Spherical Astronomy
3.6 Other Forms of Energy
Fig. 3.39 The European Herschel launched in 2009 was the most efficient recent infrared satellite.
The diameter of its main mirror, 3.5 metres, was bigger than in any earlier astronomical satellite. The final shape to the highly curved mirror was given by the Opteon company at the Tuorla observatory in Finland. The observations ended in 2103 when the liquid helium used to cool down the camera ran out.
(Photos ESA, Rami Rekola)
86 3 Observations and Instruments Fig. 3.39 (Continued)
The most energetic cosmic rays give rise to secondary radiation, mainly muons, when they hit molecules of the atmosphere. This secondary radiation can be observed from the ground. Just like particles produces by gamma rays particles of the secondary radiation can move faster than light in the atmosphere emitting ˇCerenkov radi- ation. Thus the cosmic rays can also be detected by their ˇCerenkov radiation.
The Pierre Auger telescope in Argentina ob- serves both secondary particles and ˇCerenkov ra- diation to detect cosmic rays with energies at least of the order of 1018eV. Since only about one such particle hits an area of one square kilometre in a century, the detectors have been positioned in an area of 3000 square kilometres.
Primary cosmic rays can only be directly ob- served outside the atmosphere. The detectors used to observe cosmic rays are similar to those used in particle physics. Since Earth-based accel- erators reach energies of only about 1012eV, cos- mic rays offer an excellent “natural” laboratory for particle physics. Many satellites and space- craft have detectors for cosmic rays.
Neutrinos Neutrinos are elementary particles with no electric charge. Earlier it was thought that
they are massless, but currently they seem to have a tiny rest mass. They interact with other mat- ter only through the weak nuclear force, which is very much weaker than the electromagnetic or strong nuclear force.
Most neutrinos are produced in nuclear reac- tions within stars and in supernova explosions.
Since they react very weakly with other matter, they escape directly from the stellar interior.
Due to their weak interaction neutrinos are very difficult to observe; trillions of them pass through this page every second without any no- ticeable consequences.
The first method of detection was the radio- chemical method. As a reactive agent, e.g. tetra- chloroethene (C2Cl4) can be used. When a neu- trino hits a chlorine atom, the chlorine is trans- formed into argon, and an electron is freed:
37Cl+ν→37Ar+e−.
The argon atom is radioactive and can be ob- served. Instead of chlorine, lithium and gallium might be used to detect neutrinos. The first gal- lium detectors have been running in Italy and Russia from the end of the 1980’s.
Another observation method is based on the Cerenkov radiation produced by neutrinos in ex-ˇ
Fig. 3.40 The Sudbury neutrino observatory (SNO) is at the depth of two kilometres in a still operational mine. The spherical vessel contains liquid where neutrinos produce Cherenkov radiation. 9600 photomultipliers are used to register the flashes of the radiation. Earlier the liquid was heavy water but that was later replaced with alkyl benzene.
The observations have already solved the solar neutrino problem (Sect.13.1). The new SNO+will be able to ob- serve neutrinos of even lower energies. (Photos SNOlab, H. Karttunen)
tremely pure water. The flashes of light are regis- tered with photomultipliers, and thus it is possi- ble to find out the direction of the radiation. This method is used e.g. in the Japanese Kamiokande detector and in the Sudbury neutrino observatory in Canada (Fig.3.40).
The largest neutrino detector is the IceCube in the Antarctic, finished in 2010. It consists of one cubic kilometre of ice.
Neutrino detectors must be located deep under the ground to protect them from the secondary radiation caused by cosmic rays. Also radioactive decay inside the Earth cases background noise.
The detectors have observed neutrinos from the Sun, and the Supernova 1987A in the Large Magellanic Cloud was also observed in 1987. In the future, also the dark matter and other neutrino producing phenomena will be studied.
Gravitational Radiation Gravitational astron- omy is as young as neutrino astronomy. The first attempts to measure gravitational waves were made in the 1960’s. Gravitational radiation is emitted by accelerating masses, just as electro- magnetic radiation is emitted by electric charges in accelerated motion. Detection of gravitational waves is very difficult, and they have been ob- served only very recently.
The first type of gravitational wave antenna was the Weber cylinder. It is an aluminium cylin- der which starts vibrating at its proper frequency when hit by a gravitational pulse. The distance between the ends of the cylinder changes by about 10−17m, and the changes in the length are studied by strain sensors welded to the side of the cylinder.
Another type of modern gravity radiation de- tectors measures “spatial strain” induced by grav- ity waves and consists of two sets of mirrors in directions perpendicular to each other (Michel- son interferometer), or one set of parallel mir- rors (Fabry–Perot interferometer). The relative distances between the mirrors are monitored by laser interferometers. If a gravity pulse passes the detector, the distances change and the changes can be measured. The longest baseline between the mirrors is in the American LIGO (Laser Inter-
88 3 Observations and Instruments
Fig. 3.41 The LIGO Livingston Observatory seen from the air. (Photo LIGO/Caltech)
ferometer Gravitational-wave Observatory) sys- tem (Fig. 3.41). LIGO made the first scientific observations in 2002. In early 2016 the first pos- itive detection of gravitational waves was an- nounced. The observation was made in Septem- ber 14, 2015.
Although the gravitational detectors are iso- lated from their surroundings as carefully as pos- sible, many perturbations can make them oscil- late. A positive detection requires that the same phenomenon is observed by different instruments located far away from each others.
Box 3.1 (Diffraction by a Circular Aperture) Consider a circular hole of radius R in the xy plane. Coherent light enters the hole from the direction of the negative z axis (see fig- ure). We consider light rays leaving the hole parallel to the xz plane forming an angle θ with the z axis. The light waves interfere on a screen far away. The phase difference be- tween a wave through a point(x, y)and a wave going through the centre of the hole can be calculated from the different path lengths s= xsinθ:
δ=s
λ2π=2πsinθ λ x≡kx.
Thus, the phase differenceδdepends on the x coordinate only. The sum of the amplitudes of the waves from a small surface element is proportional to the area of the element dxdy.
Let the amplitude coming through the centre of the hole be da0=dxdyˆi. The amplitude com- ing from the point(x, y)is then
da=dxdy
cosδiˆ+sinδjˆ .
We sum up the amplitudes coming from dif- ferent points of the hole:
a=
Aperture da
= R
x=−R √
R2−x2 y=−√
R2−x2
coskxiˆ+sinkxjˆ dydx
=2 R
−R
R2−x2
coskxˆi+sinkxjˆ dx.
Since sine is an odd function (sin(−kx)=
−sin(kx)), we get zero when we integrate the second term. Cosine is an even function, and so
a∝ R
0
R2−x2coskxdx.
We substitute x =Rt and define p=kR = (2π rsinθ )/λ, thus getting
a∝ 1
0
1−t2cosptdt.
The zero points of the intensity observed on the screen are obtained from the zero points of the amplitude,
J (p)= 1
0
1−t2cosptdt=0.
Inspecting the function J (p), we see that the first zero is atp=3.8317, or
2π Rsinθ
λ =3.8317.
The radius of the diffraction disk in angular units can be estimated from the condition
sinθ=3.8317λ
2π R ≈1.22 λ D, whereD=2Ris the diameter of the hole.
In mirror telescopes diffraction is caused also by the support structure of the secondary mirror. If the aperture is more complex and only elementary mathematics is used calcula- tions may become rather cumbersome. How- ever, it can be shown that the diffraction pattern can be obtained as the Fourier transform of the aperture.