Hubble Space Telescope
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- LIGHT-GATHERING POWER
- Figure 5.10 Sensitivity
- Figure 5.11 Mauna Kea Observatory
- Figure 5.12 VLT Observatory
- Figure 5.13 Resolving Power
- RESOLVING POWER
- Figure 5.14 Resolution
- Concept Check
- Figure 5.15 Atmospheric Turbulence
- Figure 5.16 European Southern Observatory
- Figure 5.17 CCD Chip
- Figure 5.18 Image Processing
- Figure 5.19 Active Optics
- Figure 5.20 Adaptive Optics
5.2 Telescope Size 
 Astronomers generally prefer large telescopes over small ones, for two main reasons. The first has to do with the amount of light a telescope can gather—its light-gathering power. The second is related to the amount of detail that can be seen—the telescope’sresolving power.
LIGHT-GATHERING POWER One important reason for using a larger telescope is simply that it has a greater collecting area, which is the total area of a telescope capable of gathering radiation. The larger the telescope’s reflecting mirror (or refracting lens), the more light it collects, and the easier it is to measure and study an object’s radiative properties. Astronomers spend much of their time observing very distant—and hence very faint—cosmic sources. In order to make detailed observations of such objects, very large telescopes are essential. Figure 5.10 illustrates the effect of increasing telescope size by comparing images of the Andromeda Galaxy taken with two different instruments. A large collecting area is particularly important for spectroscopic work, as the received radiation in that case must be split into its component wavelengths for further analysis.
The observed brightness of an astronomical object is directly proportional to the area of our telescope’s mirror and therefore to the square of the mirror diameter. Thus, a 5-m telescope will produce an image 25 times as bright as a 1-m instrument because a 5-m mirror has 52 = 25 times the collecting area of a 1-m mirror. We can also think of this relationship in terms of the length of time required for a telescope to collect enough energy to create a recognizable image on a photographic plate. Our 5-m telescope will produce an image 25 times faster than the 1-m device because it gathers energy at a rate 25 times greater. Put another way, a 1-hour exposure with a 1-m telescope is roughly equivalent to a 2.4-minute exposure with a 5-m instrument. Until the 1980s the conventional wisdom was that telescopes with mirrors larger than five or six meters in diameter were simply too expensive and impractical to build. The problems involved in casting, cooling, and polishing a huge block of quartz or glass to very fine tolerances (typically less than the width of a human hair) were just too great. However, new high-tech manufacturing techniques, coupled with radically new mirror designs, make the construction of telescopes in the 8- to 12-m range almost a routine matter. Experts can now make large mirrors much lighter for their size than had previously been believed feasible and can combine many smaller mirrors into the equivalent of a much larger single-mirror telescope. Several large-diameter instruments now exist, and many more are planned. Currently, the largest operating optical telescopes are the twin Keck instruments atop Mauna Kea in Hawaii (see Figure 5.11), administered jointly by the California Institute of Technology and the University of California. Each telescope combines 36 hexagonal 1.8-m mirrors into the equivalent collecting area of a single 10-m reflector. The first Keck telescope became fully operational in 1992; the second was completed in 1996. The high altitude and large size of these devices make them particularly well suited for detailed spectroscopic studies of very faint objects, in both the optical and infrared parts of the spectrum. (Mauna Kea’s 4-km altitude minimizes atmospheric absorption of infrared radiation, making this site one of the finest locations on Earth for infrared astronomy.)
Numerous other large telescopes can be seen in Figure 5.11. Some are designed exclusively for infrared work; others, like Keck, operate in both the optical and the infrared. To the right of the Keck domes is the 8.3-m Subaru (the Japanese name for the Pleiades) telescope, operated by the National Astronomical Observatory of Japan. It saw “first light” in 1999. In the distance is the 8.1-m Gemini North instrument, also completed in 1999 by a consortium of seven nations including the U.S. Its twin, Gemini South, in the Chilean Andes, should become operational in 2001. In terms of total available collecting area, the largest telescope currently available is the European Southern Observatory’s optical-infrared Very Large Telescope (VLT), located at Cerro Paranal, in Chile (Figure 5.12). It consists of four separate 8.2-m mirrors that can function as a single instrument. The third mirror was completed in 2000, the fourth in 2001.
RESOLVING POWER A second advantage of large telescopes is their finer angular resolution. In general, resolutionrefers to the ability of any device, such as a camera or telescope, to form distinct, separate images of objects lying close together in the field of view. The finer the resolution, the better we can distinguish the objects and the more detail we can see. In astronomy, where we are always concerned with angular measurement, “close together” means “separated by a small angle on the sky,” so angular resolution is the factor that determines our ability to see fine structure. Figure 5.13 illustrates how the appearance of two objects—stars, say—might change as the angular resolution of our telescope varies. Figure 5.14 illustrates the result of increasing resolving power with views of the Andromeda Galaxy at several different resolutions. What limits a telescope’s resolution? One important factor is diffraction, the tendency of light, and all other waves for that matter, to bend around corners. (Discovery 3-1) Because of diffraction, when a parallel beam of light enters a telescope, the rays spread out slightly, making it impossible to focus the beam to a sharp point, even with a perfectly constructed mirror. Diffraction introduces a certain “fuzziness,” or loss of resolution, into the optical system. The degree of fuzziness—the minimum angular separation that can be distinguished—determines the angular resolution of the telescope. The amount of diffraction is proportional to the wavelength of the radiation divided by the diameter of the telescope mirror. As a result we can write, in convenient units, where 1 µm (1 micron) = 10-6 m (see Appendix 2). Thus, for a given telescope size, the amount of diffraction increases in proportion to the wavelength used. Observations in the infrared or radio range are often limited by its effects. For light of any given wavelength, large telescopes produce less diffraction than small ones.
For a given telescopic size, the amount of diffraction increases in proportion to the wavelength used. Observations in the infrared or radio range are often limited by its effects. For example, according to the formula above, in an otherwise perfect observing environment, the best possible angular resolution of blue light (with a wavelength of 400 nm) that can be obtained using a 1-m telescope is about 0.25"  (0.4/1) = 0.1". This quantity is known as the diffraction-limited resolution of the telescope. But if we were to use our 1-m telescope to make observations in the near infrared, at a wavelength of 10 µm (10,000 nm), the best resolution we could obtain would be only 2.5". A 1-m radio telescope operating at a wavelength of 1 cm would have an angular resolution of just under 1º.
For light of any given wavelength, large telescopes produce less diffraction than small ones. A 5-m telescope observing in blue light would have a diffraction-limited resolution five times finer than the 1-m telescope just discussed—about 0.02". A 0.1-m (10-cm) telescope would have a diffraction limit of 1" and so on. For comparison, the angular resolution of the human eye in the middle of the visual range is about 0.5'. 
 Concept Check
Give two reasons why astronomers need to build very large telescopes. 5.3 High-Resolution Astronomy
ATMOSPHERIC BLURRING As we observe a star, atmospheric turbulence produces continual small changes in the optical properties of the air between the star and our telescope (or eye). The light from the star is refracted slightly, and the stellar image dances around on the detector (or on our retina). This continual deflection is the cause of the well-known “twinkling” of stars. It occurs for the same reason that objects appear to shimmer when viewed across a hot roadway on a summer day. On a good night at the best observing sites, the maximum amount of deflection produced by the atmosphere is slightly less than 1". Consider taking a photograph of a star. After a few minutes of exposure time (long enough for the intervening atmosphere to have undergone many small random changes), the image of the star has been smeared out over a roughly circular region an arc second or so in diameter. Astronomers use the term seeing to describe the effects of atmospheric turbulence. The circle over which a star’s light (or the light from any other astronomical source) is spread is called the seeing disk. Figure 5.15 illustrates the formation of the seeing disk for a small telescope.*
Atmospheric turbulence has less effect on light of longer wavelengths—ground-based astronomers generally “see” better in the infrared. However, counteracting this improvement in image quality is the fact that the atmosphere is wholly or partially opaque over much of the infrared range. (Sec. 3.3) For these reasons, to achieve the best possible observing conditions, telescopes are sited on mountaintops (to get above as much of the atmosphere as possible) in regions of the world where the atmosphere is known to be fairly stable and relatively free of dust, moisture, and light pollution from cities. In the continental United States, these sites tend to be in the desert Southwest. The U.S. National Observatory for optical astronomy in the Northern Hemisphere, completed in 1973, is located high on Kitt Peak near Tucson, Arizona. The site was chosen because of its many dry, clear nights. Seeing of less than 1" from such a location is regarded as good, and seeing of a few arc seconds is tolerable for many purposes. Even better conditions are found on Mauna Kea, Hawaii (Figure 5.11), and at numerous sites in the Andes Mountains of Chile (Figures 5.12 and 5.16), which is why many large telescopes have recently been constructed at these exceptionally clear locations.
An optical telescope placed in orbit about Earth or on the Moon could obviously overcome the limitations imposed by the atmosphere on ground-based instruments. Without atmospheric blurring, extremely fine resolution—close to the diffraction limit—can be achieved, subject only to the engineering restrictions of building or placing large structures in space. The Hubble Space Telescope (HST; named for one of America’s most notable astronomers, Edwin Hubble) was launched into Earth orbit by NASA’s space shuttleDiscovery in 1990 (see Discovery 5-1). This telescope has a 2.4-m mirror, with a (blue-light) diffraction limit of only 0.05", giving astronomers a view of the universe as much as 20 times sharper than that normally available from even much larger ground-based instruments. IMAGE PROCESSING Computers play a vital role in observational astronomy. Most large telescopes today are controlled either by computers or by operators who rely heavily on computer assistance, and images and data are recorded in a form that can be easily read and manipulated by computer programs. It is becoming rare for photographic equipment to be used as the primary means of data acquisition at large observatories. Instead, electronic detectors known as charge-coupled devices, or CCDs, are in widespread use. Their output goes directly to a computer. A CCD (Figure 5.17) consists of a wafer of silicon divided into a two-dimensional array of many tiny picture elements, known as pixels. When light strikes a pixel, an electric charge builds up on the device. The amount of charge is directly proportional to the number of photons striking each pixel—in other words, to the intensity of the light at that point. The charge buildup is monitored electronically, and a two-dimensional image is obtained. A CCD is typically a few square centimeters in area and may contain several million pixels, generally arranged on a square grid. As the technology improves, both the areas of CCDs and the number of pixels they contain continue to increase. Incidentally, the technology is not limited to astronomy—many home video cameras contain CCD chips similar in basic design to those in use at the great astronomical observatories of the world.
CCDs have two important advantages over photographic plates, which were the staple of astronomers for over a century. First, CCDs are much more efficient than photographic plates, recording as many as 90 percent of the photons striking them, compared with less than five percent for photographic methods. This means that a CCD image can show objects 10 to 20 times fainter than a photograph made using the same telescope and the same exposure time. Alternatively, a CCD can record the same level of detail in less than a tenth of the time required by photographic techniques, or record that detail with a much smaller telescope. Second, CCDs produce a faithful representation of an image in a digital format that can be placed directly on magnetic tape or disk, or even sent across a computer network to an observer’s home institution. Computers are also widely used to reduce background noise in astronomical images. Noise is anything that corrupts the integrity of a message, such as static on an AM radio or “snow” on a television screen. The noise corrupting telescopic images has many causes. In part, it results from faint, unresolved sources in the telescope’s field of view and from light scattered into the line of sight by Earth’s atmosphere. It can also be caused by electronic “hiss” within the detector. Whatever the origin of noise, its characteristics can be determined (for example, by observing a part of the sky where there are no known sources of radiation) and its effects partially removed with the aid of high-speed computers, allowing astronomers to see features that would otherwise remain hidden. Using computer processing, astronomers can also compensate for known instrumental defects and even correct some effects of bad seeing. In addition, the computer can often carry out many of the relatively simple but tedious and time-consuming chores that must be performed before an image (or spectrum) reaches its final “clean” form. Figure 5.18 illustrates how computerized image-processing techniques were used to correct for known instrumental problems in the Hubble Space Telescope, allowing much of the planned resolution of the telescope to be recovered even before its repair in 1993.
NEW TELESCOPE DESIGN The latest techniques for producing ultrasharp images take these ideas of computer control and image processing several stages further. By analyzing the image formed by a telescopewhile the light is still being collected, it is now possible to adjust the telescope from moment to moment to avoid or compensate for the effects of mirror distortion, temperature changes in the dome, and even the effects of atmospheric turbulence. By these means, some recently constructed telescopes have achieved resolutions very close to their theoretical (diffraction) limits. Even under conditions of perfect seeing, most telescopes would not achieve diffraction-limited resolution. The temperature of the mirror or in the dome may fluctuate slightly during the many minutes or even hours required for the image to be exposed, and the precise shape of the mirror may change slightly as the telescope tracks a source across the sky. The effect of these changes is that the mirror’s focus may shift from minute to minute, blurring the eventual image in much the same way as atmospheric turbulence creates a seeing disk (Figure 5.15). At the best observing sites, the seeing is often so good that these tiny effects may be the main cause of image blurring. The collection of techniques aimed at controlling these environmental and mechanical fluctuations is known as active optics.
The first telescope designed to incorporate active optics was the New Technology Telescope (NTT), constructed in 1989 at the European Southern Observatory in Chile (NTT is the most prominent instrument visible in Figure 5.16). This 3.5-m instrument, employing the latest in real-time telescope controls, achieved a resolution of about 0.5" by making minute modifications to the tilt of its mirror as its temperature and orientation changed, thus maintaining the best possible focus at all times. Figure 5.19(a) shows how active optics can improve image resolution. Active optics techniques now include improved dome design to control airflow, precise control of the mirror temperature, and the use of actuators (pistons) behind the mirror to maintain its precise shape at all times (Figure 5.19b). All of the large telescopes described earlier include active optics systems, improving their resolution to a few tenths of an arc second. NTT itself was upgraded in 1997 and has reported resolution as sharp as 0.15" under the best conditions. With active optics systems in place, Earth’s atmosphere once again becomes the main agent limiting a telescope’s resolution. Remarkably, even this problem can now be addressed, using an approach known as adaptive optics. This technique actually deforms the shape of the mirror’s surface, under computer control, while the image is being exposed, in order to undo the effects of atmospheric turbulence. Adaptive optics presents formidable theoretical and technological problems, but the rewards are so great that they have been the subject of intense research since the 1980s. Declassified SDI (“Star Wars”) technology in the 1990s provided an enormous boost to this effort. In the experimental military system shown in Figure 5.20(a), lasers probe the atmosphere above the telescope, returning information about the air’s swirling motion to a computer that modifies the mirror thousands of times per second to compensate for poor seeing. Astronomical adaptive optics systems generally do not use lasers to gauge atmospheric conditions. Instead, they monitor standard stars in the field of view, constantly adjusting the mirror’s shape to preserve those stars’ appearance. Figure 5.20(b) compares the results of a pair of observations of a nearby double star called Castor. Spectroscopic observations long ago revealed the double nature of this object, which appears to be a binary-star system. The image on the left shows the star system as seen through an ordinary, moderate-sized telescope—an oblong blur combining the light from the two stars, spread over several arc seconds. On the right, with the adaptive optics system turned on, one star is clearly distinguishable from the other. The two stars are separated by less than an arc second; the resolution, with adaptive optics turned on, is about 0.1".
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Even large telescopes have limitations. For example, according to the discussion in the preceding section, the 5-m Hale telescope should have an angular resolution in blue light of around 0.02". In practice, however, it cannot do better than about 1". In fact, apart from instruments using special techniques developed to examine some particularly bright stars, noground-based optical telescope built before 1990 can resolve astronomical objects to much better than 1". The reason is Earth’s 
