## ASTR 511 (O'Connell) Lecture Notes

CCD'S IN ASTRONOMY

Subaru CCD Mosaic 8 x (2k x 4k)

Charged coupled devices (CCD's) have been used in astronomy since the late 1970's. They are now nearly ubiquitous in observations made at wavelengths between the near-IR (~1 µ) and the X-ray. They have transformed the way astronomy is done.

## I. REFERENCES FOR CCD'S

General

CCD Data Reduction and Applications

## II. GENERAL DETECTOR CHARACTERIZATION

 QUANTUM EFFICIENCY QE = percentage of photons incident on detector which produce measurable signals Strong wavelength dependence (e.g. threshold activation cutoffs set by workfunction/band gap) Typical peak values: Eye: 1-2% Photographic plate: 1-2% Photomultiplier tube: 20-30% CCD: 70-90% IR array (HgCdTe): 30-50% Schematic QE curves for various classes of detector
• "Detective quantum efficiency" is defined as [(SNRout)/(SNRin)]2, where "in" and "out" refer to the input and output signals to/from the detector, respectively. DQE combines basic QE with the noise introduced by the detector. This quantity is really what matters in comparing detectors, but it is so dependent on specific details of operations/applications that it is rarely used.

• SPECTRAL RANGE

• Wavelength region over which QE is sufficient for operation

• DYNAMIC RANGE

• Definition: ratio of maximum to minimum measurable signal

• E.g. maximum number of events in a CCD pixel is determined by photoelectron "full well'' capacity or digitization maximum (typically 2 bytes); minimum is determined by dark current/readout noise

• Applies to a single exposure; effective dynamic range can be increased with multiple exposures

• Typical values: 100 (Pg); 104 (CCD); 105 (PMT)

• Related concepts:

• Linear Range: range of signals for which [Output] = k x [Input], where k is a constant. Generally smaller than calibrateable range

• Threshold: minimum measurable signal. Influenced by detector noise or other intrinsic characteristics (e.g. fog on Pg plates)

• Saturation point: level where detector response ceases to change with signal

• NOISE CHARACTERISTICS

• RESOLUTION

• Temporal

• PMT: ~0.0001 ms
• CCD: ~10 ms per pixel

• Spatial (array detectors)

• Pixel = minimum measureable area of detector surface. Typically 10-50 µ on Pg or semiconductor types. Pixels are not necessarily independent of one another.

• Resolution cell: according to the Nyquist criterion, the resolution cell is 2x the size of the minimum independent measurable area. For proper sampling of image, need at least 2 pixels per resolution cell. A lower pixel density implies "undersampled" imaging. A significantly higher pixel density does not provide more information and is a waste of pixels.

For an imaging application, the Nyquist criterion implies that 1 pixel should span ~ (FWHM)/2, where FWHM is the full width half maximum of the point spread function.

• Spectral

• Most UVOIR detectors are broad-band; inherently poor spectral resolution. Must use external elements (filters, gratings) to provide spectral resolution

• Exception: superconducting "3D" detectors, which measure photon energy as well as position; current spectral resolution is R ~ 100

## III. BRIEF HISTORY OF ASTRONOMICAL DETECTORS

PHOTOGRAPHIC ASTRONOMY:

Photographic emulsions work by storing in AgBr crystals the photoelectrons ejected following absorption of photons during exposure to light. Chemical reactions during development cause the crystals to precipitate grains of silver, which form a permanent image.

Film was the detector of choice for almost all applications in astronomy from around 1900 to 1960 and for imaging until about 1980. It is impossible to exaggerate its importance to the development of modern astronomy. Even with relatively low QE, photographic plates offered decisive advantages over visual observations:

1. Very long exposure times (up to a week in early applications), meaning limiting fluxes thousands of times fainter
2. A permanent, objective record of astronomical images and spectra
3. Geometric stability, important for astrometry and morphology
4. Large formats for wide field surveys
5. Extension of observations to EM bands beyond the "visual" band (the classic "photographic" band is centered around 4500 Å)

However, the photographic emulsion had serious limitations with which astronomers had long struggled:

1. It is relatively insensitive, with QE ~ 1%.
2. Pg materials have non-negligible thresholds (a minimum of ~4 electrons per grain), a strongly non-linear response function, and limited dynamic range. These properties plus the use of hard-to-control chemical processes for emulsion deposition & development makes them very difficult to calibrate quantitatively.
3. Each individual plate (or at the best, each co-processed "batch" of plates) has different characteristics.
4. Conversion of data to quantitative form (e.g. with microdensitometers) is slow and cumbersome.

For details on astronomical photography at low light levels, see Smith & Hoag 1979, ARAA, 17, 43.

PHOTOELECTRIC ASTRONOMY TO 1980:

In 1907, Joel Stebbins (UIll, UWisc) began testing various types of photoelectronic detectors. These, such as selenium phototubes, were largely experimental and temperamental, but use of PE devices spread after the 1920's. World War II greatly accelerated these technologies, with mass production of high quality vacuum tubes and sensitive electronics for detection of faint signals.

The key design for astronomy was the photomultiplier tube (PMT), the first widely used example of which was the RCA 1P21. The initial detector in a PMT is a photo-emissive cathode surface, made from alkali metal compounds, which ejects a single electron in response to a photon absorption. A series of other "secondary electron emissive" surfaces (the "dynode chain") amplifies this into a burst of ~106-7 electrons. (See illustration below.)

PMT's require cooling to suppress dark current and were typically operated at dry-ice temperatures (-78C). The photocathode QE's of PMT's in the optical region reached 20-40% by the 1960's, with very wide acceptance bands.

Although PMT's were initially operated in a "direct current" mode, the pulse-like character of PMT output led to the adoption in astronomy of the same kinds of high speed digital electronics that had been developed earlier for accelerator and nuclear physics. This kind of "pulse-counting" operation offers excellent noise rejection and permits the detection of individual photons from cosmic sources.

For more details on PMT's and PMT-based photometers, see the articles by Lallemand and Johnson in Astronomical Techniques, ed. W. A. Hiltner (Stars & Stellar Systems Vol II).

PMT's became the workhorses of multicolor photometry and spectrophotometry (e.g. the Johnson & Morgan broad-band UBV system) after 1950. They featured excellent linearity and stability, which implied an unprecedented capacity for accurate calibration of photometric measurements. PMT photometry routinely reached the 1% level for flux calibration. In turn, this was responsible for an explosion in quantitative astrophysics.

Despite their excellent performance and their broad wavelength response, PMT's were single element devices. They were not easily adaptable to 1-D, let alone 2-D, applications (such as imaging), despite heroic attempts such as the Palomar Multichannel Photoelectric Spectrometer (Oke, 1969).

Around 1960, a plethora of efforts began to develop robust 1-D and 2-D electronic devices suitable for astronomy. A dozen or so of these produced practical systems (e.g. the Secondary Electron Conduction Vidicon, the Intensified Dissector Scanner, the Intensified Reticon Scanner, the Image Photon Counting System, and the Multi-Anode Microchannel Array detector).

Most of these employ some type of image intensifier vacuum tube as a key element. This technology was developed for military night vision systems. Intensifiers produce easily detectable light pulses in a 2-D image field from individual incident photons by accelerating the photo-electrons to high energies and/or producing a cascade of electrons, then focussing these on a luminescent output screen. Disadvantages include the frequent use of high voltages (e.g. 25 kV) and serious image blur which, however, can be reduced by the use of pulse centroiding electronics. Fiber optic input/output plates were commonly used with intensifiers after 1970. Microchannel plate (MCP) intensifiers have often been used in space astronomy missions (including HST, GALEX, and FUSE).

 SOLID STATE ARRAY DETECTORS: But by far the most successful 2-D devices for astronomy emerging in the last 30 years have been solid state, semiconductor arrays. These are based on photo-conductive materials fabricated with embedded microelectronic integrated circuits on thin wafers by photolithography. Although lower quality devices can be mass-produced by microchip "foundries," professional grade detectors still need to be custom-processed. Solid state arrays are now employed as astronomical detectors from the X-ray to the far-infrared. The most widely used is the Charge-Coupled Device (CCD), which operates at wavelengths from the X-ray to ~1 µ. The basic CCD architecture was invented at Bell Labs in the late 1960's. By the mid-70's, CCDs were being tested as astronomical detectors. They did not come into widespread use until ca. 1980, following extensive development efforts associated with the Wide Field & Planetary Camera on the Hubble Space Telescope, led by J. Westphal, J. Gunn, and C. R. Lynds. Wafer containing a number of co-fabricated CCDs and other devices.

## IV. SEMICONDUCTORS

Semiconductors are crystalline materials which are not normally good conductors of electricity but which can be made to conduct under certain circumstances. The central useful property of semiconductors employed in astronomy is that their electrical properties change significantly after absorption of photons.

 BAND GAPS: The properties of semiconductors are manipulated by changing the structure of their internal energy levels, which are spread out into "bands" by the proximity of the component atoms of the solid. The "valence" band, corresponding to the outermost electrons in a normal, unexcited atom, is the lowest energy band where electrons are able to move between ions. However, no net conduction occurs as long as the band is full. Above this lie the "conduction" bands, which are not filled, and where electrons will move freely in response to external EM fields. The separation between the valence and first conduction band is called the "band gap energy", Eg. Different materials have a wide range of band gaps. "Conductors" have a zero gap, meaning that electrons are always available to transfer charge. "Insulators" have very large gaps, implying zero conduction except under extreme circumstances. "Semiconductors" have intermediate gaps.
Absorption of a photon can push a valence electron into the conduction band and produce a potential electrical signal. The photon energy must exceed Eg, which implies that there is a maximum wavelength for excitation given by:

Lammax = 12,400 Å/Eg(eV)

Obviously, materials with band gaps in the few eV range are of interest as potential UVOIR detectors. Band gaps and max wavelengths for some important materials are given in the following table:

Material Eg(eV) Lammax (Å)
Pure Si 1.1 11,300
GaAs 1.43 8,670
InSb 0.36 34,400
Hg1-xCdxTe 1.55x 8,000/x

 DOPING: The "elemental" semiconductors are those elements in group IVa of the Periodic Table (including Si and Ge). These have four electrons in their valence shells, half the maximum allowed. These are shared among the ions in "co-valent bonds." There are many other types of "compound" semiconductors, however, composed for instance of atoms from group IIIa and Va of the Table; two of these are listed in the table above. The electrical properties of pure semiconductors can be dramatically altered by adding ("doping with") small amounts (~1 part in 106) of an impurity. The result is called an "extrinsic" semiconductor. n-type: a material with more than 4 valence electrons is added (As, from group Va, in the illustration). The extra electrons cannot be accommodated in the valence band and so occupy the conduction band. They represent a persistent set of negative carriers p-type: a material with fewer than 4 valence electrons is added (e.g. B, from group IIIa). This has one fewer electron than normal and creates a small "vacuum" in the electron sea of the valence band. This is called a "hole." As valence electrons shift to fill it, the hole propagates like a positive charge in the opposite direction. The holes represent a persistent set of positive carriers. In pure semiconductors, conduction electrons and holes can also exist, if electrons are excited by thermal effects, for instance. But they always occur in pairs. Electrons and holes in n- and p-type materials, respectively, have no counterparts. Extrinsic material is electrically neutral but is more responsive than pure materials to a difference in electrical potential. By adjusting the amount of doping, the band gap of the semiconductor (donor/acceptor levels in diagram at right) can be customized. By layering n/p regions, the response to applied potentials can be adjusted to create a large variety of electronic devices. n-type doping p-type doping Doping affects energy levels
Photons are primarily absorbed by electrons in the valence band. For photon energies above Eg, the electron is boosted to the conduction band, leaving a hole behind. If a positive voltage is applied at one side of the semiconductor, the electron will be attracted toward it while the hole will be repelled.

## V. BASIC CCD DESIGN

Apart from sensitivity, the key design issues for solid state arrays are to localize photon-produced charges on their surfaces and then arrange to amplify and read them out without distorting the image or introducing unacceptable amounts of noise.

A CCD is a charge-transfer device. Its storage, transfer, and amplification electronics are all fabricated on a single piece of silicon (unlike most IR arrays). During an exposure, it traps released photoelectrons in small voltage wells. After the exposure, the electrons are shifted in a series of "charge-coupled" steps across the CCD surface, amplified, read out of the CCD, and stored in a computer memory. This is "destructive readout"---i.e. one cannot read the same signal again (unlike other array architectures, where each pixel is coupled to a separate amplifier).

BASIC STRUCTURAL ELEMENT: The basic element in a CCD design is a "Metal-Oxide-Semiconductor" capacitor. See the illustration above. This serves both to store photoelectrons and to shift them wholesale. The bulk material is p-silicon on which an insulating later of silicon-oxide has been grown. P-silicon can be made to have very small numbers of free electrons ("high resistivity") before exposure to light; this is important for best performance. A set of thin conducting electrodes of semitransparent "polysilicon" are applied.

Before exposure, the central electrode is set to a positive bias while the two flanking electrodes are set negative. This creates a "depletion" region under the central electrode containing no holes but a deep potential well to trap electrons. The region shown is about 10µ thick.

 OPERATION SEQUENCE: During exposure (controlled by a separate shutter), light enters through the "front-side" electrodes. Photoelectrons generated under the central electrode will be attracted toward the electrode and held below it. The corresponding holes will be swept away into the bulk silicon. Good performance requires little diffusion of electrons away from the potential well. The surface of the CCD is covered with MOS capactitors in a pattern like that at right. In this particular design, there are three electrodes per pixel. A single pixel is shown shaded in the diagram. Typical pixel sizes are 10-40 µ. The "parallel shift" registers are shown as rows running across the whole face of the CCD. These are separated by insulating "channel stops." At one end is a column of "serial shift" electronics and an output amplifier. Note that there is only one amplifier in this design. Contemporary large chip designs involve several amplifiers (but always many fewer than the number of pixels!). At the end of the exposure, readout of the collected electrons is accomplished by cycling ("clocking") the voltages on the electrodes such that the charge is shifted bodily to the right along the rows. The figure at the right shows how this is done. Good performance depends on near-100% transfer of the electrons to/from each electrode with no smearing or generation of spurious electrons. Each parallel transfer places the contents of one pixel in each row into the serial register column. This column is then shifted out vertically through the output amplifier and into computer memory before the next parallel transfer occurs.
BUCKET BRIGADE: The resulting transfer and readout process is illustrated in the animation below:

ADU CONVERSION: For storage in memory, the electrical signal generated by the amplifier must be digitized. This is done by an "analog-to-digital converter". This is normally adjusted such that one digital unit corresponds to more than one photo-electron. Typical values of this conversion are 2 to 8 electrons per stored digital unit.

The stored values are called "ADU's", for analog-to-digital-unit. The corresponding constant of transformation, normally quoted in units of "electrons per ADU", is often called the "Gain" (although this is confusing nomenclature because a larger Gain results in reduced ADU values).

Note that the use of such a conversion importantly affects the statistical properties of the recorded signal. If x is the recorded signal in ADU's, y is the original signal in photo-electrons, and G is the gain, then from Lecture 8 we see that:

Var(x) = Var(y)/G2

## VI. CCD DESIGN ISSUES

CCDs have undergone a long optimization process since 1980. Contemporary designs have excellent performance but still require careful calibration in order to overcome inherent limitations. There is also only a handful of reliable manufacturers of professional-grade chips.

Here are some of the issues affecting electronic design and manufacture of CCDs:

 INTRINSICALLY LOW BLUE RESPONSE (< 4500 Å): Caused by absorption in bulk Si and by electrode structures in "Frontside-Illuminated" chips. Mitigation: Use special thin, polysilicon material for electrodes. But cannot be completely transparent. Special Coatings: "Anti-reflection" coatings trap photons, causing multiple reflections as in Fabry-Perot etalons, and therefore enhance absorption. "Downconverter" coatings are phosphors which absorb blue photons but emit green photons at wavelengths where the CCD QE is higher (e.g. "mouse milk," coronene, lumogen). "Thinned, Backside-Illuminated" chips: shave off silicon subtrate, leaving only ~10µ deep unit; illuminate from backside; greatly improves blue response. For techniques, see Mike Lesser's UofA website.
Difficulties with thinned chips:

• Fragile
• Non-uniform thinning
• Surface trapping by SiO2 layer of photo-electrons produced nearby (shorter wavelengths)
• Interference effects if wavelength ~ chip thickness (i.e. in IR). Strong spatial modulation of response = "Fringing". Especially serious for night sky emission lines. Can be well calibrated for narrow-band filters or for broad-band filters. Hard for intermediate band filters.
• Thinning reduces red response. For good response 5000--10000 Å, prefer thick (~500µ) front-illuminated chips.

• CHARGE TRANSFER EFFICIENCY (CTE)

• CTE can be better than 99.999% per transfer, but has to be, since throughput of chip with 2048 required shifts = CTE2048.

• Radiation damage to CCD's in space seriously decreases CTE over several years' time.

• Mitigation:

Operational: add (electronically or with diffuse light source) a pedestal background signal (a "fat zero") over whole chip to increase mean electron density per pixel. However, adds additional noise and not suitable for very faint source applications.

Technical design: change number phases, clocking cycles; add "minichannels."

## VII. ADVANTAGES OF CCD DETECTORS

 HIGH QUANTUM EFFICIENCY: To 80-90% at peak in optical. Much effort was expended to reach these high levels. This had tremendous impact on astronomical imaging & spectroscopy. It meant the detection threshold with any instrument was extended 4-5 magnitudes over film and that therefore a 1-m telescope could now pursue the kind of science previously possible only with 4-m class telescopes. A key corollary: since we are already near 100% QE, at least in the optical region, achieving significantly lower threshold fluxes requires larger telescopes rather than better detectors. NB: "Quantum yield" can be over 100% for far-UV and X-ray photons (i.e. more than one photoelectron can be generated but fewer than 100% of photons produce responses). Sample CCD QE curves (ESO)
• LINEAR RESPONSE: Until approach full-well capacity (typically 200,000 e/pixel). This implied much improved flux calibrations and much higher precision for flux measurements at all levels.

• EXCELLENT DYNAMIC RANGE: Typically 104.

• WIDE WAVELENGTH RESPONSE : Intrinsically sensitive from X-ray to ~1µ. Other materials (e.g. InSb2, HgCdTe) with similar architectures are usable in the IR.

• GEOMETRICALLY STABLE: good for astrometry

• ONLY LOW VOLTAGES REQUIRED (~5-15 v)

• RELATIVELY CHEAP, AVAILABLE, SIMPLE: Compared to other digital 2-D systems. Standard chips cost ~\$2-200 K.

• RELATIVELY LOW NOISE : Compared to many other classes of astronomical detectors, e.g. Pg plates, Reticons, SEC vidicons, etc. But noise is not negligible. Typical read noise now 2-10 e/pixel, and dark current is largely suppressed by cooling.

• SMALL PIXELS : Typically 10-30 µ. Usually an advantage, but should match pixel size to 1/2 of smallest resolvable picture element in optical system.

• SPECIAL FORMAT/READOUT DESIGNS: By changing electrode structure and clocking cycles, one can arrange for many different integrate/readout modes.

Rapid clocking/video: inherent in CCDs intended for TV application. For bright sources, readout rates of 100 MHz are possible.

Drift Scanning:

Basic technique is to clock the chip slowly along columns while moving the telescope to keep a target centered on the moving electron cloud. Continuously read out the chip to produce a strip-image of the sky. The integration time is set by the drift time across the chip.

The simplest approach is to align the CCD columns east-west, keep the telescope fixed on the sky, and clock the chip westward at the sidereal rate. A sample application of drift scanning is described here.

Drift-scanning is now a standard method for wide-field CCD sky surveys, most notably the Sloan Digital Sky Survey

"Nod and Shuffle": this technique takes advantage of the capability to shift an image wholesale on the CCD without reading out to obtain much better sky subtraction (e.g. in near-IR where atmospheric OH emission is very bright and variable).

On-chip binning: change clocking to combine electrons from several pixels together before reading out through amplifier. E.g. combine a 2x2 pixel region on the chip into a single output pixel. This reduces the effect of amplifier readout noise on each pixel in the final data. Also reduces memory and storage requirements. Adds considerable practical flexibility to CCD systems. Obviously, however, reduces the resolution of the output image. Is useful for applications such as imagery of very faint, extended sources (e.g. galaxy halos), low spatial resolution spectroscopy, photometry of point sources under poor seeing conditions, etc.

• UBIQUITOUS : Now almost universally used in astronomy (amateur & pro). Photographic materials and older electronic detectors being phased out.

• IMMEDIATE DIGITAL CONVERSION OF DATA: The other advantages of array detectors notwithstanding, it is the immediate conversion of astronomical signals into a form capable of computer storage/analysis which has so dramatically transformed UVOIR astronomy since 1980. The practice, pace, & scope of UVOIR astronomy are entirely different now than in the "photographic" era that preceded widespread use of CCDs and other array devices. Digital conversion of images and spectra has enabled quantitative analysis of observations on a scale not possible before.

Among other things, rapid digital processing allows much improved use of observing time---notably in surveys (e.g. for variable sources in MACHO and supernova searches; moving targets such as asteroids/Kuiper Belt objects; or combined imaging/spectroscopic surveys such as SDSS, 2dF).

• SMALL SIZE : Individual chips typically less than 7 cm square. Cover only ~20% of high quality imaging field on typical modern telescope. Size limited by small fabrication yield of large, defect-free chips. Largest routinely available chips are 4096x2048.

• However, MOSAICKING technology now well developed. Typical mosaic now uses 4096x2048 chips.

• Examples: (click on images below for enlargements):

• Sloan Digital Sky Survey: 54 chip mosaic for drift scanning in 5 bands simultaneously

• Canada-France-Hawaii-Telescope MegaPrime mosaic: 40 4096x2048 chips covering 1 degree FOV

 Sloan DSS Camera Sloan DSS Camera CFHT MegaPrime Mosaic

• CRYOGENIC COOLING REQUIRED: To reduce dark noise, cooling to below -100 C necessary. Thermoelectric coolers usually not adequate, so cryogens (e.g. liquid N2) required. Introduces many practical complications. Optimal low-T chips differ from commercial types used at room temperature in digital cameras, etc.

• READOUT NOISE: Produced by variations in amplifier gain. Much effort to reduce. Now typically 2-10 e/pix.

What matters is not the noise per pixel but rather the total noise per image area, which can extend over many pixels, depending on the application.

Even at these very low levels, RON can compromise some types of observations (spectroscopy, surface photometry), see Lecture 7.

Can reduce RON effects in some applications using on-chip binning (see above). In some cases, it may be more advantageous to use "pulse counting" detectors, which can unambiguously detect individual photons, than CCD's.

• FULL WELL CONSTRAINTS: Bright sources which over-fill pixels can produce "blooming" or "bleeding" along columns, making parts of the chip other than the immediate vicinity of the source useless. Best solution is operational.

 RESPONSE NONUNIFORMITIES ("FLAT FIELD" EFFECTS): Caused by small variations in masks used to manufacture chips, deposition irregularities, thinning variations, etc. Typically 5% pixel-to-pixel. Color dependent. Requires extensive calibration, with color-matching to targets. Use "dome flats" or "twilight flats." Special observing procedures to reduce flat field effects include: Drift Scanning: see above. Multiple Offset: Break exposure into 4-5 parts, offset ~50 pixels between exposures. Combine exposures during data reduction. "Dithering" is a related technique involving smaller offsets to achieve sub-pixel spatial resolution. These methods result in reduced field of view because not all parts of the original field will have uniform exposures. SENSITIVITY TO COSMIC RAYS: Especially thick chips. Must remove effects during processing. CR's are a major problem for spacecraft detectors (e.g. HST/WFPC2). Requires that exposures be broken into multiple parts (called "CRSPLITs" on HST) so that CR events can be detected. CR hits can be removed from data frames, but this always leaves "holes" which have less exposure than other parts of the image. SENSITIVITY TO X-RAYS: An advantage for X-ray astronomy, but some materials in vicinity of CCD's, e.g. special glasses used for windows, can produce a diffuse background of X-rays which add noise to observations. CCD Flatfield Frame (AAO)   CR's on HST/WFPC2 2400 sec image (extract)
• CHARGE TRANSFER EFFICIENCY: Good CTE is often possible only for signal levels above threshold ~10-50 e/pix. For low light levels, require adding "fat zero" (typ. 1000 e/pix) electronically or by preflash. Creates added noise (10-30 e/pix).

• AVAILABILITY HOSTAGE TO COMMERCIAL MARKET: CCD availability has always been driven more by commercial and military applications than science. Scientific CCD manufacture represents only about 10% of the overall \$1B CCD market. A serious near-term issue, since industry is moving to "active pixel sensors," a different technology with amplifiers incorporated in each pixel. Requirements for good performance at low (astronomical) light levels are considerably different than for room-temperature, short-exposure, mass-produced equipment.

• DATA GLUT! CCDs typically produce 2 x Npix bytes of data per readout, where Npix is the number of pixels. For a 8096x8096 mosaic, this is 130 MB. Data storage/manipulation was a serious problem when CCD's were first introduced, and this influenced the style of data processing systems such as IRAF. Disk and tape storage are now much improved, but long-term stability of massive amounts of data now being produced is a non-trivial issue. [All that data is really important and worth saving...isn't it?]

## IX. EXAMPLE CCD SYSTEMS

• HST: WIDE FIELD PLANETARY CAMERA 2 (1992)

• 4 CCD's, optically mosaicked with beam-splitter
• Loral 800x800, 15µ pixel chips
• Front-illuminated, lumogen phosphor coating for UV response
• RON 5 e/px; QE (6000 Å) 35%

• HST: WIDE FIELD CAMERA 3 (2001)

• 2 CCD's, physically mosaicked, contiguous
• Marconi 4096x2048, 15µ pixel chips
• Thinned, back-illuminated; anti-reflection coating for enhanced blue/UV response
• Minichannel for improved CTE
• RON 3 e/px; QE (6000 Å) 70%

• SLOAN DIGITAL SKY SURVEY (1997)

• Drift scanning mosaic (not contiguous) containing 54 chips
• Tektronix/SITe 2048x2048, 24µ pixel chips in 5x6 array; for simultaneous broad-band photometry. Both frontside and thinned backside with AR coating. QE (6000 Å) ~60%; RON < 5 e/px
• Tektronix/SITe 2048x400, 24µ pixel chips; for astrometry & focus check. Frontside illuminated, RON < 15 e/px.

• FAN MOUNTAIN CCD SYSTEM (Gen II)

## X. CCD HIGH PRECISION CALIBRATION PROCEDURE

A. DATA REQUIRED

• Bias frames
• Dark frames (optional)
• Flat field frames
• Sky flat/fringe frames (optional)
• Flux calibrator fields
• Target frames

B. REDUCTION PROCEDURE

• SUBTRACT BIAS FRAME : Bias frames (zero exposure time) taken with chip unexposed to light from telescope. These measure electronic pedestal of chip. For high precision, average many bias frames, before and after observations. Check for bias drift during night. With some chips, can determine electronic bias level from overscan region on chip. For optical "fat zero" subtract average of many frames.

• Optional: SUBTRACT DARK FRAME: Median filter many long (~30 min) dark exposures; note possible LED activity of CCD electrodes. Scale result to integration time of each data frame before subtracting.

• DIVIDE BY TWILIGHT OR DOME FLAT FIELD: To remove residual pixel-to-pixel variations (typical 5%). Make exposures of twilight sky (good diffuse source, but tricky to get exposure level right; only 2 chances per night). Or make many exposures of diffusely illuminated screen in dome (disadvantage: often not uniformly illuminated). Must be done for each filter used (color sensitive effects).

[Note: photon statistics must yield S/N significantly in excess of final desired S/N---i.e. if desire S/N = 100, require over 10,000 photons per pixel in net flat field image.]

• Optional: SUBTRACT SKY FLAT/FRINGE frame: Remove night sky emission line fringing effects (worse in near-IR) by observing uncrowded field in night sky. Take several exposures, moving telescope by, say, 25 arc-sec between them. Use "adaptive modal filter" technique to zap star images and create mean sky frame. Scale to target frames and subtract. Resulting flat field as good as 0.1% can result. Require a sky flat determined this way for each filter.

• USE MULTIPLE OFFSETS, DITHERING OR DRIFT SCANNING FOR TARGET FRAMES: For faintest possible photometry, use multiple-offset exposure technique---e.g. 500 sec exposures shifted by about 20 arc-sec each--- to reduce flat field errors. Always want more than one exposure anyway for "reality check," empirical determination of photometric errors, cosmic ray removal, etc.

• REGISTER AND COMBINE TARGET FRAMES: Re-register offset frames to sub-pixel accuracy (e.g. "Drizzle" softare). Median filter. Trim off "lost" parts of image.

• CALIBRATE FLUX SCALE: Observe "standard stars" in recommended CCD calibrator fields (e.g. star clusters) to determine nightly atmospheric extinction and telescope throughput. No need to take on non-photometric nights (variable clouds), but no flux calibration either. Important that calibrator stars overlap targets of interest in color. These frames are reduced exactly like target frames, apparent fluxes extracted, and correction factors to determine true fluxes of targets are obtained.

• APPLY GOOD PHOTOMETRIC REDUCTION SOFTWARE: for source IDs, flux measurements (point or extended source).

=====> RESULT : PHOTOMETRIC CALIBRATION GOOD TO 0.005 MAG

C. EXAMPLE SCIENTIFIC APPLICATIONS

• Subpixel registration of dithered HST imaging by A. Fruchter's "drizzling" technique. On left is one original frame (I-band, 2400 sec exposure). On right is result of drizzled processing of 12 such frames. Combined image has limiting magnitude of I ~ 28. It would be impossible to reach such levels with photographic detectors.

• Example of improvement in color-magnitude diagram for star cluster 47 Tuc. On left is 1977 state-of-the-art CMD based on widefield photographic images with photoelectric calibrations (Hesser & Hartwick, 1977). On right is a 1987 CMD derived from CCD photometry (Hesser et al. 1987). The greatly improved photometric precision reveals new features of astrophysical interest: e.g. the thinness of the main sequence near turnoff, which places strong limits on the range of ages present in the cluster.

 Photographic CMD CCD CMD

## XI. NON-UVOIR USE OF CCD & RELATED DEVICES

A. X-RAY: CHANDRA/ACIS DETECTOR

ACIS is a 10-CCD (1024x1024 chips) focal plane array used on Chandra for both imaging and spectroscopy. It uses both back-illuminated and front-illuminated versions.

It is operated in a pulse-detection mode, unlike the standard procedure at UVOIR wavelengths.

Each X-ray photon releases more than one electron in the CCD, in fact, the mean number released is ~ EXR/3.7 eV. Since Chandra operates at ~ 5 keV, the average electron cloud corresponding to one photon has ~ 1000 electrons.

The standard operation sequence is to expose for 3.2 seconds, then rapidly read out the array in 40 ms. The resulting image is analyzed by on-board software to catalogue the x,y position and the pulse amplitude of each valid photon pulse.

Because the pulse amplitude is proportional to the photon energy, ACIS achieves a spectral resolution of R ~ 10-50.

A difficulty with the ACIS design is that if more than one photon strikes the same pixel during the exposure time, the counting analysis becomes distorted, and the photon flux is underestimated. This is called "pileup." Fortunately, most X-ray sources are faint enough that this is not a problem.

B. INFRARED ARRAYS
The table in Sec. IV above shows that pure silicon photoconductor arrays will not work at IR wavelengths, but there are a number of other materials that will.

There are many varieties of IR detectors in use today. Some of these are monolithic, i.e. fabricated on single subtrates like CCD's, and some are hybrids in which the readout electronics are fabricated separately from the photon-detection devices.

Hybrids typically use silicon wafers for the readout electronics. Some actually use CCD-type architecture. The readout is connected to the photon-sensing material using conducting "bump-bonded" indium studs. If the wafers ultimately produce readout though a small number of output amplifiers, they are called "multiplexers" or "MUX's".

Here is a cross-section of the HST/NICMOS infrared detector, which uses a HgCdTe ("mercad-telluride") photon detector.

IR arrays will be discussed further in the guest lecture by Mike Skrutskie and in ASTR 512 (Spring 2004).

CCD-World (forum for discussions about CCDs and their use)

Basics:

Example Systems & Development Efforts:

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