CHARGE COUPLED DEVICES TECHNOLOGY

Roberto Bartali

ABSTRACT The aim of this project is the description and the application of Charge Coupled

Devices as light detectors of different wavelengths: from UV to near infrared part of the EM spectrum. The reader is firstly introduced to the technologies involved, then, follows, a detailed description of how these sensors works and their advantages and drawbacks, sometimes with the comparison to other technologies. With such a detailed description, the reader will be able to understand and select the right sensor for some specific observing purpose; even more he/she will be able to operate it correctly. This work is directed to all people who want to use the state of the art technology in image detection and wish to know, also, what there is behind and inside the black-box, called imaging device, he/she is placing on the prime focus of a telescope. Key words: CCD, Image Processing, Instrumentation.

1 - INTRODUCTION Astronomy is a science based on light detection, with this in mind we can recognize five “Periods in the History of Astronomy” each one characterized by the kind of light observed or the sensor used, but not all these periods are well delimited, like, for example, geologic eras; some times they overlap each other.

• Period 1. From prehistory to second half of XIX century. The only sensor available is the human eye. All data collected by the naked eye is recorded on cavern’s walls, monolithic monuments, stones, paper, etc. During this period the knowledge of electromagnetic (EM) spectrum is reduced to visible light.

• Period 2. From the application of photography (near 1870) to radio wave detection. The main sensor is the photographic plate and the EM spectrum is a little wider, including near UV light. The greatest collecting power of telescopes and the capability to integrate photons for hours, expand the known Universe which includes now several millions of stars and galaxies.

• Period 3. From radio wave detection (1932) to gamma and x ray detection (near 1964). The EM spectrum includes now radio waves. It is clear that celestial objects emit light of many different wavelengths, not just visible.

• Period 4. From high energy photons detection to the application of CCD to Astronomy (1974). The most energetic radiations of the EM spectrum are available to astronomers. Gamma rays and X rays emitted by stars are also observable and measurable.

• Period 5. From the CCD to now. An enormous increment in sensitivity and efficiency of the sensors expands the limit of the known Universe almost to its full size. Full EM spectrum is observable thanks to space telescopes capable to detect photons from extremely energetic gamma rays to UV, bands impossible to study from the surface of our planet due to their absorption by the atmosphere. Infrared light is almost fully blocked by the atmosphere, telescopes

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must kept very cool, placed at high altitude or on very cold places like Antarctica, so space telescopes orbiting far from the planet are a better solution.

Our actual knowledge of the Universe is based on the development and enhancement of photon sensing and measuring technologies. As more sensitive and large are sensors, more distant objects, in space and time, can be observed and studied. As spectral response of detectors increase, more detailed observations of some particular phenomenon can be performed. Silicon based sensors are now available for almost all wavelengths: from gamma rays to far IR We will describe in this work the technology of light detection from UV to near IR part of the EM spectrum.

2 – CHARGE COUPLED DEVICES (CCD) 2.1 - CCD BASICS

Charge Coupled Devices (CCD) are electronic imaging sensors based on the Photoelectric Effect (PE). Like any other electronic device they are mainly fabricated using

some semiconductor material, like Silicon (Si) and Germanium (Ge).

A semiconductor atom has four electrons in the valence band [1][2], if we can input some, and enough energy, one or more electrons are forced to jump to the conduction band, converting the semiconductor atom into a conductor one, because electrons in the conduction band can freely move. But this is possible only if there is enough energy, otherwise, after a few nanoseconds, the electron return to its normal energy state (back to the valence band). This

excitation energy depends on the size of the band-gap between the valence and the conduction band (figure 1a). The valence band is the outermost filled energy band in the atom, instead, the conduction band is the innermost energy level that can be occupied by energized electrons, so the difference between a conductor, a semiconductor and an insulator material is the width of the valence-conduction band separation (figure 1b).

To convert the semiconductor material to a conductor or an insulator we have to dope it, this is the process of introducing some impurities (atoms of different material) into it [Wikipedia, 2006]. To convert a semiconductor into a conductor, we have to dope it with

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some material whose atoms have more than 4 electrons in the valence band (called donors), normally we can use the group 15 atoms listed in the Periodic Table of Elements [PTE]. These atoms, all, contain 5 electrons in their valence band and are: Nitrogen, Phosphorous, Arsenic, Antimony and Bismuth. The extra electron (the fifth) can be moved around easily in the presence of an external electric field. Semiconductors with impurities like those listed are called n-type materials.

Conversely, to reduce the conductivity of a semiconductor, making it to behave like an insulating material, we have to dope it with some atoms with less than 4 electrons in their valence band. This way, there is an excess

of positive charge (a hole) and electrons can easily fill that energy level, so there is no electric current flow possible. These materials are called acceptors and they are listed

in the periodic table of elements as group 13 atoms. All, they have 3 electrons in the valence band, there are five atoms sharing this property: Boron, Aluminium, Gallium, Indium and Thallium. Semiconductors doped with some of the above materials are called p-type materials. The right combination of both (P and N types) in the right place (semiconductor geometry) gives us a working electronic device (figure 2).

After the above introduction about semiconductors, we can now see the structure of a CCD and how it works. Basically, the structure of a pixel in the CCD is, as we can see in figures 2, 3 and 3b, conformed by a bulk P-type silicon substrate and a thin N-type layer

above it. Another thin isolating oxide layer, separating the electrodes from the N-type silicon, prevents the trapping of electrons by the electrodes. This structure is really a small capacitor. A positive voltage on one electrode, induce an electric field which is able to create electron-holes immediately below it, holes are moved down, deeper in the P-type silicon, this way a depletion zone is generated. When photons arrive and penetrate the surface of the CCD, they can produce the so called photo-

electrons, when they are absorbed by silicon atoms. These photo-electrons are confined in the depletion zone below the positive electrode, this area is called the pixel well. Electrodes on each side of the well, are negatively biased, or with a voltage much less positive, so they can repel electrons (both, photo generated and thermal electrons), preventing their diffusion and recombination in the bulk P-type silicon. The pixel well is then a storage area and, until the chip is exposed to light, it is filled by photo generated electrons. Each pixel is isolated from it neighborough by a thin insulating region called the channel stop (figure 3B), this

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prevents the overflow of electrons from one pixel to the next, otherwise, we can loose the spatial resolution and the final reconstructed image became overexposed (blooming effect).

Photons, coming from the object of interest, strikes the surface of the sensor and are able to penetrate in the silicon up to certain depth, depending on their wavelength (figure 5). High energy photons (shorter wavelengths) are

absorbed near the surface, instead, lower energy ones (longer wavelengths), can travel more and are absorbed deeper in the silicon (figure 2). When one photon (or more) is absorbed by the semiconductor atom, the latter frees an electron. The freed electron, that is generated

in some part of the p-type silicon, is moved (by the electric field of the most positive electrode) and stored in the well. This process continues until the device is exposed to the incoming light, but there is only a limited quantity of room, depending on the thickness of the p-type silicon, the voltage applied to the electrode and to the size of the pixel. Largest pixels can accumulate more photoelectrons than small ones (figure 4). The quantity of possible electrons stored, is called the well capacity. As we have seen, only photons with some range of wavelengths are absorbed by the device (figure 5), but it depends on the kind of semiconductors used and on the physical structure of the CCD, so, that graph is only representative. We can see (figure 2) that short wavelength photons, less than 400 nm. are reflected by the surface (not absorbed) and long wavelength, more than 1000 nm., simply pass through the semiconductor (in other words, it is transparent to that radiation).

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When the exposure time ends, applying the right voltages to electrodes, at a very precise sequence, we can move all the electrons stored in each pixel well to the adjacent one in a process called shifting. This process must be repeated as many times as the number

of rows and columns are present in the CCD. Electrons are transported this way from the upper row, downside (vertical shifting) until they reach the lowest row, here they are discharged into another shift register and they are moved horizontally (horizontal shifting) until they reach the charge node where

they are measured and converted to a voltage, then they are sent to the output amplifier. This voltage is available on the output pin of the CCD for its subsequent processing. This analogue signal (voltage) is proportional to the number of electrons and must be transformed to a digital number by an analogue to digital converter circuit (ADC), some times an interface, like an emitter follower or an amplifier, is connected from the CCD output to the ADC. Now the information from each pixel can be introduced in the computer, stored and processed in order to reconstruct the image of the object. This image, thanks to the Internet, can be shared world-wide to scientific community and to general people for subsequent analysis.

There are two kinds of CCD, depending on the pixels arrangement: a single line or a matrix. A linear CCD has only one row of pixels (figure 6A) and a matrix CCD has an n by

m array of sensitive picture elements (figure 6B). Linear CCD are used sporadically in Astronomy, an example is the camera onboard of the Viking Mars Lander. Some imaging techniques like Time Integration Delay and Drift Scan can use linear arrays. For most imaging tasks in Astronomy, a matrix CCD is better.

To better understand how a CCD works, we can make a rain-bucked analogy (figure 7). Rain drops are photons, buckets are storage wells and conveyor belts represent shift registers. During the exposure time, each bucket is filled with rain drops. When exposure ends, each bucket is emptied and fills the adjacent one; conveyor belts, move buckets vertically and horizontally toward the last one, which act as the charge node and is where the water (electrons) is weighted and made available to the external circuit.

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2.2 - DEVELOPMENT OF THE CCD The silicon technology era begin with the invention of the transistor in 1948

[Massey 2005]. Soon the property of silicon to detect light, converting photons to electrons, was discovered; but very high fabrication cost of semiconductors, avoided the further

development of silicon light sensors. During the decade of 1960, some devices were fabricated for military and industrial purposes, but, their low efficiency and the need for very expensive and complex electronics, led these devices mainly as a “Technological curiosities”. The invention of integrated circuits, increase the interest in silicon based imaging devices because the complexity and size of the support electronic circuits were dramatically reduced. In the late of the 60th and in the beginning the decade of 1970, many photosensitive transistors were integrated in a linear image array, a short time after, a matrix image sensor was built. 1974 was the year when the first CCD was placed at the focal plane of a telescope, it was a 100 by 100 pixel sensor made by Fairchild [Fairchild Imaging]. In figure 8 we can see the first astronomical image ever taken (the full Moon). That year, a new revolution in Astronomy begins, like the one when Galileo Galilei in 1610 observed for the first time through an eyepiece of a telescope.

Resolution and sensitivity of the first CCD sensors were poor, if we compare images in figure 8 and 9, it is clear the difference, the image of M51 galaxy taken recently by the Hubble Space Telescope is order of magnitude better. To reach this kind of quality, a little more than 30 years of development was needed.

Astronomy is the science of collecting light from distant objects and, many physical processes can be identified and investigated because each one emits light at different wavelength. Until the advent of solid state imaging, the only way to take pictures of the Universe was using films and photographic

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plates. Physical and chemical properties of silver halide grains (used in photography), make them very sensitive to blue and near UV part of the electromagnetic spectrum [Ferreri 1977]. Photographic plates, are much more efficient than the eye because they can integrate the image during a long exposure time (figure 10). Silicon is much more sensitive than silver in the red and near IR region of the electromagnetic spectrum, this, and a better quantum efficiency (capacity of converting photons to electrons), give us the possibility to imaging, for example, hydrogen emission lines with a very short exposure time (10 to 100 times less).

2.3 - FRONT AND BACK ILLUMINATED CCD

Observing figure 5, we can see that photons of different energy can penetrate into silicon, until they are absorbed, up to some depth. It is clear that we can expose the semiconductor device to light on both sides (figures 11a, 11b, 11c). If we let the light

(represented by black arrows) to strike the sensitive area from the electrode side, as depicted in figure 2, 3 and 11a, we have a Front Illuminated CCD.

Obviously the electrodes must be transparent to the wavelength of interest and very thin. If light reach the silicon from the opposite side (figure 11b and 11c), electrodes do not interfere, in this case we have a Back Illuminated CCD. Both types have advantages and drawbacks one respect to the other, now we will describe both in some details. 2.3.1 - Front Illuminated CCD These are the first developed devices, they are also extensively used in consumer electronics (digital still and video cameras) and in industrial control, because they are less complex and cheap. A typical structure of a pixel is depicted in figures 2 and 3. With this in mind and, comparing that structure with data in figure 5, we can see that only photons with relatively long wavelength can be able to reach the photo-electrons collecting area, because they have to pass through electrodes and the isolating oxide layer. UV and most of blue photons are absorbed near the surface of the pixel or they are even reflected back. This gives to the CCD a great sensitivity on the red because the thickness of the sensitive area is typically 600 or more microns. The overall quantum efficiency is good, but not as high as the one of a Back side illuminated CCD (figure 12). Most photo-electrons collected are formed below the positive biased electrode. Low end CCD of this type are manufactured by Texas Instruments, Kodak, Atmel, Thomson, Sony,

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Fairchild among others. Professional CCD are manufactured by Marconi, SITe, EEV, among others. 2.3.2 - Back Illuminated CCD

These kinds of sensors are especially designed to improve the sensitivity in the blue part of the spectrum and to have a much more quantum efficiency (figure 12) respect to the front illuminated ones. Photons enters the CCD from the back, so they do not encounter any obstacle, but to reach the storage well, they have to travel a great distance, so the probability to recombine with holes is very high. To avoid this, the silicon substrate must be reduced by a process called “thinning” (figure 11c). This is a very expensive and difficult task, because we have to reduce to the minimum possible the depth of the silicon and it must be as uniform as possible, typically that thickness is less than 15 micron. A small difference in the dimensions means that the sensitivity and the efficiency may be variable. Most short wavelength photons are collected and stored easily, but the CCD became semi-transparent or fully transparent to long wavelength light, because those photons can pass through the sensitive area without interact with silicon atoms. The available area for photon interaction is greater than the available in a front illuminated CCD, so the quantity of photons collected is much more and results in very high quantum efficiency. Current technology can reach values of 90% or more (at the centre of the visible spectrum). This means that almost every incoming photon can release a photo-electron. These type of sensors are very expensive.

There are another two problems with a back illuminated CCD: fringes and fragility. Fringes are due to multiple internal reflections of photons and depend on the wavelength and the depth of the silicon. This is very difficult to avoid and make the CCD almost

unusable for spectroscopy. Due to the reduced depth, handling and mounting of a back illuminated CCD is difficult and must be dome with extreme care, otherwise the chip can be broken or bended.

Even taking into account all the intrinsic problems they have, this kind of CCD are the most used in professional Astronomy. Amateur astronomers and reduced budget observatories and institutions, uses front illuminated CCD due to their lower cost. Marconi, Loral, SITe, EEV are some of the manufacturers of back side illuminated CCD for professional uses. In order to avoid or reduce problems generated by thinning the CCD, maintaining blue sensitivity and improve red and near infrared

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response, two techniques were developed: deep depletion (figure 13a) and virtual phase (figure 13b). Deep depletion are essentially a back side illuminated CCD, but the sensitive silicon is thicker than the one of a thinned CCD, but less than of a front side illuminated CCD. This gives to the red incoming photons a better chance to be absorbed, the pixel well is also wider. To avoid the recombination in their path toward the storage well, a higher voltage is applied to electrodes. Blue photons can be directed easily to the storage well, by the enhanced electric field produced by the electrode. A virtual phase CCD, is a front illuminated CCD, but, instead to have three electrodes, it has just one. This reduces the blocking area encountered by photons. Blue light can, then, produce photo-electrons which can be directed to the well because they are not absorbed by the electrode structures, like in a normal front side CCD. This is the approach developed by Texas Instruments. Figure 13a and 13b, represent the spectral response of a deep depletion and a virtual phase CCD respectively. The horizontal axis is approximately of the same scale. We can see a nearly flat response from 350 to 800 nm in Texas Instruments technology, but deep depletion CCD has more quantum efficiency in the visible region of the spectrum. 2.4 - CCD AS IMAGING DEVICES

Astronomers used photographic plates as the main imaging system for about 150 years, a great effort to enhance their properties was done by manufacturers like Kodak, Agfa and Ilford. The increased necessity of higher resolving power, longer exposure times for capturing fainter objects and wider spectral response, was filled by some products like Kodak 103, III, II and technical Pan series of films and plates; FP and Delta series by Ilford and APX serie by Agfa. Results obtained were almost spectacular and plates were used until a few years ago. As an example we have the Palomar Digital Sky Survey, an almost full sky atlas that contains thousands of plates [DSS]. In Palomar Sky Survey, two different plates were used, a Kodak IIIaF for images in red light and the Kodak IIaJ for blue light. These images are in use today and are available, on the WEB to every one, thanks to the fact that they were digitised. Many new discoveries and a lot of science can be done today by the analysis of Palomar images.

There is a lot of discussion about the obsolescence of photographic plates and films in front of electronic detectors. There are many advantages and disadvantages of one technology respect to the other, but it seems that in Astronomy, CCD are the main and, in a short period of time, will be, the unique imaging sensors. In the table below (table 1), we will show the main differences between both types of imaging systems: photographic plates and CCD [Bartali 2003].

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Table 1: CCD and Photographic plates properties comparison.

FEATURE TO BE COMPARED FILM CCD Type of reaction Chemical, physical Physical

Quantum Efficiency <10% >80% Resolution 10 to 25 micron 6 to 24 micron

Pixel matrix size 1200x1800 (24x36 format) 512x512 to 8192x8192

Spectral response 350 to 650 nm

can be extended from 250 to 950 nm

400 to 900 nm can be extended from 100 to 1100

nm Linearity poor excellent

Time from the end of the exposure to image

>30 minutes (minutes to hours)

10 seconds to 10 minutes (for larger format CCD)

Dynamic range <16 bits >=16 bits (65536 gray levels) Equipment cost Low High

Auto guiding during exposure Not possible Yes Direct image processing Only after scanning the image Yes

Remotely image acquisition Not possible Yes Automatic image capturing (no

need for operator) Not possible Yes

Need for cooling system No Yes Special chemical and physical processes and procedures for

sensitivity increment Yes No need

Special environment for developing Yes No

Interferometric telescope connection capability Not possible Yes

Automatic correction of images with adaptive and active optics Not possible Yes

Automatic protection against saturation Not possible Yes

Antiblooming capability Not possible Yes Reciprocity effect failure Yes No

Loss of sensitivity if not cooled Yes but low Yes but high Binning capability Only after scanning the image Yes

Resolution increasing Not possible Yes Special chemical treatments before

exposure Yes No need

Cooling of the image support before and after the exposure Yes Not apply

Cooling during the exposure Yes but moderate (Around -20ºC)

Yes but high (-10 to – 110 ºC)

Duration of the image until degradation or unusability decades

Almost infinite (disregarding storage technology

obsolescence) Cost for sharing or duplicate

images Very high Very low (near zero)

Image acquisition Easy (few steps)

Complex (many steps)

Image processing requirements Moderate Very complex Exposure time to reach the same

limit magnitude Very long

(many Hours) Short

(Minutes to a few hours)

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As we can see in figure 14, there is a great difference between a photographic plate image (figure 14-a, 14-b), a CCD image from Earth surface (figure 14-c) and a CCD image from space (figure 14-d) of the same object.

Imaging with a CCD is not a straightforward task. Until the final image is ready to be analysed or printed for the public, many intermediate images must be acquired.

First of all, we have to take a Dark Frame; this is an image with the shutter closed and with the same exposure time as the science image and at the same temperature, this way only the thermal generated signal is available. To have a better figure of the dark (thermal) signal we have to take several images and average them; this is called a Master Dark. A good Master Dark is not necessarily made each night if the telescope and CCD conditions do not change. If the CCD is cooled to very low temperatures (100ºC below zero) dark current is so low that is negligible, so there is no need for a dark frame, but at higher temperature we must have it, because the dark current level is important (figure 15).

A zero exposure time image, called Bias Frame, is then taken. In this image we have all the electrons generated by the internal electronics,

electroluminescence( figure 24a), pixel defects like hot spots (figure 21) and black pixels (figure 22), bad columns and cosmic rays (figure 16, 25). A good Bias is the average of several images (master Bias). A professional grade CCD shows an almost uniform bias frame (figure 16a), instead a less quality CCD shows many defects (figure 16b).The Flat Frame is an exposure of a very uniform illuminated source. This can be done by taking a picture of an evenly illuminated screen inside the dome, a light box placed over the telescope, the twilight sky or a starless patch of the sky [Bartali 2005]. Exposure time for the flat frame must be short enough to optimize the very expensive telescope time (we want to take as

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many as possible science images, not spend all the time taking calibration frames), but long enough to reach almost a 30 to 50 % of the full well capacity (half saturation). The

temperature of the CCD during flat frame imaging must be the same as that of the raw image. Flats must also be corrected with dark and bias frames. Averaging a set of Flats is a normal technique to have a Master Flat. Flat frames should be taken each observing session. Flat frames shows basically the difference in sensitivity of pixels, all defects due to dust in the optics (CCD and

telescope), vignetting and fringes (figure 17a and 17b). Finally, the raw image is the exposure of the object of interest (figure 18). It must be taken with the telescope and the CCD at the same condition of temperature and exposure time of the auxiliary frames: dark,

flat and bias (remembering that bias is a zero exposure frame). Some times to reduce the noise, several exposures (with much less exposure time) of the same object are averaged

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together in a technique called Stacking. Now, all the images are stored in the memory of the computer and are ready to be processed. As we can see in figure 20 we have firstly to correct the raw image and the flat frame. To do this, we have to subtract dark frames (or Master dark) and bias (or Master bias) from the Flat (or Master Flat) and form Raw frames. Then, divide the resulting corrected Raw by the corrected Flat. The resulting image (figure 19) is called a Science Image and is theoretically free of defects and ready to be further analysed. A great difference is shown between raw and science images

Even when, at first sight, it seems to be a very difficult task, taking an image with a CCD is relatively easy due to the increased (and increasing) computer power and automation technologies. Today images can be taken and retrieved, just by sending a command to a remote control computer and the telescope camera (remotely operated and robotic telescopes.

2.5 - SOURCES OF CCD ERRORS Until now, we have talked about how good is a CCD, but we have to talk, also,

about the many problems arises when we want to uses them. Thanks to the digital form of the image, all errors can be corrected using specifically developed software routines.

There are many sources of problems that arise when using a CCD, some are intrinsic, derived from the fabrication process and some others are due to external factors during the operation. Now we will explain both cases and we also give the better, or possible, solution.

2.5.1 - MANUFACTURING PROCESS INDUCED ERRORS It is nearly impossible to manufacture a perfect electronic component due to some impurities in the material or in the atmosphere where the crystal is manipulated, machine imperfections, and many other factors. This is especially true for CCD. In the following list, we can explain these possible problems and the way we can try to eliminate or reduce them to minimum. Problem: Sensitivity difference from pixel to pixel.

Pixels are not identical in their spectral response nor in sensitivity (figure 17a,b), this is because of the many procedures during manufacturing process. Doping each pixel with exactly the same amount of atoms, growing the crystal uniformly is not as easy as it may appear. During fabrication, many chemical etching phases are needed, so if a very little error or imperfection arise, it is transferred and perhaps, incremented at each step. Solution: If the Flat Field exposure is done properly, we can virtually eliminate non uniformity, but this is true only in theory because it is very difficult to obtain a perfect Flat. Practically we can reach a satisfactory unification of spectral response and sensitivity.

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Problem: Hot spots Some pixels can be near their saturation value, or at some specific value, even when

they are not exposed to light, in this case we call them hot spots; they appears in the raw image as bright dots like stars (figure 21-a). They appear in all calibration frames and, of

course, in the raw image. If a star light falls exactly over a hot spot, no usable information is available. Possibly, a certain number of photo-electrons increase the level of that pixel, and, in this case it is possible to rescue some information, but it is not reliable. Bright

columns can also appears (figure 21-b) produced by the leakage of charges during vertical shifting, some electrons are left behind and increase the charge accumulated in the pixel well of the row below the defective pixel. Solution: subtract or setting to zero the value of that pixel. If we know that a star fall over that pixel, a second image with a slight position shift is taken and subtracted from the other. Problem: bad, buried or black spots

Some individual pixels or a group of adjacent ones, called bad pixels, can never be used to collect photo-electrons because they have no sensitivity, nor the capacity to collect photo-electrons. There is also the possibility that the well is not formed because the electrode is not connected to the rest of the circuit or its resistivity is higher than normal, so the positive voltage (applied to electrodes) is lower than the required. If a star falls over a black spot, we will never see it in the image. If we have a black pixel, the final image shows a black streak (dark column) starting at the defective pixel position (figure 22), because of the vertical shifting (downward). As of the

above case, this defect is visible on every image (raw and calibration), so it is easy to locate and cancel. Solution: Taking a second image of the same field, but with some position shift, help to cancel the effect, after image subtraction.

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Problem: Blooming If photo-electrons number are near the full well capacity, there is the risk of an overflow of some of them into adjacent pixels because the insulating regions (the separation between pixels) are very small, this effect is known as blooming. The raw image contains vertical white streaks. The length and wide of the streak depends on the exposure time and on the brightness of the star, both can produce too many photo-electron and an overflow occurs (figure 23). Tis overflow can extends over many pixels. Solution: reducing the exposure time to the minimum possible to have a good signal to noise ratio. Taking more images of the same field with some position shift and then subtract one from the other, this way objects hided by the

overexposed or bloomed pixels can be visible again. CCD with great well capacity are less affected by this problem. If, for example, a very faint object is too near a much brighter one, like a satellite close to the parent planet, two sets of images must be taken. One, over-exposing the planet in order to enhance the visibility of the satellite and a second, with a much shorter exposure, to pick up planet details (surely satellites are not visible on this image), then stacking both images, give us all objects visible. Problem: Transport inefficiency After the exposure, the pixel well content must be shifted (transported) vertically downward until it reaches the horizontal shift register, from where it is horizontally shifted toward the charge node. But, some imperfections in the geometry of the CCD, impurities in the semiconductor materials, non uniformity of electrode resistivity, among other factors, create a non perfectly uniform displacement of photo-electrons from one pixel to the next. Some are left during each transfer stage. Some photoelectrons can also be captured by atoms and a recombination occurs. This is more visible if the well potential is near its full capacity or the well is too large, because lateral photoelectrons can easily recombine. A professional CCD has a 99.9999% of transport efficiency. Solution: unfortunately there is no possible solution for this problem until we buy a better device. With a small CCD, the effect is almost negligible, but in large CCD it can be a serious problem. To avoid the effect, and to slower the readout time, some manufacturers opted to divide the CCD into two or four symmetrical areas. Each one has its proper horizontal shift register, this way each pixel can travel only a 25% of the full path. Problem: Glowing There are two main causes for glowing in a CCD: radioactivity and heat produced by the output amplifier. The former is because some material used to manufacture the CCD or the glass window, may be weakly radioactive, so this radioactivity can induce the release of some photoelectrons producing a glow in some section of the image (figure 24-a). The position of the internal output amplifier is very important, because it is one of the parts which is working at higher temperature due to it complexity. If it is placed under, or too close to the imaging array, it can produce more dark current on those closer pixels (figure 24-b).

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Solution: Calibration frames can eliminate this problem, dark frame shows glowing like the raw images, so, when they are subtracted, glowing disappears. Unfortunately, if it rise because of radioactivity, it is not always visible. Another solution may be to turning off the CCD and wait a few seconds, then, turn on again and re-expose the image. For this reason, manufacturers place the

output amplifier as far as they can from the image area. Turning off the output amplifier during exposure time (procedure available for most CCD) eliminate the problem if it is thermal, but, if is produced by material radioactivity, there is no way to anticipate when and where it can be appears. Problem: Reflection CCD electrodes are metallic and metals are very reflective, even when they are treated to be semitransparent to let photons been able to pass through, so some quantity of photons can never reach the sensitive area because they are reflected off the sensor. The protecting glass window over the sensitive area is also reflective. Solution: antireflective coating over the glass windows and electrodes, help to reduce the number of reflected photons; this is done during the manufacturing process. Problem: poor blue and UV sensitivity High energy photons, like UV and blue ones, are absorbed near the surface of the CCD, in other words they can not travel too deep inside silicon. In a front illuminated CCD the well is placed under the electrodes, so photons, in order to interact with silicon and release photo-electrons, must pass through the electrodes. Most of them are absorbed before they can reach the usable sensitive area. Solution: If the application needs a high blue or near UV response, the solution is to select a thinned, a deep depletion or a virtual phase CCD. Another solution it is to coat the device with some material that absorb short wavelength light and re-emit photons at longer wavelengths. Normally UV is converted to yellow-green light, so these photons can be collected by the CCD, because they falls in its most sensitive region. Some material used are Coronene and Lumigen [Kitchin 1998, Howell 2000]. Problem: transparent to IR radiation: Long wavelength photons, like IR ones, are absorbed very deep inside the CCD. These lower energy photons can also pass undisturbed through the CCD, making it completely transparent above some wavelengths, like in thinned CCD (figure 2, 11) . The sensitive area must be thick but, if the photon is absorbed too much far from the well, the probability for the new released photo-electron to be absorbed by other atom is very high an then it can never reach the well, so it is not stored. Solution: Thick or front side illuminated CCD are best for imaging red objects, but, as we saw above, they have poor response to blue light. If the application asks for a better flat

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spectral response, a compromise (to have both responses, in the blue and the red) is to use a deep depletion CCD or a virtual phase device. Problem: Internal reflections (fringes)

Thinned devices suffers from multiple internal reflections at some wavelength. This is because the thickness of the silicon substrate is proportional to the wavelength of the incident light (figure 17-b). Solution: Flat field frames can help to reduce this effect, most of the time they are able to cancel it. If the CCD is to be used for spectroscopy, the resulting image is modulated and this interference pattern occult spectral lines, so it is not useful for that application. If the device is used for spectroscopy, it is better to select a front illuminated CCD. 2.6.2 - OPERATION INDUCED ERRORS

During the operation, imaging exposure, read out process, camera assembly, mounting and maintenance, we can have a series of problems as we can see in the following list. A possible set of solutions is also given. Problem: Dark current Atoms in the CCD, like in any other device, are bouncing because they are not at the temperature of the absolute zero (-273.15ºC). This vibration is the cause of some heat generated by friction. During collisions, some electrons are released and they can be attracted by the positive potential of the well, this way they are stored and mixed with photo-electrons generated by the radiation coming from the object of interest. There is no way to identify and separate ones from the others, so the raw image is the sum of internally generated electrons and those generated by incoming photons. The number of thermally generated electrons (called dark electrons) is proportional to CCD temperature, environment temperature and exposure times. The raw image show some (not necessarily) uniform background illumination that masks fainter objects (figure 15). Solution: The only way to reduce to the minimum dark current is to work at lower temperature. This implies the addition, to the CCD, of a cooling system, either a thermoelectric or liquid Nitrogen ones. But there is a lower limit for cooling (approximately 110ºC below zero); at lower temperature, electron mobility is greatly reduced, so accumulation and transport process is affected. Problem: Cosmic rays

Cosmic rays are very energetic particles (mainly protons) generated by stars in an enormous quantity. Because of their high energy, they can pass through almost any material loosing just a part of their energy. When they reach the sensitive area of the CCD, they can produce an avalanche of electrons which are stored in the well. This way they are added to those generated by incoming photons, but there is no way to distinguish them from ones produced by a cosmic ray. The raw image show white spots like stars (figure 25). Where and when cosmic rays strikes the CCD is absolutely unpredictable. The number of collected cosmic rays is proportional to the exposure time and the CCD area.

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Solution: The analysis of the Point Spread Function (PSF) of the image can show if it is a real star or a cosmic ray effect. PSF of a star tend to be Gaussian, instead a cosmic ray has a sharper profile. The probability to have cosmic rays at the same position in all frames (calibration and raw) is nearly zero, so during the processing phase, they can easily eliminated. If a real star falls under a cosmic ray, the only way to see it is to re-expose the image. Problem: Smear or blurring

If the CCD is not covered by a shutter after the exposure time ends, during the time needed for the transport of charges, incident photons continues to strikes the sensitive area, so more photoelectrons are added to each pixel. But previously accumulated photo-electrons are not in the same place, because they are now in the shifting process, so the image is distorted, smeared or blurred. This effect is proportional to the size of the CCD because larger is the pixel quantity, higher is the readout time. Solution: Closing the shutter immediately after the exposure time elapsed, eliminate this problem. Some

manufacturers divide the CCD into two separated regions, one is used to collect light and the other (masked by an opaque layer of aluminium) is used as a temporally storage area. When the exposure ends, the image is shifted at high speed from the collecting area to the protected area, and then, at slower speed, it can be shifted out from the CCD. During this time, a new image can be integrated in the collecting area. This configuration is called frame transfer CCD (figure 26). There are two other solutions: shifting at higher speed, but this increase the noise, or use devices with more than one horizontal shift registers. In this case the CCD is divided into (for example) four regions, each one with its proper shift register. Each fourth part of the image is read simultaneously, so the time needed for the read out process is reduced four times. Problem: Over and under exposition.

If we over expose, we can easily reach the full well capacity of the pixel (saturation) this way the image is just a white spot. Also, an overflow to nearest pixels is possible. Fine

details of an extended object like planet surface are occulted (figure 27). Under expose, means that the image has not enough dynamic range, so low contrast details are not visible. The worst case is when we want to take a picture of a faint object near a much brighter one, this is a very difficult task. We can not have both at the same time on the same frame. Solution: In this case, we can expose two times the same object, the first with an extended period of time, to capture fainter object, and, the second with a much

shorter exposure time, to capture the brighter object, late, during the image processing phase, both images are added (stacked) together. To avoid over or under expositions, we can take some test images.

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Problem: Focusing Focusing an image on a CCD is not an easy task; any small deviation can produce poor quality images. The full resolution of the optical system can not be achieved and some precise observation like those needed for photometry and astrometry can not be done. If we are exposing the CCD with the addition of filters, the correct focus is not necessarily the same, due to the transmission properties of the filters and its mechanical construction (filter thickness) and mounting. Solution: Taking test images and process them until the best focus is obtained. We have to re-focus the system each time we use a different filter, in addition, all calibration frames must be taken for each filter. Several methods are used for focusing purposes, but, for all, the process is the same: take a short exposure, process the image and move the focuser until the sharpest image is obtained. If we know haw much the system is unfocused, we can apply some mathematical algorithm to correct it for. Problem: Over sampling and/or undersampling.

If the star image is spread over too many pixels (figure 28) its brightness is lower than the real and it is difficult to determine the exact position of the centroid, because too many pixels in the centre are at the same value.

The opposite, of the above, is when the image occupy just one or two pixels, no available information is

obtained from the image. The bright point may be a star or just a cosmic ray hit. Undersampled image is when, a less than necessary, number of pixels are covered by the star image (figure 29), so the star is not a circle but an irregular staircase. Atmospheric blurring can also cause an over-sampled image. Solution: Depending on the field of view, pixel size and exposure time, it is necessary to obtain star images with almost nine pixels in order to obtain a Gaussian like point spread function. A correct sampled image of a star can be 2 or 2.5 times the pixel scale in order to obey the Nyquist criterion. Problem: Mounting

CCDs are very fragile devices, they must be manipulated with extreme care, especially large format thinned ones. During mounting, a small bend can change their optical and electronic properties, in the worst case, a prolonged fatigue, can produce a rupture. If they are not mounted correctly, the telescope focal plane may be not remain parallel to the CCD surface, so the image can be out of focus and a point like star will be seen as a comet. Solution: mounting and manipulating with extreme care and only by skilled people. Problem: Dust

Dust particles can be present over all optical surfaces, from the primary mirror to the CCD window. It is impossible to avoid them. In the flat frame image they are like unfocused donuts with different diameter depending on their size (figure17-a, 17-b).

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Solution: Flat field frames show dust particles, but they disappears during image processing. Problem: Vignetting

This is a very common aberration in astronomical images, it consists of a progressively darkening from the centre to the corner (figure 30). It is caused by the

obstruction of the image cone, formed by the objective lens or the primary mirror, by some mechanical part (filter, focuser, mirror support, CCD support, etc), so the full image is formed on the focal plane. It is also caused by the poor alignment of the CCD respect to the focal plane. If the CCD size do not match the dimension of the image on focal plane, vignetting appears. Solution: first of all we have to know or calculate the image size on the focal plane of the telescope, short focal length produce smaller images than long focal

lengths telescopes. Second, we have to know the exact shape of the image cone, but this is difficult in large telescopes due to the great quantities of secondary and auxiliary mirrors presents. Also, each of these mirrors and they support, must match the size of the image cone. Filters dimensions and filter support can play a significant role in vignetting, so they have to be placed as close as possible to the CCD in order to avoid any light obstruction.

3 - COOLING SYSTEMS

Dark current (electrons generated by thermal excitation) in CCD is a limiting factor to reach very faint magnitude, because it represent a bias level added to normal sky background. As we saw earlier, the only way to reduce or eliminate as much as possible the dark level, is to reduce the temperature of the sensor. There are two main techniques for obtaining this: thermoelectric and liquid nitrogen cooling. The former is much less expensive but it is able to reduce the temperature only by 30 to 50 degrees below ambient. The latter can reach temperature of 100ºC and 110ºC below zero.

3.1 - THERMOELECTRIC COOLING This is a technique based on the Peltier effect discovered more than a century ago (figure 31). An electric current applied to the Peltier junction increase the temperature on one side and reduces it to the other side. The cold side is attached to the CCD, so the sensor can be cooled. To avoid the destruction of the Peltier device, we must extract as much heat as possible from the hot side. This is done, in the simplest way, with a cooling fan, but a high efficiency heat sink and forced ventilation works much better (figure 32). If the temperature differential is pushed near the limit of the capacity of the device, a liquid heat exchanger can also be applied, so water flowing help to take the heat off the sensor and the Peltier device. Normally with a Peltier cooler we can reach 30ºC below zero, to have more cooling, a technique called stacking is used. This is when we place the cold side of one module in contact to the hot side of a second Peltier junction, this way it is possible to lower the temperature to some 50 degree below zero (figure 33).

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The electronic system interface must be able to maintain the desired temperature as constant as possible (0.1ºC or better) during the exposure time.

This cooling method is simple in operation and installation, we only need a programmable current regulated power supply, a temperature sensor placed as near as possible to the CCD and a control electronic circuit. Power consumption and overall dimensions are small. Due to the relatively low cost of such a system, it is mainly used by amateur astronomers. 3.2 - LIQUID NITROGEN COOLING

Modern giant telescopes require very sophisticated ancillary equipment to get results as good as expected from a so much expensive instrument. CCD sensors for professional Astronomy are very sensitive and, need to be cooled to the lowest possible temperature to avoid dark currents. But, due to semiconductor properties and the atom behave, the lowest temperature at which we can cool a CCD is 110ºC below zero. At lower

temperature the mobility of electrons are greatly reduced, so they can not be transported from one pixel to another, efficiently, by shift registers.

Nitrogen is in gaseous form at ambient temperature, but at very low temperature it became liquid. A very expensive and sophisticated electric and hydraulic system is needed to operate a liquid Nitrogen Cooler (figure 34). The nitrogen is placed in a reservoir thermally isolated from the outside environment and it is placed inside a chamber

maintained in vacuum. This chamber is made of glass or aluminium. Cold is conducted to the sensor by a copper wick. Sometimes, the external circuit is also contained and cooled inside the dewar to prevent noise generation. A tube, with a two way valve, is used to both

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refill and exhaust. Heat generated by the sensor, is enough to boil nitrogen and it is then expulsed through this tube. A glass window, not in contact with the CCD, protect the sensor and prevent the formation of frost. A separated chamber, between the glass window and the CCD is held at very low pressure, near a perfect vacuum. To maintain the temperature at constant level during exposure time, there is also an electrical heater (normally a resistor), this way, temperature variations are limited to +/- 0.1ºC or less.

The process of lowering the temperature take some time, no more than a few degree per minutes, otherwise a mechanical stress, due to different expansion coefficient of materials, can destroy the sensor. The CCD must be maintained as long as possible at low temperature, precisely to avoid stresses, but this imply the constant refill of liquid nitrogen in the dewar, normally every 24 hours. The following table (table 2) summarizes characteristics of both cooling systems. Table 2: Cooling system comparison. Characteristic to be evaluated Thermoelectric cooling Liquid Nitrogen Cooling

Temperature reached 10ºC to 50ºC below ambient 100ºC to 130ºC below ambient

Temperature resolution and stability

0.1ºC to 0.01ºC Depending on the control

electronics

0.1ºC to 0.01ºC Depending on the control

electronics. Time to reach lowest operation

temperature Less than 15 minutes More than 1 hour approximately

Main application Visible spectrum observations UV, Visible, near IR and IR spectrum observations

Cost Low, less than 300 US$ High, thousands US$ Implementation easy complex

Portability yes no Main Use Amateur and small observatories professional

Overall Dimensions Small, less than 2 cubic decimetres Bulk, tens of cubic decimetres

Maintenance Virtually free periodic Weight Low, less than 1 kg High, tens of kg

Auxiliary equipments fan, heat sink, power supply, control electronics, sensors

Vacuum chamber, pump, sensors, control electronics,

hydraulic and pneumatic circuits Operation simple complex

Power consumption Low, less than 300 W High, in the order of kW Installation Easy Need a skilled technician

Sensors used Current, temperature Temperature, vacuum, flow, current

Control electronics Microcontroller based (one or two cards) Several microcontroller cards

Remote controlled operation possible Possible with some restrictions Failure probability Very low possible

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4 – CCD APPLICATIONS In this section we will describe the main features of CCDs specifically designed for the detection of three wavelength range, in this case Visible (400 to 700 nm), Infrared (800nm to 24 micron) and Ultraviolet (10nm to 380nm). 4.1 - CCD APPLICATION FOR VISIBLE SPECTRUM CCD spectral sensitivity in the visible and near IR spectrum is very high and makes them ideal for normal photography and obviously for astronomy. Hydrogen atomic

emission falls in the red, so this is a great advantage over photographic plates and films. Their application on the focal plane of a telescope is

straightforward, because no needs of special filters or coatings over the

optics are needed. But if we have to produce near true colour images of celestial objects, a set of filters that matches the eye perception is needed. In Astronomy a special set of filters called UBV and UBVRI are used and their spectral response are depicted in figure 35. With these filters it is possible to do precise photometry and crate colour magnitude diagrams, fundamental for calculating temperature and evolution of stars. Unfortunately, not all manufacturers offers filters with identical response. In figure 36 we can appreciate the eye spectral response and we can easily compare it with a typical CCD response (figure 10, 12, 13a, 13b).

From the first 100x100 pixel array, some 30 years ago, we have now devices with 100 million of photosites (figure 37) in a sigle chip. This CCD offers a sensible surface of 95 by 95 mm. These kind of devices try to solve

one of the most restricting problem of their application in modern large telescopes: extremely large field of view.

CCDs pixel size, ranging from 4 to 24 microns wide, let the user to obtain images of different resolutions, small pixels are useful where very fine details, like planetary surface or double star measurements, are needed. Instead, for large and wide field survey of

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galaxies and to search for asteroids, mid or large size pixels can be used. Large telescopes have an image area, measured at the focal plane, of several centimetres wide. To get full advantage of that, for wide field imaging, we have to form a mosaic of several devices (figure 38), but there are, many times, restrictions in their fabrication, because of the internal control circuit position, so, at most we can have a three size buttable CCD. In other words, a mosaic camera offers the possibility to have almost any number of CCDs and billions of pixels available, but the resulting image is not continuous, it shows gaps (figure 39). To avoid these gaps it is necessary to take a second image of the same field, but with the camera rotated 90 degree or the telescope position slightly shifted from the original one. This imply to double the time needed to obtain the image and, a lot more processing time. Computer memory, processing software and power needs are also incremented. 4.2 - CCD FOR IR SPECTRUM APPLICATION IR Astronomy is a relatively new science because infrared light is almost fully blocked by the atmosphere, just a few windows in the spectrum are available. But, to observe these bands, a special detector and telescope technology shuld to be implemented. Heat emitted by the environment and instrumentation can be higher than IR radiation we receive from stars, so we have to cool, as much as possible, every part of the telescope (including detectors and electronics) and to observe from high altitude places or very cold ones, like the Antarctic continent. The better place to do IR astronomy is obviously from space, where low temperature and no blocking atmosphere are present. Telescope optics for IR observations are different from the used for the visible. For visible wavelengths mirrors are coated with a thin layer of aluminium and, to prevent oxidation, a layer of quartz is deposited over the aluminium. Optics for IR are coated with different materials, depending on the wavelength of interest, cost, performance and application. Some of this materials are [II-VI Inc, 2006]:

• Zinc Selenide (ZnSe) • Zinc Sulfide (ZnS) • Germanium (Ge) • Gallium Arsenide (GaAs) • Silicon (Si) • Cadmium Telluride (CdTe) • Copper (Cu) • Aluminum (Al) • Molybdenum (Mo) • Gold (Au)

Gold and copper are the most used for they high reflectivity in the IR. Observing in the IR is very important because we can see different phenomena, like protostars in the development phase, protoplanetary disks and in general all objects occulted inside and behind thick dust nebulae. This is because shorter wavelength photons can not pass through dust, but they produce the heating of dust particles when they strike them. This way dust particles reemit photons in the IR spectrum (figure 40). Dust act like a screen and let us see the scene behind it in indirect form. Cool and low mass objects like red dwarf stars, methane dwarfs, brown dwarfs, hot Jupiter exoplanets and, in general, all objects with temperature lower than 2,000ºC emits photons with too low energy to be seen

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in the visible (blackbody peak emission depends on the object temperature, lower the temperature, longer the wavelength emitted). CCDs for IR observation are more difficult to build because different materials,

other than Silicon, are needed in order to have the required spectral sensitivity. Another difficult consists of their capability to trap photons with much less energy than those of visible light. For this reason, until a few years ago, IR CCDs were built with a reduced number and relatively large size pixels. Today we have IR CCDs large enough (but less than those for the visible spectrum) and we

can build mosaic cameras like the one depicted in figure 38. For visible light observation, CCD temperature must be maintained as low as possible to avoid dark current generation and it is independent to the wavelength observed, but, for IR detection, the working temperature depends also on the wavelength and the sensitive materials used. We will briefly describe most common material used and their performance for IR CCD arrays [Hoffman, 2004]. IR sensors are made in hybrid form, they have some electronics circuit inside, normally the read out interface, to avoid the increment of noise. This approach is similar to CMOS imaging sensor technology, widely used in consumer and industrial applications, but not used in professional Astronomy (with some exceptions). 4.2.1 – “Si PIN” DETECTORS

These are broad band detectors operating from the visible to the near IR with the highest quantum efficiency in the IR. Even when they are hybrid sensors, the fill factor is nearly 100% due to the position of electronic circuits outside the sensitive area. They are radiation tolerant, better than normal CCDs. Large available format are

1024x1024 pixels. In figure 41 we can see the spectral response (figure 41-a) and an example of these sensors (figure 41-b). These detectors can work at relatively high temperature (100 to 300 degrees Kelvin).

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4.2.2 – “HgCdTe” DETECTORS Some of the better, and most used, materials for IR detection are Mercury (Hg), Cadmium (Cd) and Tellurium (Te). Depending on their cut-off wavelength they are classified as short wave and mid wave detectors.

• Short wave: their spectral response is very uniform from 0.85 to 3.2 micron with sharp cut-on and cut-off curve; they reach a 70% to 80% of quantum efficiency, but if coated with antireflective coating

they can reach up to 95%. Modifying the quantity of Cadmium, Zinc and Tellurium in the substrate, the cut-on wavelength can be extended to the visible spectrum. The operating temperature is about 100ºK. In figure 42 there are a graph of their spectral

response and an image of them. • Mid Wave: their spectral response is very

uniform up to 5.2 micron with sharp cut-off curve; they reach an 80% of quantum efficiency, but if coated with antireflective coating they can reach up to 95%. The operating temperature is about 70ºK. In figure 43 there is a graph of the spectral response of a CCD made by Raytheon.

4.2.3 – “InSb” DETECTORS Indium (In) and Antimonium (Sb) based sensors are widely used for IR light detection. They works at very low temperature, about 30ºK. The advantage of these material is that they are sensitive over the full visible and the IR, up to 5 microns (figure 44-a). If coated, with one or several layers of antireflective materials, their quantum efficiency can be very high, near 100% at 1 micron. In figure 44-b there is an example of a 2k x 2k pixels sensor. To maintain low dark current and noise, the module is mounted on a metallic pedestal and several capacitors and resistors are bonded on the board itself. This kind of sensors were developed for NASA IR SPITZER space telescope.

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4.2.4 – “Si:As IBC” DETECTORS Detection of the longest IR wavelength up to 25 micron (figure 45-a) needs a very different combination of materials like Silicon (Si), Arsenic (As), Boron (B) and Carbon(C). Noise and operating temperature must be maintained extremely low, their

working temperature are about 10ºK. Current state of the art sensors arrays are made of 2k x 2k pixels (figure 45-b), but the largest array tested is of one million pixels (Raytheon). Their spectral response is very linear and reach 80% with one layer of antireflective coating (solid line in figure 45-a). The dashed line

represent the spectral response without the antireflective coating. This kind of sensors are selected for the James Webb Space Telescope and are operating in the SPITZER space telescope. To avoid interference from heat generated by the telescope, control electronics, motors, etc, a telescope for the IR must be maintained as cool as possible; the air inside the dome and the telescope structure must be maintained at the same temperature of the air outside the observatory, to do this, several hours before observing session, a series of ventilation windows on the dome are opened and closed by a control computer. 4.3 - CCD FOR UV SPECTRUM APPLICATION

Fortunately for our body, but not so much for astronomical observation, UV light is almost completely absorbed by the atmosphere (in the ozone layer), so the only way to observe physical phenomena that produce short wavelength light is from space. But, observing from space, there are many factors that reduce the life of the detector and related instrumentation, degrading, with time, their performance, the most important of that is radiation.

Exploring in the UV, we can understand how stars and galaxies forms and their evolution. Temperature, chemical composition and density of the interstellar medium and study of hot young stars are also a field for UV Astronomy. The combination of IR and visible spectrum observations, with UV, give us a full vision of the evolution of the Universe.

Only radiation around 300 to 400 nm can reach the surface of our planet, but UV Astronomy needs the detection of the full UV range, from 10nm (extreme UV) to 380 nm (near UV). UV spectrum is divided into 4 bands:

• Near: 320nm to 400nm. • Mid: 200nm to 320nm. • Far: 91.2nm to 200nm. • Extreme: 10nm to 91.2nm. Detection of short UV photons is very difficult because they are absorbed by silicon

near its surface, so front illuminated devices are impossible to use and back illuminated CCD have several problems like fringes, so a different approach, with hybrid technology is

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used. The best technology available today, is the so called microchannel detector. These detectors are an hybrid between a CCD and a photomultiplier tube, taking the best from each one. As we saw in the above section (about IR detectors) not all wavelengths can be detected with the same material, so we cannot have a universal sensor for the full UV spectrum, but several ones made of different semiconductor materials. 4.3.1 - SENSORS FOR NEAR UV Photon penetration in silicon based CCD, depends on their wavelength (energy), near UV photons can be detected with back illuminated CCD, deep depletion CCD and Virtual phase CCD, but for all of these technologies, the quantum efficiency is low at shorter wavelengths. The addition of some coating substances like Coronene and Lumigen can improve the probability of detection because they act as light wavelength down converters. These substances absorb UV light and emit photons with much lower energy, normally in the range of 500 to 550 nm. Quantum efficiency of CCD at these wavelengths

is very high, but we have the problem of their high sensitivity to visible light too, so UV photoelectrons are mixed with yellow-green ones produced by the visible component of the incoming light. This is the major drawback, because we want to sense UV radiation alone. A better solution is to use sensors made of Silicon Carbide (SiC) which are completely blind,

in other words transparent, to all wavelengths greater than 380 nm, making them more suitable for UV (figure 46). This property became from the fact that the band-gap of SiC semiconductors is wider than the Silicon (Si), Gallium (Ga) and Arsenide (As) semiconductors. Due to this, we need to give much more energy to excite an electron and send it to the conduction band. Visible and IR light photons can not excite electrons in the valence band of SiC semiconductors because they have not enough energy, instead, UV photons carry the right quantity of energy. SiC semiconductors can, also, operate at higher voltage, temperature and radiation levels, these last two qualities are very useful for space based telescopes. The sensitivity of SiC sensors in the near UV is 10 thousand more than the one of Ga-As or Si sensors. 4.3.2 - MICROCHANNEL TECHNOLOGY SENSORS FOR MID AND FAR UV These are sensors based on the electron multiplication effect, they can be fabricated using vacuum tubes or semiconductor technologies. For Earth based observations, both are used, but for UV space based Astronomy, the latter is the best choice due to their lower power consumption and reduced size.

To understand how they works, it is better to explain the channel vacuum tube multiplier. To find a more detailed description about photomultiplier tubes, there is a complete section devoted to them, below (section 5).

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UV photons strikes the sensitive material of the photocathode, which in turn, release a photo-electron (figure 47). This photo-electron enter into a narrow and curved semiconductive channel. During its path, toward the anode, it impact the inner

curved walls of the channel several times, every time a bend is encountered. After each impact, more electrons are released due to a near total internal reflection effect. Each electron, hit the wall at the subsequent bending, and so on. This process continue until all electrons are trapped by the anode electrode, which is held at a high positive voltage. This electron avalanche has a gain up to 100 million, depending on the number of bends in the channel. The quantum efficiency is low (< 30%), but sensitivity is very high. The channel

can be made very thin, so if we place several channels, side by side, we can detect not only photons coming from a unique specific point, but we are able to recreate the image of the emitting object. A device with this characteristics is called a microchannel array (figure 48). A semiconductor microchannel

array consist of a thin plate pierced with thin holes (channels) and each channel is, at all practical effect, a miniature photomultiplier tube with a diameter of about 25 microns. The top surface is maintained negatively charged, with respect to the bottom surface, by the application of a high voltage. This surface is also coated with some photoemitting material especially designed to have its peak sensitivity at the interested wavelength. The flux of accelerated electrons, spread out from the bottom, where they are directed to some more conventional detector. The bottom of the plate is an array of anode electrodes, positively charged. Electrons generated inside each channel are shifted out serially (like in normal CCD) and are available on the output of an amplifier. This way we can know the exact position of each detected photon. Amplification (multiplication) factor, can be incremented by facing the anode array outputs of one device to the cathode (channel entrance) of a second one. The advantages of a microchannel array are many, and they are used on all UV space telescopes, the largest ones are the detectors of the GALEX space telescope. Some of the best features of microchannel arrays are:

• Very high gain, • Compact size, • Fast time response. • Two dimensional detection. • High spatial resolution(25 microns).

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• Stable operation even in the presence of magnetic fields. • High sensitivity to high energetic particles and photons, make them a suitable

choice for gamma rays, x rays, UV and neutrons. • Low power consumption. • High sensitivity (depending on the cathode material). • Low dark current.

Their quantum efficiency is low, about 20 to 30%. In figure 49-a there is a picture of a microchannel array made by Hamamatsu, while, in figure 49-b and 49-c there are pictures of both UV detectors onboard the GALEX space Telescope.

4.4 – COMPARING CCD OBSERVATIONS To fully understand how the Universe work, it is necessary to observe each phenomenon in as many wavelengths as possible. This is done by the comparison of images taken by different instruments on Earth and from space. Each instrument can show us different details, depending on the combination of the image detector and its wavelength

range. To show an example of this, we have selected the starburst galaxy M82 in Ursa Major. Figure 50 is a full size image of that galaxy, covering the full visible spectral range; this image, taken by the Hubble Space Telescope, is a false colour image, where each colour represents the light of some specific wavelength. This image is a combination of four images taken with 4 filters with the central wavelength at 435nm (blue), 555nm (green), 658nm (red) and 814nm (red-orange) respectively. To better understand the difference in images through filters, we will enlarge a section, of figure 50, of the left side of

the central region. This way, we can have a better idea of the difference between a visible light image and a picture of the same area, but in the infrared. Very dense clouds of dust

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and gas (figure 51-a) let pass only a small quantity of the light generated by bright star clusters that lies behind and inside those clouds. If we observe in infrared light (figure 51-

b), thousand of hot stars and many young star clusters, clearly appears. Energetic UV light from hot stars heat dust particles and, as a consequence, they glow, emitting IR light. Only brightest star clusters are visible in figure 51-a. Complementing these images with others, in different wavelengths as for example far UV and X-ray, we can argue the age of those star clusters and their evolution stage. If their brightness are greater in UV and Xray, than in IR, they are young, instead, if they glow more in red and IR, they are old. The advantage from space observation is the possibility to take diffraction limit pictures, because of the absence of blurring atmosphere. For example, the gathering light

power of the Spitzer telescope is les than that of the Mt. Palomar Schmidt telescope; but the resolution of figure 52a (from Spitzer space telescope) is much more than that of figure 52-

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b and 52-c, which are digitized images of the Palomar Digital Sky Survey (DSS). These images are made using red sensitive photographic plates. These images show the enormous difference in sensitivity and quantum efficiency of CCDs over photographic plates. Hydrogen gas emitted from the galaxy nucleus is not visible in figure 52-b and it is barely visible in figures 52-c. The two sets of images (figures 52 and 53) can tell us the story of M82, we can argue that it collided with a closer galaxy (M82, not shown in pictures) and as a result a huge star formation begins. Due to the high brightness of these stars, we know that they are young (we can also calculate their age), so the collision was not so far in time. It is though that the interaction with M82 was 600 million years ago, because most of the stars observed are young and massive due to their strong UV light emission.

In figure 53-a we can see two superimposed images of the galaxy M82 in near (yellow) and far (blue) UV light (GALEX space telescope). Figure 53-b is the same galaxy, but the source is an UV telescope onboard the Astro1 satellite. Figure 53-c is a scanned

image with a blue sensitive photographic plate, that belongs to the Palomar Digital Sky Survey. Hot gas flowing from the nucleus is basely visible.

CONCLUSIONS Our knowledge of the Universe is strictly related to photon detectors for every band of the electromagnetic spectrum. When Galileo point his telescope to the heavens, the Universe, suddenly increase its dimension, but astronomical observations still were subjective, depending on the eye of the observer, no matter of the telescope dimension. Photographic plates let astronomers to permanently record observations in a objective mode, only instrumental defects and operation processes, can affect images. Precise measurements of colours, brightness, size and position of stars and galaxies are possible with CCD and PMT.

From the advent of Astrophotography, in the middle of the XIX century, until today, astronomers and engineers worked together to develop devices capable to detect visible light, UV, IR, etc. Thank to this cooperation, we can now observe the Universe in every wavelength of the electromagnetic spectrum.

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Astronomy, is today, inconceivable without electronics. Silicon based sensors, computers, satellites, robotic systems, are all tools used by astronomers to increment their knowledge of the Universe. We are able to take pictures of the very early Universe, of objects 13 billion light years away, thank to the precise understanding and application of the photoelectric effect on semiconductors materials. The invention of the transistor and the integration of millions of them into a tiny silicon chip favoured the development of the charge coupled device (CCD). Today CCDs are the light sensor for excellence for a wide range of wavelengths of the electromagnetic spectrum (from gamma rays to far infrared). They are now capable to detect almost every incoming photon and convert it to a measurable electric current. Very large scale integration technologies help to the development of larger sensors, with many millions of pixels, incrementing the resolution and the field of view of astronomical images, taking the full advantage of the optics of modern large telescopes. Excellent linearity and wide spectral response are features that let astronomers to measure photometric properties of stars and galaxies, giving clues to understand how they born, evolve and die. Bulky and high power consuming vacuum tube technologies are in the course to be replaced with tiny and low power silicon sensors having the same, and many times, a much better performance. Even when actual state of the art technology give us a near perfect sensor, many features will be improved in the future: flatness of spectral response, greater quantum efficiency, fast read out speed, low noise, large pixel count, selective read out, are characteristics that surely we will see in the next generation of photon detectors. REFERENCES CCD section Atomic structure and semiconductors technologies [1] Energy Bands: http://www.tpub.com/neets/book7/24c.htm[2] Solid State band theory: http://www.chemistry.adelaide.edu.au/external/soc-rel/content/bands.htm[3] Photoelectric effect: http://zebu.uoregon.edu/text/photoe.txt[4] Physics, Charles Sturt University: http://hsc.csu.edu.au/physics/core/implementation/9_4_3/943net.html#net2[5] Drakos N., 1999, Physics 1501 Modern Technology: http://theory.uwinnipeg.ca/mod_tech/node1.html[6] Bordes N., 1999, Photonic devices, Australian Photonics CRC, http://oldsite.vislab.usyd.edu.au/photonics/devices/index.html[7] Wikipedia 2006, Semiconductors: http://en.wikipedia.org/wiki/SemiconductorsPTE, Periodic Table of Elements: http://www.dayah.com/periodic/Images/periodic%20table.png[8] Hepburn C.J., Britney’s guide to semiconductor physics, the basics of semiconductors: http://britneyspears.ac/lasers.htmCCD fundamentals and history

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[9] Aikens R., 1991, Charge Coupled devices for quantitative electronic imaging, IAPPP communication No 44, Jun-Aug 1991. [10] Richmond M, Introduction to CCDs: http://spiff.rit.edu/classes/phys445/lectures/ccd1/ccd1.html[11] Tulloch S., 2006-1, Introduction to CCDs: http://www.ing.iac.es/~smt/CCD_Primer/Activity_1.ppt[12] Fairchild Imaging, Fairchild History: http://www.fairchildimaging.com/main/history.htm[13] Evolving towards the perfect CCD: http://zebu.uoregon.edu/ccd.html[14] Peterson C., 2001, How it works: the charge-coupled device or CCD: http://www.jyi.org/volumes/volume3/issue1/features/peterson.html[15] Massey D., 2005, Bell System Memorials - the transistor: http://www.bellsystemmemorial.com/belllabs_transistor.html[16] Ferreri W., Fotografia Astronomica, Il Castello, 1977 CCD fabrication and operation [17] Lesser M., 2006, CCD Glossary, The University of Arizona Imaging Technology Laboratory: http://www.itl.arizona.edu/Education/glossary.html[18] Brock k., 2001, Photodiodes, SPIE’s OE Magazine, august 2001. [19] CCD University, 2006, Apogee Instruments Inc.: http://www.ccd.com/ccdu.html[20] Pavesi, 2003, A primer on photodiode technology: http://science.unitn.it/~semicon/pavesi/tech2.pdf[21] Tulloch S. (2006-2), Use of CCD cameras: http://www.ing.iac.es/~smt/CCD_Primer/Activity_2.ppt[22] Tulloch S. (2006-3), Advanced CCD techniques: http://www.ing.iac.es/~smt/CCD_Primer/Activity_3.ppt[23] Tulloch S. (2006-4), Low light level CCD: http://www.ing.iac.es/~smt/CCD_Primer/LLLCCD.ppt[24] Tulloch S. (2005), Latest CCD developments: http://www.ing.iac.es/~smt/CCD_Primer/CCDlectureNov2005.ppt[25] CCD fundamentals, Princeton Instruments Acton, 2005 http://www.piacton.com/support/library.aspx[26] Abramowitz M., Davidson M.W., 2004, Concepts in digital imaging technology: http://micro.magnet.fsu.edu/primer/digitalimaging/concepts/concepts.html[27] Davenhall A.C., Privett G.J., Taylor M.B., 2001, The 2-D CCD data reduction cookbook: http://star-www.rl.ac.uk/star/dvi/sc5.htx/sc5.html#stardoccontents[28] Wong H.S. et al, TDI charge coupled devices : design and applications, IBM research and Development, 1992. [29] Rabinowitz D., Drift scanning, Michelson Summer Workshop, Caltech, 2005 [30] Gehrels T., CCD Scanning, 1986acm, 1986. Technical aspects of Drift Scanning, ESO Imaging Survey, 1997. [31] Gibson B., Hickson P., Time delay integration CCD readout technique: image deformation, 1992MNRAS..258..543G

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[32] Kodak Image Sensor Solutions: http://www.kodak.com/US/en/dpq/site/SENSORS/name/ISSHome[33] Buil C., 1991. CCD Astronomy, Willmann.Bell Inc., 1991, ISBN 0943396298. [34] Kitchin C.R.,1998, Astrophysical Techniques, IOP Publishing Ltd, 1998, ISBN 0750304987. [35] Howell S., Handbook of CCD Astronomy, Cambridge, 2000. [36] [DSS] Palomar Digital Sky Atals: http://archive.stsci.edu/dss/index.html[37] Bartali R., 2003, Do photographic plates still have a place in professional Astronomy?. [38] Kodak technical literature (CCD, photographic films and filters); www.kodak.com[39] Ilford technical literature (film): www.ilford.com[40] Agfa technical literature (film): www.agfa.com[41] Texas Instruments technical literature (CCD): www.ti.com[42] http://www.pinnipedia.org/optics/vignetting.html[43] http://www.astrocruise.com/geg.htm[44] http://www.chartchambers.com/whyln2.html[45] Atmel technical literature (CCD): www.atmel.com[46] Pfanhauser W., Application notes Roper Scientific gmbh, 2006: http://www.roperscientific.de/theory.html[47] Hoffaman A., Mega Pixel detector arrays: visible to 28 micron, Proceedings SPIE vol. 5167, 2004. [48] II-VI Inc, Optics manufacturing, 2006: http://www.iiviinfrared.com/opticsfab.html[49] Chaisson, AT405, 2004: http://138.238.143.191/astronomy/Chaisson/AT405/HTML/[50] Acreo AB, Infrared Detector Arrays for Thermal Imaging Tutorial "Infrared Detectors", 2004: http://www.acreo.se/upload/Publications/Tutorials/TUTORIALS-INFRARED-2.pdf[51] Teledyne Scientific and Imaging, Infrared and visible FPA, 2006: http://www.teledyne-si.com/infrared_visible_fpas/index.html[52] Carruthers G, Electronic Imaging: http://138.238.143.191/astronomy/topics.htm[53] Clampin M, UV-Optical CCD, STSI, 2001 [54] Bonanno G., New development in CCD technology for the UV-EUV spectral range, Catania Astrophysical Observatory, 1995. [55] Galaxy Evolution Explorer, Home page: http://www.galex.caltech.edu/[56] Spitzer space telescope, home page: http://www.spitzer.caltech.edu/[57] Hubble Space Telescope, home page: http://hubblesite.org/[58] UV Astronomy, Wikipedia, 2006: http://en.wikipedia.org/wiki/UV_astronomy[59] Electro optical component Inc, Silicon Carbide detectors, 2006: http://www.eoc-inc.com/UV_detectors_silicon_carbide_photodiodes.htm[60] Timothy J.G., Optical detectors for spectroscopy, 1983, 1983PASP..95..810T: http://adsabs.harvard.edu/cgi-bin/nph-bib_query?bibcode=1983PASP...95..810T&db_key=AST[61] O Connell R.W., Introduction to Ultraviolet Astronomy, 2006: http://www.astro.virginia.edu/class/oconnell/astr511/UV-astron-f01.html[62] Sheppard S.T., Cooper J.A., Melloch M.R., Silicon Carbide Charge Coupled Devices,: http://www.ecn.purdue.edu/WBG/Device_Research/CCDs/Index.html[63] Cree Research Inc., Silicon Carbide Semiconductors, 2003: http://www.mdatechnology.net/techsearch.asp?articleid=174

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[75] Optical Society of America, Optics Infobase, 2006: http://www.opticsinfobase.org/ocisdirectory/040_5250.cfm[76] Sakaki N, et al., Development of multianode photomultipliers for the EUSO focal surface detector, International Cosmic Ray conference, 2003: http://euso.riken.go.jp/publication/icrc28_233.pdf#search=%22PHOTOMULTIPLIERS%22[77] Breskin A., Ion-induced effects in GEM & GEM/MHSP gaseous photomultipliers for the UV and the visible spectral range, 2004 http://arxiv.org/ftp/physics/papers/0502/0502132.pdf[78] Casolino M., Space applications of Silicon photomultipliers: ground characterizations and measurements on board the International Space Station with the Lazio experiment, 2006: http://www.cosis.net/abstracts/COSPAR2006/03209/COSPAR2006-A-03209-1.pdf?PHPSESSID=41d280d7162dda45323d561244363f44#search=%22PHOTOMULTIPLIERS%22[79] Barral J., Study of silicon photomultipliers, 2004: http://www.stanford.edu/~jbarral/Downloads/StageOption-Rapport.pdf#search=%22PHOTOMULTIPLIERS%22[82] University of Pisa, Physics Department, Silicon Photomultiplier, 1995: http://www.df.unipi.it/~fiig/research_sipm.htm[83] Piemonte C., SiPM: status of the development, 2006: http://sipm.itc.it/intro/device.html[84] Ninkovic J., The avalanche drift diode: A back illuminated silicon photomultiplier, 2006: http://www.hll.mpg.de/twiki/bin/view/Avalanche/AvalancheDriftDiode IMAGE CREDITS Figure 1a, 1b Atomic energy bands: http://www.tpub.com/neets/book7/24c.htmFigure 2 CCD geometry (adapted from): http://www.ing.iac.es/~smt/CCD_Primer/Activity_2.pptFigure 3 Pixel structure: Kitchin C.R.,1998, Astrophysical Techniques, IOP Publishing Ltd, 1998, ISBN 0750304987. Figure 3b (adapted from): http://www.ing.iac.es/~smt/CCD_Primer/Activity_2.pptFigure 4 Pixel size and well capacity relationship: Bartali R., 2006 Figure 5 Silicon absorption depth graphic (adapted from): Howell S., Handbook of CCD Astronomy, Cambridge, 2000. Figure 6A, 6B Linear CCD: http://www.fairchildimaging.com/products/fpa/ccd/linear/ccd_191.htmFigure 6B Matrix CCD: http://www.fairchildimaging.com/products/fpa/ccd/area/ccd_3041.htmFigure 7

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CCD rain-buckets analogy: http://www.microscopyu.com/articles/digitalimaging/ccdintro.htmlFigure 8 First astronomical CCD image: http://zebu.uoregon.edu/ccd.htmlFigure 9 Galaxy M51: http://hubblesite.org/gallery/wallpaper/pr2005012a/800_wallpaperFigure 10 Image sensors quantum efficiency: Howell S., Handbook of CCD Astronomy, Cambridge, 2000. Figure 11a, 11b, 11c Front and back side CCD: Bartali R., 2006 Figure 12 Back and front illuminated CCD comparison: http://www.site-inc.comFigure 13a Optoelectronics Databook, 1984, Texas Instruments: http://www.ti.comFigure 13b Deep Depletion CCD (pixel structure): http://www.ing.iac.es/~smt/redsense/deep_depletion.PDFCCD42-90 CCD Datasheet, Marconi Applied Technology (QE graph): http://www.marconitech.comFigure 14a Crab Nebula in red light: http://archive.stsci.edu/cgi-bin/dss_formFigure 14b Crab Nebula in blue light: http://archive.stsci.edu/cgi-bin/dss_formFigure 14c Crab Nebula VLT: http://www.eso.org/outreach/press-rel/pr-1999/phot-40f-99-normal.jpgFigure 14d Crab Nebula HST: http://hubblesite.org/gallery/wallpaper/pr2005037a/800_wallpaperFigure 15 Dark frames examples: http://www.frazmtn.com/~bwallis/drk_tmp.htmFigure 16 Bias frames examples: http://www.eso.org/projects/odt/Fors1/images/bias.jpghttp://www.carleton.edu/departments/PHAS/astro/pages/knowledgebase/biasdark.htmlFigure 17a, 17b Flat field frame example: http://www.highenergyastro.com/CVFUN.htmlhttp://www.mso.anu.edu.au/observing/detectors/imager.phpFigure 18 Example of raw image: Bartali 2003 Figure 19 Example of science image: Bartali R., Rosner A., 2003 Figure 20 CCD imaging, basic steps: Bartali R., 2006

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Figure 21a Hot pixel: Bartali R., 2003 Figure 21b Bright column: Tulloch S., Use of a CCD camera. Figure 22 Dark Streacks: HET609 CDRom, 2006 Figure 23 Blooming example: Bartali R., 2003 Figure 24a Example of glowing: Buil C., CCD Astronomy Figure 24b CCD output amplifier: http://spiff.rit.edu/classes/phys445/lectures/ccd1/structure_1.gifFigure 25 Cosmic rays: http://spider.ipac.caltech.edu/staff/kaspar/obs_mishaps/images/cr.htmlFigure 26 Frame transfer CCD: http://www.sinogold.com/images/tc237face.gifFigure 27 Example of an overexposed image: Bartali R., 2003 Figure 28 Oversampled stellar image: http://sctscopes.net/Photo_Basics/CCD_Camera/Choosing_a_CCD_Camera/CCD_Parameters/ccd_parameters.htmlFigure 29 Undersmpled stellar image: http://sctscopes.net/Photo_Basics/CCD_Camera/Choosing_a_CCD_Camera/CCD_Parameters/ccd_parameters.htmlFigure 30 Vignetting: http://www.astrocruise.com/geg.htmFigure 31 Peltier module: An Introduction to Thermoelectrics, Tellurex Corporation, 2006. Figure 32 Cooling Peltier module: An Introduction to Thermoelectrics, Tellurex Corporation, 2006. Figure 33 Multistage module: Melcor Thermal Solutions, Melcor Corporation. Figure 34 Liquid Nitrogen cooler: top – Tellurex Corporation. Bottom – Buil C., CCD Astronomy, Willmann Bell, 1991. Figure 35 UBVRI filters: http://outreach.atnf.csiro.au/education/senior/astrophysics/photometry_colour.htmlFigure 36 Eye spectral response: http://www.marine.maine.edu/~eboss/classes/SMS_491_2003/sound/em-spectrum_human-eye_asu_380x300.gifFigure 37 100 million pixel CCD:

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http://www.dalsasemi.com/news/news.asp?itemID=252Figure 38 Mosaic camera: http://www.cfht.hawaii.edu/Instruments/Imaging/CFH12K/images/CFH12K-FP_lr.jpgFigure 39 Mosaic image: http://www.astro.indiana.edu/~vanzee/SMUDGES/field.htmlFigure 40-a Orion nebula in visible band from Harvard Observatory: http://138.238.143.191/astronomy/Chaisson/AT405/HTML/AT40506.htmFigure 40-b Orion nebula in the IR band from NASA: http://138.238.143.191/astronomy/Chaisson/AT405/HTML/AT40506.htmFigure 41-a, 41-b Si PIN sensor: Hoffman A, Mega pixel detector arrays from visible to 25 micron, 2004 Figure 42-a, 42-b HgCdTe sensors (SW): Hoffman A, Mega pixel detector arrays from visible to 25 micron, 2004 Figure 43 HgCdTe sensors (MW): Hoffman A, Mega pixel detector arrays from visible to 25 micron, 2004 Figure 44-a, 44-b InSb sensors: Hoffman A, Mega pixel detector arrays from visible to 25 micron, 2004 Figure 45-a, 45-b Si:AsIBC sensors: Hoffman A, Mega pixel detector arrays from visible to 25 micron, 2004 Figure 46 SiC response: http://www.eoc-inc.com/UV_detectors_silicon_carbide_photodiodes.htmFigure 47 Channel multiplier, adapted from: http://www.olympusmicro.com/primer/digitalimaging/concepts/photomultipliers.htmlFigure 48 Microchannel structure: Hamamatsu photomultiplier tubes, Hamamatsu Corp., 2006. Figure 49-a Microchannel detector: Hamamatsu photomultiplier tubes, Hamamatsu Corp., 2006. Figure 49-b, 49-c GALEX UV detectors: http://www.galex.caltech.edu/ Figure 50 M82 HST: http://hubblesite.org/newscenter/archive/releases/2006/14/image/a/format/large_webFigure 51-a M82 visible light (HST): http://www.seds.org/messier/m/m082.htmlFigure 51b-a M82 IR ligh (HST)t: http://www.seds.org/messier/m/m082.htmlFigure 52-a M82 Spitzer: http://hubblesite.org/newscenter/archive/releases/2006/14/image/i/format/large_web

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Figure 52-b M82 IR DSS: http://archive.stsci.edu/cgi-bin/dss_search?v=poss2ukstu_ir&r=09+55+52.19&d=%2B69+40+48.8&e=J2000&h=15.0&w=15.0&f=gif&c=none&fov=NONE&v3Figure 52-c M82 red DSS: http://archive.stsci.edu/cgi-bin/dss_search?v=poss2ukstu_red&r=09+55+52.19&d=%2B69+40+48.8&e=J2000&h=15.0&w=15.0&f=gif&c=none&fov=NONE&v3Figure 53-a M82 GALEX: http://www.galex.caltech.edu/GALLERY/GALEX-M82.jpgFigure 53-b M82 ASTRO1: http://www.seds.org/messier/Pics/More/m82a1uv.jpgFigure 53-c M82 blue DSS: http://archive.stsci.edu/cgi-bin/dss_search?v=poss2ukstu_blue&r=09+55+52.19&d=%2B69+40+48.8&e=J2000&h=15.0&w=15.0&f=gif&c=none&fov=NONE&v3=

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