Refraction

Light travels through a vacuum at its maximum speed of about 3.0 × 10⁸ m/s, and in a straight path. Light travels at slower speeds through different materials, such as glass or air. (A material through which light travels is called a medium, plural media.) When traveling from one medium to another, some light will be reflected at the surface of the new medium. The light that continues through the new medium will either speed up or slow down depending on how fast it can travel through each medium. For example, light travels more quickly through air than through water. The refractive index of a medium is the ratio of the speed of light in a vacuum to the speed of light in the medium. The higher the refractive index, the more light is slowed down by the substance.

If light enters the new medium at a right angle to the surface (along the normal), it will change speed, but not direction. If it enters at an angle, its speed and its direction will change. The direction the light takes depends on whether it travels faster or slower in the new medium. Imagine driving a car from smooth pavement onto a sandy beach. If you approach the beach straight on, the car will slow down, but not change direction. If the you approach the beach at an angle, one of the tires will be slowed down by the sand before the other is, and the car will turn in the direction of the tire that touched the sand first.

Light follows the same same principle and bends towards the normal when traveling into a medium with a higher index of refraction, and away from the normal when traveling into a medium where it can go faster. In the diagram below, light is leaving air and entering glass, so it bends towards the normal on the way in, and away on the way out of the glass.

Reflection

When light hits a surface that it can't travel through, it bounces back. If the surface is smooth, like a mirror, the light will reflect in a predictable way. If the surface is flat, the angle at which a beam of light approaches the mirror will be equal to the angle at which the beam is reflected, so i = r in the diagram below.

Lenses

Lenses form images by refraction and are typically made of either glass or plastic. They are ground so that their surfaces are either segments of spheres or planes. If a lens is convex or converging, it takes parallel light rays from a distant object and bends them so that they converge to a single point called the focal point. The distance from the lens to the focal point is called the focal length of the lens.

If a lens is concave or diverging, it takes parallel rays and bends them so that they spread out. The rays will then appear to originate from a point in front of the lens. This point is also called the focal point, and its distance is measured in negative units. 

Telescopes

The earliest telescopes, as well as many amateur telescopes today, use lenses to gather more light than the human eye could collect on its own. They focus the light and make distant objects appear brighter, clearer and magnified. This type of telescope is called a refracting telescope.

Most refracting telescopes use two main lenses. The largest lens is called the objective lens, and the smaller lens used for viewing is called the eyepiece lens.

from: http://www.astronomygcse.co.uk/AstroGCSE/Unit5/RefractingTelescopes.htm

The size of an image produced by a lens is proportional to the focal length of the lens. The longer the focal length, the larger the image. The brightness of an image from a telescope depends partly on how much light is collected by the telescope. The light-gathering power of a telescope is directly proportional to the area of the objective lens. The larger the lens, the more light the telescope can gather. Doubling the diameter of the lens increases the light gathering power by a factor of 4. Brightness of images also depends on how big an area the image light is spread over. The smaller the area, the brighter the image.

The magnifying power of a telescope is the ratio of an object's angular diameter to its naked eye diameter. This depends on the focal length of both lenses.

Magnification = focal length of objective lens/focal length of eyepiece lens

An example to try:

A small refracting telescope has an objective of focal length 100 cm. If the eyepiece has a focal length of 4.0 cm, what is the magnification of the telescope?

Answer: 

Magnification = 100/4.0 = 25 (usually written as 25×) 

Magnification might seem like the most important aspect of a telescope, but there are limits to how sharp an image a telescope can produce because of the blurring effects of the Earth's atmosphere. Magnifying a blurred image makes it bigger, but not clearer, so the priority when  telescopes are built is to have the greatest light-gathering power possible. Gathering more light makes brighter images, and brighter images make it easier to see faint details.

Galileo is credited with being the first person to use a telescope to make observations of the night sky. After hearing of the invention of the telescope in 1608, he built one of his own, called a Gallilean Telescope, in 1609 using a convex objective lens and a concave eyepiece lens. His telescope could magnify objects 3 times. Telescopes he made later magnified objects up to 30 times.

Limitations of Refracting Telescopes

Lenses create a type of image distortion known as chromatic aberration. This occurs because as light passes through a lens, different colors are bent through different angles (like in a prism) and brought to a focus at different points. Because of this, stars viewed through a simple lens are surrounded by rainbow colored halos. This can be corrected for by adding a thin lens of a different kind of glass behind the objective lens.

Lenses present other optical problems including how difficult and expensive it is to make large lenses completely free of defects. Glass also absorbs most ultraviolet light, and visible light is substantially dimmed as it passes through a lens. In addition, lenses in telescopes can only be supported around the outside, so large lenses can sag and distort under their own weight. All of these problems affect image quality and clarity.

Curved mirrors can bend light and make parallel light rays converge to a focus. This focus is directly in the path of the incoming light, so there are several ways of making images from the mirror visible. One is called a Newtonian reflector, where a flat mirror is used to point the light rays out to an eyepiece.

There are several other types of reflectors that solve the issue of where to focus the light in different ways. Cassegrain reflectors have a convex secondary mirror and a hole in the middle of the primary mirror.  Prime focus telescopes have no secondary optics and the observer or camera observes the image from near the focal point. Coudé telescopes use a convex secondary mirror like a Cassegrain and an angled mirror like a Newtonian reflector to move the light rays to a focal point away from the telescope. This arrangement is useful when optical equipment is being used that is too heavy to mount directly on the telescope.

Reflecting telescopes have many advantages over refracting telescopes. Mirrors don't cause chromatic aberration and they are easier and cheaper to build large. The are also easier to mount because the back of the mirror can be used to attach to the mount. Reflecting telescopes have a few disadvantages as well. Because they are normally open, the mirrors have to be cleaned. Also, unless the mirrors and other optics are kept at the same temperature as the outside air, there will be air currents inside the telescope that will cause images to be fuzzy.

Different reflectors use different shapes of mirrors. Parabolic mirrors will focus all incoming light rays to a single point. However, images from a parabolic mirror will have  a defect called coma, where images far from the center of the field of view are elongated. A spherical mirror surface is relatively easy to make, but different parts of a spherical mirror have slightly different focal lengths, so images will be fuzzy. Mirrors in modern telescopes area made in various shapes to correct for these errors. Some telescopes use a combination of mirrors and lenses. Schmidt-Cassegrain telescopes use a spherical mirror with a correcting plate that corrects the focus. 

LCOGT's 1.0 meter telescopes are quasi-Ritchey-Chrétien telescopes. A true Ritchey-Chrétien has a hyperbolic primary and a hyperbolic secondary mirror. In the design of LCOGT's 1.0 meter telescopes, the shape of the mirrors has been changed a bit in order to find a more optimal optical design for the system as a whole. Because the mirror shapes have changed, the 2 mirrors alone no longer are a Ritchey-Chrétien telescope in the strict definition of the design.  It is still close though, hence the name, "quasi-RC."

This animation shows how light travels in LCOGT's 2.0 meter Falukes Telescopes. 

Astronomers' view of the night sky improved dramatically as soon as photography technology had developed enough to allow telescopes to take photographs of the sky.

In 1840, John W. Draper became the first person to photograph a celestial object when he took a daguerrotype image of the Moon. As his technique improved, Draper captured more detail of craters and other features on the moon's surface. It quickly became clear that photography would revolutionize astronomy. Astronomers began to use photographs to collect precise records of the position, brightness, spectra and features of celestial objects, and no longer had to rely on their hand-drawn sketches and log notes. The first daguerrotypes and photographs were not very sensitive and were complicated to use. In the early 1870s, Richard Leach Maddox developed a new photographic medium made with a dry gelatin, cadmium bromide and silver nitrate preparation on glass plates. Eventually as the technology improved, these glass plates could be made more sensitive, and inexpensively. 

In 1880 John W. Draper's son Henry Draper used this new dry plate technology to make a 51 minute exposure of the Orion Nebula, which was essentially the first deep sky astrophoto.

Charge coupled devices, or CCDs are sensitive detectors of photons that can be used in telescopes instead of film or photographic plates to produce images. CCDs were invented in the late 1960s and are now used in digital cameras, photocopiers and many other devices. Its inventors, Willard Boyle and George E. Smith received the Nobel Prize in physics in 2009 for their work.

A CCD is a tiny microchip onto which the light that the telescope collects is focused. The microchip consists of a large grid of individual light sensing elements called pixels. There are 2048 pixels along each side of the chip in the Merope Camera in Faulkes Telescope North. Each pixel is a 13.5 micrometers(µm) square printed on a cracker sized piece of silicon 50µm thick. Tissue paper is about the same thickness. The images below are of astronomical CCDs from one of LCOGT's telescopes and shows the front and back of a CCD.

When light falls onto one of the pixels, electrons are released from atoms in the pixel. To measure the amount of light that fell onto each pixel, the number of electrons that was released has to be counted. This is done by measuring the charge on the pixel at the end of the last row in the grid. Then that charge is discarded and all the other charges in the row are made to move along to that one corner pixel. The next charge in line is then measured, and so on – until all the charges in that row have been dealt with. Then all the charges in all the remaining rows are made to move over one row, and the whole process is repeated. Amazingly, the entire chip can be "read" in less than 10 seconds. It is this method of read out that distinguishes CCDs from other devices (such as photodiodes and CMOS devices) that convert photons to electrons.

The animation below gives a visualization of how this works.

CCDs are increadibly powerful tools for astronomers because when a telescope's motion is synchronized with the Earth's rotation, the camera can “stare” at one spot in space for hours at a time. The longer the CCD is exposed to the sky, the more photons will land on it, and fainter, more distant objects can be imaged than are otherwise visible. CCD exposures are so long in astronomy (seconds, minutes or even longer) compared to digital cameras (normally a fraction of a second), that CCDs in telescopes are usually kept very cold (−50° to -100°C). Keeping the CCD at a very low temperature minimizes the effects of thermal noise. At any given temperature, a certain fraction of the electrons in the atoms of the CCD itself will will have enough thermal energy to liberate themselves. They are then indistinguishable from electrons liberated by the interaction of the CCD with incoming photons from the telescope, so they get counted as if they were light from a star.

Image Quality

Environmental Factors:

Image quality is affected by many factors some of which come from the environment. Seeing conditions, the Moon, light pollution, and clouds all affect the quality and accuracy of astronomical observations and images.

Astronomical seeing describes the conditions of the night sky and how suitable it is for astronomical observations. Turbulence and temperature variations in the Eath's atmosphere cause astronomical objects to appear to twinkle and form blurry images, and places a limit on a telescope's ability to resolve stars. These effects can come from anywhere between the air in the telescope itself to air in the high atmosphere. A seeing disk is the angular diameter of a star's image, or the region in which the star appears to be moving, which is spread out because of the motion in the air that the light traveled through to get to the CCD. At the very best astronomical sites in the world, this disk is usually less than 1 arcsec in diameter and sometimes as small as 0.25 arcsec. Bad seeing conditions can have seeing disks of 4 arcseconds or more.

Bad seeing caused by temperature differences between the telescope and the surrounding air is called tube currents. Many observatories attempt to minimize this effect by keeping the observatory air conditioned during the day to the temperature forecast for the time the observatory will open in the evening. The less of a difference between the telescope and the outside air temperatures, the less tube currents will be present. Bad seeing can also be caused by warm ground, such as hot asphalt near a telescope. Many observatories attempt to minimize these effects by planning as few roads and as little development as possible in the surrounding area. Seeing conditions are generally better at high elevations because the atmosphere is thinner, which is why most of the world's great observatories are located on high mountain peaks.

This is a gif "movie" made of 8 individual frames taken from a video of the Lunar crater Clavius. It shows the effect of our Earth's atmosphere on astronomical images. 

The Moon, especially when it is full or close to full, fills the sky with light. This extra light fills many pixels on the CCD with light from the moon, making observations of faint objects difficult or impossible. LCOGT's telescopes are often scheduled for maintenance during the few days surrounding the full Moon. Light pollution from cities also makes observations more difficult in highly populated areas for the same reason as the Moon.

Telescope Effects:

Image quality is also affected by various properties of CCDs themselves, as well as the read out process and telescope optics. CCDs are subject to variations in the sensitivity of the individual pixels as well as noise and imperfections coming from the optics of the telescope. These can cause a variety of effects such as in the images below:

Traps: Traps are pixels on a CCD that behave irregularly and cause effects that look like dark lines. A pixel behving like a trap only allows electrons to travel through it when it has a certain number itself. In these cases, a CCD will display traps when a short exposure is taken that does not allow enough photons to hit the trapping pixel. If, however, a longer exposure is used, or a star or other bright object appears over the area of the trap pixel, the pixel will behave normally and no dark lines will be seen. The dark lines on the left side of this image are the result of several traps.

Dust rings: Dust on either the filter, the window protecting the CCD, or any of the corrective optics will leave little donut shapes on an image like the one below. They appear as rings because the dust grains lie on optical surfaces above the focal plane so when they cast a shadow on the CCD, it is out of focus. Astronomers can measure the size of a dust ring and tell exactly where in the optics the grain of dust lies. Dust grains on the CCD itself leaves little dark spots. Calibration frames completely remove dust spots and rings from images.

Blooming: This effect is caused by using an exposure that is too long. As pixels fill up and can no longer hold extra charge, they spill over into adjacent pixels in the column or columns. This is a very overexposed image of Jupiter.

Diffraction Features: These features are caused by the bars that support the secondary mirror on a reflecting telescope.  The shape and angle of the diffraction features depends on the angle and orientation of these supports, which differ from telescope to telescope. In the image below, the x-shaped rays coming off the brighter stars are caused by the secondary mirror supports of the Sedgwick telescope.

Cosmic Rays: High energy particles sometimes hit a CCD during an exposure and leave bright, sharp spots or lines on an image. They are easy to tell apart from stars which cover a wider area of pixels. Some of these cosmic rays come from supernovae, black holes and other objects in the universe, and some come from the decay of radioactive atoms in the optics of the telescope. Astronomers can use computer algorithms to remove cosmic rays from thier images. They use the fact that stars and other objects have smooth edges, while cosmic rays have bright, sharp edges to identify them.

Improving Image Quality from CCDs

Astronomers use several calibration techniques to compensate for many of the irregularities mentioned above.

Dark Frames: A dark frame is an image taken for a similar length of exposure as the planned observation exposures, and the planned flat frames (see below), but with the camera shutter closed so no light can reach the CCD. This frame contains an image of the noise caused by dark current (electrons moving randomly in the CCD) and read out noise. Often the exact length of observations a telescope will be making is not known in advance, so astronomers make a master dark by averaging several long exposure dark frames with the same exposure length, and dividing by the number of seconds of the exposure. In many telescopes the CCDs are kept very cool with liquid nitrogen because dark current is mainly a consequence of heat in the CCD.

Zeros: These are also commonly known as biases or bias frames. A zero second exposure with the shutter closed records any noise in the CCD caused specifically by the read out process and the computer being used. Some of this noise will be random from exposure to exposure and some of it will have a repetative pattern. It is important to take many zeros and average them to get a useful calibration frame, called a master zero or master bias. This frame won't be able to remove all the noise, but it can lessen the effects of any pattern in the noise. With LCOGT's 2m telescopes, it takes around 22 seconds for each zero to be read out.

Flat Frames: Flat frames are images taken of a uniformly illuminated surface such as a section of the sky at twilight or a special screen made for this purpose. The flat frame exposures are timed so that each pixel will be filled to about three quarters of its capacity. This process records variations in the sensitivity of individual pixels, as well as intrinsic variability of the CCD, dust and other obstructions in the light path. It is critical to take flat fields in the same filter as the science data to be processed since different wavelengths of light behave differently as they pass through the telescope's optics and interact with the CCD.

Once all of these calibration frames have been taken, the flat frames will be averaged. The master zero (or master bias) is subtracted first, then the master dark. The resulting frame is called a master flat.

For every image taken with a CCD, a master dark is multiplied by the number of seconds of the exposure of the image. This result is then subracted from the image. The master flat is then divided into the dark-subtracted image, and the image is as accurate and free of noise patters as possible.

Astronomers often use telesopes to view very distant and faint objects. This can mean leaving the CCD exposed to the sky for many seconds or minutes, or even hours in some cases. As the Earth rotates, stars in the sky appear to move. Telescopes have to be able to compansate for this motion. Otherwise stars and other objects would appear as streaks in an image, rather than as points of light.

Equatorial Mounts

The most common type of mount for small telescopes is the equatorial mount. It works by having one axis aligned to the Earth's celestial pole, which is parallel to the axis of the Earth's rotation. It then uses a motor to to make the telescope slowly rotate about that axis at just the right speed so that when the telescope is pointed at an object, that object stays centered in the field of view.

LCOGT's 0.4 and 1.0 meter telescopes use a type of equatorial mount called a C-Ring mount.

The animation below shows a schematic of an LCOGT 1m telescope responding rapidly to a request for observations. The telescope slews to the target and then tracks it.

Altitude-Azimuth Mounts

Altitude-azimuth or altazimuth mounts are usually prefered for larger telescopes because they are easier and cheaper to build for a massive telescope than an equatorial mount would be. Altitude-azimuth mounts have two axes, a vertical and horizontal, that the telescope can rotate on. Both axes must move at variable rates to in order to compensate for the Earth's motion and keep the orientation of the field of view from rotating. This is complicated and normally done using a computer guiding system.

LCOGT's 2.0 meter telescopes use altitude-azimuth mounts.

Spectra

Spectroscopy uses the principle first suggested by Isaac Newton, that white light is a combination of many different colors of light, and a prism can be used to separate white light into a rainbow called a spectrum (plural: spectra). In 1814 Joseph von Fraunhofer shone sunlight through a prism, then magnified the spectrum. He was surprised to see over 600 fine, dark lines in the Sun's spectrum. These lines are called spectral lines.

About 50 years later, Robert Bunsen and Gustav Kirchhoff used a gas burner invented by Bunsen (a Bunsen burner) which produced a clean flame with no color of its own. They added different substances to the flame, then passed the emitted light through a prism. They soon discovered that each chemical element produces its own unique pattern of spectral lines.

When a substance is very hot, it will emit light at all wavelengths in a continuous spectrum. If a continuous spectrum of light passes through a gas cloud, atoms in that gas will absorb only photons of certain wavelengths. The rest of the light will pass right through the cloud. So in this case astronomers would see an absorption line spectrum, where certain wavelenths have been absorbed by the gas. When astronomers view a spectrum from a gas that does not have a light source behind it, they will see the wavelengths that the gas emits. This is called an emission line spectrum. Absorption lines and emission lines will be in the same place for the same gas. These lines can be explained by the behavior of electrons in an atom.

Electron Energy Levels

To understand the spectral lines they were observing, scientists needed to develop a better model of the atom than they had at the time. Niels Bohr studied the spectra of hydrogen and was the first to come up with an explanation that fit with observations. He proposed that electrons in an atom can only exist at certain distances or energy levels from the nucleus, and that because they can only exist in certain levels, electrons around the nucleus of an atom can only absorb certain wavelengths of light, which give them enough energy to move further away from the nucleus. An electron will usually stay in a higher energy level for a very short time (10^-8 seconds) before dropping down to a lower energy level and emitting a photon of the same energy it absorbed. This video from Penn State University illustrates how electrons move between different energy levels.

Spectrographs

Astronomers use spectrographs attached to telescopes to view the spectral lines of stars and other celestial objects. A spectrograph combines either a prism or diffraction grating to spread the light from a source into its spectrum. It then has a detector, usually a CCD, to record the spectrum. Astronomers use computers to analyze the spectra and create graphs such as this one from a Wolf Rayet Star:

Applications of Spectroscopy

Astronomers can compare the spectral lines they observe to the spectral lines of known elements to learn about the chemical composition and temperature of astronomical objects. They can also study the motion of astronomical objects, because objects that are moving towards an observer will have wavelengths shifted slightly towards the blue wavelengths (called blueshift) and objects moving away will have wavelengths slightly shifted towards the red wavelengths (called redshift). This is because of the Doppler effect. This principle can also be used to find stars that are binary star systems, or stars with planets orbiting them. In these cases, the star will have blueshift some of the time and redshift some of the time because the gravitational pull of the orbiting object will make the star move towards the Earth some of the time and away from the Earth some of the time.

More information

For more information on how spectroscopy works and what astronomers use spectroscopy for, please consider watching the following videos:

This is the first of a 3-part series that is a very good basic description of the use of spectroscopy in astronomy:

A fascinating 17 minute talk by Garik Israelian about what astronomers are currently discovering using spectroscopy:

Other Resources:

Very inexpensive spectrograph kits available from the Solar Center at Stanford University.

An excellent simulation of electrons absorbing photons of various wavelengths.

This animation from PhET Interactive Simulations illustrates how electrons absorb and emit photons.

• Go to the Hydrogen atom simulation.
• Click on the "Run Now!" button to start the simulation.
• In the simulation, use the selector in the top left to choose "Prediction."
• Select "Bohr."
• Turn on the power to the electron gun (click the red button on the drawing) and observe the simulation.

Radio Waves

The objects astronomers study such as stars, galaxies, quasars, pulsars, planets, supernovae and more, all emit visible light, as well as radiation that our eyes can't detect such as infrared and ultraviolet radiation. They also emit radio waves which are another part of the same electromagnetic spectrum. Radio waves have much longer wavelengths than the rest of the electromagnetic spectrum and range from several centimeters to several kilometers.

Radio Telescopes

Radio telescopes are used to study radio waves and microwaves between wavelengths of about 10 meters and 1 millimeter emitted by astronomical objects. Radio waves with wavelengths longer than about 10 meters are absorbed and reflected by the Earth's atmosphere and do not reach the ground. Many radio waves shorter than 1 centimeter are also absorbed by the Earth's atmosphere and only a few wavelength bands make it through. Wavelengths between 1 and 20 cm only experience minor distortions while traveling through the atmosphere and signal processing software can be used to correct for these effects.

Angular Resolution

Radio telescopes have to be much larger than optical telescopes because the wavelengths of radio waves are so much larger than the wavelengths of visible light. Radio wavelengths are between λ ≈ 3 km to λ ≈ 30 cm, while visible light wavelengths are between λ ≈ 4 x 10^-7(violet) and λ ≈ 7 x 10^-7(red). Angular resolution is a measure of how small details of an area in the sky can be seen. The larger the telescope, the more detail can be observed in a given wavelength.

Angular resolution (θ) of a telescope can be calculated using the wavelength of light or radio waves (λ) the telescope is being used to observe, and the diameter (D) of the telescope. 

θ = 2.5 x 10⁵ x λ/D, where θ is in arcseconds and λ and D are in meters

or

θ = 1.22 x λ/D, where θ is in radians and λ and D are in meters

So for example, one of LCOGT's 1.0 meter telescopes should have an angular resolution of approximately 0.1" when observing violet wavelengths. A 65 meter diameter radio telescope observing radio wavelengths of 5 cm would have an angular resolution of 192".

Interferometry

As you can see, the resolution achieved by a typical radio telescope at a typical radio wavelengths is not very detailed. To overcome this difficulty, radio astronomers use multiple radio telescopes at the same time, a technique called interferometry. This gives angular resolutions of 0.001" or better by effectively creating a single telescope as large as the distance between the two farthest telescopes. The light gathering power is not increased by this technique, but the angular resolution in greatly improved. The Very Large Array (VLA) in New Mexico consists of 27 radio telescopes each 25 meters in diameter, arranged in a Y shaped configuration. All 27 telescopes are used simultaneously to observe a target, then their observations are added together.

Image courtesy of NRAO/AUI 

Very Long Baseline Interferometry

The longer the distance between two telescopes, the better the resolution when they are used together. Radio astronomers sometimes use telescopes that are thousands of kilometers apart to improve the resolution of their observations. This is called very long baseline interferometry or VLBI. At such great distances, it would take too long to send information from the observations back and forth, so each telescope has its own atomic clock and records the observation. Then later the observations from the various telescopes can be synchronized and combined.

Space telescopes have the advantage of not being subject to the blurring effects of the Earth's atmosphere. In addition, there are many wavelengths from the electromagnetic spectrum that do not reach Earth because they are absorbed or reflected by the Earth's atmosphere. In fact, as you can see from the diagram below, it is only the wavelengths of visible light, a small portion of infrared, and a part of the radio wavelengths that reach Earth at all. To observe ultraviolet, x-rays, gamma rays or infrared, astronomers have had to put telescopes outside of Earth's atmosphere (normally in orbit around the Earth).

Credit: NASA/IPAC

Infrared Astronomy

Water vapor in the atmosphere absorbs much of the infrared radiation from space so the infrared observatories on Earth are located on high, dry mountains such as Mauna Kea in Hawaii. Even at high altitudes, however, the quality of observations in infrared is limited. The best solution for infrared observing is to put a telescope in orbit above the Earth and use radio to send data back to Earth. The Herschel Space Observatory was launched in May 2009 and the Spitzer Space Telescope was launched in August 2003. Herschel's primary mirror is 3.5 meters in diameter and the telescope is designed for infrared wavelengths from 55 to 672 µm. Spitzer's primary mirror is 0.85 meters across and it is designed for wavelengths of 3 to 180 µm. Infrared observatories in space must be kept very cold because otherwise infrared radiation from the telescope itself would interfere with its ability to observe infrared radiation from space. Spitzer recently exhausted its liquid helium coolant and only a few of its instruments are still being used. Herschel's liquid helium is expected to last at least 3 years.

Astronomers study the infrared wavelengths to study the early universe and to learn about objects that are too cold to generate visible light including brown dwarf stars and dust clouds.

Visible Wavelengths

Visible wavelengths make it through Earth's atmosphere, but turbulence in the atmosphere causes images of stars to be blurred and spread out by at least 0.5", even at the best observing sites in the world. The Hubble Space Telescope observes from an orbit about 559 km above the Earth at wavelengths from near infrared through the visible range and into the ultraviolet. It has a 2.4 meter primary mirror. It was put into orbit in 1990 and had a major repair in 1993. In May 2009 it was serviced again and should last until its successor, the James Webb Space Telescope (JWST), is launched in 2014. The JWST wil be optimized for infrared observation, however, and ground based observatories will be the main source of observations in the visible range when Hubble is no longer able to operate. The Kepler Space Telescope was launched in March, 2009, has a 1.4 meter diameter primary mirror, and observes in the visible and infrared range of wavelengths. It has a very large field of view and is being used to search for Earth sized planets in our galaxy.

X-ray Astronomy

X-ray telescopes make it possible to study objects with temperatures between 10⁶ and 10⁸ K (between about 1 million and 100 million °C). When atoms in a gas are this hot, they move so fast that when they collide, they emit X-ray photons with wavelengths less than 10 nm. The Earth's atmosphere blocks all X-rays from space, so space telescopes must be used to observe in these wavelengths. X-rays have such high energy that the typical reflecting telescope design used for radio, infrared and optical telescopes cannot be used as the X-rays would just penetrate into the mirror. Instead, mirrors are arranged in specially shaped tubes so that the X-rays coming into the telescope just skim off the surface of the mirror (similar to skipping a stone on the surface of a lake) and onto a detector. Two X-ray telescopes currently in space are the Chandra X-ray Observatory and the XMM-Newton.

Gamma-Ray Astronomy

Gamma-ray telescopes such as the Fermi Gamma-ray Space Telescope do not use mirrors at all and instead have special detectors to measure the energy and direction of the most energetic electromagnetic radiation in the universe, gamma-rays. The Fermi Gamma-ray Space Telescope detects gamma-rays with energies from 10 keV to 300 GeV and and has a very large field of view. It sees approximately 20% of the sky at once, and can cover the entire sky every three hours. Studying gamma-rays helps astronomers learn more about many things including active galactic nuclei, blazars, gamma-ray bursts, pulsars and solar flares.

Now that you have finished this section you are ready to try the following activities:

• How to observe using the Real Time Control Interface