Transcript Introduction to CCDs

Introduction to CCDs
Basic principles of CCD imaging
Original notes by S.Tulloch & G.Bonanno modified and integrated by GPS
(Rev.103 Nov08)
What is a CCD ?
Charge Coupled Devices (CCDs) were invented in the 1970s and originally found application as memory devices
and more recently as shift registers, analog delay lines, particle detectors.
Their light sensitive properties were quickly exploited for imaging applications (TV cameras) and they produced a
major revolution in Astronomy. They improved the light gathering power of telescopes by almost two orders of
magnitude. Nowadays an amateur astronomer with a CCD camera and a 15 cm telescope can collect as much light
as an astronomer of the 1960s equipped with a photographic plate and a 1m telescope.
CCDs work by converting light into a pattern of electronic charge in a silicon chip. This pattern of charge is converted
into a video waveform, digitised and stored as an image file on a computer.
The sensitive surface of a CCD is composed of a large number of juxtaposed square or rectangular elements (pixels)
sensitive to light; each pixel (a few microns size) creates and accumulates electrical charge proportional to the amount
of light it has received and constitutes an elementary point of the image, as does each grain of a photographic film. It is
the reading of the accumulated charges in the different pixels that allow the image to be reconstructed.
CCDs used in astronomy are all black and white detectors; it is possible to produce color images by trichromatism
(three images of the same object, each one with one of the three fundamental filters and combine these images
during analysis). CCDs have practically no detection threshold (unlike films) and the number of electrons generated
gives a direct photometric measure.
•
•
Il rumore intrinseco di un
rivelatore e’ tutto ciò che e’
introdotto dal rivelatore stesso e
non dal segnale. Determina la
soglia inferiore di rilevabilità (con
basso rumore intrinseco si
osservano sorgenti più deboli)
Il rumore di un CCD e’ dovuto a:
–
–
•
•
•
Agitazione termica (trascurabile
al di sotto di -100°C
Rumore di lettura del circuito di
uscita (ineliminabile)
Se il rivelatore non introducesse
rumore, l’unico a restare sarebbe
quello della sorgente stessa
(RSHOT = √Nfotoni).
Per fare analisi quantitative il
segnale deve superare di almeno
3 volte il rumore
Il minimo valore di segnale
misurabile e’ detto “sensibilità” e
si può considerare dato dal
rumore del rivelatore
“Dynamic range”: the ratio,
usually expressed in decibels, of
the maximum to the minimum
signal that a system can handle
(Grey levels). Used to describe
the limits of receivers.
Photoelectric Effect
Increasing energy
The effect is fundamental to the operation of a CCD. Atoms in a silicon crystal have electrons arranged in
discrete energy bands. The lower energy band is called the Valence Band, the upper band is the Conduction
Band. Most of the electrons occupy the Valence band but can be excited into the conduction band by heating
or by the absorption of a photon. The energy required for this transition is 1.26 electron volts (eV). Once in this
conduction band the electron is free to move about in the lattice of the silicon crystal. It leaves behind a ‘hole’
in the valence band which acts like a positively charged carrier. In the absence of an external electric field
the hole and electron will quickly re-combine and be lost. In a CCD an electric field is introduced to sweep
these charge carriers apart and prevent recombination.
Conduction Band
1.26eV
Valence Band
Hole
Electron
Thermally generated electrons are indistinguishable from photo-generated electrons . They constitute a
noise source known as ‘Dark Current’ and it is important that CCDs are kept cold to reduce their number.
1.26eV corresponds to the energy of light with a wavelength of 1mm. Beyond this wavelength silicon becomes
transparent and CCDs constructed from silicon become insensitive.
CCD Analogy
A common analogy for the operation of a CCD is as follows:
An number of buckets (Pixels) are distributed across a field (Focal Plane of a telescope)
in a square array. The buckets are placed on top of a series of parallel conveyor belts and collect rain fall
(Photons) across the field. The conveyor belts are initially stationary, while the rain slowly fills the
buckets (During the course of the exposure). Once the rain stops (The camera shutter closes) the
conveyor belts start turning and transfer the buckets of rain , one by one , to a measuring cylinder
(Electronic Amplifier) at the corner of the field (at the corner of the CCD)
The animation in the following slides demonstrates how the conveyor belts work.
CCD Analogy
CCD COLUMNS
PHOTONS
(VERTICAL
CONVEYOR BELTS)
PIXELS (BUCKETS)
SERIAL REGISTER
(HORIZONTAL
CONVEYOR BELT)
OUTPUT
AMPLIFIER
Exposure finished, buckets now contain samples of rain, i.e. charge has been collected in pixels.
Conveyor belt starts turning and transfers buckets. Rain collected on the vertical conveyor
is tipped into buckets on the horizontal conveyor.
SERIAL REGISTER
(HORIZONTAL
CONVEYOR BELT)
Vertical conveyor stops. Horizontal conveyor starts up and tips each bucket in turn into
the measuring cylinder .
After each bucket has been measured, the measuring cylinder
is emptied , ready for the next bucket load.
`
OUTPUT
AMPLIFIER
A new set of empty buckets is set up on the horizontal conveyor and the process
is repeated.
Eventually all the pixels have been measured, the CCD has been read out.
CCD main functions
Exposure
1. Charge generation: photoelectric effect
2. Charge collection: potential walls (pixels)
Shutter closes
3. Charge transfer: CCD columns shift
4. Readout: serial register and amplification
Structure of a CCD 1.
The image area of the CCD is positioned at the focal plane of the telescope. An image then builds up that
consists of a pattern of electric charge. At the end of the exposure this pattern is then transferred, pixel by
pixel, through the serial register to the on-chip amplifier. Electrical connections are made to the outside
world via a series of bond pads and thin gold wires positioned around the chip periphery.
Image area
Metal, ceramic or plastic package
Connection pins
CCD: 512x512 pixels,
1024x1024, 2048x2048,
4096x4096
Sensitive surface:
Gold bond wires •TC211: 2.64x2.64mm2
(192x165 pixels, 13.75x16µm
Bond pads
size)
•Thomson 7863: 8.8x6.6mm2
(384x288 pixels, 23x23µm)
•SITe SI-003A: 24.6x24.6mm2
(1024x1024 pixels, 24x24µm)
•LORAL 2K3eb: 30.7x30.7mm2
Silicon chip
(2048x2048 pixels, 15x15µm)
On-chip amplifier
Serial register
CCD arrays may be smaller in
comparison with standard
photography (24x36mm).
“Buttable” CCDs for mosaic.
Structure of a CCD 2.
CCDs are are manufactured on silicon wafers using the same photo-lithographic techniques used
to manufacture computer chips. Scientific CCDs are very big ,only a few can be fitted onto a wafer.
This is one reason that they are so costly.
The photo below shows a silicon wafer with three large CCDs and assorted smaller devices. A CCD has
been produced by Philips that fills an entire 6 inch wafer! It is the worlds largest integrated circuit.
Don Groom LBNL
Structure of a CCD 3.
The diagram shows a small section (a few pixels) of the image area of a CCD. This pattern is reapeated.
Channel stops to define the columns of the image
Plan View
One pixel
Cross section
Transparent
horizontal electrodes
to define the pixels
vertically. Also
used to transfer the
charge during readout
Electrode
Insulating oxide
n-type silicon
p-type silicon
Every third electrode is connected together. Bus wires running down the edge of the chip make the
connection. The channel stops are formed from high concentrations of Boron in the silicon.
Structure of a CCD 4.
Below the image area (the area containing the horizontal electrodes) is the ‘Serial register’ . This also
consists of a group of small surface electrodes. There are three electrodes for every column of the image area.
Image Area
On-chip amplifier
at end of the serial
register
Serial Register
Cross section of
serial register
Once again every third electrode in the serial register is connected together.
Structure of a CCD 5.
Details of a corner of an EEV CCD
160mm
Bus wires
Serial Register
Read Out Amplifier
Edge of
Silicon
Image Area
The serial register is bent twice to move the output amplifier away from the edge
of the chip (Voltage bus). The arrows indicate how charge is transferred through the device.
Structure of a CCD 6.
Details of the on-chip amplifier of a Tektronix CCD and its circuit diagram
20mm
Output Drain (OD)
Gate of Output Transistor
Output Source (OS)
SW
R
RD
OD
Output Node
Reset
Transistor
Reset Drain (RD)
Summing
Well
R
Output
Node
Serial Register Electrodes
Output
Transistor
OS
Summing Well (SW)
Last few electrodes in Serial Register
Substrate
Electric Field in a CCD 1.
Electric potential
The n-type layer contains an excess of electrons that diffuse into the p-layer. The p-layer contains an
excess of holes that diffuse into the n-layer. This structure is identical to that of a diode junction.
The diffusion creates a charge imbalance and induces an internal electric field. The electric potential
reaches a maximum just inside the n-layer, and it is here that any photo-generated electrons will collect.
All science CCDs have this junction structure, known as a ‘Buried Channel’. It has the advantage of
keeping the photo-electrons confined away from the surface of the CCD where they could become trapped.
It also reduces the amount of thermally generated noise (dark current).
p
n
Potential along this line shown
in graph above.
Cross section through the thickness of the CCD
Electric Field in a CCD 2.
Electric potential
During integration of the image, one of the electrodes in each pixel is held at a positive potential. This
further increases the potential in the silicon below that electrode and it is here that the photoelectrons are
accumulated. The neighboring electrodes, with their lower potentials, act as potential barriers that define
the vertical boundaries of the pixel. The horizontal boundaries are defined by the channel stops.
p
n
Region of maximum
potential
Charge Collection in a CCD.
Charge packet
pixel
boundary
pixel
boundary
incoming
photons
Photons entering the CCD create electron-hole pairs. The electrons are then attracted towards
the most positive potential in the device where they create ‘charge packets’. Each packet
corresponds to one pixel
n-type silicon
Electrode Structure
p-type silicon
SiO2 Insulating layer
Charge Transfer in a CCD 1.
In the following few slides, the implementation of the ‘conveyor belts’ as actual electronic
structures is explained.
The charge is moved along these conveyor belts by modulating the voltages on the electrodes
positioned on the surface of the CCD. In the following illustrations, red electrodes
are held at a positive potential, black ones are held at a negative potential.
1
2
3
Charge Transfer in a CCD 2.
+5V
2
0V
-5V
+5V
1
0V
-5V
+5V
3
0V
-5V
1
2
3
Time-slice shown in diagram
Charge Transfer in a CCD 3.
+5V
2
0V
-5V
+5V
1
0V
-5V
+5V
3
0V
-5V
1
2
3
Charge Transfer in a CCD 4.
+5V
2
0V
-5V
+5V
1
0V
-5V
+5V
3
0V
-5V
1
2
3
Charge Transfer in a CCD 5.
+5V
2
0V
-5V
+5V
1
0V
-5V
+5V
3
0V
-5V
1
2
3
Charge Transfer in a CCD 6.
+5V
2
0V
-5V
+5V
1
0V
-5V
+5V
3
0V
-5V
1
2
3
Charge Transfer in a CCD 7.
+5V
2
0V
-5V
Charge packet from subsequent pixel enters
from the left as first pixel exits to the right.
+5V
1
0V
-5V
+5V
3
0V
-5V
1
2
3
Charge Transfer in a CCD 8.
+5V
2
0V
-5V
+5V
1
0V
-5V
+5V
3
0V
-5V
1
2
3
On-Chip Amplifier 1.
The on-chip amplifier measures each charge packet as it pops out the end of the serial register.
+5V
RD and OD are held at
constant voltages
SW
R
RD
SW
0V
-5V
OD
+10V
R
0V
Reset
Transistor
Summing
Well
--end of serial register
Output
Node
Vout
Output
Transistor
(The graphs above show the signal waveforms)
OS
Vout
The measurement process begins with a reset
of the ‘reset node’. This removes the charge
remaining from the previous pixel. The reset
node is in fact a tiny capacitance (< 0.1pF)
On-Chip Amplifier 2.
The charge is then transferred onto the Summing Well. Vout is now at the ‘Reference level’
+5V
SW
SW
R
RD
0V
-5V
OD
+10V
R
0V
Reset
Transistor
Summing
Well
--end of serial register
Output
Node
Vout
Output
Transistor
OS
Vout
There is now a few tens of µs wait while external
circuitry measures this ‘reference’ level.
On-Chip Amplifier 3.
The charge is then transferred onto the output node. Vout now steps down to the ‘Signal level’
+5V
SW
SW
R
RD
0V
-5V
OD
+10V
R
0V
Reset
Transistor
Summing
Well
--end of serial register
Output
Node
Vout
Output
Transistor
This action is known as the ‘charge dump’
OS
Vout
The voltage step in Vout is as much as
several mV for each electron contained
in the charge packet.
On-Chip Amplifier 4.
Vout is now sampled by external circuitry for up to a few tens of microseconds.
+5V
SW
SW
R
RD
0V
-5V
OD
+10V
R
0V
Reset
Transistor
Summing
Well
--end of serial register
Output
Node
Vout
Output
Transistor
OS
Vout
The sample level - reference level will be
proportional to the size of the input charge
packet.
Use of CCD Cameras
Practical considerations of using
and building CCD cameras
Nik Szymanek
Spectral Sensitivity of CCDs
Transmission of Atmosphere
The graph below shows the transmission of the atmosphere when looking at objects at the
zenith. The atmosphere absorbs strongly below about 330nm, in the near ultraviolet part of
the spectrum. An ideal CCD should have a good sensitivity from 330nm to approximately
1000nm, at which point silicon, from which CCDs are manufactured, becomes transparent
and therefore insensitive.
Wavelength (Nanometers)
Over the last 25 years of development, the sensitivity of CCDs has improved enormously, to
the point where almost all of the incident photons across the visible spectrum are detected.
CCD sensitivity has been improved using two main techniques: ‘thinning’ and the use of
anti-reflection coatings. They will be now explained in more details.
A CCD sensitivity (Quantum
Efficiency, QE) is expressed as
the number of electrons
produced per incident photon.
To improve UV sensitivity:
•Anti-reflection coating:
fluorescent coat excited by UV
light, re-emits in the green
•Thinning: Back-lit CCDs
Pixel capacity:105-106
electrons. Saturation,
“blooming” (bright stars)
(1 nm = 10 angstroms)
Incoming photons
Thick Front-side Illuminated CCD
p-type silicon
n-type silicon
625mm
Silicon dioxide insulating layer
Polysilicon electrodes
These are cheap to produce using conventional wafer fabrication techniques. They are used in
consumer imaging applications. Even though not all the photons are detected, these devices
are still more sensitive than photographic film.
They have a low Quantum Efficiency due to the reflection and absorption of light in the
surface electrodes. Very poor blue response. The electrode structure prevents the use of
an Anti-reflective coating that would otherwise boost performance.
The amateur astronomer on a limited budget might consider using thick CCDs. For
professional observatories, the economies of running a large facility demand that the detectors
be as sensitive as possible; thick front-side illuminated chips are seldom if ever used.
Anti-Reflection Coatings 1
Silicon has a very high Refractive Index (denoted by n). This means that photons are strongly reflected
from its surface.
ni
nt
Fraction of photons reflected at the
interface between two mediums of
differing refractive indices
=
[
nt-ni
nt+ni
2
]
n of air or vacuum is 1.0, glass is 1.46, water is 1.33, Silicon is 3.6. Using the above equation we can
show that window glass in air reflects 3.5% and silicon in air reflects 32%. Unless we take steps to
eliminate this reflected portion, then a silicon CCD will at best only detect 2 out of every 3 photons.
The solution is to deposit a thin layer of a transparent dielectric material on the surface of the CCD. The
refractive index of this material should be between that of silicon and air, and it should have an
optical thickness = 1/4 wavelength of light. The question now is what wavelength should we choose, since
we are interested in a wide range of colours. Typically 550nm is chosen, which is close to the middle of the
optical spectrum.
Anti-Reflection Coatings 2
With an Anti-reflective coating we now have three mediums to consider :
ni
ns
nt
Air
AR Coating
Silicon
The reflected portion is now reduced to :
In the case where
[
2
n t x n i- n s
2
nt x ni+ns
2
]
n2s = nt the reflectivity actually falls to zero! For silicon we require a material
with n = 1.9, fortunately such a material exists, it is Hafnium Dioxide. It is regularly used to coat
astronomical CCDs.
Anti-Reflection Coatings 3
The graph below shows the reflectivity of an EEV 42-80 CCD. These thinned CCDs were designed
for a maximum blue response and it has an anti-reflective coating optimised to work at 400nm. At this
wavelength the reflectivity falls to approximately 1%.
Incoming photons
Thinned Back-side Illuminated CCD
15mm
Anti-reflective (AR) coating
p-type silicon
n-type silicon
Silicon dioxide insulating layer
Polysilicon electrodes
The silicon is chemically etched and polished down to a thickness of about 15microns. Light enters
from the rear and so the electrodes do not obstruct the photons. The QE can approach 100% .
These are very expensive to produce since the thinning is a non-standard process that reduces the
chip yield. These thinned CCDs become transparent to near infra-red light and the red response is
poor. Response can be boosted by the application of an anti-reflective coating on the thinned
rear-side. These coatings do not work so well for thick CCDs due to the surface bumps created
by the surface electrodes.
Almost all Astronomical CCDs are Thinned and Backside Illuminated.
Quantum Efficiency Comparison
The graph below compares the quantum of efficiency of a thick frontside illuminated CCD and a
thin backside illuminated CCD.
‘Internal’ Quantum Efficiency
If we take into account the reflectivity losses at the surface of a CCD we can produce a graph showing
the ‘internal QE’ : the fraction of the photons that enter the CCDs bulk that actually produce a
detected photo-electron. This fraction is remarkably high for a thinned CCD. For the EEV 42-80 CCD,
shown below, it is greater than 85% across the full visible spectrum. Todays CCDs are very close to
being ideal visible light detectors!
Appearance of CCDs
The fine surface electrode structure of a thick CCD is clearly visible as a multi-coloured
interference pattern. Thinned Backside Illuminated CCDs have a much planer surface
appearance. The other notable distinction is the two-fold (at least) price difference.
Kodak Kaf1401 Thick CCD
MIT/LL CC1D20 Thinned CCD
Computer Requirements 1.
Computers are required firstly to coordinate the sequence of clock signals that need to be sent to a CCD
and its signal processing electronics during the readout phase, but also for data collection and the
subsequent processing of the images.
The CCD Controller
In this first application, the computer is an embedded system running in a ‘CCD controller’. This controller will
typically contain a low noise analogue section for amplification and filtering of the CCD video waveform,
an analogue to digital converter, a high speed processor for clock waveform generation and a fibre optic
transceiver for receipt of commands and transmission of pixel data.
An astronomical system might require clock signals to be generated with time resolutions of a few tens
of nanoseconds. This is typically done using Digital Signal Processing (DSP) chips running at 50Mhz. Clock
sequences are generated in software and output from the DSP by way of on-chip parallel ports. The most
basic CCD design requires a minimum of 7 clock signals. Perhaps 5 more are required to coordinate the
operation of the signal processing electronics. DSPs also contain several on-chip serial ports
which can be used to transmit pixel data at very high rates. DSPs come with a small on-chip memory for
the storage of waveform generation tables and software. Less time critical code , such as routines to
initialise the camera and interpret commands can be stored in a few KB of external RAM. The computer
running in the CCD controller is thus fast and of relatively simple design. A poorly performing processor here
could result in slow read out times and poor use of telescope resources. Remember that when a CCD is reading
out the telescope shutter is closed and no observations are possible. For an amateur observer using a small CCD
with a fast readout time, a slow CCD controller may not be such a disadvantage; there are not so many pixels
to process.
Computer Requirements 2.
The Data Acquisition System(DAS)
This will be typically based around a SUN SPARC workstation which is a high-end desktop computer. Pixel
data will be received from the CCD controller by way of a fibre optic. The hardware in such a system will be
cheap and ‘off-the shelf’, the only speciality item being the high speed fibre optic transceiver card.
The hardware may typically consist of a Sparc Ultra 6 workstation, 500Mb of RAM, a 9GB hard-drive
and a DAT drive. There will also be a high speed Ethernet card for connection to the observatory Local Area
Network.The software required to carry out the data acquisition task is typically developed in-house by each
observatory and represents the major cost of such a system. It will provide an easy-to-use interface (typically
graphics based) between observer and instrument. Its complexity will be further increased by the need to talk to
other telescope systems such as the Telescope Control System. This will allow information on the pointing
of the telescope to be stored alongside the pixel data as a ‘file header’.
Computer Requirements 3.
Image Processing Computers
These are used for reduction and analysis of the astronomical data. Many astronomers process their data
in real-time, i.e. they may be analysing one exposure whilst the next exposure is actually been taken. Others
will take a cursory look at their data in real time but leave the heavy image processing tasks for when they
return to their home institution. With large mosaic cameras producing very large data files, a high end system
is required.
A typical system would be :
•A PC with a 1GHz CPU
•Enough RAM for at least 2 images , using 4 bytes per pixel (for a mosaic camera this could run to 500MBytes)
•At least 100 GBytes (300GBytes would be better) of local hard disc space
If we use such a system to analyse images from a four chip CCD mosaic containing 36 Million pixels, the following
performance would be obtained :
Linearisation, bias subtraction and flat-fielding : ~150 sec
de-fringing
: ~300 sec
object detection and star/galaxy separation
: ~300 sec
Computer Requirements 4.
Image Processing Computers (Contd.)
This professional system is unusual in its high demands on disc space and RAM. The processor speed , however, is
the same as that found in current PCs costing a few thousand dollars. An amateur observer with a small 1K square
CCD camera will find a medium level PC quite sufficient for operation of the camera and for image processing.
The system specs would typically be:
•
Pentium III 500 MHz processor,
•
256 MB RAM
•
32 MB video memory
•
30GB Hard Drive
•
CD Writer
•
a 19” monitor (twin monitors are even better, one for images , one for text)
For operation of the camera the bottleneck is often the data transfer between camera and PC. For image
processing applications such as Maximum Entropy or Lucy-Richardson de-convolution
(a form of image sharpening), a high speed PC is needed.
Schematic view of CCD camera electronics
Acquisition
commands
CCD camera
Clocks
CCD
Cooling
system
I/O
interface
Computer
Amplifier
Digital
converter
Image
Image defects in CCDs
Blooming: limited capacity of a CCD pixel
Dark columns: traps
Cosmic rays
Bright columns: traps
Hot Spots: non uniformity of structure
Electronic processing
Calibration exposures
Blooming 1.
The charge capacity of a CCD pixel is limited: when a pixel is full, the charge starts to leak into
adjacent pixels. This process is known as ‘Blooming’.
pixel
boundary
Photons
pixel
boundary
Overflowing
charge packet
Spillage
Photons
Spillage
Blooming 2.
The diagram shows one column of a CCD with an over-exposed stellar image focused on one pixel.
The channel stops shown in yellow prevent the charge
spreading sideways. The charge confinement provided by
the electrodes is lower so the charge spreads vertically up
and down a column.
The capacity of a CCD pixel is known as the ‘Full Well’. It is
dependent on the physical area of the pixel. For Tektronix
CCDs, with pixels measuring 24mm x 24mm it can be as much as
300,000 electrons. Bloomed images will be seen particularly
on nights of good seeing where stellar images are more compact .
Flow of
bloomed
charge
In reality, blooming is not a big problem for professional
astronomy. For those interested in pictorial work, however, it can
be a nuisance.
Blooming 3.
The image below shows an extended source with bright embedded stars. Due to the long
exposure required to bring out the nebulosity, the stellar images are highly overexposed
and create bloomed images.
M42
Bloomed star images
(The image is from a CCD mosaic and the black strip down the center is the space between adjacent detectors)
Dark columns
Unless one pays a huge amount it is generally difficult to obtain a CCD free of image defects.
The first kind of defect is a ‘dark column’. Their locations are identified from flat field exposures.
Dark columns are caused by ‘traps’ that block the vertical
transfer of charge during image readout. The CCD shown at
left has at least 7 dark columns, some grouped together in
adjacent clusters.
Traps can be caused by crystal boundaries in the silicon of
the CCD or by manufacturing defects.
Although they spoil the chip cosmetically, dark columns are
not a big problem for astronomers. This chip has 2048 image
columns so 7 bad columns represents a tiny loss of data.
Flat field exposure of an EEV42-80 CCD
Cosmic rays, Hot spots, Bright columns
There are three more image defect types : cosmic rays, bright columns and hot spots.
Their locations are shown in the image below, which is a lengthy exposure taken in the dark (a ‘Dark Frame’)
Bright
Column
Bright columns are also caused by traps . Electrons contained
in such traps can leak out during readout causing a vertical streak
Hot Spots are pixels with higher than normal dark current. Their
brightness increases linearly with exposure times
Cluster of
Hot Spots
Cosmic rays
Cosmic rays are unavoidable. Charged particles from space or
from radioactive traces in the material of the camera can
cause ionisation in the silicon. The electrons produced are
indistinguishable from photo-generated electrons.
Approximately 1-2 cosmic rays/cm2/minute will be seen.
A typical event will be spread over a few adjacent pixels and
contain several thousand electrons.
Somewhat rarer are light-emitting defects which are hot spots
that act as tiny LEDS and cause a halo of light on the chip.
900s dark exposure of an EEV42-80 CCD
Electronic processing
Some defects can arise from the read-out processing electronics. This negative image has a
bright line in the first image row.
M51
Dark column
Hot spots and bright columns
Bright first image row caused by
incorrect operation of signal
processing electronics.
Calibration exposures
Bias frames
CCD defects, readout noise
exposure time = 0
Flat field
pixel non-uniformity, dust, …
exposure time ≈ medium (half saturation)
Dark frames
dark current, defects
exposure time = science image
Biases, Flat Fields and Dark Frames 1.
These are three types of calibration exposures that must be taken with a scientific CCD camera,
generally before and after each observing session. They are stored alongside the science images
and combined with them during image processing. These calibration exposures allow us to compensate for
certain imperfections in the CCD. As much care needs to be exercised in obtaining these images as for
the actual scientific exposures. Applying low quality flat fields and bias frames to scientific data can
degrade its quality, rather than improve it.
Bias Frames (CCD defects, readout noise)
A bias frame is an exposure of zero duration taken with the camera shutter closed. It represents the zero
point or base-line signal from the CCD. Rather than being completely flat and featureless, the bias frame
may contain some structure. Any bright image defects in the CCD will of course show up, there may also be
slight gradients in the image caused by limitations in the signal processing electronics of the camera.
It is normal to take about 5 bias frames before a night’s observing. These are then combined using an image
processing algorithm that averages the images, pixel by pixel, rejecting any pixel values that are appreciably
different from the other 4. This can happen if a pixel in one bias frame is affected by a cosmic ray event. It
is unlikely that the same pixel in the other 4 frames would be similarly affected so the resultant ‘master bias’,
should be uncontaminated by cosmic rays. Taking a number of biases and then averaging them also reduces
the amount of noise in the bias images. Averaging 5 frames will reduce the amount of read noise (electronic
noise from the CCD amplifier) in the image by the square-root of 5.
Biases, Flat Fields and Dark Frames 2.
Flat Fields (pixel non-uniformity, dust, etc.)
Some pixels in a CCD will be more sensitive than others (non-uniformity due to slightly different pixel
dimensions and depth because of manufacturing process). In addition there may be dust spots on the surface
of either the chip, the window of the camera or the colored filters mounted in front of the camera. A star
focused onto one part of a chip may therefore produce a lower signal than it might do elsewhere. These
variations in sensitivity across the surface of the CCD must be calibrated out or they will add noise to the
image. The way to do this is to take a ‘flat-field ‘ image: an image in which the CCD is evenly illuminated
with light. Dividing the science image, pixel by pixel, by a flat field image will remove these non-uniformities
and sensitivity variations very effectively.
Since some of these variations are caused by shadowing from dust spots, it is important that the flat fields
are taken shortly before or after the science exposures as the dust may move around. As with biases, it is
normal to take several flat field frames and average them to produce a ‘Master’.
A flat field is taken by pointing the telescope at an extended, evenly illuminated source. The twilight sky or
the inside of the telescope dome are the usual choices. An exposure time is chosen that gives pixel values
about halfway to their saturation level, i.e. a medium level exposure.
Dark Frames (dark current, defects)
Dark current is generally absent from professional cameras since they are operated at low temperature, using liquid
nitrogen as a coolant. Amateur systems running at higher temperatures will have some dark current and its effect
must be minimized by obtaining ‘dark frames’ at the beginning of the observing run. These are exposures with
the same duration as the science frames but taken with the camera shutter closed. These are later subtracted
from the science frames. Again, it is normal to take several dark frames and combine them to form a Master,
using a technique that rejects cosmic ray features.
Biases, Flat Fields and Dark Frames 3.
A dark frame and a flat field from the same EEV42-80 CCD are shown below. The dark frame shows
a number of bright defects on the chip. The flat field shows a cross patterning on the chip
created during manufacture and a slight loss of sensitivity in two corners of the image. Some dust
spots are also visible.
Dark Frame
Flat Field
Biases, Flat Fields and Dark Frames 4.
All raw images contain undesirable effects, inherent to the CCD
system technology used; preprocessing consists in “cleaning” the raw
image of these effects.
Science Raw Image
Dark Frame
Science
-Dark
Output Image
Science -Dark
Flat Field Frame
Flat-Bias
Flat
-Bias
Bias Frame
Bias Frames and Flat Fields 5.
In the absence of dark current, the process is slightly simpler :
Science Frame
Bias Image
Science
-Bias
Output Image
Science -Bias
Flat-Bias
Flat Field Image
Flat
-Bias
Pixel Size and Binning 1.
Nyquist Sampling
It is important to match the size of a CCD pixel to the focal length of the telescope. Atmospheric seeing
places a limit on the sharpness of an astronomical image for telescope apertures above 15cm. Below this
aperture, the images will be limited by diffraction effects in the optics. In excellent seeing conditions, a large
telescope can produce stellar images with a diameter of 0.6 arc-seconds. In order to record all the information
present in such an image, two pixels must fit across the stellar image; the pixels must subtend at most
0.3 arc-seconds on the sky. This is the ‘Nyquist criteria’. If the pixels are larger than 0.3 arc-seconds the
Nyquist criteria is not met, the image is under-sampled and information is lost. The Nyquist criteria also
applies to the digitisation of audio waveforms. The audio bandwidth extends up to 20KHz , so the Analogue
to Digital Conversion rate needs to exceed 40KHz for full reproduction of the waveform. Exceeding the
Nyquist criteria leads to ‘over-sampling’.This has the disadvantage of wasting silicon area; with improved
matching of detector and optics a larger area of sky could be imaged.
Under-sampling an image can produce some interesting effects. One of these is the introduction of
features that are not actually present. This is occasionally seen in TV broadcasts when, for example,
the fine-patterned shirt of an interviewee breaks up into psychedelic bands and ripples. In this example,
the TV camera pixels are too big to record the fine detail present in the shirt.
This effect is known as ‘aliasing’.
Nyquist/Shannon Theorem: fc≥2fmax (fc=sampling freq, fmax=signal max freq)
Aliasing
When a sinusoidal signal is measured or sampled at regular but not sufficiently close intervals, one will obtain
the same sequence of samples that would be obtained from a sinusoid of a lower frequency. Specifically, if a
sinusoid of frequency f (in cycles per second for a time-varying signal, or in cycles per centimeter for spacevarying signal) is sampled fs samples per second or per centimeter, the resulting samples will also be compatible
with a sinusoid of frequency (Nfs-f) and one of frequency (Nfs+f), for any integer N. If fs>2f , the lowest of
these image frequencies will be the original signal frequency, but otherwise it will not. In the case that fs<2f<2fs ,
the lowest image frequencies will be at (fs-f) , the lowest image frequency in a sense masquerades as the
sinusoid that was sampled and is called an alias of the sinusoid that was actually sampled, albeit inadequately
sampled.
One way to avoid such aliasing is to make sure that the signal does not contain any sinusoidal component with a
frequency equal to or greater than fs/2. This condition is sometimes called the Nyquist criterion, and is
equivalent to saying that the sampling frequency fs must be high enough; either greater than twice the highest
frequency or some other more complicated criterion.
206265
Aperture in mm X f-number
Pixel Size and Binning 2.
Matching the Pixels to the telescope
Example 1.
The William Herschel Telescope, with a 4.2m diameter primary mirror and a focal ratio of 3 is to be used
for prime focus imaging. What is the optimum pixel size assuming that the best seeing at the telescope site
is 0.7 arc-seconds ?
First we calculate the ‘plate-scale’ in arc-seconds per millimeter at the focal plane of the telescope.
Plate Scale (arc-seconds per mm) =
206265
=16.4 arc-sec per mm
Aperture in mm X f-number
(Here the factor 206265 is the number of arc-seconds in a Radian)
Next we calculate the linear size at the telescope focal plane of a stellar image (in best seeing conditions)
Linear size of stellar image = 0.7 / Plate Scale = 0.7/ 16.4 = 43 microns.
To satisfy the Nyquist criteria, the maximum pixel size is therefore 21microns. In practice, the nearest
pixel size available is 13.5 microns which leads to a small degree of over-sampling.
Pixel Size and Binning 3.
Example 2.
An Amateur telescope with a 20cm aperture and a focal ratio of 10 is to be used for imaging. The best
seeing conditions at the observing site will be 1 arc-second. What is the largest pixel size that can
be used?
Plate Scale (arc-seconds per mm) =
206265
=103 arc-sec per mm
Aperture in mm X f-number
Linear size of stellar image = 1 / Plate Scale = 1/ 103 = 9.7 microns.
To satisfy the Nyquist criteria, the maximum pixel size is therefore 5 microns. This is about the lower
limit of available pixel sizes.
Pixel Size and Binning 4.
Binning
In the first example we showed that with 13.5µm pixels the system exceeded the Nyquist Criteria even
on nights with exceptionally good sub-arcsecond seeing. If we now suppose that the seeing is 2
arc-seconds, the size of a stellar image will increase to 120µm on the detector. The image will now
be grossly over-sampled. (One way to think of this is that the image is less sharp and therefore requires
fewer pixels to record it). It would be more efficient now for the astronomer to switch to a detector with
larger pixels since the resultant image files would be smaller, quicker to read out and would occupy less
disk space.
There is a way to read out a CCD so as to increase the effective pixel size, this is known as ‘Binning’. With
binning we can increase pixel size arbitrarily. In the limit we could even read out the CCD as a single large
pixel. Astronomers will more commonly use 2 x 2 binning which means that the charge in each 2 x 2 square
of adjacent pixels is summed on the chip prior to delivery to the output amplifier. One important advantage
of ‘on-chip binning’ is that it is a noise free process.
Binning is done in two distinct stages : vertical binning and horizontal binning. Each may be done without
the other to yield rectangular pixels.
Pixel Size and Binning 5.
Stage 1 :Vertical Binning
This is done by summing the charge in consecutive rows .The summing is done in the serial register. In the
case of 2 x 2 binning, two image rows will be clocked consecutively into the serial register prior to the serial
register being read out. We now go back to the conveyor belt analogy of a CCD. In the following animation
we see the bottom two image rows being binned.
Charge packets
Pixel Size and Binning 6.
The first row is transferred into the serial register
Pixel Size and Binning 7.
The serial register is kept stationary ready for the next row to be transferred.
Pixel Size and Binning 8.
The second row is now transferred into the serial register.
Pixel Size and Binning 9.
Each pixel in the serial register now contains the charge from two pixels in the image area. It
is thus important that the serial register pixels have a higher charge capacity. This is achieved
by giving them a larger physical size.
Pixel Size and Binning 10.
Stage 2 :Horizontal Binning
This is done by combining charge from consecutive pixels in the serial register on a special electrode
positioned between serial register and the readout amplifier called the Summing Well (SW).
The animation below shows the last two pixels in the serial register being binned :
SW
1
2
3
Output
Node
Pixel Size and Binning 11.
Charge is clocked horizontally with the SW held at a positive potential.
SW
1
2
3
Output
Node
Pixel Size and Binning 12.
SW
1
2
3
Output
Node
Pixel Size and Binning 13.
SW
1
2
3
Output
Node
Pixel Size and Binning 14.
The charge from the first pixel is now stored on the summing well.
SW
1
2
3
Output
Node
Pixel Size and Binning 15.
The serial register continues clocking.
SW
1
2
3
Output
Node
Pixel Size and Binning 16.
SW
1
2
3
Output
Node
Pixel Size and Binning 17.
The SW potential is set slightly higher than the serial register electrodes.
SW
1
2
3
Output
Node
Pixel Size and Binning 18.
SW
1
2
3
Output
Node
Pixel Size and Binning 19.
The charge from the second pixel is now transferred onto the SW. The binning is now complete
and the combined charge packet can now be dumped onto the output node (by pulsing the voltage
on SW low for a microsecond) for measurement.
Horizontal binning can also be done directly onto the output node if a SW is not present but this can
increase the read noise.
SW
1
2
3
Output
Node
Pixel Size and Binning 20.
Finally the charge is dumped onto the output node for measurement
SW
1
2
3
Output
Node
Advanced CCD Techniques
CCD Imaging
Nik Szymanek
Integrating and Video CCD Cameras.
There is a difference in the geometry of an Integrating CCD camera compared
to a Video CCD camera. An integrating camera, such as is used for most astronomical
applications, is designed to stare at an object over an exposure time of many minutes.
When readout starts and the charge is transferred out of the image area, line by line,
into the serial register, the image area remains light sensitive. Since the readout can
take as long as a minute, if there is no shutter, each stellar image will be drawn out into a
line. An external shutter is thus essential to prevent smearing. These kind of CCDs are
known as ‘Slow Scan’.
A video CCD camera is required to read out much more rapidly. A video CCD may be used
by the astronomers as a finder-scope to locate objects of interest and ensure that the telescope
is actually pointed at the target or it may be used for auto-guiding. These cameras must read out
much more quickly, perhaps several times a second. A mechanical shutter
operating at such frame rates could be unreliable. The geometry of a video CCD, however ,
incorporates a kind of electronic shutter on the CCD and no external shutter is required.
These kind of CCDs are known as ‘Frame Transfer’.
Slow Scan CCDs 1.
The most basic geometry of a Slow-Scan CCD is shown below. Three clock lines
control the three phases of electrodes in the image area, another three control those in the
serial register. A single amplifier is located at the end of the serial register. The full image area
is available for imaging. Because all the pixels are read through a single output, the readout
speed is relatively low. The red line shows the flow of charge out of the CCD.
Image Area
Image area clocks
Output Amplifier
Serial Register clocks
Serial Register
Slow Scan CCDs 2.
A slightly more complex design uses 2 serial registers and 4 output amplifiers. Extra clock
lines are required to divide the image area into an upper and lower section. Further clock lines
allow independent operation of each half of each serial register. It is thus possible to read out the
image in four quadrants simultaneously, reducing the readout speed by a factor of four.
Serial clocks A
Serial clocks B
Amplifier A
Amplifier B
Upper Image area clocks
Lower Image area clocks
Amplifier D
Amplifier C
Serial clocks C
Serial clocks D
Slow Scan CCDs 3.
There are certain drawbacks to using this ‘split-frame readout’ method. The first is that each
amplifier will have slightly different characteristics. It may have a slightly different gain or
a differing linearity. Reconstructing a single image from the four sub-images can be an
image processing nightmare and unless the application demands very high readout speed,
most astronomers would wait slightly longer for an image read out through a single amplifier.
Another drawback is cost. CCDs that have all of their output amplifiers working are rare
and come at a premium price.
Most CCDs are designed with multiple outputs. Even if only one of the working outputs is actually
used, the others provide valuable backups should there be, for any reason, an amplifier failure.
Video CCDs 1.
In the split frame CCD geometry, the charge in each half of the image area could be shifted
independently. Now imagine that the lower image area is covered with an opaque mask. This
mask could be a layer of aluminum deposited on the CCD surface or it could be an external
mask. This geometry is the basis of the ‘Frame transfer’ CCD that is used for high frame rate video
applications. The area available for imaging is reduced by a half. The lower part of the image
becomes the ‘Store area’.
Image area
Image area clocks
Opaque mask
Store area clocks
Store area
Amplifier
Serial clocks
Video CCDs 2.
The operation of a Split Frame Video CCD begins with the integration of the image in the image area.
Once the exposure is complete the charge in the image area is shifted down into the store area
beneath the light proof mask. This shift is rapid; of the order of a few milliseconds for a large CCD.
The amount of image smear that will occur in this time is minimal (remember there is no external shutter).
Integrating Galaxy Image
Video CCDs 3.
Once the image is safely stored under the mask, it can then be read out at leisure. Since we can
independently control the clock phases in the image and store areas, the next image can be
integrated in the image area during the readout. The image area can be kept continuously integrating
and the detector has only a tiny ‘dead time’ during the image shift. No external shutter is required
but the effective size of the CCD is cut by a half.
Correlated Double Sampler (CDS) 1.
The video waveform output by a CCD is at a fairly low level : every photo-electron in a pixel
charge packet will produce a few micro-volts of signal. Additionally, the waveform is
complex and precise timing is required to make sure that the correct parts are amplified and
measured.
The CCD video waveform , as introduced in Activity 1, is shown below for the period of
one pixel measurement
Vout
t
Reset feedthrough
Reference level
Charge dump
Signal level
The video processor must measure , without introducing any additional noise, the Reference level
and the Signal level. The first is then subtracted from the second to yield the output signal voltage
proportional to the number of photo-electrons in the pixel under measurement. The best way to
perform this processing is to use a ‘Correlated Double Sampler’ or CDS.
Correlated Double Sampler (CDS) 2.
The CDS design is shown schematically below. The CDS processes the video waveform and outputs
a digital number proportional to the size of the charge packet contained in the pixel being read. There
should only be a short cable length between CCD and CDS to minimise noise.The CDS minimises the
read noise of the CCD by eliminating ‘reset noise’. The CDS contains a high speed analogue processor
containing computer controlled switches. Its output feeds into an Analogue to Digital Converter (ADC).
R RD OD
Reset switch
CCD On-chip Amplifier
.
Inverting Amplifier
-1
OS
ADC
Input Switch
Polarity Switch
Computer Bus
Pre-Amplifier
Integrator
Correlated Double Sampler (CDS) 3.
The CDS starts work once the pixel charge packet is in the CCD summing well and the CCD reset
pulse has just finished. At point t0 the CCD wave-form is still affected by the reset pulse and so
the CDS remains disconnected from the CCD to prevent this disturbing the video processor.
t0
t0
Output wave-form of CCD
Output voltage of CDS
-1
Correlated Double Sampler (CDS) 4.
Between t1 and t2 the CDS is connected and the ‘Reference ‘ part of the waveform is sampled.
Simultaneously the integrator reset switch is opened and the output starts to ramp down linearly.
t1 t2
t1
Reference
window
-1
t2
Correlated Double Sampler (CDS) 5.
Between t2 and t3 the ‘charge dump’ occurs in the CCD. The CCD output steps negatively by an amount
proportional to the charge contained in the pixel. During this time the CDS is disconnected.
t2t3
t1
-1
t2 t3
Correlated Double Sampler (CDS) 6.
Between t3 and t4 the CDS is reconnected and the ‘signal’ part of the wave-form is sampled. The input to
the integrator is also ‘polarity switched’ so that the CDS output starts to ramp-up linearly. The width of the
signal and sample windows must be the same. For Scientific CCDs this can be anything between 1 and 20
microseconds. Longer widths generally give lower noise but of course increase the read-out time.
t3 t4
t1
Signal
window
-1
t2 t3
t4
Correlated Double Sampler (CDS) 7.
The CDS is then once again disconnected and its output digitised by the ADC. This number , typically a
16 bit number (with a value between 0 and 65535) is then stored in the computer memory. The CDS
then starts the whole process again on the next pixel. The integrator output is first zeroed by closing
the reset switch. To process each pixel can take between a fraction of a microsecond for a
TV rate CCD and several tens of microseconds for a low noise scientific CCD.
t2 t3
t4
Voltage to be
digitised
The type of CDS is called a ‘dual slope integrator’.
A simpler type of CDS known as a ‘clamp and sample’
only samples the waveform once for each pixel.
It works well at higher pixel rates but is noisier
than the dual slope integrator at lower pixel rates.
t1
-1
ADC
Noise Sources in a CCD Image 1.
Noise level sets the smallest detectable signal (faint sources) so it must be properly controlled.
The main noise sources found in a CCD are :
1.
READ OUT NOISE
Caused by electronic noise in the CCD output transistor and possibly also in the external circuitry.
Read noise places a fundamental limit on the performance of a CCD. It can be reduced at the
expense of increased read out time. Scientific CCDs have a read out noise of 2-3 electrons RMS.
2.
DARK CURRENT
Caused by thermally generated electrons in the CCD even in the absence of lighting, increasing
with temperature and integration time.
Consequences: generates “thermal noise” limiting the detection of faint sources and can even
saturate the pixels. Eliminated by cooling down the CCD usually at ~-100ºC.
3.
PHOTON NOISE
Also called ‘Shot Noise’. It is due to the fact that the CCD detects photons. Photons arrive in an
unpredictable fashion described by Poisson statistics. This unpredictability causes noise.
R  N 
4.
PIXEL RESPONSE NON-UNIFORMITY
Defects in the silicon and small manufacturing defects can cause some pixels to have a higher
sensitivity than their neighbours. This noise source can be removed by ‘Flat Fielding’.
Noise Sources in a CCD Image 2.
Before these noise sources are explained further some new terms need to be introduced.
FLAT FIELDING
This involves exposing the CCD to a very uniform light source that produces a featureless and even
exposure across the full area of the chip. A flat field image can be obtained by exposing on a
twilight sky or on an illuminated white surface held close to the telescope aperture (for example the
inside of the dome). Flat field exposures are essential for the reduction of astronomical data.
BIAS REGIONS
A bias region is an area of a CCD that is not sensitive to light. The value of pixels in a bias region
is determined by the signal processing electronics. It constitutes the zero-signal level of the CCD.
The bias region pixels are subject only to readout noise. Bias regions can be produced by
‘over-scanning’ a CCD, i.e. reading out more pixels than are actually present. Designing a CCD with
a serial register longer than the width of the image area will also create vertical bias strips at the left
and right sides of the image. These strips are known as the ‘x-underscan’ and ‘x-overscan’ regions
A flat field image containing bias regions can yield valuable information not only on the various
noise sources present in the CCD but also about the gain of the signal processing electronics
i.e. the number of photoelectrons represented by each digital unit (ADU) output by the camera’s
Analogue to Digital Converter.
Noise Sources in a CCD Image 3.
Flat field images obtained from two CCD geometries are represented below. The arrows represent
the position of the readout amplifier and the thick black line at the bottom of each image represents
the serial register.
Y-overscan
Here, the CCD is over-scanned in X and Y
Image Area
X-overscan
CCD With Serial
Register equal in
length to the image
area width.
Image Area
X-overscan
CCD With Serial
Register greater in
length than the image
area width.
X-underscan
Y-overscan
Here, the CCD is over-scanned in Y
to produce the Y-overscan bias area.
The X-underscan and X-overscan are
created by extensions to the serial
register on either side of the image area.
When charge is transferred from the image
area into the serial register, these extensions
do not receive any photo-charge.
Noise Sources in a CCD Image 4.
These four noise sources are now explained in more detail:
READ NOISE
This is mainly caused by thermally induced motions of electrons in the output amplifier. These cause
small noise voltages to appear on the output. This noise source, known as Johnson Noise, can be
reduced by cooling the output amplifier or by decreasing its electronic bandwidth. Decreasing the
bandwidth means that we must take longer to measure the charge in each pixel, so there is always
a trade-off between low noise performance and speed of readout. Mains pickup and interference from
circuitry in the observatory can also contribute to Read Noise but can be eliminated by careful design.
Johnson noise is more fundamental and is always present to some degree.
The graph below shows the trade-off between noise and readout speed for an EEV4280 CCD.
Read Noise (electrons RMS)
14
12
10
8
6
4
2
0
2
3
4
5
Tim e spent m easuring each pixel (m icroseconds)
6
Noise Sources in a CCD Image 5.
DARK CURRENT
Electrons can be generated in a pixel either by thermal motion of the silicon atoms or by the absorption
of photons. Electrons produced by these two effects are indistinguishable. Dark current is analogous to
the fogging that can occur with photographic emulsion if the camera leaks light. Dark current can be
reduced or eliminated entirely by cooling down the CCD. Science cameras are typically cooled with liquid
nitrogen (-196ºC) to the point where the dark current falls to below 1 electron/pixel/hour (~-100ºC) where
it is essentially un-measurable. Amateur cameras cooled thermoelectrically may still have substantial
dark current. The graph below shows how the dark current of a TEK1024 CCD can be reduced by cooling.
Electrons per pixel per hour
10000
1000
100
10
1
-110
-100
-90
-80
-70
-60
Temperature Centigrade
-50
-40
Noise Sources in a CCD Image 6.
PHOTON NOISE
This can be understood more easily if we go back to the analogy of rain drops falling onto an array
of buckets; the buckets being pixels and the rain drops photons. Both rain drops and photons arrive
discretely, independently and randomly and are described by Poisson statistics. If the buckets are
very small and the rain fall is very sparse, some buckets may collect one or two drops, others may collect
none at all. If we let the rain fall long enough all the buckets will measure the same value ,
but for short measurement times the spread in measured values is large. This latter scenario is essentially
that of CCD astronomy where small pixels are collecting very low fluxes of photons.
Poissonian statistics tells us that the Root Mean square uncertainty (RMS noise) in the number of
photons/second detected by a pixel is equal to the square root of the mean photon flux (the
average number of photons detected/second).
For example, if a star is imaged onto a pixel and it produces on average 10 photo-electrons/second
and we observe the star for 1 second, then the uncertainty of our measurement of its brightness
will be the square root of 10 i.e. 3.2 electrons. This value is the ‘Photon Noise’.
Increasing exposure time to 10 seconds will increase the photon noise to 10 electrons (the square root
of 100) but at the same time will increase the ‘Signal to Noise ratio’ (SNR). In the absence of other
noise sources the SNR will increase as the square root of the exposure time (SNR1=3.1, SNR2=10).
Astronomy is all about maximising the SNR.
{Dark current, described earlier, is also governed by Poisson statistics. If the mean dark current
contribution to an image is 900 electrons/pixel, the noise introduced into the measurement
of any pixels photo-charge would be 30 electrons}
Noise Sources in a CCD Image 7.
PIXEL RESPONSE NON-UNIFORMITY (PRNU)
If we take a very deep (at least 50,000 electrons of photo-generated charge per pixel) flat field exposure ,
the contribution of photon noise and read noise become very small. If we then plot the pixel values
along a row of the image we see a variation in the signal caused by the slight variations in sensitivity
between the pixels. The graph below shows the PRNU of an EEV4280 CCD illuminated by blue light.
The variations are as much as +/-2%. Fortunately these variations are constant and are easily removed
by dividing a science image, pixel by pixel, by a flat field image.
3
% variation
2
1
0
-1
-2
-3
0
100
200
300
400
500
column number
600
700
800
Noise Sources in a CCD Image 8.
HOW THE VARIOUS NOISE SOURCES COMBINE
Assuming that the PRNU has been removed by flat fielding, the three remaining noise
sources combine in the following equation:
NOISEtotal =
(READ NOISE)2 + (PHOTON NOISE)2 +(DARK CURRENT)2
In professional systems the dark current tends to zero and this term of the equation can be
ignored. The equation then shows that read noise is only significant in low signal level
applications such as Spectroscopy. At higher signal levels, such as those found in
direct imaging, the photon noise becomes increasingly dominant and the read noise
becomes insignificant. For example , a CCD with read noise of 5 electrons RMS will
become photon noise dominated once the signal level exceeds 25 electrons/pixel.
If the exposure is continued to a level of 100 electrons/pixel, the read noise contributes
only 11% of the total noise.
Photon Transfer Method 1.
Using two identical flat field exposures it is possible to measure the read noise
of a CCD with the Photon Transfer method. Two exposures are required to remove
the contribution of the PRNU and of small imperfections in the flat fields caused by uneven
illumination.
The method actually measures the conversion gain of the CCD camera; the number of electrons
represented by each digital interval (ADU) of the analogue to digital converter, however, once
the gain is known the read noise follows straightforwardly.
This method exploits the Poissonian statistics of photon arrival. To use it, one requires an image
analysis program capable of doing statistical analysis on selected areas of the input images.
Photon Transfer Method 2.
Bias area 1
Image area 1
Flat Field
Image 1.
STEP 1
Measure the Standard Deviation in the two bias areas and average
the two values.
result= NoiseADU the Root Mean Square readout noise in ADU.
STEP 2
Measure the mean pixel value in the two bias areas and the two
image areas. Then subtract MeanBias area 1 from MeanImage area 1
Bias area 2
Image area 2 result= MeanADU ,the Mean Signal in ADU.
Flat Field
Image 2.
As an extra check repeat this for the second image, the Mean should
be very similar. If it is more than a few percent different it may be
best to take the two flat field exposures again.
Photon Transfer Method 3.
STEP 3 The two images are then subtracted pixel by pixel to yield a third image
Image 1
-
Image 2
=
Image 3
Image area 3
STEP 4
Measure the Standard Deviation in image area 3
result= StdDevADU .
The statistical spread in the pixel values in this subtracted image
area will be due to a combination of readout noise and photon noise.
STEP 5
Now apply the following equation.
Gain =
2 x MeanADU
(StdDevADU ) 2 - (2 x NoiseADU 2).
The units will be electrons per ADU, which will be inversely
proportional to the voltage gain of the system.
Photon Transfer Method 4.
STEP 6 The Readout noise is then calculated using this gain value :
Readout Noiseelectrons= Gain x NoiseADU
Precautions when using this method
The exposure level in the two flat fields should be at least several thousand ADU
but not so high that the chip or the processing electronics is saturated. 10,000 ADU
would be ideal. It is best to average the gain values obtained from several pairs
of flat fields. Alternatively the calculations can be calculated on several
sub-regions of a single image pair. If the illumination of the flat fields is not
particularly flat and the signal level varies appreciable across the sub-region on which
the statistics are performed, this method can fail. If good flat fields are unavailable,
as will be the case if the camera is connected to a spectrograph, then the sub-regions
should be kept small.
Deep Depletion CCDs 1.
Electric potential
The electric field structure in a CCD defines to a large degree its Quantum Efficiency (QE). Consider
first a thick frontside illuminated CCD, which has a poor QE.
Cross section through a thick frontside illuminated CCD
In this region the electric potential gradient
is fairly low i.e. the electric field is low.
Potential along this line
shown in graph above.
Any photo-electrons created in the region of low electric field stand a much higher chance of
recombination and loss. There is only a weak external field to sweep apart the photo-electron
and the hole it leaves behind.
Deep Depletion CCDs 2.
Electric potential
In a thinned CCD , the field free region is simply etched away.
Cross section through a thinned CCD
There is now a high electric field throughout the
full depth of the CCD.
This volume is
etched away
during manufacture
Problem : Thinned CCDs may have good blue
response but they become transparent
at longer wavelengths; the red response
suffers.
Red photons can now pass
right through the CCD.
Photo-electrons created anywhere throughout the depth of the device will now be detected. Thinning
is normally essential with backside illuminated CCDs if good blue response is required. Most blue
photo-electrons are created within a few nanometers of the surface and if this region is field free,
there will be no blue response.
Deep Depletion CCDs 3.
Electric potential
Ideally we require all the benefits of a thinned CCD plus an improved red response. The solution is to use a
CCD with an intermediate thickness of about 40mm constructed from Hi-Resistivity silicon. The increased
thickness makes the device opaque to red photons. The use of high-resistivity silicon means that there are no field
free regions despite the greater thickness.
Cross section through a Deep Depletion CCD
Problem :
High-resistivity silicon contains much lower
impurity levels than normal. Very few wafer
fabrication factories commonly use this
material and deep depletion CCDs have to
be designed and made to order.
Red photons are now absorbed in
the thicker bulk of the device.
There is now a high electric field throughout the full depth of the CCD. CCDs manufactured in this way
are known as Deep depletion CCDs. The name implies that the region of high electric field, also known as
the ‘depletion zone’ extends deeply into the device.
Deep Depletion CCDs 4.
The graph below shows the improved QE response available from a deep depletion CCD.
The black curve represents a normal thinned backside illuminated CCD. The Red curve is actual data from
a deep depletion chip manufactured by MIT Lincoln Labs. This latter chip is still under development.The blue
curve suggests what QE improvements could eventually be realised in the blue end of the spectrum once
the process has been perfected.
Deep Depletion CCDs 5.
Another problem commonly encountered with thinned CCDs is ‘fringing’. This is greatly reduced
in deep depletion CCDs. Fringing is caused by multiple reflections inside the CCD. At longer
wavelengths, where thinned chips start to become transparent, light can penetrate through and be
reflected from the rear surface. It then interferes with light entering for the first time. This can give
rise to constructive and destructive interference and a series of fringes where there are minor
differences in the chip thickness.
The image below shows some fringes from an EEV42-80 thinned CCD
For spectroscopic applications, fringing can render some thinned CCDs unusable, even those
that have quite respectable QEs in the red. Thicker deep depletion CCDs , which have a much
lower degree of internal reflection and much lower fringing are preferred by astronomers
for spectroscopy.
Advantages of CCDs for Astronomy 1.
• High Quantum Efficiency: Astronomical CCDs can reach peak quantum efficiencies of greater than 90%.
This means that nine out of ten photons hitting a pixel form an electron-hole pair that can be detected and
counted. Equally important is that CCDs have a high quantum efficiency across a wide frequency range. A
back-illuminated CCD for instance has QE > 60% for more a 500 nm waveband. CCDs are far more
sensitive than either photographic emulsions of photomultiplier tubes. This means that they are much
"faster" than photographic emulsions so that the same length exposure will reveal far fainter sources than a
photograph. Shorter exposures therefore can often be used so that many more fields can be imaged in a
single observing session.
Broad Spectral Response: Early generation CCDs were sensitive in the red part of the spectrum but less
so in the blue or ultraviolet regions. Improved technology means that current chips have a broader spectral
response.
• Large Dynamic Range: CCDs have a very large dynamic range. This is basically the range between the
lowest and highest charge values that can be detected in a well. CCDs may have a dynamic range with a
factor of 100,000 × compared with only 100 × that is typical for photographic emulsions. They are thus
useful for imaging astronomical objects where there is naturally a large dynamic range in the sources.
• Linear Response: If a CCD pixel receives twice the number of photons than another pixel, it will have
double the amount of charge than the first pixel. the number of electrons in a pixel therefore is proportional
to the number of incident photons to within about 0.1%.
Advantages of CCDs for Astronomy 2.
• Low Noise: As with any semiconductor material, the silicon in CCDs produces random "noise" due to
thermal vibrations. This noise can degrade a signal. Fortunately modern CCDs are designed to produce very
low levels of noise that can then be accounted for in subsequent analysis. One way of reducing noise is to
cool the CCD chip - the cooler it is the less vibration in the atoms. Cooling can be thermoelectric (eg Peltier
cooling) and/or by using cryogenic systems (liquid nitrogen).
• Stable: CCDs are physically very stable. They do not expand or warp due to thermal or mechanical
changes. The actual photosensitive chip is normally encased in a protective enclosure.
• Digital Readout: The information obtained in a CCD exposure is directly readout into a computer or
digital storage device at the end of the exposure. This means that analysis can start straight away without the
need for developing a plate as photography requires. It also avoids the need for scanning and digitizing
plates. The CCD data can be easily backed up on tape or disk and even transmitted via the Internet.
Problems with CCDs for Astronomy 1.
• Small Size and Field of View: Initially CCDs were quite small, a 512 × 512 pixel CCD with a total of
262,144 pixels was considered large. The problem with such small area detectors is that they give a very
small field of view, much smaller than a photographic plate achieved on the same telescope. As
manufacturing techniques have improved and the cost per pixel decreased, CCDs have grown in size. 1024
× 1024 pixels are now standard whilst chips with 4096 × 4096 (that is 16.8 million) pixels now exist.
Nonetheless the field of view is still smaller than photographic systems. To try and get around this some
observatories use several CCDs mosaiced together to give a much larger field of view.
One example of this is the wide-field imager, MegaCam which consists of 40 2048 x 4612 pixel CCDs (a
total of 340 megapixels). This covers a full 1° × 1° field-of-view with a resolution of 0.187 arcsecond per
pixel and makes the most of the 0.7 arcsecond median seeing at Mauna Kea. MegaCam sits at the prime
focus of the CFHT, the Canada-France Hawaii telescope, a 3.6 m optical telescope in Hawaii.
• Cost: Professional-grade CCD chips are expensive. The cost per pixel has decreased significantly since
their introduction but the chips used for professional astronomical CCD cameras have much higher
specification levels and finer tolerances than those used in many commercial cameras.
• Calibration: In order to extract the information accurately and fully from a CCD images many calibration
and adjustment steps are required. Additional frames or exposures such as flat-fields and dark frames must
be taken and accurately logged. A CCD's performance also changes with temperature so this needs to be
accounted for.
• Cooling: To minimize noise on an astronomical image a CCD chip must be cooled. If this involves
cryogenic materials such as liquid nitrogen, this must be added to the camera well in advance of observing
to allow sufficient time for it to cool down. Care must also be taken to ensure the camera stays thermally
stable.
Mosaic Cameras 1.
When CCDs were first introduced into astronomy, a major drawback, compared to photographic
plate detectors was their small size. CCDs are still restricted in size by the silicon wafers that
are used in their production. Most factories can only handle 6” diameter wafers. The largest
photographic plates are about 30 x 30cms and when used with wide angle telescopes can
simultaneously image a region of sky 60 x 60 in size. To cover this same area of sky with a smaller
CCD would require hundreds of images and would be an extremely inefficient use of the telescope’s
valuable time. It is unlikely that CCDs will ever reach the same size as photographic detectors, so for
applications requiring large fields of view, mosaic CCD cameras are the only answer. These are
cameras containing a number of CCDs mounted in the same plane with only small gaps between
adjacent devices.
Mosaic CCD cameras containing up to 30 CCD chips are in common use today, with even larger
mosaics planned for large survey telescopes in the near future. One interesting technical challenge
associated with their design is in keeping all the chips in the same plane (i.e. the focal plane of the
telescope) to an accuracy of a few tens of microns. If there are steps between adjacent chips then
star images will be in focus on one chip but not necessarily on its neighbors.
Most new CCD are designed for close butting and the construction of mosaics. This is achieved
by using packages with electrical connections along one side only, leaving the other three sides free
for butting. The next challenge is to build CCDs which have the connections on the rear of the package
and are buttable on 4 sides! This would allow full unbroken tiling of a telescopes focal plane and the best
possible use of its light gathering power.
Mosaic Cameras 2.
The pictures below show the galaxy M51 and the CCD mosaic that produced the image.
Two EEV42-80 CCDs are screwed down onto a very flat Invar plate with a 50 micron gap
between them. Light falling down this gap is obviously lost and causes the black strip
down the centre of the image. This loss is not of great concern to astronomers, since it
represents only 1% of the total data in the image.
Mosaic Cameras 3.
Another image from this camera is shown below. The object is M42 in Orion.
This false color image covers an area of sky measuring 16’ x 16’. The image
was obtained on the William Herschel Telescope in La Palma.
Mosaic Cameras 4.
A further image is shown below, of the galaxy M33 in Triangulum. Images from
this camera are enormous; each of the two chips measures 2048 x 4100 pixels. The
original images occupy 32MB each.
Nik Szymanek
Mosaic Cameras 5.
The Horsehead Nebula in Orion.
The mosaic mounted in its camera.
Mosaic Cameras 6.
This colossal mosaic of 12 CCDs is in operation at the CFHT in Hawaii. Here is an
example of what it can produce. The chips are of fairly low cosmetic quality.
Picture : Canada France Hawaii Telescope
Mosaic Cameras 7.
This mosaic of 4 science CCDs was built at the Royal Greenwich Observatory. The positioning
of the CCDs is somewhat unusual but ultimately all that matters is the total area covered . A smaller
fifth CCD on the right hand side is used for auto-guiding the telescope. An example of this camera’s
output is shown on the left.
M13
Camera Construction Techniques 1.
The photo below shows a scientific CCD camera in use at the Isaac Newton Group. It is
approximately 50cm long, weighs about 10Kg and contains a single cryogenically cooled CCD.
The camera is general purpose detector with a universal face-plate for attachment to various
telescope ports.
Mounting clamp
Pre-amplifier Pressure Vessel
Vacuum pump port
Camera mounting
Face-plate.
Liquid Nitrogen
fill port
Camera Construction Techniques 2.
The main body of the camera is a 3mm thick aluminium pressure vessel able to support
an internal vacuum. Most of the internal volume is occupied by a 2.5 liter copper can
that holds liquid nitrogen (LN2). The internal surfaces of the pressure vessel and the external
surfaces of the copper can are covered in aluminized mylar film to improve the
thermal isolation. As the LN2 boils off , the gas exits through the same tube that is used for
the initial fill
The CCD is mounted onto a copper cold block that is stood-off from the removable end plate
by thermally insulating pillars. A flexible copper braid connects this block to the LN2
can. The thickness of the braid is adjusted so that the equilibrium temperature of the
CCD is about 10 degrees below the optimum operating temperature. The mounting
block also contains a heater resistor and a Platinum resistance thermometer that are used to
stabilize the CCD temperature. Without using a heater resistor close to the chip for thermal
regulation , the operating temperature and the CCD characteristics also, will vary with
the ambient temperature.
The removable end plate seals with a synthetic rubber ‘o’ ring. In its center is a fused
silica window big enough for the CCD and thick enough to withstand atmospheric
pressure. The CCD is positioned a short distance behind the window and radiatively cools
the window. To prevent condensation it is necessary to blow dry air across the outside.
Camera Construction Techniques 3.
The basic structure of the camera is that of a Thermos flask. Its function is to protect the
CCD in a cold clean vacuum. The thermal design is very important, so as to maximize the hold
time and the time between LN2 fills. Maintenance of a good vacuum is also very important,
firstly to improve the thermal isolation of the cold components but also to prevent contamination
of the CCD surface. A cold CCD is very prone to contamination from volatile substances such
as certain types of plastic. These out-gas into the vacuum spaces of the camera and then condense
on the coldest surfaces. This is generally the LN2 can.
On the back of the can is a small container filled with activated charcoal known as a ‘Getter’.
This acts as a sponge for any residual gases in the camera vacuum. The getter contains a heater
to drive off the absorbed gases when the camera is being pumped.
This camera is vacuum pumped for several hours before it is cooled down. The pressure at the
end of the pumping period is about 10 -4 mBar. When the LN2 is introduced, the pressure will
fall to about 10-6mBar as the residual gases condense out on the LN2 can.
When used in an orientation that allows this camera to be fully loaded with LN2, the boil
off or ‘hold time’ is about 20 hours. The thermal energy absorbed by 1g of LN2 turning to
vapor is 208J, the density of LN2 is 0.8g/cc. From this we can calculate that the heat load
on the LN2 can, from radiation and conduction is about 6W.
Camera Construction Techniques 4.
A cutaway diagram of the same camera is shown below.
(To avoid frost it is necessary that the CCD be place in an area devoid of
water vapor: either in vacuum or in a dry nitrogen atmosphere).
Thermally
Insulating
Pillars
Electrical feed-through
Vacuum Space
Pressure vessel
Pump Port
Telescope beam
Face-plate
CCD
Focal Plane
of Telescope
Optical window
...
CCD Mounting Block Thermal coupling
Boil-off
Nitrogen can
Activated charcoal ‘Getter’
Camera Construction Techniques 5.
The camera with the face-plate removed is shown below
CCD
Retaining
clamp
Temperature servo circuit board
Aluminised Mylar
sheet
Gold plated copper
mounting block
Top of LN2
can
Platinum resistance
thermometer
Pressure
Vessel
‘Spider’.
The CCD mounting
block is stood off from
the spider using
insulating pillars.
Location points (x3)
for insulating pillars
that reference the CCD
to the camera face-plate
Signal wires to CCD
Camera Construction Techniques 6.
A ‘Radiation Shield’ is then screwed down onto the spider, covering the cold components but not
obstructing the CCD view. This shield is highly polished and cooled to an intermediate temperature
by a copper braid that connects it to the LN2 can.
Radiation Shield
Camera Construction Techniques 7.
Some CCDs cameras are embedded into optical instruments as dedicated detectors.
The CCD shown below is mounted in a spider assembly and placed at the focus of a
Schmidt camera.
CCD Signal connector (x3)
Copper rod or ‘cold finger’
used to cool the CCD. It is
connected to an LN2 can.
‘Spider’ Vane
CCD Clamp plate
Gold plated
copper CCD
mounting
block.
FOS 1 Spectrograph
CCD Package