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 represent the third of the three most important
technical innovations for observational UVOIR astronomy in the second
half of the 20th century (the others being space telescopes
and desktop/personal computers). 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
- Buil, CCD Astronomy
- Howell, Handbook of CCD Astronomy
- Kitchin, Astrophysical Techniques
- Rieke, Detection of Light : general introduction to detector physics
- LLM Sec. 7.3.3.
MacKay, Ann Rev Astr Ap, 24, 255, 1986
Timothy, PASP, 95, 810, 1983: discussion of many detector types
CCD Data Reduction and Applications
Situ CCD Testing"
- Massey, A User's Guide to CCD
Reductions with IRAF (1997)
JOSAA, 3, 2131, 1986: high precision technique
Dithering and CR Rejection in Undersampled HST Data"; also see
PASP, 106, 250, 1994 and references therein. Widely-used ALLFRAME
and DAOPHOT point-source CCD photometry programs.
Franx et al. AJ, 98, 538, 1989: sample application to
Bertin & Arnouts, A&AS, 317, 393, 1996 SExtractor---source
ID and extraction software for digital images. Obtain software
II. GENERAL DETECTOR CHARACTERIZATION
- QE = percentage of photons incident on detector which produce
- Strong wavelength dependence (e.g. threshold cutoffs set
by work function/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
- "Detective quantum efficiency" is defined to be
[(SNRout)/(SNRin)]2, where SNR
is the signal-to-noise ratio and "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.
- Wavelength region over which QE is
sufficient for operation
- 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 the 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
increase with the signal
- 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, we
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.
- Most UVOIR detectors are broad-band and have inherently poor spectral
resolution. One must use external elements (filters, gratings) to provide
- 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 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:
However, the photographic emulsion had serious limitations with which
astronomers had long struggled:
- Very long exposure times (up to a week in early applications),
meaning limiting fluxes thousands of times fainter
- A permanent,
objective record of astronomical images and spectra
- Geometric stability -- important for astrometry and morphology
- Large formats for wide field surveys
- Extension of observations to EM bands beyond the
"visual" band. The classic "photographic" band is centered around 4500
For details on astronomical photography at low light levels, see
Smith & Hoag
1979, ARAA, 17, 43.
- It is relatively insensitive,
with QE ~ 1%.
- Pg materials have non-negligible thresholds (a minimum of ~4
photoelectrons 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
- Each individual plate (or at the best, each co-processed "batch" of
plates) has different response characteristics.
Conversion of data to quantitative form (e.g. with microdensitometers) is
slow and cumbersome.
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
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
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,
where the average induced current was measured, 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
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 after 1950, e.g. the Johnson & Morgan broad-band UBV system
(see Lecture 14 for details).
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 efforts such as the Palomar Multichannel Photoelectric Spectrometer (Oke,
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
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
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 50 years have been solid
state, semiconductor arrays. These are based on
photo-conductive materials fabricated with embedded
microelectronic integrated circuits on thin wafers
Although lower quality devices can be mass-produced by microchip
"foundries," professional grade detectors still need to be
The image at the right shows a wafer containing a number
of co-fabricated CCDs and other devices.
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
Wide Field & Planetary Camera on the Hubble Space Telescope, led
by J. Westphal, J. Gunn, and C. R. Lynds.
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
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:
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.
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
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"
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: 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.
Doping affects energy levels
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. 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
layer of silicon-oxide has been grown. P-silicon can be made to have
very few 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 10 µ thick.
"BUCKET BRIGADE": The resulting
transfer and readout process is illustrated in the animation
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
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
The surface of the CCD is covered with MOS capactitors in a pattern
like that at the 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. 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
Each parallel transfer places the contents of one pixel from 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.
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
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
Difficulties with thinned chips:
- Use special thin, polysilicon material for electrodes; but these 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 a ~10 µ deep unit and illuminate from
backside; greatly improves blue response. The technique was pioneered
Lesser at the University of Arizona.
- 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).
Produce strong spatial modulation of response or "Fringing".
Especially serious for night sky emission lines. Can be well
calibrated for narrow-band filters or for broad-band filters. Harder
for intermediate band filters.
- Thinning reduces red response. For good response
at 5000-10000 Å, thick (~500 µ)
front-illuminated chips are preferred.
- CHARGE TRANSFER EFFICIENCY (CTE)
- CTE can be better than 99.999% per transfer, but it has to be, since
the throughput of a chip with 2048 required shifts = CTE2048.
- Radiation damage to CCD's in
decreases CTE over several years' time.
Operational: add (electronically or with diffuse light source) a
pedestal background signal (a "fat zero") over the whole chip to increase
the mean electron density per pixel. However, this adds additional noise and
compromises results on very faint sources.
Technical design: change number of phases, clocking cycles; add
VII. ADVANTAGES OF CCD DETECTORS
- HIGH QUANTUM EFFICIENCY: Reaching peak values of 80-90% in
the optical band. Much effort was expended to reach these high levels.
Sample response curves for three CCD's from ESO are shown below
- 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 detections
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.
- LINEAR RESPONSE across the whole range of operation until an
exposure approaches full-well capacity (typically 200,000 e/pixel).
This allows much improved flux calibrations and much higher precision
for flux measurements at all levels than other 2-D systems.
- EXCELLENT DYNAMIC RANGE: Typically 104.
- WIDE WAVELENGTH RESPONSE: Intrinsically sensitive from X-ray to
~1 µ. 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.
Dark current is largely suppressed by cooling. But noise is not
negligible. Typical equivalent read noise now is 2-10 e/pixel,
- SMALL PIXELS: typically 10-30 µ. Usually an advantage
in devising instrument/detector combinations, but should match pixel
size to 1/2 of smallest resolvable picture element in optical
- SPECIAL FORMAT/READOUT DESIGNS: By changing electrode structure
and clocking cycles, one can arrange for many different integrate/readout
Rapid clocking/video: inherent in CCDs intended for TV
applications and useful in some astronomical applications
(e.g. speckle interferometry). For bright sources, readout rates of
100 MHz are possible.
The basic drift-scanning technique
et al. 1994) is to clock the chip slowly along its
parallel registers while moving the telescope to keep a target
centered on the moving electron cloud as it builds up. The chip is
continuously read out 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 parallel registers
east-west, keep the telescope fixed on the sky, and clock the chip
westward at the sidereal rate. A sample application of drift scanning
Drift-scanning is now a standard method for wide-field CCD sky
surveys, most notably the Sloan Digital
"Nod and Shuffle": this
technique takes advantage of the capability to shift an image
wholesale on the CCD without it reading out to obtain much better sky
subtraction (e.g. in the near-IR where atmospheric OH emission is very
bright and variable).
On-chip binning is possible by changing clocking to combine
electrons from several pixels together before reading out.
E.g. combine a 2x2 pixel region on the chip into a single
output pixel. This suppresses the effect of amplifier readout noise on
each pixel in the final data and also reduces memory and storage
requirements. It adds considerable practical flexibility to CCD
systems. Obviously, however, binning reduces the resolution of the
output image. Binning 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
- UBIQUITY: CCDs are now almost universally used in
astronomy (amateur & pro). Photographic materials and older
electronic detectors are 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 near-real-time
assessment and hence 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).
VIII. LIMITATIONS/DISADVANTAGES OF CCD'S
- SMALL SIZE: Individual chips typically are less than 7 cm
square. They cover only a small fraction of the high quality imaging
field on typical modern telescopes. Device size is limited by the small
fabrication yield of large, defect-free chips. The largest routinely
available chips are 4096x2048.
- However, mosaicking technology is now well developed. Typical
mosaics are now constructed from 4096x2048 chips.
- The largest CCD mosaic camera,
containing 201 CCD's, is being built for
the Large Synoptic Survey
- Other 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 is necessary. Thermoelectric coolers are usually not
adequate, so cryogens (e.g. liquid N2) are required.
Introduces many practical complications. Chips with excellent low-T
performance differ from commercial types used at room temperature in
digital cameras, cellphones, etc.
- READOUT NOISE: Produced by variations in amplifier gain.
There has been much effort to reduce readout noise, which is now
typically equivalent to 2-10 e/pix.
Even at these very low levels, RON can
compromise some types of observations (spectroscopy, surface
photometry), see Lecture 7.
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.
RON effects in some applications can be reduced by 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, rather than CCD's.
- FULL WELL CONSTRAINTS: Bright sources which over-fill
pixels can produce "blooming" or "bleeding" along columns, making
parts of the chip even far from the source useless. The best solution is
- 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 special exposures ("dome flats"
or "twilight flats"). Special observing procedures to suppress 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 (see below). These methods
result in reduced field of view because not all parts of the original
field will have uniform exposures.
- FRINGING: see above. Caused by interference effects within
- SENSITIVITY TO COSMIC RAYS: High energy particles produce
strong, localized signals in CCDs, especially in thick chips. Their
effects must be removed during processing. CR's are a major
problem for spacecraft detectors (e.g. on HST). Mitigation
requires that exposures be broken into multiple parts (called
"CRSPLITs" on HST) so that CR events can be identified. CR hits can be
removed from processed 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
CCD Flatfield Frame (AAO)
CR's on HST/WFPC2 2400 sec exposure
- CHARGE TRANSFER EFFICIENCY: Because of "traps" within the
CCD's, good CTE is often possible only for signal levels above a
threshold of ~10-50 e/pix. For good efficiency at low light levels,
this requires adding a "fat zero" (typically 1000 e/pix)
electronically or by "preflash" from a diffuse light source. 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 $2B CCD market. A serious issue, since "active pixel
sensors," a different technology with amplifiers incorporated in each
pixel, now dominate the larger imaging array industry. Requirements
for good performance at low (astronomical) light levels are
considerably different than for room-temperature, short-exposure,
- 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 solid-state storage are now much improved, but
long-term stability of the massive amounts of data now being produced
is a non-trivial issue.
Cf. the LSST. [All that data is
really, 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, 15 µ pixel chips
- Front-illuminated, lumogen phosphor coating for UV response
- RON 5 e/px; QE (6000 Å) 35%
- SLOAN DIGITAL SKY SURVEY (1997)
- Drift scanning mosaic (not contiguous) containing 54 chips
- Tektronix/SITe 2048x2048, 24 µ pixel chips in 5x6
array; for simultaneous broad-band photometry. Both frontside and
thinned backside with AR coating. QE (6000 Å) ~60%; RON < 5
- Tektronix/SITe 2048x400, 24 µ pixel chips; for
astrometry & focus check. Frontside illuminated, RON < 15 e/px.
- HST WIDE FIELD CAMERA 3 (2001)
- 2 CCD's, physically mosaicked, contiguous
- Marconi 4096x2048, 15 µ pixel chips
- Thinned, back-illuminated; anti-reflection coating for enhanced blue/UV
- Minichannel for improved CTE
- RON 3 e/px; QE (6000 Å) 70%
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 AND DATA-TAKING PROCEDURES
- SUBTRACT BIAS FRAME: Bias frames (zero exposure time) are
taken with the chip unexposed to light. These measure the electronic
pedestal of the chip. For high precision, average many bias frames,
before and after observations. Check for bias drift during the night.
With some chips, you can determine the electronic bias level from the
"overscan" region on the chip.
- 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: This removes the
residual pixel-to-pixel variations (typical 5%). Make exposures of
the 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 because of
color sensitive effects.
[Note: photon statistics must yield a SNR significantly in excess of
the final desired SNR -- i.e. if you want SNR = 100, the net
flat field image must contain well over 10,000
photons per pixel.]
- Optional: SUBTRACT SKY FLAT/FRINGE FRAME: Remove night sky
emission line fringing effects (worse in near-IR) by observing an
uncrowded field in the night sky. Take several exposures, moving
telescope by, say, 25 arc-sec between them. Use the "adaptive modal
filter" technique to zap star images and create a mean sky frame.
Scale to target frames and subtract. Resulting flat field as good as
0.1% can result. A sky flat determined this way is needed for each
- TARGET FRAMES: USE MULTIPLE OFFSETS, DITHERING OR DRIFT
SCANNING: For faintest possible photometry, use the
multiple-offset exposure technique -- e.g. 500 sec exposures shifted by
about 20 arc-sec each -- to reduce flat field errors and
identify cosmic rays.
You always want more than one exposure anyway for "reality
checks" and empirical determination of photometric errors and
- REGISTER AND COMBINE TARGET FRAMES: Re-register the offset
frames to sub-pixel accuracy (using, e.g., "Drizzle" softare),
creating an image "stack."
- REMOVE COSMIC RAY EVENTS: Median filter the signals in each
pixel in the stack. Outlying signals caused by cosmic rays and
other transients are thereby rejected from the combined image. You
must track the remaining "good" exposure time for each pixel.
- COSMETIC TREATMENT: trim off the "lost" parts of the
image. Depending on the application, you can interpolate over or mask
out bad pixels.
- 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 (e.g. variable clouds), but no flux
calibration for those, either. Calibrator stars should overlap
targets of interest in color. These frames are reduced exactly like
target frames, apparent fluxes extracted, and the multiplicative
factors needed to convert to 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. See image pair below. On the
left is one original frame (I-band, 2400 sec exposure). Most of the
"features" in this image are cosmic ray tracks, not cosmic objects.
On the right is the result of drizzled processing of 12 such frames.
The combined image has a limiting magnitude of I ~ 28. It would be
impossible to reach such levels with photographic plates or other
less-stable and calibrateable systems..
- Example of improvement in a color-magnitude diagram for
the globular star cluster 47 Tuc. On the left is a 1977
state-of-the-art CMD based on widefield photographic images with
photoelectric calibrations (Hesser & Hartwick, 1977). On the 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. (Later, yet higher precision, CCD imaging with HST in
the near-ultraviolet bands revealed the existence of multiple
populations with similar ages but different metal and He abundances in
many globular clusters.)
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
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".
a cross-section of the HST/NICMOS infrared detector, which uses
a HgCdTe ("mercad-telluride") photon detector.
ADDITIONAL WEB LINKS
CCD-World (forum for
discussions about CCDs and their use in astronomy)
Example Systems & Development Efforts:
June 2017 by rwo
Bandgap images from Britney's Guide to Semiconductors.
MOS capacitor illustration from Molecular Expressions. Bucket brigade
animation and front/back illumination drawing by C. Tremonti (UWisc).
CCD design drawings from C. MacKay, Annual Reviews (1986). Most other
images from public observatory sites. Text copyright © 1989-2017
Robert W. O'Connell. All rights reserved. These notes are intended
for the private, noncommercial use of students enrolled in Astronomy
511 at the University of Virginia.