Emission Line Surveys Lecture 1 Mauro Giavalisco Space Telescope Science Institute University of Massachusetts, Amherst 1 1 From January 2007

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

Disclaimer We wrote these lectures from the point of view of the “observer” 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 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 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 It also reflects our personal tastes and bias Readers are strongly encouraged to do further, comparative research in any specific subject discussed here Readers are strongly encouraged to do further, comparative research in any specific subject discussed here

Why Emission Line Surveys To effectively look for a specific class of sources in some pre-assigned volume of space and/or at some pre-assigned point in time To effectively look for a specific class of sources in some pre-assigned volume of space and/or at some pre-assigned point in time “effectively”: with high yield (low contamination) and in large numbers “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 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 Traditional flux selection plus follow-up spectroscopy highly inefficient to cull special classes of sources from the general counts

Notations, Definitions, Reminders and World Model. I Throughout these lectures, we use: Throughout these lectures, we use:  F: flux, in units of erg/s/cm 2  f : flux density, in units of erg/s/cm 2 /Hz  f : flux density, in units of erg/s/cm 2 /Å  f  f |d /d  = f c/ 2  1 Å = cm  c = cm/s

Notations, Definitions, Reminders and World Model. II Throughout these lectures, we use: Throughout these lectures, we use:  L: luminosity, in units of erg/s  l : luminosity density, in units of erg/s/Hz  l : luminosity density, in units of erg/s/Å  f = l (1+z) / 4  D L 2 (z)  f = l / 4  D L 2 (z) (1+z)  F = L / 4  D L 2 (z) D L (z) = D L (z; H 0,  m,   ) : luminosity distanceD L (z) = D L (z; H 0,  m,   ) : luminosity distance z is the redshift defined as z = a(t 0 )/a(t) – 1z is the redshift defined as z = a(t 0 )/a(t) – 1 t is the cosmic time and t 0 is the age of the universet is the cosmic time and t 0 is the age of the universe

Notations, Definitions, Reminders and World Model. III Throughout these lectures, we use: Throughout these lectures, we use:  AB magnitudes:  m AB = -2.5 Log 10 (f ) (Oke 1974; Oke & Gunn 1977)(Oke 1974; Oke & Gunn 1977)  ST magnitudes  m ST = -2.5 Log 10 (f ) (Walsh 1995)(Walsh 1995)  World Model (when needed):  H 0 = 70 km/s/Mpc   m = 0.3;   = 0.7

CCD and near-IR Detectors Most common devices used in emission line surveys Most common devices used in emission line surveys Photon counting devices: Photon counting devices:  DN = G N   DN: Calibrated Data Number, I.e. what we read from the detector after calibrations  G: inverse gain  N  : number of photons, in a finite wavelength interval  Detectors add their own “signal” and noise: Detectors add their own “signal” and noise:  DN obs = DN + K +    is removed during calibration (bias + d.c. + …)   is a random variable with < <  <    ron 2 rms + d.c. 2 rms + …<    ron 2 rms + d.c. 2 rms + … Typical values:Typical values: –[ron 2 rms ] 1/2 ~ a few (as low as ~1) to a few 10 e - /pix –[d.c. 2 rms ] 1/2 ~ 0.01 to a few e - /sec/pix Let’s assume G=1 in the following Let’s assume G=1 in the following

The Finite Resolution element The smallest spatial scale or wavelength interval the instrumentation can resolve: 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) arcsecGood (bad) seeing: 0.6 (2) arcsec HST resolution (V band): 0.03 arcsecHST 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)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 If pixel size is well matched to resolution element (Nyquist sampling): FWHM (of PSF or LSF) covered by 4 pixels

S/N: Signal-to-Noise Ratio Most important metric to asses sensitivity. Most important metric to asses sensitivity. S/N in some finite wavelength interval  either the passband width or the spectral resolution element S/N in some finite wavelength interval  either the passband width or the spectral resolution element since we detect (count) photons, uncertainty on photon counting is simply  = N 1/2, and thus: since we detect (count) photons, uncertainty on photon counting is simply  = N 1/2, and thus:  S/N = S  / [S  + B  + N 2  ] 1/2  S  : number of photons from source  B  : number of photons from background  N  : equivalent number of photons from additional sources of noise (typically detector)

Width of Emission Lines The finite width of an emission line along the wavelength axis. 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: Commonly measured by the Full Width at Half Maximum (FWHM). For a gaussian line profile:   ~ 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 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   / =  v / c If source is at redshift z, wavelengths are “stretched” by (1+z), thus observed FWHM and rest-frame FWHM related by: 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)

Equivalent Width of Emission Lines Metric to asses the strength of an emission line. 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 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 It represents the wavelength range over which the continuum luminosity equals the line luminosity  W = L / l = F / f  W = L / l = F / f Unaffected by extinction (line and continuum extinct by equal amount) 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 W and rest-frame W related by If source is at redshift z, wavelengths are “stretched” by (1+z), but luminosity (number of photons) is conserved. Thus, observed W and rest-frame W related by  W (obs) = W (rest) (1+z)

How to Detect Emission Lines Directly: observing the spectra of some class of candidates 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) 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)

How to Detect Emission Lines: Spectroscopy

Spectroscopy Ly  z=5.65 Vanzella et al., in prep.

How to Detect Emission Lines: Photometry

Finding galaxies at high- redshift: color selection B 435 V 606 z 850 Unattenuated Spectrum Spectrum Attenuated by IGM B 435 V 606 i 775 z 850 z~4 1.Color selection is very efficient in finding galaxies with specific spectral types in a pre-assigned redshift range 2.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”

How to Detect Emission Lines: Photometry

Source Selection

Emission-Line Sources

Weeding Out Interlopers

Spectroscopy Follow-up: the Interlopers

Spectroscopy Follow-up: the Targets