Instrumentation Concepts Ground-based Optical Telescopes

1. IAG-USP (Keith Taylor) Instrumentation ConceptsGround-based Optical Telescopes Keith Taylor (IAG/USP) Aug-Nov, 2008

2. Imaging considerations • Trading field of view vs. angular resolution • A large field at coarse spatial resolution or smaller field of view at high fidelity? • Generally detector defines fixed pixel format • Constraints also driven by seeing and pixel scale (or camera f-ratio) • Photometric accuracy • Simple morphological discrimination or accurate flux measures as well? • Definition of passbands and central wavelengths • Multiple simultaneous passbands?

3. Astronomical CCD Imaging • Simplest astronomical instrument (in principle) • Undispersed 2D images of field of view • Generally use filters to limit spectral band-pass • Detector itself may supply band-width • Polarimetric capabilities? • Rapid reads can give limited time resolution information

4. Standardized Filter Systems • Variety of different filter systems prevalent in optical/IR domain. eg: • UBVRI / JHKLM – Johnson/Cousins (UV/optical) (NIR) • u g r i z - Sloan Digital Sky Survey (SDSS) optical bands • Extras and modifications • y - UKIDSS IR band • K’ and K* - modifications to K to avoid thermal radiation • HST, Spitzer etc. defined by wavelength rather than name (revolutionary!)

5. Examples of CCD Imaging CCD = Charge-coupled device More sensitive than photographic plates by factors ~50 Data can be read directly into computer memory, allowing digital enhancement and manipulation Negative image can enhance contrasts False-color image to visualize brightness contours

6. cf: Photographic Imaging(eg: AAT c1990) • 14inch (=350mm) • ~5 Giga Pixels • cf: Largest CCD currently: • 8 Mega Pixels • (4k-by-2k) • CCD arrays up to • ~0.1 Giga Pixels • Generally seeing-limited • How do we obtain higher spatial definition?

7. Adaptive Optics Computer-controlled “adaptive” mirror adjusts the mirror surface (many times per second) to compensate for distortions induced by atmospheric turbulence And yet further spatial resolution?

8. Interferometry Recall: Resolving power of a telescope depends on diameter D: min = 1.22 /D. This holds true even if not the entire surface is filled out. • Sparsely filled aperture: • Combine the signals from several smaller telescopes to simulate one big mirror  • Interferometry

9. Spectroscopic considerations • What kind of spectral feature are of interest? • Emission or absorption lines; continuum shapes • Broad, narrow or spectrally unresolved • Low or high contrast with continuum • Spectral Energy Distributions (SEDs) • Line centres ; equivalent widths ; line shapes ; kinematic mapping? • and/or precise spectrophotometry? • One or many targets simultaneously?

10. The simplest spectrograph Using a prism (or a grating), light can be split up into different wavelengths (colors!) to produce a spectrum. Spectral lines in a spectrum tell us about the chemical composition and other properties of the observed object

11. Typical grating spectrograph • Simple grating spectrograph • Spectrum extracted along a slit so ‘imaging’ in one dimension • Off source light along slit used to measure and subtract sky background

12. What you get • Optical long slit spectrum of a galaxy • Minimal data reduction in displayed spectral image • Can see galaxy, bad pixels, cosmic ray hits and sky lines • Need off source signal to measure and extract target (sky subtraction) Sky lines Target

13. Considerations for Spectroscopy • Basic parameters - resolution and central wavelength for spectrum • Slit width (if selectable) affects resolution • Wavelength range • Set by combination of detector geometry and spectral resolution • Some spectrographs provide large -range at low-R; others provide only a few 1,000kms-1 range, so centering on a critical line of interest (eg: H) But, what if you need both high-R and large -range?

14. High Resolution and lots ofSpectrum • X-dispersed echelle grating spectrometers allow high resolution and lots of spectral coverage • Achieve this by having two orthogonal gratings • One gives the high resolution (in y-axis) the other spreads the spectrum across the detector(in x-axis) • However, the slit is consequently much shorter

15. STELES echelle spectrograph(for SOAR) Primary disperser (echelle grating) Secondary (orthogonal) disperser (VPHG) Blue channel Red channel

16. Multiobject Spectroscopy • To get spectra for lots of objects at once. Can use two approaches • Multislit - have several slits in the image plane and get spectra for all of them • Use fibres to pipe light from different parts of the focal plane while reformatting them along the spectrograph slit • Both techniques were developed in the 80s and perfected in the 90s

17. Fibre Fed Systems • AAT 2dF (now replaced by AAOmega) • Pickoff fibres positioned by robot • Include sky fibres for each object

18. Multi-slit spectroscopy • Example of multislit spectral image • Easier to achieve at telescope (can use holes in a mask) but preparation and reduction can be more complex • Need to ensure spectra don’t overlap

19. LDSS-2mask superimposed on sky image Great care has to be taken in selecting objects to study so that they don’t overlap in wavelength direction. Also need objects of similar brightness so the SNRs are similar. Mask optimization is NOT trivial! Field acquisition is NOT trivial But what if you want images and spectroscopy simultaneuously?

20. Integral Field Spectroscopy • Extended (diffuse) object with lots of spectra • Use “contiguous 2D array of fibres or ‘mirror slicer’ to obtain a spectrum at each point in an image Tiger SIFS MPI’s 3D

23. Large-field imagingspectrographs • Narrow band filters • Image a field in a single narrow band • Use enough narrow bands and you have very low res. spectroscopy • Fabry-Perot • Effectively acts as a narrow tunable filter • Can thus image a field in emission lines of choice (eg. TTF)

24. Fabry-Perot • Light enters etalon and is subjected to multiple reflections • Transmission spectrum has numerous narrow peaks at wavelengths where path difference results in constructive interference • need ‘blocking filters’ to use as narrow band filter • Width and depth of peaks depends on reflectivity of etalon surfaces: finesse

25. Fourier Transform Spectrometer • As translation mirror scans an interference pattern is produced that is the FT of the source spectrum • Scan distance defines the resolution of the spectrum • • Advantage - get spectrum of whole field • • Disadvantage - get broad band noise

26. IFTS for NGST

27. Detectors for Opical/near-IR(current) • Photon Counters: • Image tube + TV camera + real-time discrimination (not solid state) • eg: IPCS - c70s to c80s • CCDs now dominate - Hi QE but … • Integrate signal on detector – no time resolution • Finite read-noise • Finite read-time • EMCCDs – new generation of Photon Counters • CCD-like QEs • V. high frame-rates

28. DQE - the key to gooddetectors • Detector quantum efficiency - the fraction of incident photons detected - is the key measure for the effectiveness of a detector; • Traditional photographic plates, while large in size, have DQE of only about 10% • CCDs and similar semiconductor devices can have DQE as high as 90% (though wavelength dependent) • Like having a telescope with 9 times the collecting area

29. CCDs • CCDs combine photon detection with integration and multiplexing • Incident photons excite charge carriers which are stored and integrated in a capacitor • CCDs are also uniquely effective in transferring charge from 2D to 1D • charge ‘clocked’ from pixel to pixel and read out at fixed point • ideal for multiplexing

30. CCD Array Camera • Semiconductor fabrication limits the size of a CCD detector • To get a large area need to mosaic detectors together Subaru Mosaic CCD Camera

31. Near-IR Detectors • CCDs use Silicon as their substrate • Valance to conduction bandgap in silicon is 1.1eV so restricted to detecting photons with wavelength < 1 micron • Need different materials for infrared • InSb for 1 to 5 micron, HgCdTe for 1 to 2.5 micron • Detector elements bonded to Si CCD system to provide multiplexing readout

33. IR Arrays vs. Optical • IR arrays are smaller, more expensive (by factor of ~10/pixel) • Readout has to be faster because of higher backgrounds • Use of different materials can push to longer wavelengths • More difficult to work with, less helpful characteristics, more expensive • At longest wavelengths have to stress the detector to produce lower energy band gaps

34. Lecture 1 (2 sessions) - Synopsis • Fundamental Principles • Introduction • Information Theory • Seeing-limited observations • Diffraction-limited observations • Signal-to-noise estimates

35. IAG/USP (Keith Taylor) Lecture 2 Fundamentals

36. UVOIR Astronomy Definition: • UVOIR = the "UV, Optical, Near-Infrared" region of EM spectrum • Shortest wavelength: 912 Å (or 91.2 nm) --- Lyman edge of H I; interstellar medium is opaque for hundreds of Å below here • Longest wavelength: ~3µm (or 3000 nm) --- serious H2O absorption in Earth's atmosphere above here • Ground-based UVOIR: • 0.3µm (or 300nm) <  < 2.5µm (or 2,500nm)

37. UVOIR Astronomy Uniqueness: • Best developed instrumentation; • Best understood astrophysically; • Highest density of astrophysical information; • Provides prime diagnostics on the two most important physical tracers. ===> UVOIR observations/identifications are almost always prerequisites to a thorough understanding of cosmic sources in other EM bands.

38. Proof Plasma (to 105K) Stars

39. Observational Priorities • Astronomy driven by discoveries rather than theoretical insights • Direction of field shaped by observations in about 3/4 of instances. • Few important astronomical discoveries were predicted; many were actually accidental

40. Accidental Discoveries • Uranus • Expanding universe • Pulsars • Supermassive black holes/AGN‘s • Large scale structure • Dark matter in spiral galaxies • X-ray emitting gas in clusters of galaxies • Gamma ray bursts • Extra-solar planets • High redshift evolution of galaxies • HST contributions were actually hindered by theoretical prejudice. A deep pencil-beam survey was delayed by 5 years.

41. Counterexamples: theory-driven discoveries • Neptune • General relativistic distortion of space-time near Sun • 21 cm line of HI • Helioseismology • Cosmic microwave background Conclusions Is Observational Astronomy a Science? (Build, don’t Think)

42. Technology drives Discovery • Key technology development for UVOIR astronomy: • 17th century: telescopes • 19th century: spectroscopy, photography, quality lens making, large structure engineering • 20th century: large mirror fabrication, electronic detectors, digital computers, space astronomy • Since 1980: array detectors

43. Telescope size: determines ultimate sensitivity • Diameter doubling time ~45 years • Largest telescopes now 8-10m diameter • Collecting area of 10-m is 4*106 that of the dark-adapted eye • In planning: 15-m to 40-m • For a given technology, cost  D2.6 • Cost is roughly proportional to mass • Even using new technologies, next generation of large ground-based telescopes will cross the $1 billion threshold.

44. The Future? • NB: Number of ground-based telescopes is NOT inversely proportional to their size • Almost as many 8m telescopes as there are 4m telescopes • How many 30m telescopes are there going to be in the next 50 years? (at US$1B a pop)

45. Review of some Basics • c = ν = 3.1010 cm/s • E = hν (ergs) • F = L/4d2 • G = 6.67.10-8 (c.g.s) • c = 3.1010 cm.s-1 • k = 1.38.10-16 • h = 6.626.10-27 • mH ~ mproton = 1.67.10-24 grams • me = 0.91.10-27 grams • eV = 1.602.10-12 ergs • Luminosity of Sun = 4.1033 ergs/sec • Mass of the Sun = 2.1033 grams

46. Flux measurements • Signal-to-Noise Ratio • "Sensitivity"---i.e. the faintest source measurable---is not simply a matter of the size of the photon collector. • It is instead a signal-to-noise (SNR) issue: • SNR = measured value / uncertainty and is dependant on many things, including: • Structure of source (point vs. extended) • Nature of luminous background & surroundings • Foreground absorption • Telescope & instrument throughput • Characteristics of detectors (quantum efficiency, noise)

47. SNRs in Astronomy • Fundamental limit set by photon statistics: • SNR < √ N, where N = no. of detected source photons • Typical SNR's in Astronomy: • Measures of astronomical EM fluxes: • Best precision: SNR ~ 1000 (0.1% error) • Low by lab standards! Problems: difficulty of calibration; faintness of interesting sources. • Typical "good" measures: SNR ~ 20-30 • Threshold detections: SNR ~ 5-10

48. Noise Sources • Detector Noise (CCDs) • Read-noise (rms ~3-10e-1/read) • Dark noise (3.10-4e-1/s/pixel) • Determined by Temperature of detector • Background Noise (Diffuse) • Artificial light pollution • Earth's atmosphere • Ecliptic scattered sunlight • Scattered Galactic light • Background Noise (Discrete) • Exclusion zone around bright stars caused by scattered light within instrument • Source "confusion" caused by diffractive blending of multiple faint sources

49. Magnitude System • An ancient and arcane, but compact and by now unchangeable, way of expressing brightnesses of astronomical sources. • Magnitudes are a logarithmic measure of spectral flux density (not flux!) • Monochromatic Apparent Magnitudes • m = −2.5 log10 f − 21.1 where f is in units of erg.s−1.cm−2.Å−1 • A system of “monochromatic magnitudes per unit wavelength”

50. Magnitude Normalization • Normalization is chosen to coincide with the zero point of the widely-used “visual” or standard “broad-band” V magnitude system: • i.e. m(5500Å) = V • Zero Point: fluxes at 5500Å corresponding to m (5500Å) = 0, are (Bessell 1998) • f0= 3.63.10−9 erg.s−1.cm−2.Å−1; or • fν0= 3.63.10−20 erg.s−1.cm−2.Hz−1; or = 3630 Janskys • f0/hν = 1005 photons.s−1.cm−2.Å−1 is the corresponding photon rate per unit wavelength