Radio telescopes
5.12 Radio telescopes
These are used for the wavelength range that typically ranges from 0.75 m (408 MHz) down to 1 cm (22 GHz). This is the range of wavelengths that reach the ground through the ‘radio window’ limited at the bottom end by the ionosphere and at the upper end by absorption by water vapour in the atmosphere. A great advantage of the fact that radio waves are far longer than optical wavelengths is that they are not absorbed by dust. Optical astronomers can only see ∼10% of the way towards the centre of our Milky Way Galaxy due the thick dust lanes that lie along its plane. Radio astronomers really can observe parts of the universe
inaccessible to optical astronomers! The majority of modern telescopes are either prime focus paraboloids, where
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the centre of the primary refl ecting surface to a point where the radio receivers are located. They tend to have very short focal ratios as in Figure 5.1.
5.12.1 The feed and low noise amplifi er system
At the focus, a feed collects the radio waves that have been focused by the primary refl ector (and possibly a secondary refl ector) and passes the signal to a low noise amplifi er made using transistors with inherently low noise levels. However, at room temperature, the electrons within the amplifi er produce a signifi cant amount of thermal noise which can be reduced by cooling the amplifi ers to tem- peratures near to absolute zero. The amplifi ers are mounted within a cryostat that is cooled by a gaseous helium refrigeration system. To minimize the heat reaching the amplifi er from its surroundings, the interior of the cryostat is evacuated to eliminate heat transfer by convection, a radiation shield of highly polished metal reduces transfer of heat by radiation and the wires to supply current to the ampli- fi ers are carried by fi ne wires wrapped around the coldest part of the refrigeration system to extract heat within them. The receivers then reach temperatures of ∼ 12 K and the thermal noise is reduced to minimal levels.
However, the noise produced within the low noise amplifi er is only part of the total noise that the radio telescope will receive even when not directly pointing at
a radio source. The total noise is determined by the ‘system noise temperature’, which is the temperature of a cavity in which a noiseless receiver is placed which would give the same signal strength as the real system whose noise contributions come from a number of causes (Figure 5.20).
In summary, these are: (1) The noise produced by the cooled amplifi er; this might be of order 8 K.
(2) Noise left over from the Big Bang. This is called the Cosmic Microwave Background, and has an equivalent temperature of ∼3 K; it will be discussed in detail in Chapter 9. (3) Noise from our own Milky Way Galaxy, caused by electrons travelling close to the speed of light (called relativistic electrons) spiralling around the magnetic fi eld of the Galaxy. It is called synchrotron radiation as it was fi rst observed being emitted from particle accelerators called synchrotrons.
(4) Radiation from molecules in the atmosphere, water vapour being a particular problem. The contribution will be less when the humidity is low, and high dry locations will always give a lower contribution. Higher frequencies are particularly affected, which is why Jodrell Bank’s radio
telescopes observing the Cosmic Microwave Background at 30 GHz (∼0.7 cm wavelength) are located on the fl anks of Mount Teide on Tenerife
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Figure 5.20 The contributions to the total system noise which is typically 25 K.
in the Canary Islands. Antarctica is a superb location from which to observe at very short wavelengths as the water vapour is frozen out of the atmosphere!
(5) The ground is at ∼290 K and so radiates as a black body at this temperature. Of course the radio telescope does not look at the ground, but some radiation can be diffracted around the edge of the dish, and the struts that support the secondary refl ector will also scatter radiation into the feed system of the telescope.
The higher the frequencies transferred down coaxial cables from the focus the greater their loss, so the frequencies are translated downwards to lower frequencies (called down-conversion to intermediate frequencies) before transfer along cables to the laboratories where the data acquisition systems are located. In the more modern systems, the radio frequency data is sampled after initial amplifi cation and then transferred along fi breoptic cables carried by modulated infrared laser light.
5.12.2 Radio receivers
The simplest form of radio receiver detects the radio signal and measures the received signal strength. Such systems initially plotted out the data on chart recorders but now, of course, computers digitize and record the data. These can
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be used to carry out surveys of the radio sky. As the radio telescope scans across
a strong radio source, the received signal strength will rise to a maximum and then fall, plotting out a typically Gaussian shape on the chart. This is the beam width that can be calculated from the same equation that was used to calculate the resolution of an optical telescope. If we consider a typical scenario, that of using the 76-m Lovell Telescope at a wavelength of 21 cm then:
Beam width ⫽ ∆θ ⫽ 1.22 λ/D
⫽ 1.22 ⫻ 0.21/76 rad ⫽ 3.4 ⫻ 10 ⫺3 rad ⫽ 3.4 ⫻ 10 ⫺3 ⫻ 57.3 ⫻ 60 arcmin ⫽ 11.6 arcmin
You can immediately see that the effective resolution of even the third largest fully steerable telescope in the world is ∼12 times less than the human eye at visible wavelengths! In fact this is an underestimate of the actual beam width as the feed of the antenna will not collect waves refl ected from towards the edge of the dish as effi ciently as from its centre. This means that the effective diameter of the dish is smaller than its actual size, and this can be compensated for by using a larger constant in the equation for the beam width. A value of 1.6 rather than 1.22 is appropriate, giving a beam width for the Lovell Telescope at 21cm
wavelength of ∼15 arcmin.
5.12.3 Telescope designs
The Lovell Telescope is a prime focus design, where the receiver is placed directly at the focus of the primary refl ector (Figure 5.21). However, most modern telescopes use a Cassegrain confi guration as in Figure 5.22. One advantage is that in the region of the focal point, which is behind the primary refl ector, it is possible to locate
a ‘carousel’ supporting a number of receivers each of which can be brought to the focal point by simply and quickly rotating the carousel. This gives what is termed ‘frequency fl exibility’. It should be pointed out that a Cassegrain design is not well suited for low frequencies as the size of the secondary refl ector has to become prohibitively large. In some telescopes the secondary mirror can be replaced by
a primary focus box, and in others a dielectric ‘lens’ is used above the feed at the Cassegrain focus to improve the effi ciency. At very low frequencies, a primary focus system is mandatory; which is why NASA, whose space tracking radio telescopes are of Cassegrain design, asked for the use of the (prime focus) Lovell Telescope in attempts to receive signals directly from two space probes, lost at Mars, which should have been transmitting at the frequencies just above 400 MHz.
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Figure 5.21 The 76-m Lovell Telescope – a prime focus design. Image: Ian Morison.
The shape of large parabolic antennas such as the 100-m Effelsberg Tele- scope in Germany and the 105-m Green Bank Telescope in West Virginia, USA will distort as their elevation changes due to changes in the gravitational forces acting on the structure (Figure 5.23). Both telescopes employ tech- niques to mitigate these effects. The Effelsberg Telescope uses a novel method which allows the shape to distort, but always to a surface with an accurate parabolic shape. The position of the focus will vary in elevation so the feed is moved to follow the focal point. In the Green Bank Telescope each of the 2004 panels that make up the surface area is mounted on actuators located at each of their corners. A laser system, mounted close to the focus, is used to measure the precise position of each panel and to adjust them as necessary to compensate for the gravitational deformations. (This is equivalent to the tech- nique of active optics that is used for optical telescopes.) The swan neck holds the secondary mirror and receiver systems away from the path of radio waves
falling on the primary reflector. This is called an offset-feed system. It prevents the supports that hold the secondary reflector and receivers scattering noise into the system.
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Figure 5.22 The 32-m Cambridge Telescope – a Cassegrain design. The cone houses
a carousel supporting four receivers. Image: Ian Morison.
5.12.4 Large fi xed dishes
Currently the world’s largest radio telescope is the 305-m Arecibo Telescope in Puerto Rico (Figure 5.24). It is built within a natural depression in the ground and has a spherical surface. High above the dish is a complex system of secondary and tertiary mirrors in the receiver complex to correct for the inherent spherical aberration. The secondary mirror is itself 25 m in diameter! As a radio source moves across the sky it is ‘tracked’ by moving the secondary mirror and receiver complex along an arm suspended above the surface.
A telescope called FAST is being built in China which will have an overall diameter of 500 m. Each of the 4600 panels that will make up the surface will
be mounted on actuators so that, as a radio source moves above the telescope, part of the surface, approximately 300 m across, is formed into a parabolic sur- face to refl ect incident radio waves up to a focus above the surface. The focus box is moved above the surface to follow the focal point by controlling the length of the support cables. This provides a coarse adjustment, with a laser controlled system to provide fi ne control of the feed position within the focus box. This innovative telescope should be completed by 2013.
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Figure 5.23 The 105-m Robert C. Byrd Green Bank Telescope. Image: National Radio Astronomy Observatory.
Figure 5.24 The 305-m Arecibo Telescope. Image courtesy of the NAIC – Arecibo Observatory,
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5.12.5 Telescope arrays
The resolution provided at radio wavelengths by even the largest single dishes is far below that of even a small optical telescope. By linking a number of
antennas into an array it is possible to synthesize a telescope whose effective size is given by the greatest separation of antennas in the array. Consider a line of antennas 5 km long whose central antenna was located at the North Pole. Looking down from above, this line of telescopes would be seen to rotate and, in 12 h, would fi ll out a circular area of diameter 5 km. If the data received by all of the antennas during the 12 h period is processed in an appropriate way an image of the sky above can be produced whose resolution is equal to that of
a 5-km telescope. As this depends on the rotation of the Earth the technique is termed ‘Earth Rotation Aperture Synthesis’. It is obviously not practical to locate an array at the North Pole, but providing one is not too near the equator the technique is viable.
At each instant of observation, each pair of telescopes is measuring some information about the structure of the radio source depending on their separation and orientation relative to the radio source. Due to the rotation of the Earth, the information about the radio source changes throughout the observing period and this enables an image of the source to be created. If there are N antennas in the array there are N(N-1)/2 combinations. The signals from each antenna are brought to a central ‘correlator’. Here delays have to be intro- duced into the signal path of each telescope so that the signals are combined
coherently. Three signifi cant arrays are:
• The Very Large Array (VLA) in the plains of New Mexico (Figure 5.25). It comprises three arms of nine 25-m antennas in a ‘Y’ formation. The array can be confi gured in a number of ways with a greatest overall separation of
36 km. Waveguides are currently used to carry the radio signals to the central ‘ correlator’ where the signals from the individual telescopes are combined in pairs. It can operate at wavelengths of 400 cm down to 0.7 cm.
• The MERLIN array across the UK has an overall size of 217 km and normally incorporates fi ve 25-m antennas and one 32-m antenna. For observations where high sensitivity is required the 76-m Lovell Telescope can be incorpo- rated into the array. It can operate at wavelengths of 74 cm down to 1.3 cm.
• The Giant Metre Radio Telescope (GMRT) is located at a site near Pune in India. The GMRT consists of thirty 45-m diameter dishes separated by up to
25 km. Fourteen of the 30 dishes are located in a compact central array in
a region of about 1 km 2 with the remaining 16 dishes spread out along the three arms of a ‘Y’-shaped confi guration. It can operate over a wavelength
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Figure 5.25 The Very Large Array in New Mexico, USA. Image: Wikipedia Commons.
The resolution of these arrays depends on the overall diameter and operating wavelength. One signifi cant attribute of the MERLIN array is that its resolution is comparable with that of the HST. At the wavelength of 5 cm the resolution is given by:
Resolution ⫽ ∆θ ⫽ 1.22 λ/D
⫽ 1.22 ⫻ 0.05/217 000 rad ⫽ 2.8 ⫻ 10 ⫺7 rad
⫽ 2.8 ⫻ 10 ⫺7 ⫻ 57.3 ⫻ 3600 arcsec ⫽ 0.06 arcsec
This is 1/15th of an arcsecond, comparable with that of the HST resolution at visible wavelengths and so allows images in the optical and radio parts of the
electromagnetic spectrum to be directly compared.
5.12.6 Very Long Baseline Interferometry (VLBI)
The highest resolutions available in any part of the electromagnetic spectrum are provided by arrays of radio telescopes that cross continents in what are called
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stretching from the UK across to Poland and from Finland to Sicily. It is also possible to include two antennas in China, one in South Africa and the Arecibo Telescope in the Caribbean. As a radio source is observed simultaneously at each observatory, the data is digitized and ‘time stamped’ using the observatory’s atomic clock. The data are then stored on arrays of hard disks. At the end of the observing session the hard disks are shipped to Holland where the data are combined in a correlator at JIVE – the Joint Institute for VLBI in Europe. It is now also possible to transfer data directly to JIVE using the internet and this allows immediate analysis of the data which is a great asset when planning observations of transient events. Due to the great size of the array, resolutions of less than 1 milliarcsec are possible allow- ing, for example, the expanding shell of a supernova in a galaxy 12 million light years to be imaged and its expansion rate measured (as shown in Figure 5.26). The white spots in the radio image, lower right, are supernovae remnants, one of which has been imaged by the EVN over a 12-year interval.
Figure 5.26 The starburst galaxy M82. Upper left: Hubble Space Telescope image of M82 in visible light.
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Two further VLBI arrays exist; the US Very Long Baseline Array (VLBA) that stretches from Hawaii to St Croix in the Caribbean, and the Asia Pacifi c Telescope (APT) which incorporates antennas in Australia, Japan, China and South Africa to provide high resolution studies of radio sources in the southern hemisphere.
5.12.7 The future of radio astronomy
The international radio astronomical community is now building and planning the radio telescopes that will dominate the observations of the twenty-fi rst century – the Atacama Large Millimetre Array (ALMA) and the Square Kilometre Array (SKA). ALMA is being constructed in the Atacama Desert region of Chile at a height of 5000 m (Figure 5.27). Here there is little water vapour in the atmo- sphere above so enabling observations from 10 mm down to as little as 0.3 mm wavelengths. The 50 or more individual antennas can be spaced up to several
kilometres apart giving unprecedented resolution at these short wavelengths. ALMA will be able to observe the clouds of gas and dust from which all the stars and galaxies in the universe form and help us to more fully understand the process of star formation.
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The SKA will be built in either South Africa or Australia and will have an effective collecting area of 1 km 2 – hence its name. It would thus have the same collecting area as 130, 100-m radio telescopes! The collecting area of the main array will, however, be made up arrays of small antennas, expected to be 15 m in size. This requires the use of more electronic systems but overall, the required collecting area can be achieved more cheaply than by the use of larger individual antennas.
The SKA will consist of an inner core of antennas making up a large part of the array’s total collecting area, together with outer stations arranged in a log- spiral pattern extending out to distances of 3000 km. It will thus combine high sensitivity allied to high resolution to give an instrument whose performance will greatly exceed radio telescope arrays currently in use. The SKA will be a highly fl exible instrument designed to address many of the most fundamental questions in astronomy. It will be the only instrument capable of observing the universe in the period, called the ‘dark ages’, when the hydrogen and helium formed in the Big Bang had yet to form the fi rst stars. Close to the central core will be a ‘phased array’ of receiving elements capable of observing much of the sky at one time – ideal for detecting transient events – and also able to carry out several observing tasks simultaneously. The artist’s impression of the SKA in Figure 5.28 shows the central core at the upper right, some of the outer stations in the foreground and middle left, the phased array. Phase 1 of the project, with over 10% of the total collecting area, is hoped to be completed by 2015 with fi nal completion of the array in 2020.
Figure 5.28 Artist’s impression of the central region of the SKA. Image: Xilos Studios,
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