Matter and the radiation it produces or annihilates are of vital importance to astronomy. The physical nature of various astronomical bodies dominates the efforts of astronomers, and it is through the radiation emitted by these bodies that we are able to say what they are like. To a great extent, work in astronomy today is what has been called "the practice of radiation diagnostics," that is, the collection, analysis, and interpretation of the radiation emitted all across the electromagnetic spectrum by astonomical bodies. Therefore, we want to continue in this chapter our investigation into the nature of matter, radiation, and its detection that we began in Chapter 4.
Some of the earliest philosophical speculations, such as that
by Thales, were concerned with what the material world is made
of. Is each substance, such as rock or wood, infinitely divisible
so that its subdivisions always yield the same properties as the
whole substance? Or is there some level of structure below which
the subdivisions will show new properties and forms? The Greek
philosophers Democritus (460?-?370 B.C.), Leucippus (c. 440 B.C.),
and later Epicurus (341-270 B.C.) suggested that the material
objects of our experience are actually made up of fundamental
units. They called these units atoms, and they visualized
them as indestructible, indivisible, infinite in number and variety,
and capable of being assembled into various forms and shapes.
However, the most important aspect was the concept that matter
was never born from nothing, but sprung only from new combinations
of atoms. As pointed out in Chapter 1, they laid the foundation
for the concept of transformability as one of the great themes
in science. Along these lines, the Roman poet Lucretius (99?-?55
B.C.) suggests that, "All nature then, as it exists by itself,
is founded on two things: there are bodies and there is void in
which these bodies are placed and through which they move about...."
Thus physical existence is composed of two realities; everlasting
particles, atoms, and what we may loosely call vacuum, or the
absence of particles.
5.1.1. Atoms, Molecules, and Nuclei
Although, a number of early conceptual schemes were based on small indivisible particles, it was not until the period between Robert Boyle (1627-1691), British chemist-physicist, and the death of the French chemist Antoine Lavoisier (1743-1794) that the concept of the chemical element was to arise and give a new, vital, and precise meaning to atoms. The concept of the atom was not widely accepted until the English scientist John Dalton (1766-1844) developed his atomic theory of gases around 1800. Dalton stated that, "Matter, though divisible in an extreme degree, is nevertheless not infinitely divisible. That is, there must be some point beyond which we cannot go in the division of matter....I have chosen the word atom to signify these ultimate particles....[which for] all homogeneous bodies are perfectly alike in weight, figure, etc. In other words, every particle of hydrogen is like every other particle of hydrogen...." Dalton also saw matter as composed of combinations of atoms of a finite number of chemical elements. And, he reaffirmed the concept that transformation involves the rearrangement of atoms, but neither their creation or destruction.
Modern science has shown that even atoms can be subdivided into more fundamental units. In 1897, the English physicist J. J. Thomson (1856-1940) identified the electron, which carries a unit of negative electrical charge, as a constituent of atoms. Also in England early in this century, Ernest Rutherford (1871-1937) demonstrated that most of the atom is empty space and that nearly all its mass is concentrated in the nucleus. The principal constituent of the nucleus is the positively charged proton, which was identified by Rutherford about 1919. By 1932, a second particle having no electrical charge and called a neutron was also shown to reside in the nucleus by the British physicist James Chadwick (1891-1974).
The current picture of atoms, based on the work of Danish physicist Niels Bohr (1885-1962) and known as the Bohr atom, is that of a nucleus made of protons and neutrons, with electrons in orbits surrounding the nucleus. The chemical identity of each atom of an element is determined by the number of protons in its nucleus, which in turn establishes the element's atomic number. The simplest nucleus is that of hydrogen, with one proton and atomic number 1; the atomic number of helium is 2, that of lithium is 3, and so on. As indicated above, protons have a unit of positive electrical charge equal and opposite to the negative charge of electrons; the mutual attraction between positive and negative charges holds atoms together. Neutral atoms have as many electrons as protons (Figure 1). When there are fewer electrons than protons, the atom is known as a positive ion. Atoms bond together to form molecules, with the number of atoms varying from two (as in the oxygen molecules we breathe) to many millions (as in complex hydrocarbons that compose biological life). This conceptual model, with some additional refinements, accounts satisfactorily for the periodicities in the physical and chemical properties used to arrange the elements in what is called the periodic table (Appendix 4).
Atomic nuclei are made up of from 1 to about 260 protons and neutrons. The atom's mass is approximately that of the nucleus, since the mass of either protons or neutrons is almost 2000 times greater than that of electrons.
Although the nucleus of any particular element contains a fixed number of protons, the number of neutrons may vary from a few more to a few less than the number of protons. These nuclei with different numbers of neutrons and consequently different masses are called isotopes. All the elements in the periodic table, except for 20, possess two or more isotopes, so that their atomic weights depend on the relative abundance of each isotope in nature.
The search for the ultimate indivisible constituents of matter
that began before Democritus continues even today. Physicists
have discovered more then 100 subatomic particles. Evidence
exists that some of these subatomic particles, such as protons
and neutrons, are themselves composed of even more basic particles,
called quarks. Is there an ultimate end to the divisibility
of matter? On a conceptual level, most scientists would like
to think so, but the historical record of this line of investigation
seems to cast some doubt that such a limit may ever be found.
5.1.2. States of Matter
Matter exists in three states: In solids atoms are bound to permanent positions relative to each other; in liquids the particle bonds are weak and temporary; and by contrast, in gases there is no significant bonding between atomic particles and the particles have no permanent positions relative to each other. The particles of a gas can be molecules (which consist of two or more atoms), atoms themselves, or ions and electrons.
Most of the matter in the Universe is either in the form of gas
or a gas composed of free electrons and positive ions called a
plasma. Each atom in a plasma can be stripped of one electron--so
that there are two independent particles per atom--or stripped
of two electrons--three particles per atom--and so on. Only some
of the atoms in a plasma may lose an electron, or in the extreme
case, all the atoms may lose all their electrons. Such a high
degree of ionization is generally the result of very high temperatures.
5.1.3. Temperature of Gases and Plasmas
In the realm of atoms (as in a gas) gravity does not cause changes in motion, as it does in the macroscopic world. The atomic world is dominated by electromagnetic forces. As with the force of gravity, the intensity of electric and magnetic fields weakens as the inverse square of the distance from their source. At first glance, this might suggest that Newtonian mechanics ought to describe motion in the atomic domain, gravity as a cause simply being replaced by electromagnetic forces. Such is not the case in general, however, and the mechanics of the atom is called quantum mechanics. Its details go beyond our needs in this book, so we only point out that motion in the atomic world has a discrete nature rather than the continuous characteristics of our everyday experience. This radical departure in the fundamental cause-and-effect relationship from the macroscopic realm, in which we live, to the atomic realm is in the same conceptual vein as Einstein's concept of the photon. That is, the concepts underlying quantum mechanics and the photon are the same.
In gases, atomic particles dart about rapidly, colliding millions of times each second and changing their directions of motion just as frequently. Each gas particle has a kinetic energy proportional to the product of its mass and the square of its velocity. After a collision, the velocity can be either greater or smaller than it was before the collision; the kinetic energy of each particle changes in its repeated collisions. Collectively, however, the gas particles will have some average kinetic energy, that changes only when energy is added to the gas or removed from it. Another way of saying this is that the average kinetic energy changes when the gas is heated or cooled. Thus what we call heat or thermal energy is no more than the collective kinetic energies of atomic constituents.
Temperature is a measure of the average kinetic energy of gas particles. The motion of the particles composing a body, such as ice or water or water vapor, is called random thermal motion and is illustrated in Figure 2. It increases as the temperature goes up, and it decreases as the temperature goes down. Absolute zero is reached when the average kinetic energy is zero. Seen in terms of the motion of the particles in a gas, temperature is a measure of that motion: The greater the temperature of the gas, the greater the random thermal motion.
Temperatures in astronomy are usually measured on the absolute, or Kelvin (K), scale. In this system, there are 100 divisions (degrees) between the freezing point (273 K) and the boiling point (373 K) of water.
When we heat the air in a vessel, we are increasing the kinetic energy of each particle, which means a higher average kinetic energy. As a result, the gas particles move about faster, and they collide more frequently and more violently with their surroundings. If the density of the gas particles is quite large, then the hot gas can transport a great deal of thermal energy (or heat) from one place to another. The flow of energy in nature is from regions in which the energy content is high to those in which it is low; natural physical processes tend to even out the energy content.
Our understanding of the nature of matter, particularly gases
and plasmas, as to its chemical composition, temperature, density,
and motions, can be enhanced by examining the radiation either
emitted or absorbed by matter. In this and the remaining sections
we will consider how radiation conveys information about matter,
and what techniques are used by astronomers to collect and measure
radiation.
5.2.1. Spectroscopy
Until about a century ago, astronomers were concerned primarily with the positions and motions of celestial bodies. They knew little or nothing about the physical nature of these bodies and were really not able to find out. Today, however, concern with the physical nature of celestial bodies, which is the field of astrophysics, is one of the most actively pursued areas in astronomy. This change in orientation is primarily the result of two developments in of our understanding of light.
The first development was the invention of the spectroscope in the late 1850s by chemist Robert Bunsen (1811-1899) and physicist Gustav Kirchhoff (1824-1887). Earlier in 1704, Newton, in his Opticks, had described how one saw a rainbow of colors when sunlight was passed through a prism, the principal component of a spectroscope. After Newton, both William Wollaston (1766-1828) in England and Joseph Fraunhofer (1787-1826) in Bavaria used prisms to investigate the colors emitted by various chemical elements. Thus the spectroscope is essentially a device used to separate white light into its component colors (Figure 3).
The second development was the recognition that under certain circumstances each chemical element emits a specific set of colors that is peculiar to it, much like a person's fingerprints. As early as the 1830s this fact was suggested in connection with the presence, identity, and abundance of different elements in ores. The real beginnings of the field of spectroscopy occurred in the last half of the nineteenth century in the laboratories of Bunsen and Kirchhoff at the University of Heidelberg.
From his experiments, Kirchhoff was able to formulate three empirical laws of spectroscopic analysis. These laws describe the physical conditions under which matter will produce light having one of three different spectra of colors. One of the first astronomical applications of these laws was in trying to determine the chemical composition of the Sun and stars.
By 1864, the English astronomer Sir William Huggins (1824-1910) had identified nine elements in the bright star Aldebaran. Four years later, in 1868, Sir Norman Lockyer (1836-1920) detected an element in the solar spectrum that was unknown on the Earth at that time. It was later found in natural gas, but it still carries its solar name--helium. In the years following, more elements were identified in stars. In addition, it was also discovered that the spectra of colors in the white light coming from stars contain sufficient similarities that the stars can be arranged into spectral classes. There are, however, small but important differences in the spectra of colors from star to star.
By the beginning of this century, an important tie had been developed
between the researcher in the laboratory and the astronomer in
the observatory. From this collaboration came a new way of perceiving
nature, and this has fundamentally altered our conceptual view
of the Universe. The extension of spectrum analysis to radiation
in parts of the electromagnetic spectrum other than the visible
and the ability to move above the Earth's obscuring atmosphere
to view radiation coming from the depths of the Universe are the
Rosetta stone of today's astronomy.
5.2.2. Types of Spectra
Just as sound waves of various wavelengths are transported simultaneously through the same region of air, electromagnetic waves of different wavelengths can move through the same point in space and superimpose to form composite waves, or white light. In turn, white light can be dispersed, or separated, into its component colors, or wavelengths, to form a spectrum. The study of spectra is called spectroscopy. Let us briefly explain how we can accomplish the dispersion of composite light by using a triangular piece of glass called a prism.
In the refraction of a ray of light, the angle through which light is refracted depends on wavelength; the angle of refraction is greater for shorter wavelengths than it is for longer ones. Consider white light passing through the slit of a narrow diaphragm and then through a glass prism, as in Figure 3. The light separates into its component wavelengths, since short-wavelength light is refracted through larger angles than is long-wavelength light. Thus waves of different wavelengths go off in different directions. The result, when imaged on a screen, is a rainbow-colored sequence of images of the slit containing an image for each wavelength present in the white light.
Another means of dispersing white light to produce a spectrum
is the diffraction grating. Unlike the ordinary glass
prism, which is transparent only to visible and infrared radiation,
the grating is useful over a broad spectrum, from X-ray to infrared
wavelengths. In its simplest form, the diffraction grating is
a plate containing a very large number of very narrow parallel
slits uniformly spaced at distances that are only a few times
the wavelength of light. (By "large number," we mean
many thousands of slits per centimeter.) The spectrum is viewed
in the direction of the light source, as in Figure 3. Since
the amount of bending, or diffraction, of electromagnetic waves
at each slit depends on wavelength, composite light is separated
into its component colors. (For more details on these dispersing
devices as they are used in an instrument called a spectroscope,
see Section 4.)
5.2.3. Kirchhoff's Laws: The Nature of Light Sources
When we analyze white light from various astronomical sources for its color composition, we do not always find a continuous rainbow-colored sequence of wavelengths. Spectra can be classified and interpreted according to laws formulated by Kirchhoff more than a century ago. The three basic types of spectra--continuous, emission, and absorption--and the physical conditions under which they are formed are given by Kirchhoff's laws and are illustrated in Figure 4:
Kirchhoff's First Law (Continuous Spectrum): The spectrum of a radiating solid, liquid, or highly pressurized gas is an uninterrupted sequence of wavelengths known as a continuous spectrum.
Kirchhoff's Second Law (Emission Spectrum): The spectrum of a radiating rarefied gas is a set of discrete or isolated wavelengths whose appearance is a series of bright-colored lines that form a pattern characteristic of the chemical composition of the gas and is known as an emission or bright-line spectrum.
Kirchoff's Third Law (Absorption Spectrum): Light from a radiating source producing a continuous spectrum will, if it passes through a cooler gas, have certain specific wavelengths characteristic of the cooler gas removed from the spectrum. The spectrum appears continuous except where it is crossed by dark lines, which indicate that these wavelengths have been removed, and it is known as an absorption or dark-line spectrum.
Illustrations of the different types of spectra are given in Figure 4. There are many common examples of light sources whose spectra are one of the three basic types. For example, the glowing filament of an electric light bulb produces a continuous spectrum. A neon sign is an example of an emission spectrum. The spectrum of a gas composed of molecules is many sets of very closely spaced spectral lines known as emission bands. And, as a final example, the spectrum of the Sun and most stars is an absorption spectrum.
We will see additional examples of each type of spectrum many
times in the remaining chapters. The important point to remember
is that the type of spectrum for a light source tells us something
about the conditions in and around that source.
5.2.4. Identifying the Chemical Elements
An astronomical light source, such as a star or a gaseous nebula, contains a mixture of chemical species, each either emitting or absorbing its own set of wavelengths of electromagnetic radiation. By knowing what wavelengths are emitted by different chemical elements from laboratory studies, astronomers can identify individual elements in the light source from the measured wavelengths of its spectral lines, regardless of whether they are emission or absorption lines.
Identification is done in the following way: Light from a celestial body is collected by a telescope and then passed through a spectrograph in order to disperse the white light from the source and form its spectrum. The photographic plate on which the spectrum is recorded is called a spectrogram. As a standard against which unknown wavelengths in the astronomical spectrum can be measured, an emission spectrum of a known gas, such as neon or vaporized iron or titanium, is placed above and below the astronomical spectrum. (The mechanism for placing the laboratory spectrum on the astronomical spectrogram is a part of the telescope and spectrograph.) With these comparison lines of known wavelength the astronomer can determine the unknown wavelengths of the astronomical object's spectral lines. An example of a star's absorption spectrum is given in Figure 5. The absorption spectrum is gray with black absorption lines, and the comparison spectrum of neon shows white emission lines on a black background.
Kirchhoff's laws of spectrum analysis tell us about the general physical conditions of the light source. And if the spectrum of the light source contains absorption or emission lines, we can measure their wavelengths and identify the chemical elements that are present.
Can more detailed information about the light source be found? Suppose we want to know the temperature of the light source. Can this be done? Yes it can, for special types of light sources known as ideal radiators, or blackbodies. In Chapter 14 we shall outline the basis for obtaining more detailed information about stars, such as density and rotation.
5.3.1. Radiation Detectors
Before discussing the instruments used with optical telescopes to obtain quantitative measurements of radiation, let us consider briefly the most important component of these instruments: the radiation detector. Telescopes are capable of collecting light over a wide range of wavelengths, but it is the radiation detector that actually determines what the telescope "sees." One radiation detector with which we are all familiar is the human eye. Since the eye is so familiar to us, we can use it to illustrative those properties of radiation detectors which are of interest. Such properties are:
5.3.2. Properties of Radiation Detectors
The eye is sensitive to the narrow-wavelength region between about 3500 and 7000 A, as shown schematically in Figure 6. However, the eye does not respond equally to all colors in the visible spectrum. It is most sensitive to the middle of the wavelength region, the green region, and sensitivity drops to zero toward either the violet (short wavelengths) or the red (long wavelengths) regions.
The variation and range of detector response is the way in which the eye responds to one photon or to a flood of photons. Experience tells us that the eye does not respond in the same way to both. For the eye, a minimum number of photons is required, depending on their wavelength, to make it respond, as illustrated in Figure 6. In other words, there is a limit to how faint a light source we can see, and that visibility limit depends on whether we are looking at violet, green, or red light.
All of us have experienced the loss of response when the eye is exposed to a very bright light. In such cases, the eye saturates--it no longer responds--and no scene is visible, just an intense and painful brilliance. To be useful, the range between minimum and saturation of visibility should be quite large, say, a factor of 100 or 1000. If anywhere between the lower and upper limits of the eye's response, we double the number of photons from a light source, do we observe that the light is twice as bright? The answer in general is no. By and large, over the eye's range of response, doubling the stimulus does not double the response; in other words, we say that the response is nonlinear. This concept of linearity is important because in seeking the amount of radiant energy emitted by an astronomical source, astronomers usually compare the unknown light source to one of known energy output. Thus they have to know how their radiation detector responds to increasing or decreasing numbers of photons. Let us now switch and consider two other radiation detectors, the photographic emulsion and the photoelectric device.
Photographic emulsions record photons by undergoing a photochemical change that will ultimately deposit silver on a glass plate or acetate film. Various photographic emulsions can be manufactured such that they will respond to different wavelength regions within and beyond either end of the visible spectrum, which makes the photographic emulsion more versatile than the eye. Photographic emulsions, like the eye, are nonlinear in their response; they have a rather complicated response depending on position in their response range. A simulated wavelength response is shown in Figure 6 for a fictitious photographic emulsion sensitive to infrared photons.
Photographic emulsions possess a significant advantage over the eye in that they build up the image by storing their response over time. Thus time exposures allow the astronomer to collect information on a photographic plate about very faint light sources that cannot be seen with the eye through the same telescope. How faint a star can we photograph? The telescope's aperture sets the initial limit. Ultimately, however, the limit is set by a weak illumination coming from the whole night sky. This background radiation comes from two sources: starlight scattered by the Earth's atmosphere and the airglow emitted by the atmosphere itself. The inherent disadvantage of photographic emulsions is that its photon-capturing efficiency is low; it can record only 1 or 2 percent of the incident photons (those which activate the light-sensitive coating). Facing this inefficiency, astronomers have found other types of radiation detectors to improve the performance of telescopes.
The photoelectric device is an application of the photoelectric effect. The basic principle is to use photons to eject electrons from a metal surface by exposing it to a beam of light and then to measure the number of electrons liberated with electronic circuitry. A photoelectric device, like photographic emulsions, can be made to respond to different wavelength regions by using different metals for the surface of a device. The biggest advantage of the photoelectric device is that it has a very large range in its response. In addition, its response is linear to the number of incident photons, as shown in Figure 6 for a fictitious device. With modern electronics it is possible to adapt the photoelectric device to count individual photons or to use a mosaic of devices to form a picture much as a photographic plate does.
As an illustration of the importance of photoelectric devices as radiation detectors, only about 14 percent of the nights of observing on the 5-m Hale telescope are devoted to photographic work. On 85 percent of the nights, some kind of photoelectric detecting device is being used.
The photographic plate and the photoelectric device enhance our ability to detect light from astronomical sources of different brightnesses, but they are not basically analyzing instruments. To analyze light, astronomers must equip an analyzing instrument with either of these detectors and attach it to a telescope. The two basic types of analyzing instrument are the spectrograph and the photometer.
The spectrograph disperses composite light from the source into its component wavelengths so that we can, for example, determine the elements that compose the light source. Spectroscopy, which is the study of the spectra of light sources, is astronomy's fundamental interpretive tool.
A prism or grating spectrograph (Figure 7) receives concentrated light from the telescope's objective on a narrow rectangular entrance slit. The light diverging past the slit enters a collimator, whose purpose is to deliver a beam of parallel rays to the dispersing device. After these rays of composite light either pass through a prism or reflect off a grating dispersing device, they will have been separated into their constituent wavelengths. The dispersed light is focused finally by a camera system on a radiation detector (a photographic plate or a photoelectric device) as individual color images of the entrance slit. Each wavelength forms a distinct image of the slit. The different color images of the slit are arrayed in an orderly progression of colors from red to violet to create a spectrum of the composite light.
Whereas the spectrograph is used to examine the spectral composition of radiation, the photometer is an analyzing instrument used to measure the amount of radiation coming from an astronomical object (Figure 8). It measures the amount of radiant energy on either a relative or an absolute scale at one wavelength or in a band of wavelengths. The radiation detector is generally today a photoelectric device, and thus the photoelectric photometer is much like an exposure meter on a camera: Incident light is converted proportionally into an electric current. One can use a variety of techniques to define the wavelength region for the photometer, such as color filters. Another way is to use the photometer in conjunction with a spectrograph, where the photometer is made to scan the spectrum formed by the spectrograph.
The photometer is usually limited to measuring only one light source, such as a star, at a time. But the limitation is compensated for by the photometer's very great accuracy. Because of its quick response to changes in amounts of light, a photoelectric photometer is particularly useful in continually monitoring the change in brightness of an object whose emission of radiant energy varies over time (for example, a number of stars are known to be variable light sources).
5.3.7. Viewing Problems Through the Atmosphere
The Earth receives electromagnetic radiation of all wavelengths from various directions in space, but most of the electromagnetic spectrum is screened out by the atmosphere well above the Earth's surface (Figure 9). Wavelengths from only two regions of the electromagnetic spectrum, however, are able to penetrate the atmosphere to any extent. These two spectral windows in the atmosphere through which astronomers observe the Universe are called the optical window--from about 3000 to 10,000 A, or roughly the visible-wavelength region--and the radio window--which includes the wavelength region from about 1 mm to 30 m. The telescopes astronomers build to take advantage of these two atmospheric windows are thus logically called optical telescopes and radio telescopes. We discussed optical telescopes in the last chapter, and we will consider radio telescopes in the next section.
With the advent of the space age, astronomers have been able to use aircraft, balloons, rockets, and now primarily satellites to extend their vision of the Universe by going above part or all of the Earth's veiling atmosphere. Astronomers are aghast at what the space-platformed telescope has revealed through radiation in the ultraviolet, X-ray, gamma-ray, and infrared regions. These results will be discussed in later chapters.
The theoretical resolving power of any optical telescope is never fully realized because the lower layers of our atmosphere are unsteady or turbulent, and turbulence blurs and distorts a star's image and causes it to twinkle, or scintillate. The rapid scintillations break the starlight into many dancing specks of light, which in long exposures merge to form the fuzzy stellar images we see in photographs. Even under the best conditions, optical images are no sharper than about 1 second of arc--the angle subtended by a dime at a distance of 1 mile--and are typically several seconds across. The less the atmospheric turbulence, the less the stars twinkle, and the better is the seeing, as astronomers refer to it. A planet, however, shines with a steady light because each point on the tiny disk twinkles out of step with neighboring points; thus we see an average of all the twinkling points.
A technique called speckle photography, can be used with large telescopes to get around the image smearing that comes from atmospheric turbulence. Several extremely short time exposures (less than 0.01 s) can be made, and in each the star image on the photographic plate appears as a cluster of sharp specks of different brightness. The information from each photograph can then be fed into a computer that reassembles the several photographic images into a single unsmeared image of the star. Figure 10 shows the reconstructed image of Betelgeuse (Alpha Orionis), demonstrating what speckle photography can do in resolving minute sources of light, such as the disk of a large star.
Other nuisances hamper our observation of the heavens. The night sky's transparency varies as smog, dust, and atmospheric haze cloud it. The upper atmosphere is also suffused with a faint light called airglow. It arises when atmospheric molecules absorb ultraviolet photons from sunlight and then reradiate the energy in a few wavelengths of the green, red, and infrared spectral regions. On long exposures, airglow fogs a photograph and reduces the contrast between the faintest images and the sky background.
Another problem is that starlight penetrating the atmosphere is bent increasingly toward the vertical so that a star appears to be slightly closer to the zenith (the point directly above the observer) than it really is (Figure 11). This atmospheric refraction effect is greatest near the horizon (about 0.5o), for there the light's path through the air is the longest. The consequences of this are that when we observe the rising or setting Sun, it is really below our horizon, since refraction has raised the Sun's image above the horizon by a 0.5o, which is the Sun's own angular diameter.
In 1800, William Herschel detected the infrared component of solar radiation by positioning thermometers beyond the red end of the Sun's visible spectrum. His discovery foreshadowed the astronomy of those portions of the electromagnetic spectrum to which the human eye is not sensitive. What we have seen over the last 15 years in those spectral regions whose wavelengths are longer than the visible spectrum--infrared, microwave, and radio--has revolutionized our concept of the Universe. In this section, we want to discuss the detection of radiation in those long-wavelength regions.
A large portion of the infrared spectrum does not reach ground level because of absorption by water vapor, carbon dioxide, and molecular oxygen that lie between the ground and about 15 km of altitude. Consequently, airplanes, balloons, rockets, and satellites have been used to lift infrared telescopes above the veiling atmosphere. However, astronomers can also locate infrared observing facilities on mountaintops, such as the one in the Hawaiian islands shown in Figure 12.
The modern infrared telescope is designed and functions much like a reflecting telescope used for visual observations with one major difference. Because bodies that are warmer than their surrounding, that is contain more thermal energy than the surroundings, will attempt to equalize the energy content with those surroundings, they will radiate electromagnetic energy. If these bodies are only slightly warmer than their surroundings, say a few tens or hundreds of degrees Kelvin, they will radiate in the infrared portion of the spectrum and not the visible. Thus the major difference between a visual and an infrared reflector is that care must be given to the design so that the infrared reflector does not collect infrared radiation from warm objects nearby. Thus infrared reflectors have no tube around the telescope, have smaller secondary mirror supports, have been shield from hot electronic equipment, etc. The infrared detector, which generates an electrical signal when heated by infrared photons, work best when cooled to temperatures near absolute zero, such as with liquid helium. With a liquid-helium-cooled infrared detector on the appropriate analyzing instruments, the modern infrared telescope is a powerful research tool.
A number of new telescopes have been designed and built for infrared astronomy only. A national observatory for infrared astronomy is located high on the 4200-m inactive Hawaiian volcano Mauna Kea (Figure 12). A 3.0-m infrared telescope constructed by NASA and the University of Hawaii is in operation there along with a 3.8-m infrared telescope belonging to the United Kingdom. Other major infrared telescope facilities are the Multiple-Mirror Telescope in Arizona, the University of Wyoming facility, and Mexico's 2.1-m reflector.
5.4.2. Discovery of Celestial Radio Waves
In 1931, a Bell Telephone engineer, Karl Jansky (1905-1950), was trying to find where the interference disrupting transatlantic radiophone circuits came from. He discovered that some of the radio noise was not from the Earth--it was extraterrestrial. The primary source was the center of the Milky Way, in the constellation of Sagittarius. In 1936, an Illinois radio engineer, Grote Reber (b. 1911), pursued the phenomenon farther. He built the first parabolic radio telescope (9.5-m in diameter) and made the first radio map of the sky. The strongest signals he found came from the star clouds in Sagittarius and from several discrete sources toward the center of our Galaxy. The next major discovery was in 1942, by British radar operators and scientists tracking down suspected radar jamming during World War II; they discovered that the interference was radio emission from the Sun.
At first astronomers did not grasp just how significant Jansky's work was; they were preoccupied with their observations of the Universe through the optical window of the Earth's atmosphere. After World War II, however, radio astronomy came into its own when physicists, radio engineers, and astronomers joined forces to build larger and more efficient radio telescopes. Radio astronomy since then has led to startling discoveries, such as interstellar molecules, pulsars, and the enigmatic quasars. Today our concept of any celestial body is based on its appearance all across the electromagnetic spectrum, and the radio region is an extremely important component.
Because the physical nature of a radio wave is exactly the same as that of a light wave, except for the longer wavelength, the problem of designing a radio telescope is similar in theory to that of designing an optical telescope. In practice, however, there are some differences. Radio waves pass through most materials without any interaction; thus it is not possible to design a "lens" for radio waves that will focus them in a refracting telescope. But any metal will reflect radio waves, so a dish-shaped metal mirror will focus radio waves, just as a glass mirror focuses light waves. The reflecting surface of the dish can be an open, fine-wire mesh or a solid metal with a parabolic shape. Radio waves are reflected from the surface and converge toward a focal point, where a small collector aerial absorbs the concentrated energy, as shown in Figure 13, turning it into an electric current. From there the current or signal is carried by an electrical cable to the receiving equipment, which processes the signal just as in your home radio receiver.
After amplification, the signal variations are recorded in one of several ways (Figure 13). With the final step being to fed the recorded signal into a computer for analysis. When the computer has done its job, a formerly invisible part of the Universe is revealed, as shown by the radio map in Figure 14.
The sensitivity of radio telescopes can be increased, resulting in a greater pointing accuracy and higher resolving power, by enlarging the collecting area of the dish or by improving the capabilities of the receiver. With the largest radio telescopes we can obtain a resolution approaching a few seconds of arc, comparable with that of large optical telescopes. The most powerful radio telescopes can detect energy from sources whose power is comparable with that of a terrestrial FM broadcast station ten thousand billion km (1018 cm) away.
The radio telescope is remotely controlled by the astronomer from an electronic console, just as is the large optical telescope. Moderate-sized radio dishes, up to about 100 m or so in diameter, are steerable and have equatorial mountings that follow the rotating celestial sphere just as optical telescopes do. Larger and consequently heavier dishes, however, use an altazimuth mounting. This minimizes the distortion in the shape of the dish due to changing the orientation of the dish in the Earth's gravitational field.
Even larger and more unwieldy antennas are fixed, pointing upward always, while the rotating Earth sweeps much of the sky by the antenna's field of view. Arecibo, Puerto Rico has the biggest fixed antenna, a metal dish 305 m across contoured out of a natural bowl in the ground (Figure 15). It can survey the sky to within 20o of the zenith, allowing coverage of about 40 percent of the entire sky.
Astronomers have searched endlessly for better resolving power. It can be achieved by building bigger telescopes or by observing at shorter wavelengths for a chosen aperture or by using the phenomenon of interference discussed in Chapter 4. Interferometry is a technique involving two or more radio telescopes (optical telescopes can be used to make a visual interferometer also). Radio radiation from an astronomical source received at the individual telescopes is combined to obtain data that have a spatial resolving power equal to that of a single telescope as large as the distance between the individual receivers. With interferometry an astronomer can obtain details about the spatial structure of a given celestial object that a single radio telescope could never reveal.
The separation between the individual telescopes of the interferometer is limited only by the ability to correlate the results from each radio telescope, because the technique depends on combining, at the same instant, the signals received by the separate telescopes. With the advent of the atomic clock (a clock governed by the vibrations of certain atoms), it became possible to record the signals received by the different telescopes, along with the precise time, and to compare them later. This allowed the individual telescopes to be greatly separated, even on opposite sides of the Earth. The technique is called very long baseline interferometry (VLBI).
In the Very Large Array (VLA) radio interferometer recently put into operation in New Mexico (Figure 16) signals from each of 27 individual radio telescopes are combined by a computer. Each dish is 25 m in diameter, and the 27 individual telescopes are moved along railroad tracks arranged in the shape of an enormous 21-km Y. Nine dishes can be located on each branch of the Y, and the system can provide a total of 351 interferometer pairs of antennas. The energy-collecting power of the VLA is roughly equivalent to a single 122-m telescope, and it has a spatial resolution smaller than 1 second of arc or better than that of the 5-m Hale optical telescope.
A proposal has been made that the VLA become the heart of an array of ten 25-m radio telescopes widely spread from Hawaii across the United States to Puerto Rico. Known as the Very Long Baseline Array (VLBA), it would operate as a single instrument, with each antenna directly controlled from the main operation center in Socorro, New Mexico. This would give the VLBA the ability to ferret out the structure of radio sources as small as 0.0003 second of arc in diameter. Although it would operate as primarily a ground-based facility, a proposal has also been made to complement VLBA by putting a single large antenna in Earth orbit to work with it, making a truly immense radio interferometer.