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Emission Line Surveys Lecture 1

This lecture provides an overview of emission line surveys, covering topics such as survey types, observational techniques, and historical notes. It also discusses the importance of emission line surveys and their potential future prospects.

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Emission Line Surveys Lecture 1

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  1. Emission Line SurveysLecture 1 Mauro Giavalisco Space Telescope Science Institute University of Massachusetts, Amherst1 1From January 2007

  2. Outline • Definitions • Why emission lines • Types of surveys and methodology • Target surveys • Blind surveys • Sensitivity • Narrow-band imaging • Slit spectroscopy • Slitless spectroscopy • Observational techniques • Results from Emission Line Surveys • Historical notes • Discussion of recent and ongoing surveys • Methodology • Results • Future prospects

  3. Disclaimer • We wrote these lectures from the point of view of the “observer” • They do not aim at providing a complete review of emission line surveys and their results • Rather, the choice of material is aimed at maximizing pedagogical value, illustrating current interesting problems, and at helping potential observers planning and designing their own emission line surveys • It also reflects our personal tastes and bias • Readers are strongly encouraged to do further, comparative research in any specific subject discussed here

  4. Why Emission Line Surveys • To effectively look for a specific class of sources in some pre-assigned volumeof space and/or at some pre-assigned point in time • “effectively”: with high yield (low contamination) and in large numbers • Exploit the presence of emission line in the spectral energy distribution of most astrophysical sources • Traditional flux selection plus follow-up spectroscopy highly inefficient to cull special classes of sources from the general counts

  5. Notations, Definitions, Reminders and World Model. I • Throughout these lectures, we use: • F: flux, in units of erg/s/cm2 • fn: flux density, in units of erg/s/cm2/Hz • fl: flux density, in units of erg/s/cm2/Å • fl = fn• |dn/dl| = fn• c/l2 • 1 Å = 10-8 cm • c = 2.9979 • 1010 cm/s

  6. Notations, Definitions, Reminders and World Model. II • Throughout these lectures, we use: • L: luminosity, in units of erg/s • ln: luminosity density, in units of erg/s/Hz • ll: luminosity density, in units of erg/s/Å • fn = ln• (1+z)/ 4p• DL2(z) • fl = ll/ 4p• DL2(z) • (1+z) • F = L / 4p• DL2(z) • DL(z) = DL(z; H0, m, L) : luminosity distance • z is the redshift defined as z = a(t0)/a(t) – 1 • t is the cosmic time and t0 is the age of the universe

  7. Notations, Definitions, Reminders and World Model. III • Throughout these lectures, we use: • AB magnitudes: • mAB = -2.5 • Log10(fn) - 48.595 • (Oke 1974; Oke & Gunn 1977) • ST magnitudes • mST = -2.5 • Log10(fl) - 21.1 • (Walsh 1995) • World Model (when needed): • H0 = 70 km/s/Mpc • m = 0.3; L = 0.7

  8. CCD and near-IR Detectors • Most common devices used in emission line surveys • Photon counting devices: • DN = G • Ng • DN: Calibrated Data Number, I.e. what we read from the detector after calibrations • G: inverse gain • Ng: number of photons, in a finite wavelength interval Dl • Detectors add their own “signal” and noise: • DNobs = DN + K + e • K is removed during calibration (bias + d.c. + …) • e is a random variable with • <e> = 0 • < e2 > = ron2rms + d.c.2rms + … • Typical values: • [ron2rms]1/2 ~ a few (as low as ~1) to a few 10 e- /pix • [d.c.2rms]1/2 ~ 0.01 to a few e-/sec/pix • Let’s assume G=1 in the following

  9. The Finite Resolution element • The smallest spatial scale or wavelength interval the instrumentation can resolve: • Spatial (PSF): the seeing (ground) or diffraction limit (space) • Good (bad) seeing: 0.6 (2) arcsec • HST resolution (V band): 0.03 arcsec • Depends on the size of the telescope, wavelength and… luck! • Poor image quality spreads photons over a large area, adds noise (2x seeing = 4x noise) • Spectroscopic (resolution): the spectral resolution element • Depends on the dispersion of the spectral element (prism, grism, grating) and on the slit aperture • If pixel size is well matched to resolution element (Nyquist sampling): FWHM (of PSF or LSF) covered by 4 pixels

  10. S/N: Signal-to-Noise Ratio • Most important metric to asses sensitivity. • S/N in some finite wavelength interval Dl, either the passband width or the spectral resolution element • since we detect (count) photons, uncertainty on photon counting is simply dN = N1/2, and thus: • S/N = Sg / [Sg + Bg + N2g]1/2 • Sg : number of photons from source • Bg : number of photons from background • Ng : equivalent number of photons from additional sources of noise (typically detector)

  11. Width of Emission Lines • The finite width of an emission line along the wavelength axis. • Commonly measured by the Full Width at Half Maximum (FWHM). For a gaussian line profile: • s ~ 0.425 • FWHM • The line width reflects the kinematics of the emission region (kinematics of the gas or of the individual sources in the case of integrated emission). If v is a measure of the velocity field within the emission region • Dl / l = Dv / c • If source is at redshift z, wavelengths are “stretched” by (1+z), thus observed FWHM and rest-frame FWHM related by: • FWHM(obs) = FWHM(rest) • (1+z)

  12. Equivalent Width of Emission Lines • Metric to asses the strength of an emission line. • The width of a top-hat emission line of equal luminosity and peak value equal to the continuum at the line wavelength • It represents the wavelength range over which the continuum luminosity equals the line luminosity • Wl = L / ll = F / fl • Unaffected by extinction (line and continuum extinct by equal amount) • If source is at redshift z, wavelengths are “stretched” by (1+z), but luminosity (number of photons) is conserved. Thus, observed Wl and rest-frame Wl related by • Wl(obs) = Wl(rest) • (1+z)

  13. How to Detect Emission Lines • Directly: observing the spectra of some class of candidates • Indirectly: comparing the photometry of the line through narrow-band passbands (on-band images) to that of the continuum through either narrow or broad-band passbands (off-band images)

  14. How to Detect Emission Lines: Spectroscopy

  15. Spectroscopy Lya z=5.65 Vanzella et al., in prep.

  16. How to Detect Emission Lines: Photometry

  17. Unattenuated Spectrum Spectrum Attenuated by IGM B435 V606 z850 Finding galaxies at high-redshift: colorselection B435V606i775z850 • Color selection is very efficient in finding galaxies with specific spectral types in a pre-assigned redshift range • Wide variety of methods available, targeting a range of redshifts, galaxies’ SEDs: • Lyman and Balmer break (Steidel et al., GOODS) • BX/BM • (Adelberger et al., COSMOS) • DRG • (van Dokkum et al., GOODS) • BzK • (Daddi et al.) • Photo-z • (Mobasher et al) • Here, the case of • “Lyman-break galaxies” z~4

  18. How to Detect Emission Lines: Photometry

  19. How to Detect Emission Lines: Photometry

  20. Source Selection

  21. Source Selection

  22. Emission-Line Sources

  23. Weeding Out Interlopers

  24. Spectroscopy Follow-up: the Interlopers

  25. Spectroscopy Follow-up: the Targets

  26. Spectroscopy Follow-up: the Targets

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