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 devicesand 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 amajor revolution in Astronomy. They improved the light gathering power of telescopes by almost two orders ofmagnitude. 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 convertedinto 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 imagesduring analysis). CCDs have practically no detection threshold (unlike films) and the number of electrons generatedgives 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

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 heatingor 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.

Incr

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Valence Band

Conduction Band

1.26eV

Thermally generated electrons are indistinguishable from photo-generated electrons . They constitute anoise 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 1m. Beyond this wavelength silicon becomes transparent and CCDs constructed from silicon become insensitive.

photon phot

on

Hole Electron

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 thebuckets (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.

PHOTONS

PIXELS (BUCKETS)

CCD COLUMNS (VERTICAL CONVEYOR BELTS)

SERIAL REGISTER (HORIZONTAL CONVEYOR BELT)

OUTPUT AMPLIFIER

CCD Analogy

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 conveyoris 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 intothe measuring cylinder .

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After each bucket has been measured, the measuring cylinderis 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.

Connection pins

Gold bond wires

Bond pads

Silicon chip

Metal, ceramic or plastic packageImage area

Serial register

On-chip amplifier

CCD: 512x512 pixels, 1024x1024, 2048x2048, 4096x4096

Sensitive surface:•TC211: 2.64x2.64mm2

(192x165 pixels, 13.75x16µm size)•Thomson 7863: 8.8x6.6mm2

(384x288 pixels, 23x23µm)•SITe SI-003A: 24.6x24.6mm2

(1024x1024 pixels, 24x24µm)•LORAL 2K3eb: 30.7x30.7mm2

(2048x2048 pixels, 15x15µm)

CCD arrays may be smaller in comparison with standard photography (24x36mm). “Buttable” CCDs for mosaic.


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


One pixel

Channel stops to define the columns of the image

Transparenthorizontal electrodesto define the pixels vertically. Also used to transfer the charge during readout

Plan View

Cross section

The diagram shows a small section (a few pixels) of the image area of a CCD. This pattern is reapeated.

ElectrodeInsulating oxiden-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.


On-chip amplifierat end of the serial register

Cross section ofserial register

Image Area

Serial Register

Once again every third electrode in the serial register is connected together.

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.


The serial register is bent twice to move the output amplifier away from the edgeof the chip (Voltage bus). The arrows indicate how charge is transferred through the device.

Edge

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Image Area

Serial Register

Read Out Amplifier

Bus

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Details of a corner of an EEV CCD


OD

OS

RDRSW

Output Node

Substrate

Output Transistor

Reset Transistor

SummingWell

20mOutput Drain (OD)

Output Source (OS)

Gate of Output Transistor

Output Node

R

Reset Drain (RD)

Summing Well (SW)

Last few electrodes in Serial Register

Serial Register Electrodes

Details of the on-chip amplifier of a Tektronix CCD and its circuit diagram

Electric Field in a CCD 1.

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).

n p

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Potential along this line shownin graph above.

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Cross section through the thickness of the CCD

Electric Field in a CCD 2.

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 definethe vertical boundaries of the pixel. The horizontal boundaries are defined by the channel stops.

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Region of maximum potential

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Charge packetp-type siliconn-type silicon

SiO2 Insulating layer

Electrode Structure

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Charge Collection in a CCD.

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

Charge Transfer in a CCD 1.

In the following few slides, the implementation of the ‘conveyor belts’ as actual electronicstructures is explained.

The charge is moved along these conveyor belts by modulating the voltages on the electrodespositioned on the surface of the CCD. In the following illustrations, red electrodesare held at a positive potential, black ones are held at a negative potential.

123

123

+5V

0V

-5V

+5V

0V

-5V

+5V

0V

-5V

Time-slice shown in diagram

1

2

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123

+5V

0V

-5V

+5V

0V

-5V

+5V

0V

-5V

1

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123

+5V

0V

-5V

+5V

0V

-5V

+5V

0V

-5V

1

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Charge packet from subsequent pixel entersfrom the left as first pixel exits to the right.

123

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0V

-5V

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+5V

0V

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1

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On-Chip Amplifier 1.

OD

OS

RDRSW

Output Node Output

Transistor

Reset Transistor

SummingWell

+5V

0V

-5V

+10V

0V

R

SW

--end of serial register

Vout

Vout

The on-chip amplifier measures each charge packet as it pops out the end of the serial register.

The measurement process begins with a resetof the ‘reset node’. This removes the charge remaining from the previous pixel. The resetnode is in fact a tiny capacitance (< 0.1pF)

RD and OD are held at constant voltages

(The graphs above show the signal waveforms)


OD

OS

RDRSW

Output Node Output

Transistor

Reset Transistor

SummingWell

+5V

0V

-5V

+10V

0V

R

SW


Vout

Vout

The charge is then transferred onto the Summing Well. Vout is now at the ‘Reference level’

There is now a few tens of µs wait while externalcircuitry measures this ‘reference’ level.


OD

OS

RDRSW

Output Node Output

Transistor

Reset Transistor

SummingWell

+5V

0V

-5V

+10V

0V

R

SW


Vout

Vout

This action is known as the ‘charge dump’

The voltage step in Vout is as much as several V for each electron contained in the charge packet.

The charge is then transferred onto the output node. Vout now steps down to the ‘Signal level’


OD

OS

RDRSW

Output Node Output

Transistor

Reset Transistor

SummingWell

+5V

0V

-5V

+10V

0V

R

SW


Vout

Vout

Vout is now sampled by external circuitry for up to a few tens of microseconds.

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

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)

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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.

(1 nm = 10 angstroms)

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)

Thick Front-side Illuminated CCD

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 detectorsbe as sensitive as possible; thick front-side illuminated chips are seldom if ever used.

n-type silicon

p-type silicon

Silicon dioxide insulating layerPolysilicon electrodes

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Anti-Reflection Coatings 1

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.

Silicon has a very high Refractive Index (denoted by n). This means that photons are strongly reflected from its surface.

nt-ni

nt+ni

ni

nt

Fraction of photons reflected at the interface between two mediums of differing refractive indices = [ ]2


The reflected portion is now reduced to :

In the case where 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.

With an Anti-reflective coating we now have three mediums to consider :

nt

Airni

ns AR Coating

Silicon

[ ]nt x ni-ns2

nt x ni+ns

2

2

ns nt

2=


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%.

Thinned Back-side Illuminated CCD

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.

n-type silicon

p-type silicon

Silicon dioxide insulating layerPolysilicon electrodes

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Anti-reflective (AR) coating

15m

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 showingthe ‘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-colouredinterference 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 CCDand its signal processing electronics during the readout phase, but also for data collection and the subsequent processing of the images.

The CCD ControllerIn this first application, the computer is an embedded system running in a ‘CCD controller’. This controller willtypically 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. Clocksequences 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 portswhich 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 computerrunning in the CCD controller is thus fast and of relatively simple design. A poorly performing processor herecould result in slow read out times and poor use of telescope resources. Remember that when a CCD is readingout the telescope shutter is closed and no observations are possible. For an amateur observer using a small CCDwith a fast readout time, a slow CCD controller may not be such a disadvantage; there are not so many pixelsto process.


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 AreaNetwork.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 (typicallygraphics 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’.


Image Processing ComputersThese are used for reduction and analysis of the astronomical data. Many astronomers process their datain 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 followingperformance would be obtained :

Linearisation, bias subtraction and flat-fielding : ~150 secde-fringing : ~300 secobject detection and star/galaxy separation : ~300 sec


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 squareCCD 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

Clocks

CCD

Amplifier

Digitalconverter

Cooling system

I/O interface

Computer

CCD cameraAcquisition commands

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

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Phot

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Blooming 1.

The charge capacity of a CCD pixel is limited: when a pixel is full, the charge starts to leak intoadjacent pixels. This process is known as ‘Blooming’.

Phot

onsOverflowing

charge packet

Spillage Spillage

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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 chargespreading sideways. The charge confinement provided bythe electrodes is lower so the charge spreads vertically upand down a column.

The capacity of a CCD pixel is known as the ‘Full Well’. It isdependent on the physical area of the pixel. For TektronixCCDs, with pixels measuring 24m x 24m it can be as much as300,000 electrons. Bloomed images will be seen particularly on nights of good seeing where stellar images are more compact .

In reality, blooming is not a big problem for professionalastronomy. For those interested in pictorial work, however, it canbe a nuisance.

Flow

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Blooming 3.

Bloomed star images

The image below shows an extended source with bright embedded stars. Due to the longexposure required to bring out the nebulosity, the stellar images are highly overexposedand create bloomed images.

(The image is from a CCD mosaic and the black strip down the center is the space between adjacent detectors)

M42

Dark columnsUnless 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.

Flat field exposure of an EEV42-80 CCD

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 inadjacent 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.

Cosmic rays, Hot spots, Bright columns

Bright columns are also caused by traps . Electrons containedin such traps can leak out during readout causing a vertical streak

Hot Spots are pixels with higher than normal dark current. Theirbrightness increases linearly with exposure times

Cosmic rays are unavoidable. Charged particles from space orfrom radioactive traces in the material of the camera cancause ionisation in the silicon. The electrons produced areindistinguishable 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.

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’)

Cosmic rays

Cluster ofHot Spots

BrightColumn

900s dark exposure of an EEV42-80 CCD

Electronic processing

M51

Some defects can arise from the read-out processing electronics. This negative image has a bright line in the first image row.

Dark column

Hot spots and bright columns

Bright first image row caused byincorrect operation of signalprocessing electronics.

Calibration exposures

Bias framesCCD defects, readout noiseexposure time = 0

Flat fieldpixel non-uniformity, dust, …exposure time ≈ medium (half saturation)

Dark framesdark current, defectsexposure 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 imagesand combined with them during image processing. These calibration exposures allow us to compensate forcertain imperfections in the CCD. As much care needs to be exercised in obtaining these images as forthe 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 framemay 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 reducesthe 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.


Flat Fields (pixel non-uniformity, dust, etc.)Some pixels in a CCD will be more sensitive than others (non-uniformity due to slightly different pixeldimensions and depth because of manufacturing process). In addition there may be dust spots on the surfaceof 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-uniformitiesand 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 orthe 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 effectmust be minimized by obtaining ‘dark frames’ at the beginning of the observing run. These are exposures withthe same duration as the science frames but taken with the camera shutter closed. These are later subtractedfrom 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.

Dark Frame Flat Field

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 dustspots are also visible.


Flat Field Frame

Bias Frame

Flat-Bias

Science -Dark

Output Image

Flat-BiasScience -Dark

Dark Frame

Science Raw Image

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.

Bias Frames and Flat Fields 5.

Flat Field Image

Bias Image

Flat-Bias

Science -Bias

Output Image

Flat-BiasScience -Bias

Science Frame

In the absence of dark current, the process is slightly simpler :

Pixel Size and Binning 1.

It is important to match the size of a CCD pixel to the focal length of the telescope. Atmospheric seeingplaces 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 offeatures 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)

Nyquist Sampling

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 space-varying 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.

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Aperture in mm X f-number


Example 1.The William Herschel Telescope, with a 4.2m diameter primary mirror and a focal ratio of 3 is to be usedfor 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) = =16.4 arc-sec per mm

(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 nearestpixel size available is 13.5 microns which leads to a small degree of over-sampling.

Matching the Pixels to the telescope

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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) = =103 arc-sec per mm

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.

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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’. Withbinning 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 advantageof ‘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 withoutthe other to yield rectangular pixels.

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.

Stage 1 :Vertical Binning

Charge packets


The first row is transferred into the serial register


The serial register is kept stationary ready for the next row to be transferred.


The second row is now transferred into the serial register.


Each pixel in the serial register now contains the charge from two pixels in the image area. Itis thus important that the serial register pixels have a higher charge capacity. This is achieved by giving them a larger physical size.


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 :

Stage 2 :Horizontal Binning

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SW

Output Node


Charge is clocked horizontally with the SW held at a positive potential.

123

SW

Output Node


123

SW

Output Node


The charge from the first pixel is now stored on the summing well.

123

SW

Output Node


The serial register continues clocking.

123

SW

Output Node


123

SW

Output Node


The SW potential is set slightly higher than the serial register electrodes.

123

SW

Output Node


123

SW

Output Node


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 voltageon 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.

123

SW

Output Node


Finally the charge is dumped onto the output node for measurement

123

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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 astronomicalapplications, 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 usedby 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 outmuch 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 clocks

Serial Register clocks Serial Register

Output Amplifier

Image Area

Slow Scan CCDs 2.

A slightly more complex design uses 2 serial registers and 4 output amplifiers. Extra clocklines 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 theimage in four quadrants simultaneously, reducing the readout speed by a factor of four.

Upper Image area clocks

Lower Image area clocks

Amplifier C

Amplifier A Amplifier B

Amplifier D

Serial clocks C Serial clocks D

Serial clocks A Serial clocks B

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 ora 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 rareand 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.

Image area clocks

Store area clocks

Amplifier

Serial clocks

Image area

Store area

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. Thismask could be a layer of aluminum deposited on the CCD surface or it could be an externalmask. This geometry is the basis of the ‘Frame transfer’ CCD that is used for high frame rate videoapplications. The area available for imaging is reduced by a half. The lower part of the image becomes the ‘Store area’.

Opaque mask

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 areabeneath 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 beintegrated in the image area during the readout. The image area can be kept continuously integratingand the detector has only a tiny ‘dead time’ during the image shift. No external shutter is requiredbut 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 pixelcharge 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 leveland the Signal level. The first is then subtracted from the second to yield the output signal voltageproportional 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.

OD

OS

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. ADC-1

Pre-Amplifier

CCD On-chip Amplifier

Inverting AmplifierIntegrator

Reset switch

Input Switch Polarity Switch


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).

Com

pute

r Bus

t0 t0


Output wave-form of CCD

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.

-1

Output voltage of CDS

t1 t2 t1 t2

-1


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.

Referencewindow

t1 t2 t3t3t2

-1


Between t2 and t3 the ‘charge dump’ occurs in the CCD. The CCD output steps negatively by an amountproportional to the charge contained in the pixel. During this time the CDS is disconnected.

t3 t4 t1 t2 t3 t4

-1


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.

Signal window

t1 t2 t3 t4

ADC

-1


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.

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. V

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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 CURRENTCaused by thermally generated electrons in the CCD even in the absence of lighting, increasingwith temperature and integration time.Consequences: generates “thermal noise” limiting the detection of faint sources and can evensaturate the pixels. Eliminated by cooling down the CCD usually at ~-100ºC.

3. PHOTON NOISEAlso called ‘Shot Noise’. It is due to the fact that the CCD detects photons. Photons arrive in anunpredictable fashion described by Poisson statistics. This unpredictability causes noise.

4. PIXEL RESPONSE NON-UNIFORMITYDefects 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’.

NR


Before these noise sources are explained further some new terms need to be introduced.

FLAT FIELDINGThis involves exposing the CCD to a very uniform light source that produces a featureless and evenexposure 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 REGIONSA 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 witha serial register longer than the width of the image area will also create vertical bias strips at the leftand 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 variousnoise sources present in the CCD but also about the gain of the signal processing electronicsi.e. the number of photoelectrons represented by each digital unit (ADU) output by the camera’sAnalogue to Digital Converter.


Flat field images obtained from two CCD geometries are represented below. The arrows representthe position of the readout amplifier and the thick black line at the bottom of each image represents the serial register.

CCD With SerialRegister equal inlength to the image area width.

CCD With SerialRegister greater inlength than the image area width.

Image Area

Image Area

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X-o

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X-u

nder

scan

Y-overscan

Y-overscan

Here, the CCD is over-scanned in X and Y

Here, the CCD is over-scanned in Yto 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 extensionsdo not receive any photo-charge.

Noise Sources in a CCD Image 4.These four noise sources are now explained in more detail:

READ NOISEThis 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 alwaysa 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.

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Time spent measuring each pixel (microseconds)

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DARK CURRENTElectrons can be generated in a pixel either by thermal motion of the silicon atoms or by the absorptionof 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 liquidnitrogen (-196ºC) to the point where the dark current falls to below 1 electron/pixel/hour (~-100ºC) whereit is essentially un-measurable. Amateur cameras cooled thermoelectrically may still have substantialdark current. The graph below shows how the dark current of a TEK1024 CCD can be reduced by cooling.

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-110 -100 -90 -80 -70 -60 -50 -40

Temperature Centigrade

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PHOTON NOISEThis can be understood more easily if we go back to the analogy of rain drops falling onto an arrayof buckets; the buckets being pixels and the rain drops photons. Both rain drops and photons arrivediscretely, 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 collectnone 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 essentiallythat 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/secondand we observe the star for 1 second, then the uncertainty of our measurement of its brightnesswill 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 measurementof any pixels photo-charge would be 30 electrons}


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 sensitivitybetween 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.

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HOW THE VARIOUS NOISE SOURCES COMBINEAssuming that the PRNU has been removed by flat fielding, the three remaining noisesources combine in the following equation:

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 levelapplications such as Spectroscopy. At higher signal levels, such as those found indirect 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 contributesonly 11% of the total noise.


NOISEtotal = (READ NOISE)2 + (PHOTON NOISE)2 +(DARK CURRENT)2

Photon Transfer Method 1.

Using two identical flat field exposures it is possible to measure the read noiseof a CCD with the Photon Transfer method. Two exposures are required to removethe 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.


Bias area 1Image area 1

Bias area 2Image area 2

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 result= MeanADU ,the Mean Signal in ADU.

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.

Flat Field Image 1.

Flat Field Image 2.


Image area 3

STEP 4

Measure the Standard Deviation in image area 3result= 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.

STEP 3 The two images are then subtracted pixel by pixel to yield a third image

Image 1 - Image 2 = Image 3

Gain = 2 x MeanADU

(StdDevADU ) 2 - (2 x NoiseADU 2).

The units will be electrons per ADU, which will be inverselyproportional to the voltage gain of the system.


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 ADUbut not so high that the chip or the processing electronics is saturated. 10,000 ADUwould be ideal. It is best to average the gain values obtained from several pairsof 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 whichthe 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-regionsshould be kept small.

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Potential along this line shown in graph above.

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Cross section through a thick frontside illuminated CCD

Deep Depletion CCDs 1.

The electric field structure in a CCD defines to a large degree its Quantum Efficiency (QE). Considerfirst a thick frontside illuminated CCD, which has a poor QE.

In this region the electric potential gradient is fairly low i.e. the electric field is low.

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-electronand the hole it leaves behind.

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Cross section through a thinned CCD


In a thinned CCD , the field free region is simply etched away.

There is now a high electric field throughout the full depth of 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.

This volume is etched away during manufacture

Problem : Thinned CCDs may have good blue response but they become transparent at longer wavelengths; the red responsesuffers.

Red photons can now passright through the CCD.

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Cross section through a Deep Depletion CCD

Deep Depletion CCDs 3.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 40m 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.

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.

Red photons are now absorbed in the thicker bulk of the device.

Problem :High-resistivity silicon contains much lower impurity levels than normal. Very few waferfabrication factories commonly use thismaterial and deep depletion CCDs have to be designed and made to order.


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 froma 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.


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 bereflected 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 thosethat have quite respectable QEs in the red. Thicker deep depletion CCDs , which have a muchlower degree of internal reflection and much lower fringing are preferred by astronomersfor 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 smallerCCD would require hundreds of images and would be an extremely inefficient use of the telescope’svaluable time. It is unlikely that CCDs will ever reach the same size as photographic detectors, so forapplications 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 challengeassociated 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 achievedby using packages with electrical connections along one side only, leaving the other three sides freefor 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 gapbetween them. Light falling down this gap is obviously lost and causes the black stripdown 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 imagewas 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.

Picture : Canada France Hawaii Telescope

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.

Mosaic Cameras 7.

This mosaic of 4 science CCDs was built at the Royal Greenwich Observatory. The positioningof 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.

Pre-amplifier Pressure Vessel Vacuum pump port

Liquid Nitrogen fill port

Camera mountingFace-plate.

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 varioustelescope ports.

Mounting clamp


The main body of the camera is a 3mm thick aluminium pressure vessel able to supportan 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 externalsurfaces 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 forthe 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 LN2can. 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 mountingblock 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 fusedsilica window big enough for the CCD and thick enough to withstand atmosphericpressure. 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.


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 contaminationof the CCD surface. A cold CCD is very prone to contamination from volatile substances suchas certain types of plastic. These out-gas into the vacuum spaces of the camera and then condenseon 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 theend of the pumping period is about 10-4 mBar. When the LN2 is introduced, the pressure willfall 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.


. . .

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 ofwater vapor: either in vacuum or in a dry nitrogen atmosphere).

Thermally Electrical feed-through Vacuum Space Pressure vessel Pump PortInsulatingPillars

Foca

l Pla

ne

of T

eles

cope

Tele

scop

e be

am

Optical window CCD CCD Mounting Block Thermal coupling Nitrogen can Activated charcoal ‘Getter’

Boil-off

Face-plate


CCDTemperature servo circuit board

Platinum resistancethermometer

Gold plated copper mounting block

Top of LN2can

PressureVessel

Signal wires to CCDLocation points (x3)for insulating pillarsthat reference the CCDto the camera face-plate

Aluminised Mylarsheet

Retainingclamp

The camera with the face-plate removed is shown below

‘Spider’.The CCD mountingblock is stood off from the spider usinginsulating pillars.

Camera Construction Techniques 6.A ‘Radiation Shield’ is then screwed down onto the spider, covering the cold components but notobstructing the CCD view. This shield is highly polished and cooled to an intermediate temperatureby a copper braid that connects it to the LN2 can.

Radiation Shield


‘Spider’ Vane

CCD Clamp plate

CCD Signal connector (x3)Copper rod or ‘cold finger’used to cool the CCD. It isconnected to an LN2 can.

Gold plated copper CCD mounting block.

CCD Package

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.

FOS 1 Spectrograph

Introduction to CCDs

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