What are Photoelectric devices

PHOTOELECTRIC DEVICES 
The photoelectric device is an application of the pho­toelectric effect. The basic principle is to liberate elec­trons from a metal surface by exposing it to photons in a light beam and then to measure the number of electrons with electronic circuitry. The photoelectric device, like the photographic emulsion, can·be made to respond to different wavelength regions by varying the metals used in making the surface of a device. The biggest advantage of the photoelectric device is that it can be manufactured to have a very large dynamic range of response; in addition its response is linear to the number of incident photons 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 photoelectric device's im­portance as a radiation detector, only about 15 per­cent of the nights of observing on the 5.1-meter Hale telescope are devoted to photographic work. On 85 percent of the nights some kind of photoelectric de­tecting device is being used.

What is Phootographiic emulsion?

PHOTOGRAPHIC EMULSION 
The photographic emulsion records photons by un­dergoing a chemical change (a photochemical effect) that will ultimately deposit silver on a glass plate or acetate film. The photographic plate can be made to respond to different wavelength regions within and beyond either end of the visible spectrum, which makes it much more versatile than the eye. Also, its response over a wavelength interval can be made much more uniform than that of the eye. The photo­graphic plate, like the eye, is nonlinear in its re­sponse; it has a rather complicated response de­pending upon the position in its dynamic range.
The photographic plate has a strong advantage over the eye since it will build up a response by storing the image. Thus time exposures allow the astronomer to collect information on a photographic plate about very faint light sources that cannot be detected by 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 the weak illumination in the night sky. This background interference comes from starlight scattered by the earth's atmosphere and from diffuse radiation in the atmosphere (airglow). Unfortunately, the photo­graphic plate's photon-capturing efficiency is low. The emulsion can record only 1 or 2 percent of the inci­dent photons (those that activate the light-sensitive coating). Facing this inefficiency, astronomers have found other types of radiation detectors to improve the telescope's performance.

What are radiation detectors and what are their properties?

Radiation Detectors and Their Properties
RADIATION DETECTORS
Before discussing the accessory instruments used with optical telescopes, let us consider briefly the most important component of these instruments: the radiation detector. The telescope is capable of col­lecting light over a very wide range of wavelengths, but it is the radiation detector that determines what the telescope sees. One radiation detector with which we are all familiar is the human eye. It possesses most of the properties of radiation detectors in general and is thus illustrative of the points we wish to make about them.
The properties of interest are the wavelength re­gions to which the detector is sensitive, the differing response of the detector over that wavelength region, and the natu re and range of detector response. Using the human eye, we can briefly illus­trate each of these properties.
PROPERTIES OF RADIATION DETECTORS 
The eye is sensitive to the narrow wavelength region between about 3500 and 7000 angstroms. 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 wavelengths, and the sensi­tivity drops to zero toward either the violet (short wavelengths) or the red (long wavelengths).
The nature and the range of detector response are expressed by the ways in which the eye responds to one photon and to a tremendous flood of photons. Common experience tells us that the eye does not respond in the same way for both. There is some threshold number of photons, depending upon their wavelength, necessary to make the eye respond. In other words, there is a limit to how faint a light source we can see, and that vis­ibility limit depends upon whether we are looking at violet, green, or red light.
All of us have experienced the loss of response of the eye when we try to look at too bright a light. That is, the eye saturates-it no longer responds-and no scene is visible to us, just an intense and painful bril­liance. To be useful, the radiation detector's dynamic range between threshold and saturation of visibility should be quite large, say, a factor of 100 or 1000. Now we may ask, "What is the response of the eye to doubling the number of photons in between the low­er and upper limits of threshold and saturation?" If we double the number of photons, do we observe that the light is twice as bright? The answer in general is no. By and large, over the dynamic range of response of the eye, 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 against one of known energy output. Thus they have to know how their radiation detector responds to increasing or de­creasing numbers of photons.
Now we look at two other radiation detectors, the photographic emulsion and the photoelectric device.

Necessary Instruments for Telescopes

Accessary Instruments for Telescopes
RADIATION DETECTORS
Before discussing the accessory instruments used with optical telescopes, let us consider briefly the most important component of these instruments: the radiation detector. The telescope is capable of col­lecting light over a very wide range of wavelengths, but it is the radiation detector that determines what the telescope sees. One radiation detector with which we are all familiar is the human eye. It possesses most of the properties of radiation detectors in general and is thus illustrative of the points we wish to make about them.
The properties of interest are the wavelength re­gions to which the detector is sensitive, the differing response of the detector over that wavelength region, and the natu re and range of detector response. Using the human eye, we can briefly illus­trate each of these properties.
PROPERTIES OF RADIATION DETECTORS
The eye is sensitive to the narrow wavelength region between about 3500 and 7000 angstroms. 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 wavelengths, and the sensi­tivity drops to zero toward either the violet (short wavelengths) or the red (long wavelengths).
The nature and the range of detector response are expressed by the ways in which the eye responds to one photon and to a tremendous flood of photons. Common experience tells us that the eye does not respond in the same way for both. There is some threshold number of photons, depending upon their wavelength, necessary to make the eye respond. In other words, there is a limit to how faint a light source we can see, and that vis­ibility limit depends upon whether we are looking at violet, green, or red light.
All of us have experienced the loss of response of the eye when we try to look at too bright a light. That is, the eye saturates-it no longer responds-and no scene is visible to us, just an intense and painful bril­liance. To be useful, the radiation detector's dynamic range between threshold and saturation of visibility should be quite large, say, a factor of 100 or 1000. Now we may ask, "What is the response of the eye to doubling the number of photons in between the low­er and upper limits of threshold and saturation?" If we double the number of photons, do we observe that the light is twice as bright? The answer in general is no. By and large, over the dynamic range of response of the eye, 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 against one of known energy output. Thus they have to know how their radiation detector responds to increasing or de­creasing numbers of photons.
Now we look at two other radiation detectors, the photographic emulsion and the photoelectric device.
PHOTOGRAPHIC EMULSION
The photographic emulsion records photons by un­dergoing a chemical change (a photochemical effect) that will ultimately deposit silver on a glass plate or acetate film. The photographic plate can be made to respond to different wavelength regions within and beyond either end of the visible spectrum, which makes it much more versatile than the eye. Also, its response over a wavelength interval can be made much more uniform than that of the eye. The photo­graphic plate, like the eye, is nonlinear in its re­sponse; it has a rather complicated response de­pending upon the position in its dynamic range.
The photographic plate has a strong advantage over the eye since it will build up a response by storing the image. Thus time exposures allow the astronomer to collect information on a photographic plate about very faint light sources that cannot be detected by 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 the weak illumination in the night sky. This background interference comes from starlight scattered by the earth's atmosphere and from diffuse radiation in the atmosphere (airglow). Unfortunately, the photo­graphic plate's photon-capturing efficiency is low. The emulsion can record only 1 or 2 percent of the inci­dent photons (those that activate the light-sensitive coating). Facing this inefficiency, astronomers have found other types of radiation detectors to improve the telescope's performance.
PHOTOELECTRIC DEVICES
The photoelectric device is an application of the pho­toelectric effect. The basic principle is to liberate elec­trons from a metal surface by exposing it to photons in a light beam and then to measure the number of electrons with electronic circuitry. The photoelectric device, like the photographic emulsion, can·be made to respond to different wavelength regions by varying the metals used in making the surface of a device. The biggest advantage of the photoelectric device is that it can be manufactured to have a very large dynamic range of response; in addition its response is linear to the number of incident photons 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 photoelectric device's im­portance as a radiation detector, only about 15 per­cent of the nights of observing on the 5.1-meter Hale telescope are devoted to photographic work. On 85 percent of the nights some kind of photoelectric de­tecting device is being used.
IMAGE INTENSIFIERS
Electronic image intensifiers do as their name implies, they intensify, or amplify, the light from weak sources of radiation. I n one such system photons from the telescope are focused onto a photocathode surface, which ejects electrons. The electrons are increased in number, accelerated, and focused by means of electric and magnetic fields onto a phosphorescent screen, which emits a spark of light for each electron that strikes it. Thus the faint light from the astrono­mical source is amplified by the device into light sufficient to record the image on a photographic plate. Alternatively a computer circuit can be used to count the electrons during the exposure. Still other image-intensifying techniques are in use or in devel­opmental stages; these techniques can reduce ex­posure times by factors of 50 to 100 over those for photographic systems.
SPECTROGRAPHS
The photographic plate and the photoelectric device enhance our ability to detect light from different as­tronomical sources, but they are not basically analyz­ing instruments. We can equip an accessory instrument with either of these detectors and attach it to the telescope to analyze light. The two basic types of ana­lyzing instrument are the spectrograph and the photo­meter.
The. spectrograph disperses the 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 receives the concentrated light coming from the telescope's objective on an entrance slit. The light diverging past the slit enters a collimator, which deliv­ers a beam of parallel rays to the dispersing device. Then these rays pass through either a prism or reflect off a grating, which separates the light into its constit­uent wavelengths. The dispersed light is focused by a camera system onto 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 images of the slit in the different wavelengths are arrayed in an orderly progression of colors from red to violet to create the observed spectrum of the composite light falling on the entrance slit.
PHOTOMETERS
The photometer is an accessory device that the astron­omer attaches to the telescope at the focal position of the objective to measure the amount of radiation com­ing from the astronomical object. Where the spec­trograph is used to examine the spectral composition of radiation, the photometer can be made to scan the spectrum formed by the spectrograph. 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 photometer is much like an ex­posure meter on a camera: Incident light is converted into an electrical current. One can use a variety of techniques to define the wavelength region for the photometer, such as a spectrograph or color filters. And the radiation detector is generally today a photo­electric device.
The photoelectric 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 photoelectric photometer's very great accuracy. Be­cause of its quick response to changes in amounts of light, the photoelectric photometer is particularly use­ful in continually monitoring the change in brightness of an object whose emission of radiant energy varies with time (for example, a number of stars are known to be variable light sources).
Infrared Devices
In 1800 William Herschel detected the infrared com­ponent of solar radiation by positioning thermome­ters beyond the red end of the sun's visible spectrum and thus foreshadowed the astronomy of invisible spectral regions. What we have seen over the last 15 years in these regions has revolutionized our concept of the universe.
INFRARED TELESCOPES
We can subdivide the infrared spectrum into three segments. A large part of the infrared spectrum is not visible at ground level because of absorption by water vapor, carbon dioxide, and molecular oxygen, which lie between the ground and about 15 kilometers alti­tude. Consequently, airplanes, balloons, rockets, and satellites are extensively used to lift the infrared tele­scope above the veiling atmosphere. Astronomers can also locate infrared facilities on mountaintops, such as
the one in the Hawaiian islands to make ground-based infrared observations.
The liquid-helium-cooled infrared detector can be used with the appropriate analyzing instruments on an ordinary optical telescope to study the cosmos. But because of the longer wavelength of infrared radi­ation, the image-producing quality of the telescope objective need not be so fine as it must be for the visible region. Thus a number of new telescopes have been designed and built for infrared astronomy only. A national observatory for infrared astronomy is built high on the 4200-meter inactive Hawaiian volcano Mauna Kea. A 3.0-meter infrared tele­scope constructed by NASA and the University of Hawaii is in operation there along with a 3.B-meter infrared telescope belonging to the United Kingdom. Other major infrared telescope facilities are the MMT in Arizona, the University of Wyoming facility, and Mexico's 2.1-meter reflector.
Now that we have described the optical window, somewhat expanded to include available parts of the infrared, we should shift our focus to the radio spec­tral window in the atmosphere. Thus the next section is on radio telescopes .

Optical Telescopes - A Complete Overview

FORMATION OF AN IMAGE
In optical astronomy the object with which we work is the image of the light source formed by the principal image-forming part of the telescope, which is called an objective. The objective of the optical telescope is either a lens or a mirror. Light rays from the light source are refracted in passing through a lens and are reflected from a mirror. The image is produced where the light rays converge to a position known as the focus. The focal length of the objective is the distance behind the lens to the focus or the distance in front of the mirror to the focus. The image of a star is just a point of light, while that of an extended object, such as the moon, is inverted.
In telescopes using either mirrors or lenses an eye­piece magnifies the image much as a reading glass magnifies small print. Or a photographic plate may be inserted into the focal plane of the objective instead of the eyepiece, transforming the telescope into a giant camera. In this case the objective lens or mirror serves as the camera lens. The advantage of photography over observing with the eye is that the photograph is available for later study and time exposures can record fainter sources than those the eye sees.
PROPERTIES OF AN IMAGE
The image formed by either a lens or a mirror has certain properties that depend upon the diameter, or aperture, of the objective and its focal length. One property is the size of the image. Si nee the image of a star is a point, size is not an important property for it. For an extended object the image size depends upon the angular size of the light source on the sky and upon the weal length of the objective.
Image brightness is important since it determines whether the image is above the threshold of visibility or how long it would take to photograph. The bright­ness of an image of a point source, such as a star, depends on how much light is intercepted by the ob­jective. Hence its brightness is proportional to the area of the objective or to the square of the aperture. Doubling the aperture but leaving the focal length the same increases the area of the objective or its light­gathering power by four times, concentrating four times as much light into the same-size image.
When we photograph an extended object, the sur­face brightness of the image depends on the amount of radiant energy per unit area of the image. The objective's area (or the square of its aperture) still determines the total amount of energy collected, but the total energy is distributed over the entire image. Thus the larger the image's area, the smaller the en­ergy per unit of area. The image size of an extended object increases in proportion to the focal length; so for a given telescope aperture the surface brightness of the image decreases as the focal length is made longer.
How well a telescope discriminates between two adjacent objects or shows fine details is called its re­solving power. Because of the wave natu re of light, the image of a point source produces a diffraction pattern; it appears as a bright central spot, called a diffraction disk, surrounded by progressively fainter rings. When the diffraction patterns of two stars that are close together no longer overlap, we can see sep­arate stellar images. The larger the telescope's aperture, the smaller is the diffraction disk of each image. A large aperture therefore im­proves the resolution of closely adjoining features by making the diffraction effect of adjacent objects over­lap less. We define resolving power as the smallest angle between two close objects whose images can just be separated by a telescope. This critical angle is directly proportional to the wavelength of the ob­served radiation and inversely proportional to the ap­erture of the objective.
VIEWING PROBLEMS
The theoretical resolving power of any optical tele­scope is never fully realized. The lower layers in the earth's atmosphere are unsteady and turbulent; this turbulence blurs and distorts the star's image and makes it twinkle, or scintillate. The rapid scintillations break the starlight into many dancing specks of light, which in long exposu res merge to form the fuzzy stel­lar images we see in photographs. When the atmo­spheric turbulence is low, the stars twinkle, or scintil­late, less, and the so-called seeing is improved. A planet, on the other hand, shines with a steady light because each point on the tiny disk twinkles out of step with neighboring points; we see an average of all the twinkling points.
A technique called speckle photography, which can be used with large telescopes, can get around the smearing and wiggling of the image that comes from atmospheric turbulence. In the exposure of the pho­tographic plate for an extremely short time (less than 0.01 second), each star image appears as a cluster of sharp specks of different brightness. Then the photo­graph is run through a high-speed light-sensing de­vice which measures the variations in brightness across each speck. The information from the assem­blage of specks in each of many photographed images is fed into a computer that is programmed to analyze and reassemble the information into the unsmeared image of the 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. Atmospheric atoms and molecules absorb the ultraviolet photons in sunlight and 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 entering the at­mosphere is bent increasi ngly toward the vertical so that we see a star slightly closer to the zenith (the point directly above the observer) than it really is. This atmospheric-refraction effect is great­est near the horizon (about OS), for there the light's path through the air is the longest. When we observe the rising or setting sun, it is really below our horizon,
but refraction raises the sun's image above the hori­zon by an amount equal to its apparent diameter, which is OS.
Other viewing problems are related to the geo­graphical location of the observatory. An ideal site for an optical observatory is a mountaintop where the air is dry, transparent, and steady, and the sky is dark. An observatory also needs a minimum amount of wind and relatively easy access. The southwestern part of the United States satisfies most of these conditions and has many clear days and nights. Kitt Peak National Observatory is located there, about 65 miles south­west of Tucson, Arizona.
REFLECTING AND REFRACTING TELESCOPES 
Telescopes that use lenses for the objective are known as refracting telescopes, while those that employ a mirror are called reflecting telescopes. The objectives of the early refracting telescopes could not form sharp images because of a condition known as spherical ab­erration; single lenses also failed to bring all colors to a common focus, a failure called chromatic aberration. A compound lens, or two lenses of different types of glass cemented together was invented to minimize these aberrations in refracting telescopes.
Spherical aberration also occurs in a reflecting tele­scope. If the surface of the mirror is parabolic rather than spherical, that aberration is eliminated although some minor deficiencies still remain.
Why are the big modern telescopes of the re­flecting type? Reflecting telescopes have many advantages over refractors: The reflecti ng telescope is free from chromatic aberration, making it ideal for all­purpose photography and spectroscopy. Also, since a lens must be supported by its edges, there is a mate­rial limit to how large a lens system can be. But a mirror can be held both at its edges and from the back, the supports allowing a wide range of sizes for mirror systems. The largest refractor has an aperture slightly larger than 1 meter, butthe largest reflector is 6 meters in diameter.
There are other advantages to reflectors: The glass for the mirror in a reflecting telescope need not be so optically pure and homogeneous as that required for a large lens because the light reflects off the front surface and does not pass through the mirror, as it does through a lens. And the mirror has only one surface that must be painstakingly ground-the ach­romatic lens has four. To counter changes in tem­perature that would affect the focal length of the reflector, large mirrors are constructed of fused quartz or of a zero-expansion pyroceramic material. The mirror's surface is coated with a thin layer of highly reflecting aluminum that is replaced many times during the life of the telescope.
FOCAL POSITION FOR REFLECTING TELESCOPES
Reflecting telescopes can be designed for many kinds of astronomical work through choice of the focal ar­rangement to suit the type of obser­vation. For photography, photometry, and spec­troscopy of faint objects the prime focus is best because its small focal length lessens the exposu re time required. The Newtonian focus, most useful for small telescopes, is now little used by professional astronomers. In both these arrangements the ob­server works at a considerable distance above the ob­servatory floor since both focal positions are near the entrance of the telescope.
I n the Cassegrain focal arrangement a convex sec­ondary mirror positioned at the top in front of the focus slows the rate at which light rays converge, ef­fectively increasing the telescope's focal length. The secondary mirror reflects the converging rays to the bottom of the telescope and through a hole in the objective mirror to focus behind the objective. This is a much more convenient observing position since it is near the floor and behind the telescope. Of all the observations made with the 5.1-meter Hale telescope on Palomar Mountain 75 percent are from the Cas­segrain focus.
We might think that putting the secondary mirror and its supports or the observer's cage for the prime focus into the path of the light rays would obscure part of the image; but the only effect is to cut down the amount of light reaching the objective; the loss is small, and the quality of the image is not affected.
Equipment that is too heavy and bulky to be at­tached to the back of the primary mirror or is sensitive to changing gravitation as the telescope moves can be placed in a room below the observatory floor. An aux­iliary flat mirror diverts the long converging beam down the hollow polar axis around which the tele­scope rotates, and with this coude focal arrangement the focus can remain stationary no matter which way the telescope points.
TELESCOPE MOUNTINGS
An optical telescope, in order to follow an object as the earth's rotation carries it across the sky, must be free to move. To track stars accu rately and to permit a telescope to be conveniently pointed in any direction, the equatorial mounting system is used for most tele­scopes. This system has two axes of rota­tion: The telescope can rotate in an east-west direc­tion, called hour-angle, around its polar axis, which is aligned with the earth's axis of rotation; another al­lows the telescope to swing in a north-south direction about the declination axis, which is perpendicular to the polar axis.
Large telescopes are usually positioned by a com­puter from an operating console and guided to the exact location with hand controls. Once a large tele­scope is properly set, the computer operating a clock drive slowly turns it westward around its polar axis at the same rate as the earth turns eastward, keeping the stellar images locked in position in the field of view. The great simplicity in equatorial mounting is that tracking requires continuous motion about only one of its two axes. The disadvantage, which obtains in the largest telescopes now in operation and planned for the future, lies in the stresses on the polar axis due to gravity. The polar axis is inclined in the earth's gravitational field and must rotate on one edge of its end. For a very large telescope that is a difficult engineering problem.
One means of removing some of the stress from the primary axis is to align it with gravity. Such a mounting is known as altazimuth mounting; with it a telescope rotates about a vertical axis and about a horizontal axis. This mounting's disadvantage is that, to track a star, it must turn continuously about both axes at the same time. When the telescope approach­es the area of the sky directly overhead, continuous tracking becomes virtually impossible. Even with this disadvantage the altazimuth mounting will be the pri­mary mounting for very large telescopes to be con­structed in the future.
OTHER APPROACHES TO MAKING TELESCOPES
The principal problems in building very large tele­scopes on the earth's surface today are costs and con­struction time. A new 5-meter Hale telescope would now cost about 25 million dollars and take 10 years to build, while a 10-meter telescope would cost 200 mil­lion dollars and take 20 years to build, and a 25-meter telescope would require about 3 billion dollars and 50 years to construct. Clearly some dramatic changes in design are needed to lower cost and construction time.
A new telescope design, called the Multiple Mirror Telescope (MMT), which is well suited for infrared observations, has been installed on Mount Hopkins in Arizona. It uses a mosaic of inde­pendent mirrors of small size to collect and focus light in order to simulate the collecting ability of a large­aperture single mirror. The MMT consists of a circular array of six identical 1.8-meter mirrors on an al­tazimuth mounting; the array has light-gathering power equivalent to that of a 4.5-meter single mirror.
The six mirrors are not thick solid ones but are of a new lightweight design. They are partially hollow, which requires a smaller mechan­ical structure to move them; thus they save money and construction time. The six images from the six mirrors may be either superimposed to form a single image or aligned along a spectrographic slit, one on top of the other, to take full advantage of slit geome­try. The pointing directions of the six mirrors are locked together by laser beams. This instrument has been successful in demonstrating the practicality of the multiple-mirror concept, and it may be the fore­runner of telescopes that are equivalent to a 25-meter (82-foot) telescope. Under consideration is an MMT consisting of eight 5-meter lightweight mirrors, having the light-gathering power of a 14-meter telescope, the angular resolution of a 22-meter telescope, and (it is hoped) the cost of a 4-meter telescope.
The MMT is not the only new design which shows an artist concep­tion of these new designs for future large telescopes, but it is not now certain whether any of them will ever be built. The success of the Space Telescope, a 2.4-meter conventional reflector that is to be put into orbit in early 1985, will not lessen the need for a mam­moth new telescope on the ground but will probably increase it. Since Space Telescope does not have to contend with light losses produced by the atmos­phere, a very large telescope will be required on the ground to observe in visible wavelengths what Space Telescope is able to observe at shorter wavelengths, where it will primarily operate.

Telescope and Formation of Image

Formation of Image

In optical astronomy the object with which we work is the image of the light source formed by the principal image-forming part of the telescope, which is called an objective. The objective of the optical telescope is either a lens or a mirror. Light rays from the light source are refracted in passing through a lens and are reflected from a mirror. The image is produced where the light rays converge to a position known as the focus. The focal length of the objective is the distance behind the lens to the focus or the distance in front of the mirror to the focus. The image of a star is just a point of light, while that of an extended object, such as the moon, is inverted.
In telescopes using either mirrors or lenses an eye­piece magnifies the image much as a reading glass magnifies small print. Or a photographic plate may be inserted into the focal plane of the objective instead of the eyepiece, transforming the telescope into a giant camera. In this case the objective lens or mirror serves as the camera lens. The advantage of photography over observing with the eye is that the photograph is available for later study and time exposures can record fainter sources than those the eye sees.

The Properties of an Image Formed by Telescope

Properties of an Image 

The image formed by either a lens or a mirror has certain properties that depend upon the diameter, or aperture, of the objective and its focal length. One property is the size of the image. Si nee the image of a star is a point, size is not an important property for it. For an extended object the image size depends upon the angular size of the light source on the sky and upon the weal length of the objective.
Image brightness is important since it determines whether the image is above the threshold of visibility or how long it would take to photograph. The bright­ness of an image of a point source, such as a star, depends on how much light is intercepted by the ob­jective. Hence its brightness is proportional to the area of the objective or to the square of the aperture. Doubling the aperture but leaving the focal length the same increases the area of the objective or its light­gathering power by four times, concentrating four times as much light into the same-size image.
When we photograph an extended object, the sur­face brightness of the image depends on the amount of radiant energy per unit area of the image. The objective's area (or the square of its aperture) still determines the total amount of energy collected, but the total energy is distributed over the entire image. Thus the larger the image's area, the smaller the en­ergy per unit of area. The image size of an extended object increases in proportion to the focal length; so for a given telescope aperture the surface brightness of the image decreases as the focal length is made longer.
How well a telescope discriminates between two adjacent objects or shows fine details is called its re­solving power. Because of the wave natu re of light, the image of a point source produces a diffraction pattern; it appears as a bright central spot, called a diffraction disk, surrounded by progressively fainter rings. When the diffraction patterns of two stars that are close together no longer overlap, we can see sep­arate stellar images. The larger the telescope's aperture, the smaller is the diffraction disk of each image. A large aperture therefore im­proves the resolution of closely adjoining features by making the diffraction effect of adjacent objects over­lap less. We define resolving power as the smallest angle between two close objects whose images can just be separated by a telescope. This critical angle is directly proportional to the wavelength of the ob­served radiation and inversely proportional to the ap­erture of the objective.

What viewing problems a telescope can have?

Viewing Problems in Telescopes
The theoretical resolving power of any optical tele­scope is never fully realized. The lower layers in the earth's atmosphere are unsteady and turbulent; this turbulence blurs and distorts the star's image and makes it twinkle, or scintillate. The rapid scintillations break the starlight into many dancing specks of light, which in long exposu res merge to form the fuzzy stel­lar images we see in photographs. When the atmo­spheric turbulence is low, the stars twinkle, or scintil­late, less, and the so-called seeing is improved. A planet, on the other hand, shines with a steady light because each point on the tiny disk twinkles out of step with neighboring points; we see an average of all the twinkling points.
A technique called speckle photography, which can be used with large telescopes, can get around the smearing and wiggling of the image that comes from atmospheric turbulence. In the exposure of the pho­tographic plate for an extremely short time (less than 0.01 second), each star image appears as a cluster of sharp specks of different brightness. Then the photo­graph is run through a high-speed light-sensing de­vice which measures the variations in brightness across each speck. The information from the assem­blage of specks in each of many photographed images is fed into a computer that is programmed to analyze and reassemble the information into the unsmeared image of the 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. Atmospheric atoms and molecules absorb the ultraviolet photons in sunlight and 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 entering the at­mosphere is bent increasi ngly toward the vertical so that we see a star slightly closer to the zenith (the point directly above the observer) than it really is. This atmospheric-refraction effect is great­est near the horizon (about OS), for there the light's path through the air is the longest. When we observe the rising or setting sun, it is really below our horizon,
but refraction raises the sun's image above the hori­zon by an amount equal to its apparent diameter, which is OS.
Other viewing problems are related to the geo­graphical location of the observatory. An ideal site for an optical observatory is a mountaintop where the air is dry, transparent, and steady, and the sky is dark. An observatory also needs a minimum amount of wind and relatively easy access. The southwestern part of the United States satisfies most of these conditions and has many clear days and nights. Kitt Peak National Observatory is located there, about 65 miles south­west of Tucson, Arizona.

What should be the focal position for reflecting telescopes?

Focal Position for Reflecting Telescopes
Reflecting telescopes can be designed for many kinds of astronomical work through choice of the focal ar­rangement to suit the type of obser­vation. For photography, photometry, and spec­troscopy of faint objects the prime focus is best because its small focal length lessens the exposu re time required. The Newtonian focus, most useful for small telescopes, is now little used by professional astronomers. In both these arrangements the ob­server works at a considerable distance above the ob­servatory floor since both focal positions are near the entrance of the telescope.
I n the Cassegrain focal arrangement a convex sec­ondary mirror positioned at the top in front of the focus slows the rate at which light rays converge, ef­fectively increasing the telescope's focal length. The secondary mirror reflects the converging rays to the bottom of the telescope and through a hole in the objective mirror to focus behind the objective. This is a much more convenient observing position since it is near the floor and behind the telescope. Of all the observations made with the 5.1-meter Hale telescope on Palomar Mountain 75 percent are from the Cas­segrain focus.
We might think that putting the secondary mirror and its supports or the observer's cage for the prime focus into the path of the light rays would obscure part of the image; but the only effect is to cut down the amount of light reaching the objective; the loss is small, and the quality of the image is not affected.
Equipment that is too heavy and bulky to be at­tached to the back of the primary mirror or is sensitive to changing gravitation as the telescope moves can be placed in a room below the observatory floor. An aux­iliary flat mirror diverts the long converging beam down the hollow polar axis around which the tele­scope rotates, and with this coude focal arrangement the focus can remain stationary no matter which way the telescope points.

What is Telescope Mounting?

Telescope Mounting

An optical telescope, in order to follow an object as the earth's rotation carries it across the sky, must be free to move. To track stars accurately and to permit a telescope to be conveniently pointed in any direction, the equatorial mounting system is used for most tele­scopes. This system has two axes of rota­tion: The telescope can rotate in an east-west direc­tion, called hour-angle, around its polar axis, which is aligned with the earth's axis of rotation; another al­lows the telescope to swing in a north-south direction about the declination axis, which is perpendicular to the polar axis.
Large telescopes are usually positioned by a com­puter from an operating console and guided to the exact location with hand controls. Once a large tele­scope is properly set, the computer operating a clock drive slowly turns it westward around its polar axis at the same rate as the earth turns eastward, keeping the stellar images locked in position in the field of view. The great simplicity in equatorial mounting is that tracking requires continuous motion about only one of its two axes. The disadvantage, which obtains in the largest telescopes now in operation and planned for the future, lies in the stresses on the polar axis due to gravity. The polar axis is inclined in the earth's gravitational field and must rotate on one edge of its end. For a very large telescope that is a difficult engineering problem.
One means of removing some of the stress from the primary axis is to align it with gravity. Such a mounting is known as altazimuth mounting; with it a telescope rotates about a vertical axis and about a horizontal axis. This mounting's disadvantage is that, to track a star, it must turn continuously about both axes at the same time. When the telescope approach­es the area of the sky directly overhead, continuous tracking becomes virtually impossible. Even with this disadvantage the altazimuth mounting will be the pri­mary mounting for very large telescopes to be con­structed in the future.

World's Window - The Telescope

The Telescope

The earth receives electromagnetic radiation of all wavelengths from various directions in outer space; yet wavelengths from only two regions of the electro­magnetic spectrum are able to penetrate the earth's atmosphere freely. Most of the electromagnetic spec­trum is screened out by the atmosphere well above the earth's surface. The two spectral windows in the atmosphere through which we can observe the universe from the earth's surface are called the optical window-from about 3,000 ang­stroms to about 10,000 angstroms, or roughly the visible-wavelength region-and the radio window­which includes the wavelength region from about 1 millimeter to 30 meters. The telescopes we build on the earth's su rface to take advantage of these two windows are thus logically called optical telescopes and radio telescopes.
With the advent of the space age, astronomers have been able to take advantage of ai rcraft, balloons, rockets, and satellites to extend our vision of the uni­verse by going above part or all of the earth's veiling atmosphere. Astronomers are aghast at what space telescopes carried by these vehicles have revealed through radiation in the ultraviolet, X-ray, gamma-ray, and infrared regions.

Electromagnetic radiation, Information in the spectrum of light, Photons and Emission and absorption of radiation

Electromagnetic radiation. 
 Electromagnetic radiation is made up of electric and magnetic fields of force continually interacting and propagating in the form of waves at the speed of light. Electromagnetic waves possess a continuous range of wavelengths from short-wavelength gamma rays through X rays, ultra­violet, visible, infrared and microwaves to long-wave­length radio waves. Light displays all the properties of wave phenomena. White (composite) light is com­posed of a range of wavelengths. The apparent bright­ness of light source varies inversely as the square of its distance.
Information in the spectrum of light.
The wave­lengths, or spectral composition, of composite waves contain information about the physical nature of the light source and its environment. The study of the spectra, produced by dispersing white light into its constituent wavelengths, is called spectroscopy. Kirchhoff's laws are diagnostic in that one can infer the physical conditions of the light source by the type of spectrum its light forms: continuous, emission, ab­sorption. Spectral lines in emission and absorption spectra uniquely identify the chemical composition of the emitting or absorbing gas. Three radiation laws­Planck's, Stefan-Boltzmann, and Wien's-can be ap­plied to the analysis of the continuous spectrum of a blackbody to determine its temperature. Thermal sources of radiation like the sun and stars emit radiant energy much like a blackbody. gained much information about the universe from the radiation it emits. And as we develop an even greater understanding of the nature of radiation-its for­mation, propagation, interaction with matter, and destruction-we can explore more deeply the dim sources in the outer reaches of the cosmos, almost back to the beginning of time.
Photons.
Electromagnetic radiation has properties showing it to be discrete as well as wavelike. Photons are discrete units of electromagnetic energy that have no inertia (are massless), are electrically neutral, and move at the speed of light. The energy content of a photon is inversely proportional to its wavelength. Photons are created by taking energy from atoms and destroyed by transferring their energy to the internal energy of atoms.
Emission and absorption of radiation.
The atomic processes responsible for the emission and absorp­tion of photons are summarized in a model for the atom known as the Bohr atom. In hydrogen (and sim­ilarly for other atoms), the electron can occupy only a selected number of allowed orbits; that is, not all pos­sible radius values are permitted for electron orbits; the electron normally resides in the lowest energy orbit, which is the one closest to the nucleus. Elec­trons make transitions to larger allowed orbits when atoms absorb photons and transitions to smaller al­lowed orbits when atoms emit photons. An electron can remain in a higher energy orbit for a very short time before it spontaneously drops to a lower energy state emitting a photon. All electron transitions begin­ning for absorption or ending for emission in an al­lowed energy state constitute a spectral series, such as the Balmer series in hydrogen. diffraction dispersion

The Forces of Nature and Quantum Theory

Quantam Theory and Nature of Forces
Gravity was the first force in nature to be understood in at least a mathe­matical sense. Newton's theory shows that, even though separated by enor­mous distances, pieces of matter can influence each other's state of mo­tion. Although less familiar in many respects, the electric and magnetic forces have been known since ancient times. Like gravity they both weaken as the square of the distance away from their source. In 1873 James
Clerk Maxwell (1831-1879) showed that a relationship exists among elec­tricity, magnetism, and light-an amazingly unifying step.
In 1924 the French physicist Louis de Broglie (1892- ) pointed out that, like light, subatomic particles also have a wave nature, as well as a dis­crete natu reo This has been verified experimentally many times. It is now an accepted fact that matter and radi­ant energy have dual natures in that they show both wave and discrete properties. Taking de Broglie's idea, Erwin Schrodinger (1887-1961), an Austrian physicist, and Werner Hei­senberg (1901-1975), a German phys­icist, independently constructed mathematical theories for atomic structure at about the same time (1925). Their theories were consoli­dated by Paul Dirac, an English phys­icist, into the mathematical formu­lation called quantum mechanics, the most rational and logical approach so far for understanding a vast variety of atomic phenomena. In reality there are no discrete electron orbits like those of planets in the solar system. Within the hydrogen atom, for exam­ple, are spherical regions surrounding the proton. In these regions the elec­tron is spread into a pattern of stand­ing waves, whose distribution corre­sponds to a discrete energy state of the atom.
All atomic properties are known to be the consequence of the electrical interaction between the nucleus and the electrons surrounding it. This electromagnetic interaction is re­sponsible for the characteristic struc­ture of each atomic species. These characteristic structures are re­sponsible for the basic forms of matter, from simple rocks and crystals to flowers and even human beings. The electromagnetic force between the electron and the nucleus is 1039 times stronger than the gravitational force between them; no one has de­tected, nor is there any prospect of detecting, the effects of gravity within atoms or molecules.
By 1932 it was known that the nu­cleus was composed of protons and neutrons. This raised the problem of what force holds the nucleus together against the mutual electrical repulsion of the protons for each other. The solution of this question was the dis­covery of the strong nuclear force of attraction. It is about a hundred times more powerful than the electromagnetic force, but of very short range, and is capable of holding to­gether nuclei with as many as a hun­dred or so protons.
Finally, a fourth force was discov­ered around 1935, the weak nuclear force, which is about 10-; times as strong as the strong nuclear force, or about a thousandth as strong as the electromagnetic force. This force is responsible for some changes in the nature of the nucleus that occur in ra­dioactive decay. It is also a very-short­range force. There is recent evidence that suggests that the electromagnetic force, the weak nuclear force, and possibly the strong nuclear force are actually different manifestations of
the same force acting differently at different distances between particles. Linking all four forces into one uni­versal expression, the so-called unified field theory, still eludes us.
The reason we are familiar with the gravitational and electromagnetic forces is that they operate on the scale of our experiences. The other two, the strong and weak nuclear forces, are confined to the nuclear scale of existence. The gravitational force increases its intensity with in­creasing mass, whereas the other forces are independent of mass. In the cosmos, as we shall see in later chapters, gravity dominates. Gravity is responsible for motion and form in the cosmic realm.

Creation and Destruction of Photons - A Core Concept

Creation and Destruction of Photons
THE BOHR ATOM
One of the most perplexing problems for early­twentieth-century physicists was why the atom emits a discrete pattern of spectral lines. By 1913, when the structure of atoms was reasonably well known, Niels Bohr (1885-1962), a Danish physicist, proposed a the­ory for the structure of the simplest atom, hydrogen, whose one electron orbits around a proton. He sug­gested that the electron can occupy only a selected number of prescribed concentric orbits about the nu­cleus, rather than having an unlimited and unspec­ified orbital distance. Also, the electron normally re­sides in the lowest energy orbit, which is the one closest to the nucleus. Orbits representing higher lev­els of energy are increasingly farther from the nu­cleus. The diameter of the first orbit corresponds to the normal size of the hydrogen atom, about
10-8 centimeter in diameter.
When the atom absorbs energy, it is said to be excited, and the electron appears in one of the outer orbits, which have successively higher energies than the lowest orbit does. The electron's change (up or down) from one allowed orbit to another is called an electron transition. An atom in a gas may acquire the internal energy necessary to excite an electron by ran­dom thermal collisions with other gas atoms, col­lisions with subatomic particles such as free electrons, or absorption of a photon traveling through the gas.
Of all the photons striking the atom only those possessing an amount of energy equal to the energy difference between a higher energy orbit and the one in which the electron is located will be absorbed and excite the atom. For example, in the hydrogen atom it takes 10.2 electron volts, or 1.63 x 10-11 erg, of energy to raise the electron from its lowest energy orbit to the next higher energy orbit. Photons with energies be­low 10.2 electron volts will not be absorbed, and con­sequently the electron will not be excited. Photons with energies in excess of 10.2 electron volts cannot raise the electron to the second orbit, but they may, if they have the right amount of energy, excite the atom by lifting the electron to even higher energy orbits.
How long will an excited atom remain that way? If a gas atom is excited, then in about a hundred­millionth of a second it will rid itself of any energy in excess of that of the lowest energy orbit by emitting the energy in the form of one or more photons. Some­what like a ball bouncing down a staircase, the elec­tron will arop in succession into one or maybe several lower energy orbits on its way to the lowest energy level, where it can reside indefinitely. With each downward transition a photon of electromagnetic radiation is emitted. This photon represents the en­ergy difference between the two orbits between which the electron makes the transition. The greater the energy difference, the greater is the amount of energy released in the form of a photon and, con­sequently, the shorter is the wavelength of the pho­ton. Bohr was led to such a model for the atom as the most straightforward means of accounting for the dis­crete amounts of energy contained in photons. (Labo­ratory experiments had already shown that light really possessed the properties of a discrete phenomenon.)
Two examples from the countless electron transi­tions possible in the hydrogen atom are shown in Figure 4.14. Besides emitting energy spontaneously, an excited atom, before it can emit a photon, may collide with another atom in the gas and transfer en­ergy to it as kinetic energy of motion. In this case no photon will be emitted.
SPECTRUM OF HYDROGEN ATOM
In addition to the model of the atom that represents it by its electron orbits, we can make a model using the energy of each allowed electron orbit. Such an energy-level model for hydrogen appears in Figure 4.16. The number of the energy levels corresponds in a one-to-one fashion with the number of the electron orbits. The distance between successive electron or­bits increases with higher orbit numbers, but the dif­ferences in energy between successive orbits grows smaller as the orbit number increases.
Suppose a hydrogen atom is excited so that its elec­tron is in the third energy level. Then it may de-excite directly to the ground state, emitting one photon. The photon would have a wavelength of 1026 angstroms,
which corresponds to the energy difference between these two orbits in the hydrogen atom. Or it may de-excite to the second energy level, emitting a pho­ton with a wavelength of 6562 angstroms, and then de-excite from the second energy level to the ground state, emitting a photon with a wavelength of 1216 angstroms. The total energy emitted in both cases is the same, but the wavelengths that result and hence the spectral lines differ.
The hydrogen-line spectrum in the visible region, known as the Balmer series, is prominent in the ab­sorption spectra of most stars. It arises from electron transitions originating on the second energy level of the atom. In the same way all the possible transitions from the ground level are known as the Lyman series, which is in the ultraviolet part of the electromagnetic spectrum. Those transitions from the third level up to higher energy levels are the Paschen series, which is in the infrared; and so on for the remaining series, whose lines appear in the far infrared on out to the microwave region of the spectrum.
Each series of spectral lines comes to a limit toward shorter wavelengths. The uppermost levels, repre­senting the electron's highest energy orbits, crowd together toward a series limit, which represents the point beyond which the proton can no longer bind the electron to it. In this case the electron has been re­moved from the atom (it has been ionized) and is free to take on any energy.
If the electron is given enough energy, either by collision or absorption of a photon, it can escape the electrical attraction of the nucleus. The atom is then ionized and is in the form of a positive ion. The ion­ized hydrogen atom cannot absorb or reradiate en­ergy in the form of discrete lines until it captures a free electron. It can execute the captu re because of the electrical force of attraction between the positively charged nucleus (the proton) and the negatively charged electron. Note that it converges toward its series limit at approximately 3646 angstroms.
SPECTRA OF OTHER ElEMENTS
In the Bohr atom, besides limits on the size of electron orbits, there is a limit to the number of electrons that may occupy a given orbit. These allowed orbits with a prescribed number of electrons in them are called electron shells.
In general, as one goes through the periodic table, electrons are added to balance the number of protons in the nucleus by filling the shells from the one closest to the nucleus outward. In hydrogen there is one elec­tron in the innermost shell, which has room for a maximum of 2 electrons. Helium's 2 electrons fill, or close, the shell so that for the element lithium the third electron must start a new shell, which is the next innermost. In the second shell there is room for only 8 electrons; in the third, 18 electrons; in the fourth, 32; and so on.
Each element has a unique set of energy levels. Consequently, the wavelength of the spectral lines originating from electron transitions between various energy levels is also unique for each element-a clear fingerprint of the element.
The amount of energy needed to ionize an atom varies from one element to the next depending on the number and "position" of the electrons. For example, to remove the outermost electron from helium takes five times as much energy as it does to do the same for sodium. Also, for a given element each additional ionization takes more energy to free an electron from an inner orbit than from an outer one because the inner one is more tightly bound to the nucleus. Thus, with carbon as an example, it takes more than twice as much energy to remove the second electron than it does to remove the first; fou r times more for the third electron than the first; almost six times more for the fourth electron; thirty-five times more for the fifth electron; and a whopping forty-four times more for the innermost sixth electron than for the outermost first electron.
Multiple ionization of a carbon atom brings a corre­sponding readjustment of the energy levels because of the altered electrical attraction between the posi­tive nucleus of the carbon atom and the reduced num­ber of electrons. Altering the energy corresponding to each allowed orbit produces different spectral lines with each succeeding ionization of the carbon atom. So we see not only different wavelengths for absorp­tion lines or emission lines between different un­ionized elements but for the same element a different spectrum after each ionization. That is, the spectrum of singly ionized carbon differs from the spectrum of neutral carbon; doubly ionized carbon differs from singly ionized; triply ionized differs from doubly ion­ized; and so on.
Using the properties of electromagnetic radiation, the atom's structure, the interaction between matter and energy, and spectrum analysis, astronomers have gained much information about the universe from the radiation it emits. And as we develop an even greater understanding of the nature of radiation-its for­mation, propagation, interaction with matter, and destruction-we can explore more deeply the dim sources in the outer reaches of the cosmos, almost back to the beginning of time.

Photons - The Discrete Nature of LIght

The Discrete Nature of Light
From his theoretical study of the emission of radiation by blackbodies Planck concluded that they do not emit or absorb energy in a continuous fashion but only discontinuously in discrete units, which later were called photons. This means that the energy transported by an electromagnetic wave is not con­tinuously distributed over the wave front; it is located at discrete points (the photons) along the wave and moves with the wave. In 1905 Einstein used Planck's idea of a discrete nature for the emission of light to explain a phenomenon discovered in 1887, known as the photoelectric effect. This effect cannot be under­stood if light has only a wave nature. Since that time an extensive body of experimental and theoretical ev­idence has been collected to validate the photon con­cept of light.
What are some of the properties of photons? They move with the speed of light, have no inertia, are electrically neutral, and are massless. Picture a radi­ating body as emitting photons of differing discrete amounts of energy in all directions. The photons retain their energy while traveling through space. Their arrival rate, or flux, at any point in space decreases with the square of their distance from the radiating source. Hence the inverse-square law of light can be understood in terms of numbers of photons (brightness of radiation) as well as in terms of electromagnetic waves.
The energy of each photon is inversely propor­tional to its wavelength. The shorter the wavelength, the more energetic is the photon; the longer the wavelength, the less energetic is the photon. That is why, for example, very-short-wavelength X-ray and gamma-ray photons can destroy molecular structures in living tissue while photons of visible light cannot. If you find that talking about the wavelength of a photon while saying that photons are the localizations of en­ergy in an electromagnetic wave seems paradoxical, then you are perfectly normal. Just as in our dis­cussion of relativity in Chapter 3, where we encoun­tered a reality not evident in our human existence, so it is with our conceptual picture of light as both wave and particle. Wavelength is a characterization of the wavelike properties of light, while the energy content of a photon refers to its discrete nature. The fact that swe can link wavelength and energy content in a math­ematical equation is strong validation of our concep­tual picture.
Photons may be absorbed by an atom, scattered by particles of matter, or converted into matter by inter­action with other photons. They are created inside atoms and in violent collisions between subatomic particles. When they lose their identity, they transfer their energy to some other physical system; and when they are created, they obtain their energy from some other physical system. Thei r creation and destruction is a classic example of the conservation of energy. The concept of light as being simultaneously both discrete photons and continuous waves is not contradictory but is a reality not borne out in our everyday life. Yet experiments are designed to inquire about either light's wave nature or its corpuscular nature; no ex­periment will simultaneously yield the discrete and the wave properties of light.
The theory of the discrete nature of light began a conceptual revolution in twentieth-century physics and astrophysics. It was used by Niels Bohr to formu­late a new model for the atom that can be used to understand how light is created and destroyed inside the atom.
Creation and Destruction of Photons
THE BOHR ATOM
One of the most perplexing problems for early­twentieth-century physicists was why the atom emits a discrete pattern of spectral lines. By 1913, when the structure of atoms was reasonably well known, Niels Bohr (1885-1962), a Danish physicist, proposed a the­ory for the structure of the simplest atom, hydrogen, whose one electron orbits around a proton. He sug­gested that the electron can occupy only a selected number of prescribed concentric orbits about the nu­cleus, rather than having an unlimited and unspec­ified orbital distance. Also, the electron normally re­sides in the lowest energy orbit, which is the one closest to the nucleus. Orbits representing higher lev­els of energy are increasingly farther from the nu­cleus. The diameter of the first orbit corresponds to the normal size of the hydrogen atom, about
10-8 centimeter in diameter.
When the atom absorbs energy, it is said to be excited, and the electron appears in one of the outer orbits, which have successively higher energies than the lowest orbit does. The electron's change (up or down) from one allowed orbit to another is called an electron transition. An atom in a gas may acquire the internal energy necessary to excite an electron by ran­dom thermal collisions with other gas atoms, col­lisions with subatomic particles such as free electrons, or absorption of a photon traveling through the gas.
Of all the photons striking the atom only those possessing an amount of energy equal to the energy difference between a higher energy orbit and the one in which the electron is located will be absorbed and excite the atom. For example, in the hydrogen atom it takes 10.2 electron volts, or 1.63 x 10-11 erg, of energy to raise the electron from its lowest energy orbit to the next higher energy orbit. Photons with energies be­low 10.2 electron volts will not be absorbed, and con­sequently the electron will not be excited. Photons with energies in excess of 10.2 electron volts cannot raise the electron to the second orbit, but they may, if they have the right amount of energy, excite the atom by lifting the electron to even higher energy orbits.
How long will an excited atom remain that way? If a gas atom is excited, then in about a hundred­millionth of a second it will rid itself of any energy in excess of that of the lowest energy orbit by emitting the energy in the form of one or more photons. Some­what like a ball bouncing down a staircase, the elec­tron will arop in succession into one or maybe several lower energy orbits on its way to the lowest energy level, where it can reside indefinitely. With each downward transition a photon of electromagnetic radiation is emitted. This photon represents the en­ergy difference between the two orbits between which the electron makes the transition. The greater the energy difference, the greater is the amount of energy released in the form of a photon and, con­sequently, the shorter is the wavelength of the pho­ton. Bohr was led to such a model for the atom as the most straightforward means of accounting for the dis­crete amounts of energy contained in photons. (Labo­ratory experiments had already shown that light really possessed the properties of a discrete phenomenon.)
Two examples from the countless electron transi­tions possible in the hydrogen atom are shown in Figure 4.14. Besides emitting energy spontaneously, an excited atom, before it can emit a photon, may collide with another atom in the gas and transfer en­ergy to it as kinetic energy of motion. In this case no photon will be emitted.
SPECTRUM OF HYDROGEN ATOM
In addition to the model of the atom that represents it by its electron orbits, we can make a model using the energy of each allowed electron orbit. Such an energy-level model for hydrogen appears in Figure 4.16. The number of the energy levels corresponds in a one-to-one fashion with the number of the electron orbits. The distance between successive electron or­bits increases with higher orbit numbers, but the dif­ferences in energy between successive orbits grows smaller as the orbit number increases.
Suppose a hydrogen atom is excited so that its elec­tron is in the third energy level. Then it may de-excite directly to the ground state, emitting one photon. The photon would have a wavelength of 1026 angstroms,
which corresponds to the energy difference between these two orbits in the hydrogen atom. Or it may de-excite to the second energy level, emitting a pho­ton with a wavelength of 6562 angstroms, and then de-excite from the second energy level to the ground state, emitting a photon with a wavelength of 1216 angstroms. The total energy emitted in both cases is the same, but the wavelengths that result and hence the spectral lines differ.
The hydrogen-line spectrum in the visible region, known as the Balmer series, is prominent in the ab­sorption spectra of most stars. It arises from electron transitions originating on the second energy level of the atom. In the same way all the possible transitions from the ground level are known as the Lyman series, which is in the ultraviolet part of the electromagnetic spectrum. Those transitions from the third level up to higher energy levels are the Paschen series, which is in the infrared; and so on for the remaining series, whose lines appear in the far infrared on out to the microwave region of the spectrum.
Each series of spectral lines comes to a limit toward shorter wavelengths. The uppermost levels, repre­senting the electron's highest energy orbits, crowd together toward a series limit, which represents the point beyond which the proton can no longer bind the electron to it. In this case the electron has been re­moved from the atom (it has been ionized) and is free to take on any energy.
If the electron is given enough energy, either by collision or absorption of a photon, it can escape the electrical attraction of the nucleus. The atom is then ionized and is in the form of a positive ion. The ion­ized hydrogen atom cannot absorb or reradiate en­ergy in the form of discrete lines until it captures a free electron. It can execute the captu re because of the electrical force of attraction between the positively charged nucleus (the proton) and the negatively charged electron. Note that it converges toward its series limit at approximately 3646 angstroms.
SPECTRA OF OTHER ElEMENTS
In the Bohr atom, besides limits on the size of electron orbits, there is a limit to the number of electrons that may occupy a given orbit. These allowed orbits with a prescribed number of electrons in them are called electron shells.
In general, as one goes through the periodic table, electrons are added to balance the number of protons in the nucleus by filling the shells from the one closest to the nucleus outward. In hydrogen there is one elec­tron in the innermost shell, which has room for a maximum of 2 electrons. Helium's 2 electrons fill, or close, the shell so that for the element lithium the third electron must start a new shell, which is the next innermost. In the second shell there is room for only 8 electrons; in the third, 18 electrons; in the fourth, 32; and so on.
Each element has a unique set of energy levels. Consequently, the wavelength of the spectral lines originating from electron transitions between various energy levels is also unique for each element-a clear fingerprint of the element.
The amount of energy needed to ionize an atom varies from one element to the next depending on the number and "position" of the electrons. For example, to remove the outermost electron from helium takes five times as much energy as it does to do the same for sodium. Also, for a given element each additional ionization takes more energy to free an electron from an inner orbit than from an outer one because the inner one is more tightly bound to the nucleus. Thus, with carbon as an example, it takes more than twice as much energy to remove the second electron than it does to remove the first; fou r times more for the third electron than the first; almost six times more for the fourth electron; thirty-five times more for the fifth electron; and a whopping forty-four times more for the innermost sixth electron than for the outermost first electron.
Multiple ionization of a carbon atom brings a corre­sponding readjustment of the energy levels because of the altered electrical attraction between the posi­tive nucleus of the carbon atom and the reduced num­ber of electrons. Altering the energy corresponding to each allowed orbit produces different spectral lines with each succeeding ionization of the carbon atom. So we see not only different wavelengths for absorp­tion lines or emission lines between different un­ionized elements but for the same element a different spectrum after each ionization. That is, the spectrum of singly ionized carbon differs from the spectrum of neutral carbon; doubly ionized carbon differs from singly ionized; triply ionized differs from doubly ion­ized; and so on.
Using the properties of electromagnetic radiation, the atom's structure, the interaction between matter and energy, and spectrum analysis, astronomers have gained much information about the universe from the radiation it emits. And as we develop an even greater understanding of the nature of radiation-its for­mation, propagation, interaction with matter, and destruction-we can explore more deeply the dim sources in the outer reaches of the cosmos, almost back to the beginning of time.

Light a discrete physical quantity

The Discrete Nature of Light
From his theoretical study of the emission of radiation by blackbodies Planck concluded that they do not emit or absorb energy in a continuous fashion but only discontinuously in discrete units, which later were called photons. This means that the energy transported by an electromagnetic wave is not con­tinuously distributed over the wave front; it is located at discrete points (the photons) along the wave and moves with the wave. In 1905 Einstein used Planck's idea of a discrete nature for the emission of light to explain a phenomenon discovered in 1887, known as the photoelectric effect. This effect cannot be under­stood if light has only a wave nature. Since that time an extensive body of experimental and theoretical ev­idence has been collected to validate the photon con­cept of light.
What are some of the properties of photons? They move with the speed of light, have no inertia, are electrically neutral, and are massless. Picture a radi­ating body as emitting photons of differing discrete amounts of energy in all directions. The photons retain their energy while traveling through space. Their arrival rate, or flux, at any point in space decreases with the square of their distance from the radiating source. Hence the inverse-square law of light can be understood in terms of numbers of photons (brightness of radiation) as well as in terms of electromagnetic waves.
The energy of each photon is inversely propor­tional to its wavelength. The shorter the wavelength, the more energetic is the photon; the longer the wavelength, the less energetic is the photon. That is why, for example, very-short-wavelength X-ray and gamma-ray photons can destroy molecular structures in living tissue while photons of visible light cannot. If you find that talking about the wavelength of a photon while saying that photons are the localizations of en­ergy in an electromagnetic wave seems paradoxical, then you are perfectly normal. Just as in our dis­cussion of relativity in Chapter 3, where we encoun­tered a reality not evident in our human existence, so it is with our conceptual picture of light as both wave and particle. Wavelength is a characterization of the wavelike properties of light, while the energy content of a photon refers to its discrete nature. The fact that swe can link wavelength and energy content in a math­ematical equation is strong validation of our concep­tual picture.
Photons may be absorbed by an atom, scattered by particles of matter, or converted into matter by inter­action with other photons. They are created inside atoms and in violent collisions between subatomic particles. When they lose their identity, they transfer their energy to some other physical system; and when they are created, they obtain their energy from some other physical system. Thei r creation and destruction is a classic example of the conservation of energy. The concept of light as being simultaneously both discrete photons and continuous waves is not contradictory but is a reality not borne out in our everyday life. Yet experiments are designed to inquire about either light's wave nature or its corpuscular nature; no ex­periment will simultaneously yield the discrete and the wave properties of light.
The theory of the discrete nature of light began a conceptual revolution in twentieth-century physics and astrophysics. It was used by Niels Bohr to formu­late a new model for the atom that can be used to understand how light is created and destroyed inside the atom.
Creation and Destruction of Photons

Atom or The Bohr Atom -

THE BOHR ATOM
One of the most perplexing problems for early­twentieth-century physicists was why the atom emits a discrete pattern of spectral lines. By 1913, when the structure of atoms was reasonably well known, Niels Bohr (1885-1962), a Danish physicist, proposed a the­ory for the structure of the simplest atom, hydrogen, whose one electron orbits around a proton. He sug­gested that the electron can occupy only a selected number of prescribed concentric orbits about the nu­cleus, rather than having an unlimited and unspec­ified orbital distance. Also, the electron normally re­sides in the lowest energy orbit, which is the one closest to the nucleus. Orbits representing higher lev­els of energy are increasingly farther from the nu­cleus. The diameter of the first orbit corresponds to the normal size of the hydrogen atom, about 10-8 centimeter in diameter.
When the atom absorbs energy, it is said to be excited, and the electron appears in one of the outer orbits, which have successively higher energies than the lowest orbit does. The electron's change (up or down) from one allowed orbit to another is called an electron transition. An atom in a gas may acquire the internal energy necessary to excite an electron by ran­dom thermal collisions with other gas atoms, col­lisions with subatomic particles such as free electrons, or absorption of a photon traveling through the gas.
Of all the photons striking the atom only those possessing an amount of energy equal to the energy difference between a higher energy orbit and the one in which the electron is located will be absorbed and excite the atom. For example, in the hydrogen atom it takes 10.2 electron volts, or 1.63 x 10-11 erg, of energy to raise the electron from its lowest energy orbit to the next higher energy orbit. Photons with energies be­low 10.2 electron volts will not be absorbed, and con­sequently the electron will not be excited. Photons with energies in excess of 10.2 electron volts cannot raise the electron to the second orbit, but they may, if they have the right amount of energy, excite the atom by lifting the electron to even higher energy orbits.
How long will an excited atom remain that way? If a gas atom is excited, then in about a hundred­millionth of a second it will rid itself of any energy in excess of that of the lowest energy orbit by emitting the energy in the form of one or more photons. Some­what like a ball bouncing down a staircase, the elec­tron will arop in succession into one or maybe several lower energy orbits on its way to the lowest energy level, where it can reside indefinitely. With each downward transition a photon of electromagnetic radiation is emitted. This photon represents the en­ergy difference between the two orbits between which the electron makes the transition. The greater the energy difference, the greater is the amount of energy released in the form of a photon and, con­sequently, the shorter is the wavelength of the pho­ton. Bohr was led to such a model for the atom as the most straightforward means of accounting for the dis­crete amounts of energy contained in photons. (Labo­ratory experiments had already shown that light really possessed the properties of a discrete phenomenon.)
Besides emitting energy spontaneously, an excited atom, before it can emit a photon, may collide with another atom in the gas and transfer en­ergy to it as kinetic energy of motion. In this case no photon will be emitted.

The Secrets of Hydrogen Spectrum

SPECTRUM OF HYDROGEN ATOM
In addition to the model of the atom that represents it by its electron orbits, we can make a model using the energy of each allowed electron orbit. Such an energy-level model for hydrogen appears in Figure 4.16. The number of the energy levels corresponds in a one-to-one fashion with the number of the electron orbits. The distance between successive electron or­bits increases with higher orbit numbers, but the dif­ferences in energy between successive orbits grows smaller as the orbit number increases.
Suppose a hydrogen atom is excited so that its elec­tron is in the third energy level. Then it may de-excite directly to the ground state, emitting one photon. The photon would have a wavelength of 1026 angstroms,
which corresponds to the energy difference between these two orbits in the hydrogen atom. Or it may de-excite to the second energy level, emitting a pho­ton with a wavelength of 6562 angstroms, and then de-excite from the second energy level to the ground state, emitting a photon with a wavelength of 1216 angstroms. The total energy emitted in both cases is the same, but the wavelengths that result and hence the spectral lines differ.
The hydrogen-line spectrum in the visible region, known as the Balmer series, is prominent in the ab­sorption spectra of most stars. It arises from electron transitions originating on the second energy level of the atom. In the same way all the possible transitions from the ground level are known as the Lyman series, which is in the ultraviolet part of the electromagnetic spectrum. Those transitions from the third level up to higher energy levels are the Paschen series, which is in the infrared; and so on for the remaining series, whose lines appear in the far infrared on out to the microwave region of the spectrum.
Each series of spectral lines comes to a limit toward shorter wavelengths. The uppermost levels, repre­senting the electron's highest energy orbits, crowd together toward a series limit, which represents the point beyond which the proton can no longer bind the electron to it. In this case the electron has been re­moved from the atom (it has been ionized) and is free to take on any energy.
If the electron is given enough energy, either by collision or absorption of a photon, it can escape the electrical attraction of the nucleus. The atom is then ionized and is in the form of a positive ion. The ion­ized hydrogen atom cannot absorb or reradiate en­ergy in the form of discrete lines until it captures a free electron. It can execute the captu re because of the electrical force of attraction between the positively charged nucleus (the proton) and the negatively charged electron. Note that it converges toward its series limit at approximately 3646 angstroms.
SPECTRA OF OTHER ElEMENTS
In the Bohr atom, besides limits on the size of electron orbits, there is a limit to the number of electrons that may occupy a given orbit. These allowed orbits with a prescribed number of electrons in them are called electron shells.
In general, as one goes through the periodic table, electrons are added to balance the number of protons in the nucleus by filling the shells from the one closest to the nucleus outward. In hydrogen there is one elec­tron in the innermost shell, which has room for a maximum of 2 electrons. Helium's 2 electrons fill, or close, the shell so that for the element lithium the third electron must start a new shell, which is the next innermost. In the second shell there is room for only 8 electrons; in the third, 18 electrons; in the fourth, 32; and so on.
Each element has a unique set of energy levels. Consequently, the wavelength of the spectral lines originating from electron transitions between various energy levels is also unique for each element-a clear fingerprint of the element.
The amount of energy needed to ionize an atom varies from one element to the next depending on the number and "position" of the electrons. For example, to remove the outermost electron from helium takes five times as much energy as it does to do the same for sodium. Also, for a given element each additional ionization takes more energy to free an electron from an inner orbit than from an outer one because the inner one is more tightly bound to the nucleus. Thus, with carbon as an example, it takes more than twice as much energy to remove the second electron than it does to remove the first; fou r times more for the third electron than the first; almost six times more for the fourth electron; thirty-five times more for the fifth electron; and a whopping forty-four times more for the innermost sixth electron than for the outermost first electron.
Multiple ionization of a carbon atom brings a corre­sponding readjustment of the energy levels because of the altered electrical attraction between the posi­tive nucleus of the carbon atom and the reduced num­ber of electrons. Altering the energy corresponding to each allowed orbit produces different spectral lines with each succeeding ionization of the carbon atom. So we see not only different wavelengths for absorp­tion lines or emission lines between different un­ionized elements but for the same element a different spectrum after each ionization. That is, the spectrum of singly ionized carbon differs from the spectrum of neutral carbon; doubly ionized carbon differs from singly ionized; triply ionized differs from doubly ion­ized; and so on.
Using the properties of electromagnetic radiation, the atom's structure, the interaction between matter and energy, and spectrum analysis, astronomers have gained much information about the universe from the radiation it emits. And as we develop an even greater understanding of the nature of radiation-its for­mation, propagation, interaction with matter, and destruction-we can explore more deeply the dim sources in the outer reaches of the cosmos, almost back to the beginning of time.

How to identify the elements from emission or absorption spectra?

IDENTIFYING THE ELEMENTS FROM EMISSION OR ABSORPTION SPECTRA
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. With the aid of laboratory spectral analysis of the different chemical elements, 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 light 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. 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.
 All objects radiate and absorb some form of electro­magnetic radiation; the wavelength region and the amount of energy depend generally on the body's temperature and physical state. From laboratory ex­periments and from theory physicists in the nine­teenth century analyzed how various bodies emit and absorb radiation as a function of temperature and wavelength. From this work they developed the con­cept of an idealized radiator called a blackbody.

What is the Bllackbody Concept?

BLACKBODY CONCEPT: 
A blackbody is an imaginary body that, when cool, absorbs all the radiant energy falling on its surface so that it is black in color; when hot, the blackbody emits energy with 100 per­cent efficiency. (Real matter is generally less than 100 percent efficient when it radiates.)
For our purposes the most important feature of the blackbody is the way in which emitted radiant energy is spread in wavelength, or the spectral energy distri­bution. Scientists have found that the distribution of energy depends only on the blackbody'S temperature and not on its chemical composition. Note how the amount of radiant en­ergy emitted by a blackbody varies with wavelength in a very recognizable way, even for different tem­peratures. The emission of radiant energy (or the brightness at each wavelength) covers a continuous range of wavelengths so that the spectrum of a black­body is a continuous spectrum; that is, there are no color bands missing from its spectrum. At room tem­perature, lampblack (a finely powdered black soot) is very close to being a blackbody because it absorbs almost all the radiation incident upon it and reflects very little. Fortunately, the radiation emitted by stars tends to be much like that emitted by a blackbody.
In 1900 the German physicist Max Planck (1858-1947) derived a mathematical expression, now called Planck's law, that describes the distribution of brightness in the spectrum of a blackbody. There are two other dis­tinguishing characteristics of the spectrum of black­body radiation: First, the energy emitted by the blackbody is greater at every wavelength as the temperature increases. Thus the total amount of radi­ant energy emitted increases with increasing tem­perature, which is known as the Stefan-Boltzmann law. Second, the greatest amount of radiation is fou nd toward shorter wavelengths (blue end of the spectrum) as the temperature increases. This is known as Wien's dis­placement law.
The significance of the blackbody-radiation laws­Planck's law, the Stefan-Boltzmann law, and Wien's law-is that bodies that emit electromagnetic radi­ation because they are hot, such as stars, do so much like a blackbody. Thus the blackbody-radiation laws are powerful diagnostic tools for measuring the tem­perature of these thermal sources of radiation. For the study of bodies that emit radiation not because they are hot (called nonthermal sources of radiation) but because of some selective physical processes, the blackbody-radiation laws are of no use. Fortunately, most of the celestial bodies-all the stars-are ther­mal sources of radiation and emit much like a black­body. Some everyday examples of thermal sources of radiation are an incandescent light bulb, the burner on an electric stove, and the flame of a cutting torch. Examples of nonthermal sources are a fluorescent light, lightning, and a television screen.