Purpose of these lectures To prepare you for observing at the Isaac Newton Telescope What the telescope + instrument can do What it can’t do What the atmospheric and lunar conditions are likely to be and their effect on the observations To choose what we’ll observe Overview of lectures Lecture 1: an overview of observing – telescope, instrument, coordinate systems, atmosphere, and moon Lecture 2: an overview of photometry and photometric data reduction Lecture 3: the TAC – “Time Allocation Committee” La Palma observing trip: preparation Lecture 1 S.C. Trager Outline Astronomical coordinate systems Telescope properties: INT plate scale, speed, type Instrument: the WFC A short recap of noise Effects of the atmosphere and background sources Coordinate systems The Celestial Sphere First we need the concept of the celestial sphere. It would be nice if we knew the distance to every object we’re interested in — but we don’t. And it’s actually unnecessary in order to observe them! The Celestial Sphere Instead, we assume that all astronomical sources are infinitely far away and live on the surface of a sphere at infinite distance. This is the celestial sphere. If we define a coordinate system on this sphere, we know where to point! Further, stars (and galaxies) move with respect to each other. The motion normal to the line of sight — i.e., on the celestial sphere — is called proper motion (which we’ll return to shortly) Astronomical coordinate systems A bit of terminology: great circle: a circle on the surface of a sphere intercepting a plane that intersects the origin of the sphere Horizon coodinates A natural coordinate system for an Earth- bound observer is the “horizon” or “Alt-Az” coordinate system The great circle of the horizon projected on the celestial sphere is the equator of this system. Horizon coodinates Altitude (or elevation) is the angle from the horizon up to our object — the zenith, the point directly above the observer, is at +90º Horizon coodinates We need another coordinate: define a great circle perpendicular to the equator (horizon) passing through the zenith and, for convenience, due north This line of constant longitude is called a meridian Horizon coodinates The azimuth is the angle measured along the horizon from north towards east to the great circle that intercepts our object (star) and the zenith. Horizon coodinates The origin of these angles (coordinates) is the observer Note that this is a left- handed coordinate system! Horizon coodinates Nearly all big telescopes (diameter ≥ 4m, telescopes built after ~1990, most radio telescopes) are in alt-az mounts This is the natural coordinate system for these telescopes But this system is dependent on the location of the observer and time of the observation: makes consistent cataloguing of objects difficult! Equatorial coordinates +90º Let’s consider a coordinate system that is tied to the astronomical objects themselves — and preferably those that don’t move! ♈ –90º Equatorial coordinates +90º In equatorial coordinates, the celestial equator is the great circle that intersects both the celestial sphere and the Earth’s equator: it’s the projection of the equator onto the celestial ♈ sphere –90º Equatorial coordinates +90º The declination δ is the celestial latitude and is measured in degrees, with 0º at the equator, +90º at the North Celestial Pole (NCP) — the intersection of the Earth’s north (rotational) pole with the celestial sphere — and ♈ –90º at the South Celestial Pole –90º Equatorial coordinates +90º The right ascension (RA) α is the celestial longitude and is measured in units of time, 0–24 hours, from west to east, with 0h at the Sun’s position when it crosses the equator from south to north, ♈ approximately at noon on 21 March. –90º Equatorial coordinates +90º The position α=0h, δ=0º is called the vernal equinox ♈ this is the sign of the constellation Aries, where the vernal equinox happened 2500 years ago ♈ The equatorial system is a right-handed system –90º Equatorial coordinates +90º Because the Earth precesses around an average direction perpendicular to the ecliptic (the plane of the Earth’s orbit around the Sun) due to the torques exerted on by the Moon, Sun, and Jupiter (more ♈ later!), the equatorial system slowly changes with time. –90º Equatorial coordinates +90º This means that the vernal equinox and the celestial equator move with respect to the distant background objects (galaxies, quasars). There we need to assign an epoch — a date — to any equatorial coordinate. ♈ (We’ll return to this shortly!) –90º The local equatorial system The local equatorial system is used to point polar-axis (or “equatorial”) mount telescopes These telescopes rotate around an axis parallel to the Earth’s rotation axis In the Northern Hemisphere, this means that the primary mount axis always points north The local equatorial system These telescopes track a star by rotations around only one axis Note that this means that the field of the image does not rotate, like it does for an alt-az telescope The local equatorial system In the local equatorial system, the hour angle HA replaces the right ascension: HA=LST–α Here LST is the local sidereal time (which we’ll define shortly!) So knowing the time of day (the LST) and the α,δ of an object, it’s very easy to locate your object with a polar-axis telescope. Note that the minus sign makes this a left-handed coordinate system! Julian Date Julian Date (JD) is an extremely useful way of keeping track of observations made over long time periods. Julian Dates are defined as the number of (Julian) days since noon on 1 January 4713 BCE (really!) Roughly now — 27 March 2014 at 08h 30m UT1 — is JD 2456743.854167 The beginning of each Julian day is defined to be at noon in Greenwich, 12h UT1 Julian Date The Modified Julian Date (MJD) is often used (because it’s shorter!): MJD=JD–2400000.5 Note that it starts at midnight in Greenwich rather than at noon (in fact, it started precisely at 00h 00m UT on Wednesday 17 November 1858!) Note also that J2000.0 is defined on the Julian day/year (=365.2500d exactly)/century (36525d) system and began at 12h (TDB) 1 January 2000 exactly, i.e., JD2451545.0 (TDB) Telescope plate scales, speeds and types y h′ θ z O D f nebula at ∞ Let’s replace our telescope with a thin lens with an equivalent focal length f Note that rays that pass through the center of a thin lens are undeviated, so that h0 = f tan ✓ f✓ ⇡ y h′ θ z O D f nebula at ∞ So the angle imaged onto a unit length on the focal plane is ✓ 1 S = = ! h0 f 206265 [00/rad] or S(00/mm) = f [mm] is the plate scale y h′ θ z O D f nebula at ∞ The Shane 3m telescope at Lick Observatory has f=15.3 m at prime focus, so S=14″/mm The full moon then images onto ≈13 cm at prime focus This is big for an optical system! Now let’s consider the size on the detector of our nebula with angular size θ For pixels (or photographic grains) of fixed size, smaller diameter h′ will put more energy into a single pixel: –2 –2 Epixel∝h′ ∝f Of course, a larger aperture (mirror) will collect more 2 photons, so Epixel∝D 2 Therefore, for diffuse (resolved) sources, Epixel∝(D/f) The ratio F=(f/D) is called the focal ratio We often speak of the f-ratio of an optical system, f/F The speed of an optical system is proportional to the energy deposited in a unit area (like a pixel), so –2 speed∝Epixel∝F , because an exposure to a given depth (the sensitivity) can be carried out more quickly in a faster system The speed of an optical system is only dependent on the focal ratio F and is independent of the diameter D A 1m telescope at f/6 is as fast as an 8m telescope at f/6, because the 1m telescope has a shorter focal length than the 8m telescope more light is concentrated into a smaller area at the focal plane Of course, sometimes we want to use a larger F — because it results in a higher magnification — and so big telescopes are useful after all! Telescope types and configurations One-mirror telescopes: A single mirror directs the light to a detector at the prime focus Usually the primary is not the appropriate shape for this configuration Need a prime focus corrector to eliminate aberrations Classical radio telescopes are often this design, but most modern (radio) telescopes don’t have a prime focus corrector Two-mirror telescopes Cassegrain, with a convex secondary A Gregorian telescope: Gregorian, with a concave secondary The INT The INT is a 2.54 m Cassegrain telescope typically used as a single- mirror telescope with an imager, the Wide-Field Camera (WFC), at the prime focus The INT has a prime-focus corrector that corrects the aberrations induced by the paraboloid primary The WFC The WFC is a mosaic of 4 thinned 2048x4096 CCDs with 13.5 µm pixels mounted at the prime focus of the INT The WFC Together with the inter- chip gaps, the mosaic is ~34’x34’ The WFC Together with the inter- chip gaps, the mosaic is ~34’x34’ The WFC Filters: six-position filter wheel with many, many filters available – see http://catserver.ing.iac.es/filter/list.php? instrument=WFC for a full list Quantum efficiency: ~60% at 380nm, ~80% at 400– 650 nm, ~15% at 950nm The INT: questions The primary mirror of the INT has a diameter of 2.54 m The corrected focal ratio of the prime focus corrector (behind which the WFC sits) is f/3.29 What’s the focal length at prime focus? What’s the plate scale of the INT at prime focus? The WFC has 13.5 µm/pixel; what is the pixel scale of the WFC? Signal-to-noise Our confidence in detecting an object or feature in astronomy is often expressed as the signal-to-noise ratio If we consider an astrophysical source from which we collect N photons in some amount of time from a detector that has its own noise — caused, say, by its own electronics — we can use the Poisson pdf to determine the total variance of the measurement If we assume that N≈⟨N⟩=µ, then the variance of the sources is also µ≈N If we call the detector’s own noise the “read noise” RN, its variance is (RN)2 Then the total variance is ! σ2 = N + (RN)2 and the “noise” — the error bar or standard deviation — is σ = N + (RN)2 p More generally, we also have a background of photons — from the night sky, contaminating sources, etc.