Most optical observatories are sited remote from human habitation, usually at the tops of mountains. This siting results from the need for dark, clear skies. At high altitudes the chances of cloudless skies are improved because the observer’s view is usually above the cloud layer.
Choosing a permanent site for an observatory can be difficult. The Royal Greenwich Observatory, for example, was originally located at Greenwich, in London, but as London pollution worsened and the number of street lights gradually increased, conditions for observing the night sky became intolerable. Between 1948 and* 1957, the observatory was moved south and re-established in the grounds of Herstmonceux Castle, near Hailsham in Sussex. Unfortunately, it was not long until atmospheric pollution began to put pressure on this establishment too. An extensive search led to the choice of a volcanic peak on the island of La Palma, in the Canaries, so that the observatory could be based high above the clouds at 7,765 feet (2,367 meters) above sea level.
An observatory is easily recognized by its characteristic cluster of domed buildings, each usually housing a telescope. Most observatories also have laboratory facilities for coating mirrors and for servicing instruments, computer rooms and workshops, as well as residential accommodations. Many observatories are so remote that they have to be totally self-sufficient for weeks at a time, often cut off by adverse weather conditions.
Most of the telescopes in observatories use reflecting optics and are of the Cassegrainian design, in which a secondary convex mirror
reflects the incoming light back through a hole in the center of the main mirror. Large refractors are less common, because of the problems of lens sag and the long length of tube required.
Many modern telescopes are fully automatic, such as the Anglo-Australian Telescope (AAT), which has a pointing accuracy of one second of arc. Most telescopes are on equatorial mounts (one axis of the mounting lies parallel with the earth’s axis of rotation). To track an object as it moves across the sky, it is necessary only to rotate the telescope slowly round the mounting axis. This method of tracking, however, leads to large, heavy, and expensive mountings. Modern mounting designs simply have a vertical and a horizontal axis (the alt-azimuth mount), both of which are driven by motors.
The heart of any reflecting telescope is its main mirror, and until recently this had to be quite thick in order to avoid flexure as the telescope moved. Far more economical designs, using thin mirrors, are now possible, allowing lightweight telescopes to be built. In these designs, deflections in the mirror are detected and corrected automatically by computer-controlled servomechanisms. Such systems, often called “rubber mirrors,” are extremely successful, an example being the United Kingdom Infrared Telescope (UKIRT) located in Hawaii.
For improved detail, as well as for studying fainter objects, it is necessary to use larger telescopes. Making single, large mirrors is a difficult and expensive task. The large telescopes of the future will probably have segmented mirrors. Several telescopes incorporating this design have already proved to be successful, namely the multiple-mirror telescope (MMT) at Mount Hopkins, near Tucson, Arizona, and the Keck Telescope at Mauna Kea, Hawaii.
Since the Apollo missions to the moon and the spectacular success of the American planetary probes, there has been little point in studying planetary objects with ground-based telescopes. Nevertheless, a precise measurement of the earth-moon distance, with an accuracy of a few inches, can be obtained. This is done by bouncing a pulse of laser light from a special reflector left by Apollo astronauts on the moon. Earth-based telescopes are used to detect the weak, reflected pulse. By timing accurately the journey of the pulse to the moon and back to earth, the exact earth-moon distance is easily evaluated.
The sun is our nearest star and, as such, has been an object of intensive study with a wide range of instruments. The solar astronomer does not suffer from the shortage of light that plagues deep-space astronomers, and therefore usually uses smaller aperture instruments than those used on the night sky. High magnification can be employed to reveal a wealth of detail on the solar surface.
The world’s largest solar telescope is the 60-inch (1.5-meter) solar telescope at the Kitt Peak National Observatory, southwest of Tucson, Arizona. Unlike other telescopes, this instrument has its main mirror buried deep in the ground, at the lower end of a tunnel parallel with the earth’s axis of rotation. The light of the sun is reflected down the tunnel by a large, driven, flat mirror at the top, open end of the tunnel. This arrangement has two advantages: the only moving part is the flat mirror, rather than the whole massive telescope; and the tunnel can be kept at a uniform temperature (because most of it is underground) and so eliminates any thermal air current that would cause a magnified image to badly “boil.”
The spectrohelioscope is a type of scanning spectroscope that enables an image of the sun to be formed in one particular wavelength, instead of the usual mixture making up white light. By examining the sun in hydrogen light, the distribution of hydrogen on the sun’s surface and in its atmosphere can be viewed. This technique makes visible the arched prominences and solar flares that could otherwise not be seen. A “narrow-band” interference filter is usually used, which passes only an extremely narrow range of wavelengths.
A more difficult observation is the study of the faint solar corona—the extended, intensely hot, outer atmosphere of the sun. The instruments used for this purpose are called coronagraphs; they are essentially telescopes with a special mask to obscure the sun’s disc. The design of coronagraphs has to be extremely skillful in order to prevent scattered stray light within the instrument from swamping the weak image of the corona. Improved images of the sun’s outer atmosphere were obtained with a type of coronagraph on board Skylab, because on the ground, the earth’s atmosphere causes so much scatter.
The photographic emulsions used at observatories still consist basically of a silver halide which is decomposed by exposure to light, although today, they are very sophisticated versions of their forerunners. Great improvements have been made in sensitivity, contrast, and grain size; the reduction of the latter allows far more detail to be recorded on a plate. Considerable effort has also gone into “gas hypersensitizing,” in which ordinary emulsions are baked in an atmosphere enriched with hydrogen, resulting in a substantial increase in sensitivity and speed.
The age of the photographic emulsion may be drawing to a close, with modern electronic detectors rapidly becoming superior in performance. Even so, a 30-minute exposure from a large telescope can produce a photographic plate that records the images of hundreds of thousands of stars and galaxies. But this information can take many weeks to analyze. The problem is that the quality of the information is considerably lower than that obtained by modern electronic detectors. The true role of photography in astronomy is really in producing sky maps and surveys. For this purpose it will probably remain useful for many years to come.
Optically fast telescopes, called Schmidt cameras, have been developed for producing large-field photographs of the sky. These cameras are made up of a large, spherical primary mirror, with glass correcting plates in front of it to compensate for the aberrations produced by the mirror. The United Kingdom’s Schmidt telescope, at the Siding Spring Observatory in Australia, has produced a 606-plate survey of the southern skies and careful examination of some of the plates led to the discovery of many new galaxies and nebulae.
It has long been realized that by studying the spectrum of an object, much information can be gained concerning its chemical composition, temperature, movement, and other physical phenomena. Astronomers have, for many decades, applied this technique to the study of stars, galaxies, and nebulae. Originally, a glass or quartz prism was used to break the incoming light into its component colors, and the resulting spectrum was photographed. Modern spectrographs use diffraction gratings rather than prisms. (Gratings are highly reflecting surfaces on which are ruled many very thin parallel lines.) The reflected light forms a spectrum that is dispersed far more than that produced by a prism, enabling more detail to be seen. These gratings have several hundreds of lines per millimeter ruled on their faces, which are commonly a thin film of aluminum deposited on an optically flat glass or quartz block. The modern spectrograph is sensitive enough to show wavelength shifts caused by the line-of-sight velocities of astronomical objects. Thus, a fairly accurate determination can be made of the speed of expansion of nebulae, rotation of galaxies, and expansion of the universe. Photographic films are still used with spectrographs because they can record the whole spectrum simultaneously. Increasingly, however, electronic detectors of various forms are replacing or supplementing film.
Second only to the spectrometer, the photometer is the astronomer’s most useful instrument. Photometric devices accurately measure the intensity of light from the stars and galaxies. They are nearly all based on highly sensitive, light-detecting electron tubes, called photomultipliers. In these devices, weak starlight causes electrons to be emitted from the surface of a photocathode, and each electron released can be made to produce many millions of secondary electrons. The output signal is easily measured and recorded and accurately represents the intensity of the light falling on the photocathode.
By measuring with color filters in front of the photomultiplier, it is possible to determine the temperature of the stars, the composition of nebulae, and the chemical processes in supernova explosions. Some modern photomultipliers are capable of detecting the individual photons that make up the incoming light. This enables the powerful technique of photon counting to be used in measuring extremely faint objects.
A particularly useful application of photometry is the plotting of the light-curves of variable stars, particularly where these stars are close binaries. Photometric observations of their varying light output reveal not only their dual nature, but also their masses, separation, relative diameters, and other information.
One of the most exciting additions to the astronomer’s laboratory is the image intensifier. For many years, telescopes were made larger, not in the search for better resolution, but to gather more light. The image intensifier is a form of electron tube in which an image is focused onto an input photocathode. The resulting emitted electrons are accelerated within the evacuated tube by a very high voltage. An electron lens focuses these electrons onto a fluorescent screen (like that in a television tube) which emits a bright flash of light when a high-speed electron strikes it. In this way, a single photon at the input releases a single electron, which is accelerated to hit the screen and causes the emission of many new photons. Such a tube has a “gain” of about 100. That is, the image on the screen is the same as that on the photocathode, but one hundred times brighter.
An image intensifier placed in the image plane of a spectrograph, with the film in contact with the screen, improves the sensitivity of the system dramatically. Often, two or three intensifiers are coupled; the input of one seeing the output of another, and so on. These “cascade” intensifiers are capable of intensifying images by many thousands of times. A three-stage intensifier (used in front of the photographic film on a 12-inch (0.3-meter) telescope) results in shorter exposure times than for plain film (with a ZOOMnch (5-meter) reflector), such as that at Mt Palomar.
Image intensifiers are used in many instruments: cameras, spectrographs, acquisition and guidance systems. Recently, a completely different and more compact type of tube, called a. microchannel plate intensifier, has become available. This does the same job as the tubes described previously, but is free from the pincushion distortion that early tubes produced.
Considerable effort is currently being expended on replacing the output screen of intensifiers with an electronic encoding device. This would permit the electron image to be read out directly into a computer for further processing, rather than converting it back into light. Results are promising, with prototype devices already being tested on telescopes in the United States and the United Kingdom. But these devices are still at too new a stage to have yet contributed significantly to our knowledge of the universe.
Since the mid-1970’s, a new type of detector has gained popularity in astronomy, called the charge-coupled device, or CCD. The CCD is a semiconductor device and is an offshoot of large-scale, silicon-integrated circuit technology. Essentially, it comprises an array of tiny detectors that each store an electric charge in proportion to the amount of light that has fallen on them. The array is then “read out” electronically into a computer memory, where the data is stored for later processing. The CCD was developed as a replacement for television camera tubes, and when used in this way, it had little to offer the astronomer. When cooled to about —238” F. (—114° C), however, it can be used in the so-called “slow scan” mode, permitting exposures of many hours to be made.
Modern CCD’s have about 1,500 separate detector elements, or “pixels,” each measuring about 25 microns square, making the active area of a CCD array only about 0.47 by 0.32 inches (12 by 8 millimeters). These devices are used as if they were photographic films that can be examined directly by a powerful computer. A further advantage of these detectors is their robustness. Even if a CCD were exposed to direct sunlight, it would be unharmed, whereas image intensifiers and photomultipliers are destroyed if too much light reaches their photocathodes.
Image photon counting system
Perhaps the one electronic detector system that has contributed more to our knowledge of astronomy than any other is the /mage photon counting system (IPCS). This detector is a four-stage, magnetically-focused image intensifier, coupled via a lens to a plumbicon television camera tube. Magnetically-focused intensifiers are virtually free from the pincushion distortion produced by tubes that employ electrostatic electron lenses. The four stages provide an extremely high amplification of the faint images. A single photon, hitting the input photocathode, produces a bright flash on the output screen, which is “seen” by the television camera. Sophisticated electronics “writes” the position of the original photon in the memory of the computer. This process is repeated for each photon detected, and an image is built up in this way in the computer memory. This then can be processed later and displayed on a video screen or recorded on magnetic tape.
This photon counting system is the most sensitive yet developed for detecting faint images. It has been used almost exclusively for spectroscopy in place of photographic film, although some direct imaging has also been done with the IPCS. Only ai few such systems exist due to their great cost, all used on the largest telescopes, such aS the Palomar 200-inch (5-meter) and the Anglo-Australian telescope. For imaging applications, the IPCS is largely replaced by the cryogenic (cooled) slowscan CCD, but for applications that require time resolution it remains unrivaled.