Telescopes & Instrumentation. Across the Infrared and Submillimetre Wavebands

Transcription

1 Telescopes & Instrumentation Across the Infrared and Submillimetre Wavebands Wayne Holland Lecture for the PPARC Introductory School for new astronomy postgraduate students Edinburgh: 6 th September 2005

2 Lecture Outline Basic principles Optical/Infrared telescope design Infrared imaging arrays Submillimetre-wave telescopes Submillimetre imagers

3 Basic Principles Figures of merit: Collecting area = πd 2 /4 Focal ratio = F/D Resolution ~ λ/d (arcsec) Plate scale ~ /F (arcsec/mm)

4 Basic Principles

5 Infrared Waveband wavelength (mm) Radio Near-IR Optical Millimetre Submm Far-IR Mid-IR Jodrell Bank JCMT UKIRT Gemini

6 Primary Mirrors How accurate does the surface error have to be? Essential requirement: The optical performance of the primary mirror should not impair the image quality permitted by atmospheric turbulence (even under the very best seeing conditions which may occur at the observatory site) Materials: Zerodur is preferred for large mirrors (low CTE) usually produced by spin casting

7 Primary Mirror Casting

8 Primary Mirrors The peak signal in the long-exposure Point Spread Function of a large groundbased telescope is given by where τ a is the transmissivity of the atmosphere, r 0 the atmospheric coherence length (typically 10-30cm at 0.5µm for a good site), τ t the transmissivity of the telescope, D its diameter The Central Intensity Ratio is defined by where, Ψ 0 is the Strehl ratio* of the telescope, solely taking into account the effect of atmospheric turbulence, and Ψ the same quantity, after the telescope errors have been taken into account. *The Strehl ratio is defined as the peak intensity in an image of a point source normalised by the peak intensity it would have if diffraction-limited

9 Practical Primary Mirrors Example: A map of the surface error of the mirror shows that the average bump is only 0.014µm, or about one five thousandths of the thickness of a human hair Subaru primary mirror If the primary mirror were the size of the Big Island of Hawaii, the average bump would only have the thickness of an ordinary sheet of paper!

10 Active Optics Mirrors deform under gravity as they move about the sky To maintain shape robotic fingers or actuators are attached to the back of the mirror A mirror that doesn t have to maintain its shape can be thin, and therefore relatively lightweight and manoeuvrable This technology is called active optics and makes the building of large aperture telescopes possible

11 Adaptive Optics Also for ground-based telescopes have to contend with seeing which distorts the image If can make the primary mirror deformable then can sense the light from a reference star and correct the incoming wavefront Curvature Sensing is a popular way to carry out wavefront sensing in which images from either side of focus are subtracted and normalised to yield a Laplacian of the wavefront which is then obtained by integration

12 Adaptive Optics

13 3.8m UKIRT 8-m Gemini IR Telescopes

14 Next Generation Challenges: - Segment fabrication and upkeep - Active optics - Adaptive optics - Secondary mirror - Enclosure?

15 GMST 30m+ Telescopes

16 The Ideal Detector Detect 100% of photons Each photon detected as a Large number of pixels Time lag for each photon Measure photon wavelength Up to 99% quantum efficiency One electron for each photon Over 400 million! Determined by frame readout Defined by filter Measure photon polarisation Defined by filter/polariser

17 Photon Detection 5 Basic Steps of Optical/IR Photon Detection: 1. Get light onto the detector - Mirrors, lenses, filters, cold stops, anti-reflection coatings 2. Charge generation - Popular materials include silicon, HgCdTe, InSb 3. Charge collection - Electric fields within the material collect photoelectrons into pixels 4. Charge transfer - For CCDs move photoelectrons to the edge of array where amplifiers are located 5. Charge amplification and digitisation - Need careful amp design as most generate excess noise

18 Practical Design Toroid Detector array Field stop Fixed flat Moving flat #1 Moving flat #2 Input flats Window Pupil stop Aspheric Filter wheels Input collimator

19 Infrared Pixel Geometry

20 Charge Generation For an electron to be excited from the conduction to the valance band: hν E g where E g is the energy gap of the material (electron volts) Conduction band Valence band E g Cut-off wavelength: λ c = / E g (ev)

21 Detector Zoology

22 Charge Collection Image intensity is generated by collecting photoelectrons onto a 2-D arrays of pixels Optical and IR focal plane arrays both collect charge via electric fields In the z-direction they use semiconductor diodes (p-n junction) to sweep the charge to the pixel collection nodes I p n

23 Charge Transfer IR detectors have an amplifier at each pixel, so no need for charge transfer CCDs must move charge across the focal plane array to the readout amplifier

24 Charge Amplification Similar for CCDs and IR detectors Both use MOSFETs (metal-oxide-silicon field effect transistors) to amplify the signal MOSFET

25 IR detector Cross-Section

26 Mid-IR Imager Michelle : Mid-IR Imager/Spectrometer Weight: 2 tonnes Cost: ~ 70 staff years ~ 4 million Operations: On UKIRT from Aug 2001 to Feb 2003 Since then its been on Gemini North telescope

27 Mid-IR Imager Michelle : Mid-IR Imager/Spectrometer Optics: All-reflective (except the front window and the filters). Detector: A pixel, Si:As IBC array. Cryogenics: Optics cooled to 26K, detector cooled to 6K.

28 Other IR Imagers WFCAM (UKIRT) MIRI Concept (JWST)

29 Far-IR/submm Wavebands wavelength (mm) Radio Near-IR Optical Millimetre Submm Far-IR Mid-IR Jodrell Bank JCMT UKIRT Gemini

30 Primary Mirrors Only the first micron or so of the surface is actually doing anything. Everything else is there just to hold it in the right place! Panels plus backing structure still the basic principle for large antennas. Space-frame no longer the only option for backing structure plate backing structures also used. CFRP is the best material for the structure low CTE, high stiffness to weight, but is expensive and it is tricky to join it to the rest of the antenna.

31 Active Control Already incorporate actuators to adjust the surface in order to get it set initially. So long as we can provide feedback, these can be used to adjust: Actuators on the GBT a) for non-homologous gravitational errors, using a look-up table based on a structural model, or possibly measurements; b) for thermal deformations based on temperature measurements; c) possibly for wind deformations based on wind or strain measurements.

32 Position Sensing Tend to use submm-wave telescopes in daylight hours so thermal gradients are a big problem! Use temperature sensors all over the dish and feed variations into a finite element model of the telescope This then generates the relevant movements of the actuators to keep the dish in shape

33 Submillimetre Telescopes At present, 10-15m class telescopes are operating with good surface quality (15-25µm rms) on good submillimetre sites JCMT JCMT 15m JCMT: Mauna Kea 10.4m CSO: Mauna Kea 10m HHT: Mount Graham CSO Current single dishes Coarse angular resolution (~10 arscec) Wide-field imaging? Limited sensitivity HHT

34 Wide-Field Design

35 Wide-Field Design Green Bank Telescope South Pole Telescope Concept

36 Airborne Telescopes Stratospheric Observatory for Infrared Astronomy (2005/6) 2.5m

37 Space Telescopes SPICA (2012?) 3.6m telescope cooled to ~4K Herschel Space Observatory (2007) 3.6m telescope

38 Interferometers Atacama Large Millimetre Array (2012?) 64 dishes, each of 12m in diameter at 5000m in the Chilean Andes

39 Detector Zoology Wavelength (microns) Velocity resolution (km/s) Bolometer Arrays Coherent SIS mixer arrays Schottky mixers Imaging FTS Photoconductor Arrays Grating Spectrometer Fabry-Perot Incoherent Frequency (GHz)

40 Practical Instrument Design Optical relay Filter ν Detector P V/I Data acquisition computer Modulated signal from telescope Instrument cryostat Low-noise amplifier Important considerations in the instrument design: Detector operating temperature Optical coupling Wavelength selection Low-noise electronic readouts Data acquisition system Environmental issues - microphonic pick-up, RFI,grounding, stray light

41 Bolometers Two components: A sensitive thermometer and high cross-section absorber temperature time Thermometer and absorber are connected by a weak thermal link to a heat sink Incoming energy is converted to heat in the absorber: T = T T 0 = E/C Temperature rise decays as power in absorber flows out to the heat sink Temperature rise is proportional to the incoming energy

42 Bolometer Operation Classical Germanium bolometer: Bias circuit with voltage source and load resistor Incident Power, Q Bias current, I A constant current, generated by the bias supply and load resistor, flows through the bolometer I Absorber Temperature T Heat capacity C V bias Load resistor R L Thermal Conductance, G Thermometer Bolometer R V HEAT SINK at temperature T 0 Provided the bias power remains constant: T = T 0 + (Q + P bias )/G The temperature rise causes a change in bolometer resistance and consequently in the voltage across it

43 Performance Parameters The Noise Equivalent Power (NEP) is the power absorbed that produces a S/N of unity at the bolometer output. Can be written as: NEP 2 = NEP 2 (detector) + NEP 2 (background) Ideal detector NEP consists of contributions from Johnson and Phonon noise. Background NEP is due to photon noise power from the sky, telescope and instrument. The overall NEP of a device can be written as: NEP 2 = 4k B TR/S 2 + 4k B T 2 G + 2Q(hν 0 +ηεkt bk ) The thermal time constant is a measure of the response time of the bolometer to incoming radiation. Can be written as: τ = C/G

44 Semiconductor Bolometers Composite design: - Metal-coated dielectric as absorber - Semiconductor resistance thermometer Used on ground-based telescopes for many years Theory and practice well understood Current state-of-the-art: Spider-Web composite bolometers 25 B150A Dark PSD 20 V n ( V rms / Hz 1/2 ) Low thermal conductivity - high sensitivity No 1/f noise down to 20 mhz Low absorber heat capacity - fast response time Low-cosmic ray cross-section (few %) Minimal suspended mass - robust Freq (Hz)

45 Bolometer Arrays Built in Edinburgh and came into service in 1997 Regarded as one of the most successful instruments ever built for a groundbased telescope

46 SCUBA Arrays But only has 128 pixels in two arrays How can we scale up to thousands of pixels?

47 Superconducting TES Bolometers Resistance R N I SQUID Amplifier R C V bias R(T) TES T c Temperature Voltage-biased on normal - superconducting transition Resistance is a very steep dependence on temperature in transition region Film held at constant voltage bias - change in resistance results in a change in current through the film Low noise, low power (~ 1nW) SQUID ammeter readout In development at Cardiff, NIST/GSFC, JPL, UC Berkeley, SRON

48 Practical TES Bolometers Normal metal Superconductor Substrate A bilayer of thin superconducting film and a thin normal metal act as a single superconductor By choosing the film thickness can reproduce TES devices with sharp (~2mK) tuneable transitions Cross-section of a proximity effect TES Mo-Cu PUD with thick normal bars X-ray detector IR/submm detector In*V (Watt/sqrt(Hz) Very sharp transition - big change in R for a small change in T Low resistance - Johnson noise negligible - high sensitivity Good low-frequency stability Signals can be readout by multiplexed SQUID amplifiers Frequency (Hz)

49 Operating Temperature For a fixed response time, the detector NEP is TC 1/2 Typical bolometers have materials which have heat capacities which are proportional to T (metals) and T 3 (dielectrics) So can write detector NEP as: NEP T 3/2 T 5/2 1E-11 1E-12 1E-13 NEP (W/ Hz) 1E-14 1E-15 NEP T 3/2 NEP T 5/2 1E-16 1E-17 1E Temperature (K)

50 Optical Coupling Size of a bolometric detector (typically a few mm) is small compared to the telescope diffraction spot at submm/mm wavelengths. There are two basic ways of coupling this radiation onto the detector: Illumination pattern Conical horn with section of cylindrical waveguide Horn defines the detector field-of-view Bolometer in integrating cavity D 2λ/D Feedhorn Gives a tapered (~Gaussian) illumination of the telescope Maximum aperture efficiency when horn diameter = 2λ/D Throughput is single-moded, i.e. AΩ=λ 2 Beam on the sky is close to the diffraction limit (~λ/d)

51 Optical Coupling Alternative is to dispense with feedhorns and simply have a bare pixel in the focal plane: Illumination pattern D Cold enclosure Bare pixel with cold aperture stop Pixel field-of-view is large (~π steradians) Illumination to the telescope is defined by a cold stop Results in a near top-hat illumination of the telescope

52 Telescope Performance Performance on the telescope is represented by the Noise Equivalent Flux Density (NEFD) This is the flux density that produces a signal-to-noise of unity in one second of integration: NEFD NEP ηη A e = τa c t e ν (mjy/ Hz) and 850 micron NEFD as a function of sky transmission NEFD depends very much on the weather and varies with sky transmission NEFD (mjy/rthz) µm 450µm On many occasions the fundamental sensitivity limit is set by sky-noise Normalised sky transmission

53 The Future... SCUBA-2 is a 10,000 pixel camera under development for the JCMT: Two cameras of TES devices each with 4 sub-arrays Multiplexed SQUID readouts Fully-sample the 850µm image plane (450µm will be undersampled by a factor of 2) Operate at 450 and 850µm simultaneously ~50sq-arcmin field-of-view Background-limited sensitivity DC coupled arrays (no sky chopping) SCUBA-2 will map large areas of sky 1000 times faster than the current SCUBA to the same S/N.

54 Pixel Architecture Deep-etched trench (10µm wide) mm Quarterwave Si Brick Transition Edge Sensor Detector Wafer SQUID MUX Wafer Nitride membrane (0.5 µm) Indium bump bonds (216,000) Not to scale

55 Detector Arrays SCUBA 850µm array (same pixel scale) Completed (1280) pixel prototype array The SCUBA-2 prototype array is the first of its type ever built!

56 In-Focal-Plane Multiplexers Input transformer Active SQUID Summing coil gradiometer ~1mm Dummy SQUID A full-sized (40 32 pixel) multiplexer wafer The output from the TES arrays are time-division multiplexed using a SQUID-based MUX developed by NIST

57 Prototype Sub-Array Module Sub-array (with cover) ~45mm Batwing PCB Niobium flex cables 1K SQUID amplifiers

58 Summary Optical and IR telescopes are currently in the 8-10m range with m class telescopes in various stages of design IR arrays are well developed and already have a million-plus pixels and continue to be developed for wide-field survey instruments Submm-wave telescopes have only just left the first generation stage interferometers and cooled-aperture telescopes in space are next Rapid developments in detector technology and readout circuitry in the submm will enable much larger arrays to be constructed the first submm/mm CCDs?