The rate at which the photons arrive at the
MPPC is also an issue as each pixel once fired can take up to
a microsecond to recover to a state at which it can be refired.
The table below compares the performance of a photomultiplier
with an MPPC. The MPPC is remarkably cheap for the
performance it delivers especially the excellent photon detection
ability. The MPPC is very compact, requires lower voltages and
appears to be more robust. However the limited dynamic range is
still a problem and the maximum size of these arrays is still too
small at 3 cm by 3 cm. Another major problem is the noise rate.
Since the diode is on the point of avalanche thermal electrons
frequently induce avalanches creating fake signals.
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The noise rate can be as high as 1MHz which is unimaginably high although this can be all but eradicated by
cooling to liquid nitrogen temperatures.
Whilst an MPPC provides excellent photon counting it does not provide spatial information.
A CCD is a device used for the movement of electrical charge, usually from within the device to an area
where the charge can be analysed, for example by conversion into a digital value. This is achieved by
moving charge between stages within the device. Often a CCD is integrated with an image sensor, such as a
photoelectric device which converts light into charge which is then read by the CCD.
Recently a Charge-Coupled device (CCD) has been incorporated into the MPPC thereby providing track
information. The figure below shows a CCD used in digital photography.
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6 Imaging systems
The detectors mentioned so far can measure energy and flux but it is also valuable to know what direction
the photons are coming from. In the optical region of the EM spectrum we would use lenses and mirrors to
make telescopes to achieve this, placing the detector at the focal plane. We cant do this for x-rays and
gamma rays as they simply penetrate the material and are absorbed rather than being refracted or reflected.
There are several alternative approaches:
Honeycomb collimator
Modulation collimator
Collimation
Coded mask collimator
Spark chamber
Compton telescope
Tracking
Wire chamber telescope
Grazing incidence telescope
6.1 Spark chamber
A spark chamber is a very basic design which has been superseded by more sophisticated detectors such as
drift chambers and silicon detectors. They are however very simple. A spark chamber consists of metal
plates placed in a sealed box filled with a gas such as helium or neon.
As a charged particle travels through the detector, it ionises the gas between the plates. A trigger system is
used to apply high voltage to the plates to create an electric field immediately after the particle goes through
the chamber, producing sparks on its exact trajectory as the ionisation trail is the most likely route along
which breakdown will occur.
By recording the ionisation trail it is possible to determine the path of the particle in the detector.
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6.2 Compton telescope
Compton telescopes are two-level instruments
typically sensitive to photons between 300 eV and
30 MeV. In the top level, the gamma-ray
Compton scatters off an electron in a scintillator.
This is usually a liquid scintillator as it is more
likely to interact via Compton scattering. The
scattered photon then travels down into a second
level of scintillator material within which it is
completely absorbed. This is usually a crystal
scintillator such as sodium iodide as the photon is
more likely to be completely absorbed. PMTs
viewing the two levels can determine the
interaction points in each layer and the amount of
energy deposited.
As shown in the figure, the line between these two
interaction points does not point back to the
direction of the incoming photon because the
direction of the photon changed when it scattered.
It is possible, however, to determine the angle of incidence the photon made with respect to this line.
If the energy lost in the top layer is E c and the energy lost in the second layer is E a then the original
photon energy E  Ea  Ec . Also the energy of the recoil photon after Compton scattering E'  E a .
'
Knowing E and E we can calculate the angle
 as:
ch
Ec mc 2
h
so substituting we can write cos   1 
(1  cos  ) but E  hf i 
i
E a ( E a  Ec )
me c
Unfortunately, whilst this angle can be determined, the only information that this can provide is that the
original gamma came from within a cone of half angle  .
 f  i 
Clearly it would be of significant value if the scattered electron could also be tracked as then the specific
origin of the gamma ray could be deduced based on conservation of momentum.
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6.3 Grazing incidence telescope
The grazing incidence telescope utilises the fact that electromagnetic radiation such as x-rays and gamma
rays at such short wavelengths behaves like ordinary light rays if it strikes surfaces at a shallow enough
angle.
X-rays mirrors only function correctly if the angle from the plane of reflection is very low (typically 10 arcminutes to 2 degrees. The idea was first put forward in 1952 by Hans Wolter, grazing incident telescope also
being known as Wolter telescopes.
6.4 Rotational modulation collimator
In the same way that putting your hand up to your
head to shield your eyes from the sun provides
information on the location of the sun, a rotational
modulation collimator consists of two sets of very
dense grids which are rotated. The detector sees a
brightening and darkening source in a regular
pattern. The rate of change in the brightness is
dependent upon the angle from the Z axis and the
orientation of the source relative to the telescope
with respect to the rotating grid position. Complex
computer reconstruction of the image provides the
source location.
6.5 Coded mask aperture collimator
The coded mask aperture collimator is a big metal sheet with
random holes in it that casts a shadow on the detector plates.
Since scientists know what the aperture mask looks like, they can
use the shadow cast by a source on the detector to pinpoint the
location of the source.
The detector array is responsible for recording gamma-ray and xray data. It is a thick semiconductor plate that measures the
energy change caused by the gamma-rays that strike it. In
INTEGRAL (see later) the coded mask is made of hexagonal
tungsten tiles, and the detector plane contains 19 germanium
semiconductor tiles.
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6.6 A honeycomb collimator
A honeycomb collimator restricts the incident angle of the incoming photons which
will impinge on the detector. They are composed of thousands of precisely aligned
holes cast in a lead block. The collimator conveys only those photons travelling
directly along the long axis of each hole. Photons emitted in other directions are
absorbed. Without a collimator in front of the imaging device, the image would be
indistinct.
6.7 A wire chamber telescope
See section 5.1.
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7. Great Observatories
The concept of a Great Observatory program was developed in the mid 1980s on the premise that many
discoveries in astronomy were due to improved reception of the electromagnetic spectrum. NASA's "Great
Observatories" program used four separate satellites, each designed to cover a different part of the spectrum
in ways which terrestrial systems could not. This perspective enabled the proposed x-ray and infrared
observatories to be appropriately seen as a continuation of the astronomical program begun with Hubble and
CGRO rather than competitors or replacements.
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7.1 ROSAT (1990-1999)
ROSAT, the ROentgen SATellite, was an X-ray observatory developed through a cooperative program
between Germany, United States, and the United Kingdom. Designed and operated by Germany but
launched by the United States in 1990, it was turned off in 1999.
Its main task was to make an all-sky survey in soft x-rays (0.1 keV-2 keV),
after which the satellite began a series of pointed observations of key areas of
interest. Its sensitivity to x-rays was over 1000 times greater than Uhuru and
the entire satellite was 4 m long and 2.4 tonnes. The x-ray mirror assembly
was a grazing incidence four-fold nested Wolter I telescope with an 84 cm
diameter aperture.
Power was supplied through 3 solar panels and through a rechargeable
battery during the shadow phase. The scientific payload consisted of an x-ray
telescope which was used in conjunction with one of the focal plane
instruments:

The Position Sensitive Proportional Counter
PSPCs are multiwire proportional counters
providing modest energy resolution and high
spatial resolution over a 2 degree diameter
field-of-view. They contained a mixture of 65%
argon, 20% xenon and 15% methane, the x-ray
photons passing through a thin plastic entrance
window.

The High Resolution Imager
This is made up of two cascaded microchannel
plates (MCPs) providing ~ 2 arcsec spatial
resolution.
Observations of the Andromeda galaxy (M31)
by ROSAT allowed detection of numerous
point sources. The satellite continues to orbit
approximately 390 km above the Earth.
Image taken from http://dx.doi.org/10.1051/0004-6361:20010495
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7.2 Hubble Space Telescope (1990)
In 1990 the Hubble Space Telescope was released into
space. Since then Hubble's ultra deep field images have
provided the most detailed visible light photos ever made
of the universe's most distant objects. Many Hubble
observations have led to breakthroughs in astrophysics,
such as accurately determining the rate of expansion of the
universe. The figure below is of the Horsehead nebula.
The HST is comprised of an infrared camera/
spectrometer, an optical camera/ spectrometer, a wide field
optical camera, and an ultraviolet spectrograph. High
energy observation astrophysics telescopes look for x-ray
and gamma ray emission so really Hubble should not be
included in this course. However no list of major space
based observatories would be complete without at least
referring to it briefly.
Hubble is the only telescope ever designed to be serviced
in space by astronauts, has been serviced 5 times and is
now expected to function until at least 2014, when its
successor, the James Webb Space Telescope (JWST) an
infrared space observatory, is due to be launched. The
JWST’s main scientific goal is to observe the most distant
objects in the universe beyond the reach of Hubble. JWST
will observe in the infrared, in order to penetrate dust at
cosmological distances. This means it will continue some
of Spitzer capabilities (see later), whilst some Hubble capabilities will be lost. New advances in ground
telescopes will however take over some visible observations.
7.3 The Compton Gamma Ray Observatory (1991)
The Compton Gamma Ray Observatory (CGRO) was the second of NASA's great observatories. CGRO, at
17 tons, was the heaviest astrophysical payload ever flown at the time of its launch in 1991 aboard the space
shuttle. It was safely deorbited and re-entered the Earth's atmosphere in 2000.
CGRO had four instruments that covered an unprecedented six
decades of the electromagnetic spectrum, from 30 keV to 30
GeV. In order of increasing spectral energy coverage, these
instruments were the Burst And Transient Source Experiment
(BATSE), the Oriented Scintillation Spectrometer Experiment
(OSSE), the Imaging Compton Telescope (COMPTEL), and the
Energetic Gamma Ray Experiment Telescope (EGRET). For
each of the instruments, an improvement in sensitivity of better
than a factor of ten was realized over previous missions.
BATSE featured 8 sodium iodide detectors mounted on all corners of the satellite which were sensitive to
energies from 20 keV to over 1000 keV. It served as the all-sky monitor for the CGRO detecting and
locating strong transient sources such as gamma-ray bursts as well as outbursts from other sources over the
entire sky.
OSSE consisted of four sodium iodide detectors, sensitive to energies from 50 keV to 10 MeV allowing
extensive measurements of the energy spectra of gamma-ray sources to be undertaken.
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COMPTEL uses two layers of gamma-ray detectors to reconstruct an image of a gamma-ray source in the
energy range 1 eV to 30 MeV. Gamma rays from active galaxies, radioactive supernova remnants, and
diffuse gamma rays from giant
molecular clouds were studied with this
instrument. COMPTEL's upper layer of
detectors were filled with a liquid
scintillator which scatters an incoming
gamma-ray photon via the Compton
effect. This photon is then absorbed by
NaI crystals in the lower detectors. The
instrument recorded the time, location,
and energy of the events in each layer
of detectors which made it possible to
determine the direction and energy of
the original gamma ray photon and
reconstruct an image and energy
spectrum of the source.
EGRET provided the highest energy gamma ray window for the Compton Observatory. Its energy range was
from 20 MeV to 30 GeV. EGRET was 10 to 20 times larger and more sensitive than previous detectors
operating at these high energies and made detailed observations of high-energy processes associated with
gamma-ray bursts, cosmic rays, pulsars, and active galaxies.
The EGRET instrument produced images
at these energies using high-voltage gasfilled spark chambers. High energy
gamma rays enter the chambers and
produce an electron-positron pair of
particles which create ionisation trails
and causing sparks. The path of the
particles is recorded allowing the
determination of the direction of the
original gamma ray. The particle
energies are recorded by a sodium iodide
crystal beneath the spark chambers
providing a measure of the original
gamma-ray energy.
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7.4 Chandra X-ray Observatory (1999)
The Chandra X-ray Observatory was launched in 1999. Chandra is
the third of NASA's four Great Observatories. The first was
Hubble Space Telescope; second the Compton Gamma Ray
Observatory, launched in 1991 although now deorbited; and last is
the Spitzer Space Telescope. Chandra is sensitive to x-ray sources
100 times fainter than any previous x-ray telescope, due primarily
to the high quality of its mirrors.
As we have seen x-rays are produced when matter is heated to
millions of degrees. Such temperatures occur where high magnetic
fields, or extreme gravity, or explosive forces are present. A vast
cloud of hot gas in a cluster of galaxies can be several million light
years across and contain enough matter to make hundreds of
trillions of stars. X-ray telescopes can also trace the hot gas from
an exploding star or detect x-rays from matter swirling as close as
90 kilometres from the event horizon of a stellar black hole.
Chandra carries four very sensitive mirrors nested inside each
other. Energetic x-rays striking the insides of the hollow shells are
focused onto electronic detectors at the end of the 9 m optical
bench. Depending on which of two detectors is used, very detailed
images or spectra of the cosmic source can be made and
analyzed.
Chandra consists of two main instruments. The first is the
Advanced CCD Imaging Spectrometer (ACIS) made up of 10
CCD chips and provides images as well as spectral information
of the object observed in the 0.2-10 keV range. The second is
the High Resolution Camera (HRC). This has two microchannel plate components and images over the range of 0.1-10
keV. Either of these can be moved into position on the focal
plane during an observation. In addition both of these
instruments can be used on their own or in conjunction with
one of the observatory's two transmission gratings. These
transmission gratings provide Chandra with high resolution
spectroscopy.
Chandra has imaged the spectacular, glowing remains of exploded stars, and taken spectra showing the
dispersal of elements. It has observed the region around the supermassive black hole in the centre of our
Milky Way, and found black holes across the universe. The figure above shows the cat’s eye planetary
nebula.
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7.5 Spitzer Space Telescope (2003)
The Spitzer Space Telescope (SST), previously called the Space
Infrared Telescope Facility is an infrared space observatory launched
in 2003. It is the fourth and final of NASA's Great Observatories. The
planned mission period was five years until the onboard liquid helium
supply, used to cool the telescope, was exhausted. This occurred in
2009.
As with the HST the SST is not a high energy photon observatory, its
principle detectors being an infrared camera, infrared spectrometer,
and far infrared detector arrays.
The figure below is an infrared image of Andromeda galaxy.
Spitzer Science produced outstanding results. Infrared observations are necessary for very distant
astronomical objects where all the visible light is redshifted to infrared wavelengths, for cool objects which
emit little visible light, and for regions optically obscured by dust. The figure shows the same area of space
photographed in the visible and infrared range. The bottom IR photo shows far greater detail, all the gas and
dust obscuring images in the visible spectrum disappearing if viewed in the infrared.
As already mentioned JWST will exceed Spitzer’s
performance in the near-infrared. The figure shows a full
scale model of JWST which uses a 6.5 m mirror and an IR
camera and spectrograph.
In addition the European Space Agency's Herschel Space Observatory (HST), launched in 2009, will exceed
Spitzer in the far-infrared and is capable of seeing the coldest and dustiest objects in space; for example,
dusty galaxies just starting to bulk up with new stars.
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7.6 Fermi Gamma ray Space Telescope (2008)
The Fermi Gamma-ray Space Telescope (FGST) shown in the figure
right, is a follow on to the CGRO and was launched in 2008. Budget cuts
however have meant that it is much smaller and its operation is more
narrowly defined. It carries only one main instrument, the Large Area
Telescope (LAT), used to perform gamma-ray astronomy including allsky survey studying active galactic nuclei, pulsars, other high-energy
sources. A secondary instrument aboard is the Gamma-ray Burst Monitor
(GBM), which is being used to study gamma-ray bursts. The GBM
includes 12 sodium iodide scintillation detectors which cover the lower
part of the energy range from a few keV to about 1 MeV, and 2 bismuth
germanate scintillation detectors which cover the energy range of ~150
keV to ~30 MeV The figure shows a sodium iodide crystal and PMT
assembly.
The LAT detects
individual gamma rays as follows. Gamma ray photons hit
thin metal sheets, converting to electron-positron pairs, via
pair production. These charged particles pass through
interleaved layers of silicon microstrip detectors, causing
ionization which produces detectable tiny pulses of electric
charge.
By combining
information from all
layers of this tracker, the path of the particles can be determined. Each
particle moves in a slightly different direction, creating an inverted "V"
that points back toward the direction in the sky from which the gamma
ray came. After passing through the tracker, the electron and positron
enter the calorimeter, which consists of a stack of caesium iodide
scintillator crystals used to measure the total energy of the particles.
The photo shows the LAT stack of metal sheets. The LAT is the bigger
and better successor to the EGRET instrument on the CRGO. It extends the Compton scattering telescope
concept to higher energy photons using pair production to indicate direction rather than Compton.
7.7 INTEGRAL (2002)
Another large, high-energy observatory is INTEGRAL, (European Space Agency’s INTErnational Gamma
Ray Astrophysics Laboratory), launched in 2002. It observes in similar frequencies to CGRO but uses a
fundamentally different telescope technology, coded-aperture masks. Thus, its capabilities are
complementary to CGRO and FGST and not a direct replacement.
The principle imager observes from 15 keV (hard x-rays) to 10 MeV (gamma rays). A 95 x 95 mask of
rectangular tungsten tiles sits around 3 metres above the detectors. The detector system contains a forward
plane of 128 x 128 Cadmium-Telluride tiles, backed by a 64 x 64 plane of caesium iodide tiles.
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7.8 Constellation-X (2016)
Constellation-X is an observatory intended to perform extremely sensitive x ray observations, beginning
around 2016. Although this is not a direct replacement for Chandra, Constellation-X may be several times to
several dozen times more sensitive than it. It will also extend further into the hard x-ray regions, giving it
some of the abilities of the CGRO. Constellation-X will feature a much larger telescope than any previous xray telescope, and it will have much higher spectral resolution.
One role for Constellation-X will be to take X-ray spectra from matter as it falls into supermassive black
holes. These spectra will measure the Doppler shift of matter as it orbits and plunges into black holes and by
applying the equations of General Relativity and comparing data with that predicted by the presence of dark
matter, the dark matter/energy relationship may be deduced.
None of the future missions discussed are designed for Shuttle launch, or manned servicing.
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8 Neutrino sources and detection methods
A neutrino is an elementary particle that usually travels close to the speed of light, is electrically neutral, and
is able to pass through ordinary matter almost undisturbed. It is a member of the lepton family which is
made up of electrons, muons, taus, electron neutrinos, muon neutrinos, tau neutrinos and the corresponding
antiparticle of each. Muons have the same charge as an electron but are 200 times heavier. Tau particles are
3,500 times the mass of an electron. Whereas the electron is stable, the muon and the tau are not and have
lifetimes of 2.2×10-6 s and 2.9×10-13 s respectively. Neutrinos have a very small, but non-zero mass which is
yet to be accurately measured. Most neutrinos passing through the Earth originate from the Sun, and more
than 50 trillion solar neutrinos pass through an average human body every second. Neutrinos are clearly
extremely difficult to detect.
In this course we will concentrate on two main sources: (i) Stars and (ii) Supernovae.
8.1 Stellar neutrinos
In β− decay, the weak interaction converts a neutron (n) into a proton (p) while emitting an electron (e−) and
n  p  e  e . (*)
an electron antineutrino ( e ) :
In β+ decay, energy is used to convert a proton (p) into a neutron (n), a positron (e+) and an electron
p  n  e    e . (*)
neutrino ( e ) :
Neutrinos are emitted at many stages of a star’s life. The main reaction sequences are:
(Image from http://www.gsi.de/forschung/kp/kp2/nuc-astro/StellarBurning_e.html)
Note that the positron in PP I immediately annihilates with an electron, and their mass energy is carried off
by two gamma ray photons.
PP I is dominant for temperatures from 10 megakelvin to 14 megakelvin. The PP II branch is dominant at
temperatures of 14 megakelvin to 23 megakelvin and above this the PP III chain is dominant.
For the Sun for example 86% of helium-4 is created via PP I, whereas 14% is created via PP II and a tiny
0.02% is created via PP III.
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For each PP chain there are distinct stages when neutrinos are emitted. It is easy to add up the atomic mass
units in a specific reaction, calculate the mass difference and therefore the amount of energy released in that
reaction. For example we can calculate the energy created in the PP I reaction step involving the release of a
neutrino below.
p  p  d  e   e (*)
The mass of a proton is 1.00728u and the mass of a deuterium nucleus is 2.0141u. The atomic mass unit u is
1.66054×10-27 kg and the mass difference must be converted into energy using E  mc 2 (*) giving the
answer in eV. (The mass of the positron can be neglected).
Mass difference is 2.01456u – 2.014102u = 4.58×10-4u = 7.61×10-31 kg.
E  mc 2  6.845 10 14 joules or 0.43 MeV.
For this step therefore within the PP I chain we find that each time a deuterium nucleus is created, 0.43 MeV
is also released plus a neutrino.
For the step in the PP II reaction chain involving the release of a neutrino below, either 0.343 MeV or 0.861
MeV is also released.
7
4
Be  e  37 Li  e
And in the PP III reaction chain again involving the release of a neutrino below, an energy of 17.98 MeV is
also emitted.
8
5
B48Be  e   e  
There is also a very small possibility that the 3He produced in the second stage of the PP I chain captures a
proton. This results in the highest energy emission of 18.77 MeV and is called the hep reaction.
3
2
He11H 24 He  e   e
There is also a very very small possibility that deuterium can be produced by a reaction between two protons
and an electron resulting in the generation of a neutrino and energy emission of 1.44 MeV. This is called the
pep reaction.
1
1
H 11H  e  12 H  e
In any reaction there must be conservation of momentum. For this reason when a reaction yields a neutrino
(tiny unmeasureable mass) and a nucleus, conservation of momentum means that the velocity of the neutrino
will be tremendously higher than that of the nucleus. The effect is therefore that the neutrino will carry away
all the kinetic energy too.
Prove this yourself by calculating the ratio of kinetic energies of two reaction products with masses m and M
travelling at u and v respectively following a reaction such as the pep one above. (Assume M>>m).
However if there are more than two bodies produced the energy can be shared in any fashion so long as
momentum is conserved. The neutrino can therefore have any energy up to the maximum of the total energy
generated in the reaction. The figure below shows the predicted neutrino spectrum from the Sun.
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If we could detect this characteristic spectrum we would have an important check on our understanding of
stellar physics. In fact there is no other way of checking what is going on in a star’s interior.
8.2 Supernova neutrinos
When a massive star at the end of its life collapses to a neutron star, it radiates almost all of its binding
energy in the form of neutrinos, most of which have energies in the range 10-30 MeV. There are two main
processes. The first process occurs as forces inside a star about to go supernova become so great that
electrons and protons are forced close enough together to interact through the weak force:
−
p  e   n  e
(*)
This is called inverse β decay.
A second and more important neutrino source is the thermal energy (100 billion kelvin) of the newly formed
neutron core, which is dissipated via the formation of neutrino-antineutrino pairs of all flavours. Most of the
energy produced in supernovas is thus radiated
away in the form of an immense burst of neutrinos.
The first experimental evidence of this
phenomenon came in 1987, when neutrinos from
supernova 1987A were detected. This again gave a
direct insight into the core of a supernova.
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8.3 Detecting neutrinos
It is not easy to detect neutrinos. They have no charge and nearly no mass and their interaction cross section
is so low that it would require 1017 m thickness of steel to ensure 50% of the incident neutrino flux were
captured!!!!! Fortunately it is also true that lots are produced. There are four main approaches to detecting
neutrinos. Two we will meet now and the other two will be covered when discussing the SNO detector.
8.3.1 Neutrino capture reactions: if a neutrino collides with a neutron in the nucleus of an atom it can
interact via negative beta decay. This changes the atom to a different one. For example
37
37
 e 17
Cl 18
Ar  e 
(*) so take a big container of chlorine and watch to see if any argon is produced.
A good example is the Homestake chlorine solar neutrino experiment which has been taking data since 1970.
The target contains 600 tonnes of
tetrachloroethylene (dry cleaning fluid) placed
1.5 km underground in a gold mine in South
Dakota.
A big target deep underground was needed to
account for the very small probability of a
successful neutrino capture, and to prevent
interference from other forms of radiation which
are absorbed by the rock such as cosmic rays. In
addition a thick water jacket shields the target
from neutrons.
37
It exploits the isotope 17 Cl which has a natural
abundance of 25%. The tank therefore contains about 2×1030 chlorine target nuclei. Neutrinos interact to
produce argon-37 in the tank.
This is allowed to accumulate for a month before the argon gas
is filtered out by bubbling with helium and the number of argon
atoms is calculated from the activity rate as it decays by β+
decay:
37
18
37
Ar 17
Cl  e   
Even with such a big target the production rate is less than one
argon atom per day and it is a huge experimental challenge to
detect a few tens of atoms in 1031.
Other examples of neutrino capture experiments are GALLEX
and SAGE.
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8.3.2 Neutrino scattering: neutrinos scattering off
electrons in a target material, give them high recoil
energy. These fast moving electrons can be detected
to obtain information on the energy of the neutrino
and its incoming direction.
The most common method is to observe Cherenkov
radiation emitted as the electron recoils at speeds
faster than light would travel in that medium. This
produces a cone of light, the cone angle depending
on the velocity.
Cherenkov light is produced when a charged particle
moves faster than the speed of light in that medium
(although of course slower than the speed of light in
vacuum). The figure below shows blue Cherenkov light produced in
a nuclear power station as β particle velocities exceed the speed of
light in water (2.3 × 108 m/s).
The big advantage of scattering detectors is that they provide
information on energy and direction whereas capture detectors just
give the total flux. However this comes at a cost of much higher
complexity and cost.
The interaction is sensitive to all neutrino flavours, but the scattering probability of an electron-neutrino is
about a factor of six larger than the scattering probability of the two other neutrino types.
An example of a neutrino scattering detector is Super Kamiokande
located 1000m underground in Kamioka mine in Japan. This is a
water Cherenkov detector made up of 50,000 m3 of ultra-pure water.
The detector is split into the target containing 32,000 m3 of water
viewed by 11,000 51 cm diameter PMTs which is surrounded by a
veto containing the rest of the water viewed by 1,900 PMTs. This veto
is designed to screen the target from the flux of cosmic rays. Since the
likelihood that the same neutrino will scatter in both the veto and
target is incredibly low, any signal which interacts in both is
disregarded as being unwanted background not worth being recorded.
The observatory was designed to study solar and atmospheric
neutrinos, and keep watch for supernovas in the Milky Way Galaxy.
Neutrino scattering with the electrons of water produces a recoiling
electron that moves faster than the speed of light in water creating a
cone of Cherenkov light, projected as a ring
onto the wall of the detector and recorded by
the PMTs. Using the timing information
recorded by each PMT, the path of the
incoming neutrino can be determined. The
figure shows a typical neutrino event.
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The detector succeeded in detecting neutrinos from a
supernova explosion which was observed in the Large
Magellanic Cloud in 1987. Super-Kamiokande was also
first to announce evidence of neutrino oscillations in 1998
(see later), which means a neutrino has non-zero mass.
In 2001, several thousand PMTs imploded, apparently in a
chain reaction as the pressure waves from each imploding
tube cracked its neighbours. The detector has been partially
restored with about 5000 photomultiplier tubes with
protective shells that will prevent the chain reaction from
recurring.
8.4 The solar neutrino problem
When Homestake began taking data in 1968 it consistently recorded fewer neutrinos than expected from the
Sun. We know the solar luminosity, the energy released in the various PP cycles, how many neutrinos are
emitted in those cycles and so should be able to calculate the flux accurately. So why did Homestake only
see 30% of the expected rate? This became known as the solar neutrino problem and many decades of
argument ensued. Was something wrong with the experiment? Was nuclear physics wrong? Perhaps the
solar model was wrong? Maybe chlorine is only sensitive to high energy neutrinos?
However, as other experiments sensitive to lower energy neutrinos came on line the same result emerged.
GALLEX and SAGE saw only 40% and Super Kamiokande saw only 50% of the expected signal. The
probability that something more fundamental was wrong began to be taken seriously.
Between 1968 and 2002 discussion raged. Within this period the muon neutrino was detected and the tau
neutrino, hypothesized in 1975, was finally detected in 2000. The accepted theory at the time was that
neutrinos were massless meaning that the type of neutrino would be fixed when it was produced. The Sun
emits only electron neutrinos and so all solar neutrinos were expected to be
electron neutrinos. All the neutrino detectors at the time were only sensitive to
electron neutrinos. However if neutrinos have mass, they could change (or
oscillate) from the type (or flavour) that they had been produced as within the
Sun's interior into muon or tau neutrinos - two types that would not be caught
by the detectors in use at the time. Thus, the "missing" solar neutrinos could
be electron neutrinos which changed into other types along the way to Earth
and therefore escaped detection.
The first direct evidence of solar neutrino oscillation came in 2001 from the Sudbury Neutrino Observatory
(SNO) in Canada. It detected all types of neutrinos coming from the Sun, and was able to distinguish
between electron neutrinos and the other two flavours. It was found that about 35% of the arriving solar
neutrinos are electron neutrinos, with the others being muon or tau neutrinos. The total number of detected
neutrinos agrees with the earlier predictions from nuclear physics, based on the fusion reactions inside the
Sun.
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8.5 Sudbury Neutrino Observatory (SNO) and how it solved the solar neutrino problem
SNO is a neutrino observatory located about 2 km underground in Ontario, Canada. The detector was
designed to detect solar neutrinos through their interactions with a large tank of heavy water. The detector
turned on in 1999, and was turned off on in 2006 although analysis of the data recorded still continues.
The SNO detector target consists of 1,000 tonnes of heavy water ( D2 O or
2
1 H 2 O ) contained in a 12 metre diameter transparent acrylic vessel. The
detector cavity outside the vessel is filled with normal water to provide
both buoyancy for the vessel and radioactive shielding. The heavy water is
viewed by approximately 9,600 PMTs mounted on a sphere.
These PMTs detect tiny flashes of light emitted as neutrinos are
stopped or scattered in the heavy water. At a detection rate of about one
neutrino per hour, many days of operation are required to provide
sufficient data for complete analysis. Because SNO uses heavy water, it
is able to detect not only electron-neutrinos through the scattering
interaction (which Super- Kamiokande relies on), but also the other
neutrino flavours through different interaction processes, namely the
charged current and the neutral current interactions.
8.5.1 Charged current interaction
In the charged current interaction, a neutrino converts the neutron in a deuteron to a proton. The neutrino is
absorbed in the reaction and an electron is produced. Solar neutrinos have energies smaller than the mass of
muons and tau leptons, so only electron neutrinos can
participate in this reaction. The emitted electron carries off
most of the neutrino's energy and is detectable. The
electrons produced in
this reaction are
emitted in all
directions, but there is
a slight tendency for
them to point back in the direction from which the neutrino came. Due
to the large energy of the incident neutrinos, the electron will be so
energetic that it will be ejected at light speed, which is actually faster
than the speed of light in water producing Cerenkov radiation. This is
detected by the PMTs, the amount of light proportional to the incident
neutrino energy. From the PMT hit patterns the energies of the
neutrinos can be determined and an angular distribution measured.
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8.5.2 Neutral current interaction
In the neutral current interaction, a neutrino dissociates the
deuteron, breaking it into its constituent neutron and
proton. The neutrino continues on with slightly less energy.
All three neutrino flavours are equally likely to participate
in this interaction. When the neutron captures on a
deuterium nucleus, gamma ray photons with roughly
6 MeV of total energy are produced. The directions of the
gamma rays are completely uncorrelated with the direction
of the neutrino.
8.5.3 Electron elastic scattering
In the elastic scattering interaction, a neutrino collides with an atomic electron and imparts some of its
energy to the electron. All three neutrinos can participate in
this interaction although this interaction is dominated by
electron neutrinos. This is the process through which
Super-Kamiokande observes solar neutrinos. The electrons
produced usually point in the direction that the neutrino
was travelling (away from the sun). Because this
interaction takes place on atomic electrons it occurs with
the same rate in both the heavy and light water.
With this ability to register interaction of all neutrino flavours with the target, SNO became the first
observatory to see the expected neutrino flux from the Sun.
8.5.4 Future possibilities for neutrino observations
 Super-Kamiokande will continue to provide new insights into the properties of neutrinos.
 SNO+, the upgrade to SNO will be completed within 8 years the main physics goals being the
measurement of pep solar neutrinos and geo-neutrinos (neutrinos from radioactive decays in the Earth's
core, mantle, and crust). In order to reach these goals liquid scintillator will be used as the target material.
Since a neutrino interaction with this liquid produces several times more light than an interaction in water,
the energy threshold for the detection of neutrinos can be lower.
 Further refinements of solar models can be expected.
 Observatories will collaborate to ensure coincident detection of any supernova events.
 New bigger observatories will be created to look for high energy neutrinos from gamma ray bursts and
active galactic nuclei.
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9 Gravitational waves
Einstein’s theory of General Relativity describes how space-time is
affected by mass. In it he describes how the fabric of space-time bends
and stretches when an object is placed in it. The traditional way of
visualising this is shown. Remember however that this is a 2
dimensional model. The curve is weak far away from the object, and
steeper nearer to it. The stretching is greatest very near to the object.
Remember too that both time and space are stretched. This distortion
becomes critical around objects of very high mass, black holes for
example forming a singularity (or very sharp spike) in the space-time
continuum.
Gravitational waves can be thought of as "ripples in space-time." Just
like a boat sailing through the ocean produces waves in the water,
moving masses like stars or black holes produce gravitational waves in
the fabric of space-time. A more massive moving object will produce
more powerful waves, and objects that move very quickly will produce more waves over a certain time
period. If you prefer, you can use the analogy of the movement of electric charge. When an object with
charge changes position, the electric field around it changes, the change moving out into space at the speed
of light. For example an oscillating charge distribution will cause electromagnetic waves to radiate (radio
transmitter). It is believed that the same occurs for gravitation but as yet no accepted theory exists and no
one has yet measured this effect.
Gravitational waves are also believed to be produced in an interaction between two or more massive objects,
such as the binary orbit of two black holes or the coalescence of galaxies. It is postulated that as these
objects orbit each other, they send out waves of "gravitational radiation" that are unobstructed by the fabric
of space and so can potentially provide information beyond the scope of photon (radio, visible, gamma)
observation. Gravitational waves could therefore be a unique source of information on collapsing and
colliding stars, supernovae, black hole formation and will validate the theory of general relativity.
Just like electromagnetic radiation, gravitational waves carry energy away from their sources and, in the case
of orbiting bodies, this is associated with a decrease in
the orbit. Consider for example the famous case of the
Hulse-Taylor Pulsar (PSR 1913+16) which is in a
binary orbit with another star and provided the first
(albeit still unconfirmed) evidence for gravitational
waves. Since 1974 the period of this pulsar has been
measured to high accuracy. General relativity tells us
that a binary system will emit energy as gravitational
waves and eventually the two objects will inspiral
towards each other and merge as the period of the orbit
gradually decreases. The figure from Weisberg and
Taylor (2004) shows the decrease in the orbital period
of the pulsar as a function of date. Although the
measured shift is only 40 seconds over 30 years, it has
been very accurately measured and agrees precisely
with the predictions from Einstein's theory of General
Relativity as shown.
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General Relativity states that the power (W) given off as gravitational waves by two masses m1 and m2 a
32 G 4 m12 m22 (m1  m2 )
distance r apart in a binary system is given by: P  
.
5 c5
r5
How much gravitational energy is given off by our Sun – Earth system?
Mass of Earth is 6×1024 kg
Mass of Sun is 2×1030 kg
G = 6.67×10-11 m3kg-1s-2
c = 3×108 m/s
Distance between is 1.5×1011 m
Gravitational waves are believed to be incredibly weak when they reach Earth and so any detector must be
sensitive enough to register very small variations in space-time. The technique of choice used to detect small
stretches (known as strains) in space-time is interferometry.
Laser beams are bounced back and forth along two long arms (a-c and b-d), being reflected by mirrors at the
ends to increase the effective distance covered by the light. The reflected beams are eventually recombined
and their interference pattern monitored using a
photodetector which locks on a specific dark fringe. As the
gravitational wave passes the strain of space is different in
the two arms as the wave propagates and so the path length
changes and the interference fringes shift. The shift is
recorded and the information generated is compared with
computer simulations based on a model of what the
incoming waveforms might look like. The greater the
distance travelled by the laser light, the more sensitive the
detector will be to small fluctuations. Improving the
sensitivity is key since expected variations in strain are
less than 10-21. These detectors also have to work across a wide frequency range as general relativity predicts
the frequency of gravitational waves to vary from 100 Hz for pulsars to 10-4 Hz for binary stars.
There are currently several ground-based detectors in
operation or under construction, including LIGO (USA),
VIRGO (Italy/France), GEO (Germany/Great Britain), and
TAMA (Japan). Detectors on Earth however are limited by
seismic vibrations, the length of the laser track (curvature of
Earth), the quality of the mirrors, and the requirement to
operate the laser in a vacuum to reduce power loss. Operating
in space may be a good way to avoid these problems. The
Laser Interferometer Space Antenna (LISA) will feature 3
spacecraft accurately positioned 5,000,000 km apart providing
3 interferometer signals and aims for strain sensitivities down
to 10-21. It is scheduled for launch in the next decade.
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10 Cosmic Ray Particles
Cosmic rays are energetic particles originating from outer space that impinge on Earth's atmosphere
seemingly from all directions. About 87% of all the incoming cosmic ray particles are simple protons, with
nearly 12% being helium nuclei, and slightly under 1% are heavier elements.
This is a similar composition to
that of the solar system. This is
interesting because it is
believed that most cosmic ray
particles originate from outside
the solar system hinting at a
fairly uniform element
distribution throughout the
galaxy. Of note are the
relatively high proportions of
lithium, beryllium, and boron,
all believed to originate from
the breakup of carbon, nitrogen
and oxygen in a process called
spallation. In this process a
carbon nuclei for example
collides with a proton,
breaking up to form helium
and beryllium maybe or any
other legitimate combination of
products.
The variety of particle energies
reflects the wide variety of
sources up to extreme energies
of around 1020 eV. (Particle
accelerators can only produce
1012 to 1013 eV). The spectrum
is not thermal (M-B
distribution) but is power law
shaped at high energies.
10.1 So where do cosmic ray particles come from?
mv 2
 Bqv .
Recall the action of a magnetic field on a charged particle is
r
So the radius of curvature is r 
mv
or more correctly since the particles are almost relativistic we should
Bq

1
 v2  2
 mv
say that r 
where   1  2  . The typical magnetic field in the solar system is 10-9 T and so a
Bq
 c 
100 GeV proton will have a radius of curvature of 3×1011 m. This is about 1 A.U. and it means for low
energies such as this, the trajectory bears no relation to the location of the source and that cosmic ray
particles of such energy could well have been created in the solar system. Actually we find that the low
energy cosmic ray flux does vary with the solar cycle.
But there are also magnetic fields throughout the galaxy typically of order 3×10-10 T. The radius of curvature
for 1014 eV protons can be shown to be around 1015 m. As the cosmic ray particles encounter such regions
their trajectories are bent, but different regions have different field directions and so the cosmic ray particles
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follow a random walk through the galaxy. Thus their origin cannot be ascertained and we have our
explanation as to why the flux is isotropic.
If we say that the Milky Way has a radius of about 3×1018 m, what is the energy of the cosmic ray proton
moving at near relativistic velocity which would be held within the Milky Way boundaries if the magnetic
field is 3×10-10 T ?
So we can loosely state that 100 GeV protons and below originate from the solar system. 1015 eV protons
perform random walks throughout the galaxy and 1017 eV protons may originate from outside the galaxy.
As to where ‘outside the galaxy’ means exactly we are not sure. The favoured model is that they are
produced in the accelerating shockwave of a supernova but who knows…
10.2 How old are cosmic ray particles?
Some cosmic ray nucleides are radioactive. Beryllium for example has a half life of 3.9×106 years.
Beryllium is produced via spallation and it decays to boron. Since this is short lived we can use the
abundance to estimate the average age of cosmic rays. Firstly we measure the cross sections for all the
different spallation reactions which make beryllium and boron. We model the process to predict the ratio of
beryllium to boron produced and then compare this to the measured ratio to see how much beryllium has
decayed. From this we can extract the mean age. This gives a value of 107 years indicating that these cosmic
ray particles have originated from within the galaxy.
10.3 What is the total flux of cosmic ray particles?
The maximum in the proton spectrum is at about 1 GeV and the flux at this energy is 2 m-2s-1sr-1(MeV)-1.
This means that a detector of 1m2 capable of accepting particles within an angular acceptance of 1 steradian
would record 2 counts per second within each 1 MeV range about this mean energy.
Attempts have been made to integrate the flux over all energies to give the total flux and from this the
energy density of all cosmic ray particles. This is estimated at 1 MeV/m3.
In order to fully appreciate the size of this we should compare it to the energy density of starlight at 0.3
MeV/m3. It can be seen that cosmic rays forma significant contribution to the Universe.
10.4 How do we detect cosmic rays?
Nuclei interact strongly with other matter, so
when the cosmic ray particles approach Earth
they begin to collide with the nuclei of
atmospheric gases mainly oxygen and nitrogen.
These collisions, in a process known as a
shower, result in the production of many pions
and kaons, unstable mesons which quickly
decay into muons.
Because muons do not interact strongly with the atmosphere many of these muons are able to reach the
surface of the Earth and even penetrate for some distance into shallow mines. Muons are ionizing radiation,
and may easily be detected by many types of particle detectors such as spark chambers or scintillation
detectors. If several muons are observed by separated detectors at the same instant it is clear that they must
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have been produced in the same shower event. The image above is a simplified picture of an air shower: in
reality, the number of particles created in an air shower event can reach in the billions, depending on the
energy and chemical environment (i.e. atmospheric) of the primary particle. All of the produced particles
stay within about one degree of the primary particle's path.
The figure below adds more detail to a typical cosmic ray particle shower.
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