Astronomical Optics

Part 3: The Astronomical Image

Telescope Optical Attributes
Light Grasp
Relative & Numerical Aperture
Magnification Selection
Exit Pupil
Apparent Field of View
True Field of View

Richest Field Telescopes
Effect of Central Obstruction

Optics of the Eye
Individual Differences
Illuminance & Luminance
Luminance Adaptation
Light Grasp
Contrast Sensitivity
True Field of View
Focus Diopters

This page introduces the optical principles necessary to understand the design and performance of astronomical telescope systems — the telescope and eyepiece used as a visual instrument with the eye included as a third component. It is one of series: previous pages explained basic optics and telescope & eyepiece combined; subsequent pages discuss optical aberrations, eyepiece designs and evaluating eyepieces.

Included at the end of each page is a list of Further Reading that identifies the sources used and points to additional information available online.

Telescope Optical Attributes


Collimation (pronounced kôl-ǝ-mā-shŭn) is the condition in which the three optical axes of the objective, eyepiece and observing eye are exactly aligned and coincident. It is the first of two critical adjustments necessary to obtain the peak optical performance from a telescope. Miscollimation arises when any of the three optical axes is out of alignment with the other two (diagram, below).

Observing through a spherical symmetrical lens from an off axis position has the same optical effect as observing on axis through a distorted or asymmetrical lens. Miscollimation therefore produces aberrations in "on axis" or centered star images where optical images are usually flawless. These aberrations most often resemble coma or astigmatism. They are easiest to produce by moving your head from side to side while viewing a star centered in the field of a wide angle eyepiece or in a small focal ratio objective.

Collimation is the single most important adjustment that most observers will make in their equipment, yet many amateur telescopes are typically found to need adjustment when tested at star parties or dark sky events. Some of those owners paid thousands of dollars to get 1/10 wave optics, but miscollimation can reduce the image quality of a telescope to 1/2 wave or worse.

Adjustment knobs or screws are usually provided for tuning the collimation in the mirror cells of the primary and secondary mirrors of a Newtonian (Dobsonian) telescope, but only in the secondary mirror of a Schmidt Cassegrain or Maksutov Cassegrain. It is not normally possible to collimate the objective lenses of a refractor, or the corrector plate of a catadioptric (Schmidt Cassegrain, Maksutov Cassegrain, Maksutov Newtonian) telescope.

The specific steps necessary to collimate a telescope vary with the telescope design and manufacturer. Nearly all commercially available telescopes are provided with illustrated collimation instructions, and generic instructions are available online. Refer to these first, and memorize the procedure so that it becomes easy to do on your own.

Take collimation very seriously. Although most commercial telescopes will hold collimation very well once properly adjusted, always check the collimation at the same time you use a star image to evaluate the seeing, and especially if the telescope has been transported in a car and possibly shaken up on the road.

Focuser Collimation. Collimation in the focuser or drawtube that holds the eyepiece is usually not something the user can adust. In some high end focusers or Newtonian (Dobsonian) telescopes the focuser is secured to a mounting plate with four screws that can be adjusted to pitch or yaw the focuser axis in its fixed position on the side of the telescope tube until the focuser axis points back, via the secondary mirror, to the center of the primary mirror. In most telescopes, including nearly all Schmidt Cassegrains, Maksutov Cassegrains and refractors, the focuser is built into or fixed onto the tube assembly and can't be adjusted.

The focuser alignment can be checked with a Cheshire eyepiece (image, below), if the focuser has adjustment screws as part of the tube mounting plate. In cases of extreme misalignment, the mounting screw holes may need to be enlarged or redrilled and the focuser remounted on the telescope tube, but this is very unlikely in nearly all commercial instruments available today.

For visual astronomy the focuser clamping mechanism (which consists of a thumbscrew that exerts pressure on a flexible metal collar inside the focuser drawtube) should always be used with eyepieces, despite the annoyance of the eyepiece undercut catching on the clamping band. Clamping ensures alignment with the drawtube and minimize variation in the eyepiece orientation. The eyepiece should first be seated so that the shoulder is in contact with the rim of the drawtube, and only light clamping pressure is required — you should be able to rotate the eyepiece in place. To ensure the most even distribution of pressure, the short end of an allen wrench held by the long end can be used to rotate the flexible metal band in its groove so that the gap in the band is located directly opposite the thumbscrew.

Alignment issues are muddied when a star diagonal (with mirror or prism) is used. In a test of six mirror diagonals available to me I found errors of up to 30 arcminutes in the alignment of the mirror to the optical axis, and to my annoyance I discovered that the most expensive mirror diagonals were not the most accurate. (One was shimmed with a putty adhesive, the other with layers of electrician's tape.)

Secondary & Primary Collimation. This principal form of collimation is the procedure of aligning the objective lens or primary and secondary mirrors with the eyepiece. It consists of two separate adjustments: the physical position of the focuser and mirror or lens mounting to each other (mechanical collimation), and the alignment of the optical axes of the mounted elements to each other (optical collimation).

In Newtonian telescopes, the main aspect of mechanical collimation concerns the rotational and longitudinal alignment of the secondary mirror in relation to the focuser drawtube (eyepiece). This alignment must be confirmed before any other adjustments can be made. There is some debate as to whether an offset of the secondary mirror is necessary, with the majority opinion usually that it has a modest effect. The centering of the primary mirror within the tube assembly usually cannot be adjusted and also does not seem critical, especially as the optical axis of the mirror may not pass through its physical center. Nils Olof Carlin argues that some of the concerns with mechanical collimation are based on myths and misunderstandings.

There is a lot of lore on the optical collimation of Newtonian telescopes and several useful web pages that describe the process step by step, for example by Gary Seronik, Bryan Greer, Tim Trott and Nils Olof Carlin (here and here). Things are simpler with a commercial Schmidt Cassegrain format, since only the secondary mirror can be adjusted, for example as described by Starizona and Thierry Legault.

A number of collimation aids are commercially available to assist the collimation procedure. The oldest and simplest but also remarkably effective is the Cheshire eyepiece (image, right) which is an open metal tube supporting centered crosshairs and capped at the observer end with a pinhole that passes through a polished diagonal plate. The crosshairs are illuminated by a side opening onto this plate, which is turned toward the sky or an observatory light. The Cheshire is placed in the focuser and the crosshairs viewed through the pinhole; their intersection is designed to show the direction of the optical axis created by the focuser drawtube and secondary mirror. In a Newtonian format, the secondary mirror is aligned with the focuser when it is centered on the crosshairs; then the primary mirror is adjusted until the secondary mirror appears centered within it.

Several different laser collimation tools are available, and these are especially popular among the owners of large aperture, small relative aperture Dobsonian telescopes where accurate collimation is more difficult to achieve. Current vendors include (among many) Baader Planetarium, Farpoint, HoTech and Howie Glatter. These tools are effective with Newtonian format telescopes (including Maksutov Newtonians), not Cassegrain format telescopes. HoTech offers a system that is specifically designed for Schmidt Cassegrains, but it is expensive ($400 to $450 at this writing). I have tried both the HoTech and Glatter SCT gadgets, and found that these tools carefully used made my SCT collimation worse rather than better.

A related tool is the autocollimator, for example the Catseye Collimation devices by Jim Fly. These rely on multiple reflections of a central spot or triangular mirror marker to guide optical alignment. Nils Olof Carlin, Don Pensack and Vic Menard have written about the procedure and its rationale — again, with attention to Newtonian collimation.

Star Collimation. There is however a general method of visual collimation that is very effective and entirely adequate in any type of moderate to large focal ratio (ƒ/6 or more) telescope: star collimation. It uses the defocused and focused image of any fairly bright star as the collimation guide.

Star collimation has three important benefits. First, it requires no special tools or equipment. Second, because the aim of good optics is to form a perfect image, you are directly using the optical output in order to tune the optical quality. Third, facility with star collimation allows you immediately to assess the collimation of any telescope unfamiliar to you or owned by another observer before you use it for any observing task.

Star collimation requires you to complete mechanical collimation and approximately center the reflected image of the secondary mirror in the primary. It also assumes that you can only adjust the secondary mirror in a Cassegrain format. In other words, you should use daylight inspection and adjustments to eliminate mechanical alignment from the process so far as possible, so that only optical fine tuning needs to be done with the star image.

On your first try at collimation, I strongly recommend that you prepare a collimation guide similar to the one I created for my SCT (diagram, right). Draw a circle to represent the eyepiece field of view and mark the top of the field in the orientation that you will view it. (This must be a standard position for both your body and head in relation to the telescope and a fixed orientation of the star diagonal if you use one, as changing the position of either one will "rotate" the field of view. See the discussion here.) This circle also represents the back of your secondary mounting or primary mirror cell, so you want to indicate the direction and size of the adjustment effect in the field of view produced by each of the three primary or secondary adjustment knobs or screws. To determine this, first center a star image in a medium power eyepiece, then tighten or loosen only one of the three primary or secondary adjustment knobs or screws (replacing the screws with Bob's Knobs makes this much easier). Return to the eyepiece, and record the direction and size of the shift that results. Note that loosening or tightening the knob has opposite effects but along the same radial direction. Return the star to the center of the field of view and do the same for the other two knobs or screws, taking care to label each one (for example, as a clock face position) so that you can identify it on the back of the mirror mounting. Refer to this diagram each time you tune the collimation.

The overall procedure is to start with a low power eyepiece at the first pass, then increase your magnification as you refine the collimation until you finish with your highest practicable magnification. After each use of the adjustment screws, center the star in the field of view and examine the image to check your progress. When the image appears approximately symmetrical at the given magnification, center the star, increase the magnification, adjust the defocus, and continue.

The diagram (right, A to E) shows the variety of focused star images, and defocused star images or aperture silhouettes, that you will encounter as collimation proceeds. In all situations you want to make the adjustment that moves either the star image or the shadow of the secondary mirror, as seen in the field of view of your telescope, toward the wider or dimmer side of the star image, defocused star disk or aperture silhouette, as indicated by the dotted arrows in the diagram.

Because shifting the star image to the edge of the field has the same effect as a collimation adjustment, some astronomers recommend first moving the star image around in the eyepiece field until it appears most nearly symmetrical, then making the mirror adjustment that will move the image back to the center of the field. I prefer to start in the center because you must evaluate the adjustment with a recentered star image anyway, and if you make a wrong adjustment you will see it directly rather than lose the star beyond the field of view. Either way, this illustrates the importance of evaluating the collimation only when the star is centered in the eyepiece field of view. A reticule eyepiece can help, but so also can an eyepiece with an easily visible edge of field (usually an apparent field of view of 50° or so). With higher magnification, backing your eye away from the eyepiece along the eyepiece optical axis is a simple test: if you can still see the star through the exit pupil when you view from a foot or more away from and squarely above the ocular, then the star is well centered.

You start the collimation by centering in the field of view a moderately bright star (magnitude 2 to 3) at a moderate zenith angle (20° to 40° from the zenith). In this orientation, the adjustment screws in either the primary or secondary mirror cells will be easily accessible. Start with a low power eyepiece (that is, with a eyepiece focal length that is about 3 times your objective focal ratio or relative aperture, for an exit pupil of de = 3) and examine the focused star image. If the focused star image shows any coma (A), adjust the mirror to move the star in the direction of the coma "tail". Or defocus the star image until the secondary shadow shows clearly (B), then make the adjustment corresponding to the dotted arrow to improve the collimation.

Next increase the magnification to an exit pupil of roughly de = 1 to 2 and defocus the star until about 5 diffraction rings are visible inside the image disk. In refractors defocusing either in or out will produce an identical image. In SCTs, turning the focusing knob clockwise will shift the objective focal point away from your eye and move the eyepiece focus in the extrafocal direction. In this position the secondary shadow will appear more pronounced (C); this is generally more useful at lower magnifications. Turning the focusing knob counterclockwise will shift the objective focus toward your eye and move the eyepiece focus in the intrafocal direction. In this position the central Poisson spot will appear more clearly (D); this is generally more accurate at higher magnifications. (SCT owners using the mirror adjustment focuser should note the SCT focusing procedure described below.) Either way, adjust the collimation in the direction of the widest or faintest part of the defocused star image.

An important clue at this stage, rarely described in collimation guides, is a perfectly even appearance in the inner rings all around the defocused image. Even if the secondary shadow or Poisson spot appears centered, an off center collimation will cause the rings to appear compressed or widened on opposite sides, as illustrated in the two images (below). The side where the rings appear maximally spaced, or where a larger number of separate rings are visible, or the rings are least bright, is the "wide" side that indicates the direction of the collimation adjustment.

End the procedure using a magnification high enough to clearly reveal the Airy disk and first diffraction ring (E). Defocus the star only until the outer, brightest ring appears around the Poisson spot: the spot should be exactly centered and the ring should be equally bright all around. Then focus the star and examine the first diffraction ring around the Airy disk. Good seeing is required to get this far, but when the opportunity presents examine the star diffraction artifact to ensure the diffraction ring is equally bright on all sides. If not, make the adjustment as shown in the diagram.

If the seeing is less than ideal, examine the focused image of a first magnitude star. Tiny threadlike granules will be visible in the flare or glare around the star image. These prismatic granules should appear to radiate away from the star image equally on all sides. Note that they will be quite sensitive to the location of your eye pupil, which must be exactly centered on the exit pupil. (This is easiest to judge when the eyepiece field edge is clearly visible on all sides, so again an eyepiece with a ~50° AFOV is helpful.) If the eye pupil is centered, then the shimmering filaments can be used to judge the last refinements in the collimation.

Most primary mirror cells are equipped with lock screws, which should be gently and gradually tightened in a repeated circular sequence until all lock screws are firmly seated (only firm pressure between lock screw and mirror cell is necessary to hold the collimation). Tightening one screw all the way before moving to the next, or tightening all the screws too far can disrupt the collimation.

In SCTs, there is usually a "mirror lock" knob that will fix the mirror in place. If you have an external focuser, it is best to perform collimation with an eyepiece that is in the middle of the focusing range required by all the eyepieces you routinely use, with your external focuser in the middle of its total travel span in and out of focus. Collimate with this eyepiece, focus the star, lock your mirror, then use the external focuser for routine focusing.

Eye Collimation. Correct positioning of the eye pupil around the exit pupil is a third source of collimation error. The variations in this error are shown in the diagram (below) as a side view of the relative position of the eyepiece (lens schematic), the eye, and the emerging pencils of light passing through the exit pupil (orange bar). Also shown is a schematic of the visual appearance of each aberration, with the always in focus edge of the eyepiece field stop indicated by dark gray, and the always defocused edge of the exit pupil by a lighter gray.

In the correct alignment of the eye (diagram, top left) two conditions apply: (1) the observer's iris is placed so that it is approximately in the same plane as the exit pupil and the eye point or center of the exit pupil is at the principal point of the observer's cornea and lens; and (2) the optical axis of the system passes through the fovea of the retina. In this position the focused edge of the field stop in a standard eyepiece is visible around the entire circumference of the image area, and any optical aberrations (astigmatism and apparent field curvature in particular) are radially balanced — of equal size and appearance around the edges of the field.

Even in this position, it is possible for the observer by very small movements of the head to create skew in the spicules of scatter around bright stars or (in short focal ratio telescopes) comatic aberrations in stars at the center of the image area. These subtle cues can aid the observer to judge the eye collimation.

In longitudinal misalignment of the eye (diagram, bottom left), the observer's pupil and fovea are centered on the eyepiece optical axis but the eye pupil point is located in front of or (more often) behind the exit pupil. In either position, the exit pupil vignettes the image area and creates a "keyhole" narrowing of the field. If the target is a bright object such as the moon, a sufficient distance clearly shows the outline of the exit pupil in the center of the eyepiece eye lens; it will appear out of focus if it is closer than about 25 cm to the eye. This misalignment is usually remedied with a flexible eye cup or an eyepiece with an adjustable eye rest.

In lateral misalignment (diagram, bottom right), the optical axis of the eyepiece and the eye are approximately parallel and the eye pupil or eye principal point are in the same plane as the exit pupil, but both the eye pupil and the fovea are displaced to one side of the optical axis; this causes a curved shadow or "shutter" (the edge of the exit pupil) to intrude into the field on the side of the eye's displacement. This can be distinguished from the edge of the eyepiece field stop because it is out of focus.

Finally, in axial misalignment (diagram, top right) the eye's optical axis is at an angle to the eyepiece optical axis and the fovea is far to one side of the of a wide apparent field, which usually requires the observer to shift the eye laterally (to one side) and longitudinally (slightly forward or backward) from the eye point. This does not produce a noticeable effect in traditional (≤55° apparent field of view) or well corrected wide field eyepieces, since the necessary angular misalignment and displacement of the eye is relatively small, although if the eye is displaced far enough the field stop edge will be eclipsed behind the edge of the exit pupil. However, in some wide field or superwide field eyepieces with spherical aberration of the exit pupil, a floating lenticular or "kidney bean" shadow appears in the field on that side, which is often outlined by the smeared out light of stars. This effect is quite annoying for some observers.

While it is a matter of physical control to find and hold the correct eye position while observing, various supports can make this task enormously easier. In particular, an adjustable observing chair is very convenient for observing with most refractors and Cassegrain format reflectors; Dobsonian and Newtonian reflectors should be used with sturdy, stable ladders or stepladders, preferably with a raised front rail that allows the observer to steady their upright posture at the eyepiece. No matter what kind of telescope is used, positions that require bending or leaning, or strain in the neck, back or legs should always be avoided as much as possible.


Focus is the second of two critical adjustments necessary to obtain the peak optical performance from a telescope. By design both the objective and eyepiece have a specific focal surface; focus is the distance adjustment along the optical axis that brings the two into coincident position.

The objective focal surface is the point of reference:

• The area on the objective side of the objective focal point is called intrafocal or "inside the focus". Moving the eyepiece out of focus in this direction causes the output rays to converge, corresponding to an increasing positive diopter setting that is appropriate for eyes that undercorrect (are farsighted).

• The area on the eyepiece side of the objective focal point is called extrafocal or "outside the focus". Moving the eyepiece out of focus in this direction causes the output rays to diverge, corresponding to an increasing negative diopter setting that is appropriate for eyes that overcorrect (are nearsighted).

Focusing Mechanisms. In a typical rack and pinion (gear and linear toothed track) focuser or Crayford (roller and cylinder) focuser, the eyepiece is adjusted by a pair of knobs on a transverse shaft below the optical axis. The rotation of the knobs necessary to move the eyepiece in or out is immediately obvious from the visible movement of the eyepiece. Rotating the right hand knob clockwise moves the eyepiece toward the objective on the optical axis, in the intrafocal direction; counterclockwise rotation moves the eyepiece away from the objective.

In a commercial Schmidt Cassegrain system, in which the focus is adjusted by moving the hidden primary mirror using a mirror cell focusing knob protruding from the back of the telescope, the necessary adjustment is less obvious (diagram, right):

• Turning the mirror cell focusing knob clockwise moves the primary mirror away from the secondary mirror, decreasing the focal length and moving the focal surface away from the observer. An eyepiece in focus is moved out of focus in the extrafocal direction, and an eyepiece out of focus on the intrafocal side is brought into focus. This is the focusing procedure that brings closer objects into focus.

• Turning the mirror cell focusing knob counterclockwise moves the primary mirror toward the secondary mirror, increasing the focal length and moving the focal surface toward the observer. An eyepiece in focus is moved out of focus in the intrafocal direction, and an eyepiece out of focus on the extrafocal side is brought into focus. This is the focusing procedure that brings more distant objects into focus, including objects at infinity (stars).

Focusing Procedure. The optimal focusing procedure is always to start focus from an extrafocal position — the eyepiece front focal surface must be in front of the objective focal point (diagram right, bottom). The reason has to do with how the eye focuses the diverging light rays from nearby objects. In the relaxed state used to view distant objects the eye lens is actually stretched into a flattened figure by zonule fibers or ligaments attached to the circumference of the lens. To focus on nearby objects, the iris muscle contracts in a way that relaxes the tension on these ligaments and allows the eye lens to bulge into a shorter focal ratio. Age hardens the lens, reducing and (around age 60) eliminating the elastic ability of the lens to respond to these changes in tension.

When an image is brought into focus from an intrafocal position, it approaches focus with a "nearsighted" convergence of rays. The eye handles these as it would an object near the eyes, and tightens the ocular muscles in an attempt to focus; some of this tension remains even after the image is brought into focus. In contrast, an image brought into focus from the extrafocal direction approaches focus with a "farsighted" divergence of rays. The eye cannot accommodate toward the image any farther than a completely relaxed state, and this remains once the image is brought into focus. If you overshoot the correct focus and err onto the intrafocal side, return to the extrafocal position and start again.

In commercial Schmidt Cassegrain telescopes, changes in the primary mirror alignment ("mirror shift" or "mirror flop") occur when focusing with the mirror adjustment knob. The recommended focusing procedure is first to turn the focusing knob clockwise to bring the eyepiece into an extrafocal defocus, then to bring the eyepiece back into focus by a counterclockwise adjustment in the intrafocal direction. This should be the same focusing procedure used during collimation (except that the star is left somewhat out of focus on the extrafocal side), so that the mirror tilt during collimation and in routine observation is the same.

The ability to produce a sharp focus is partly dependent on the optical quality and collimation of the instrument and partly dependent on the amount of atmospheric turbulence. Nevertheless, the "best" focus will always produce the visually smallest or most compact image of a star, and the sharpest edge in a planetary disk or lunar terminator, when the target is placed at the center of the eyepiece field of view, and will make visible the greatest number of faint stars in the field. Under optimal adjustment and viewing conditions the image will appear to "snap" unambiguously into the best focus — although the "snap" is most vivid in small focal ratio telescopes, and atmospheric turbulence can make this test unreliable.

When poor seeing makes focusing difficult, best focus is found by bracketing the focus. First bring the star into intrafocal defocus to a recognizable disk visual diameter, then take the star in the opposite direction, through the focal position into extrafocal defocus, until the disk reaches the same visual diameter. To finish, rotate the focusing knob back by half the travel between the two defocus positions.

Poor seeing can cause large variations in focus. Resist the temptation to continually adjust focus in order to "chase the seeing". The best procedure is not to make continual adjustments, but to find the focus that produces the most frequent intervals of focus. Patience, rather than manipulation, is the bast focusing procedure: focus should be changed by very small increments so that the quality of the new focus can be observed for several seconds.

High power focusing is assisted by the use of a moderately bright star, as the diffraction pattern (or turbulence speckles, under poor seeing) will appear compact and clear at best focus.

Depth of focus. This is the distance in front of and behind the optimal focal point within which image quality is still acceptably sharp. If the standard of image quality is that it is not degraded by more than 1/4 wavelength of light (the Rayleigh criterion), then:

DF = ±2λNo2

Depth of focus is proportional to the square of the relative aperture. Thus an ƒ/20 Cassegrain telescope has a depth of focus of ±0.44 mm at λ = 550 nm, but an ƒ/4 Newtonian has only ±0.02 mm. These tolerances are especially critical for astrophotographers, who typically use automated (software driven) focusers to find the optimal focal point at the camera image plane.

Light Grasp

(1) Magnitude vs. Brightness. The magnitude system was proposed in 1856 by Norman Pogson, who noticed that a five magnitude difference in the ancient system of naked eye magnitude ranks (attributed to Hipparchus and adopted by Ptolemy) corresponded very nearly to a 100 times difference in brightness. Therefore, the base value of the magnitude system was chosen as 1001/5 = 2.51189, so that 2.5125 = 100.023

Originally, Vega anchored the measurement of illuminance ratios as a 0.0 magnitude star and the numerator of the brightness ratio, making all magnitudes a log reciprocal illuminance (because 2.5120 = 1). A star at the naked eye limit magnitude of 6.5 then emits into the eye or a telescopic aperture only 1/2.5126.5 = 1/398 or 0.25% the light of Vega. The magnitude system can be extended to objects brighter than 0.0 by using negative magnitude values. The Sun, with an apparent magnitude of –26.74, is 1/2.512–26.74 or 4.97·1010 (almost 50 billion) times brighter than Vega.

The flux of Vega in the visual is 3.46 x 10-11 watts/m2 at the peak visual wavelength of 555 nm (Mégessier, 1995), and 2.46 x 10-6 lux on the Earth's surface — the naked eye brightness of a single candle at 640 meters distance. The chart (left) shows the visual brightness of illustrative planets and stars viewed with telescopes of aperture from 100 mm to 400 mm (white lines) at a unit exit pupil, and for the exit pupil of peak stellar brightness (de =3.5) in the 400 mm aperture (dotted white line). As reference the range of the naked eye (orange line) is also shown. The dashed blue line represents the point at which foveal resolution and color perception begins to degrade, and is consistent with the observation that naked eye star color becomes imperceptible below v.mag. ~2.0. Note that telescopic star images can be detected at much lower illuminances than star images observed with the naked eye; this is due to the decreased sky brightness and increased star diffraction area in the telescopic view at high magnification (de = 1.0; see below).

Magnitude is a exponential metric in which the magnitude units are actually the exponents applied to a constant base value that represents a fundamental proportion or ratio. Similar exponential metrics for the energy of a stimulus include the Richter Scale (earthquakes) and the decibel scale (sound).

The base value 2.51189 is the unit illuminance ratio (the "Pogson ratio"). If two stars differ in brightness by 1 magnitude, then one star emits roughly 2.5 times more light than the other. The numerical magnitude is the exponent of this unit, and by using fractional exponents (continuous magnitude values), the base value can be multiplied to express any brightness ratio. (The bolometric magnitude is the same ratio in the flux of a celestial source across the entire electromagnetic spectrum.)

Some references point out that the human visual response to luminance is actually a power function (the "Stevens power law"), not an exponent function. Perceptual or visual brightness J scales as:

[1] J = k·(II0)α

where I is the the stimulus intensity or radiant flux, I0 is the radiant flux of the visual threshold or limit magnitude, k is a scaling constant, and the exponent is usually the cube root α = 0.333 for point light sources viewed with a dark adapted eye. But that has a limited significance, because the log and power functions are fundamentally related, as:

[2] log(J) = α·log(II0)+log(k)

and the two functions coincide when plotted on log radiant flux across the visual range (magnitudes –1.5 to 7; diagram, right) with an exponent α = 2.5 (cf. formula [4], below). It's the difference between the visual and magnitude exponents, and the combination of two separate power functions for foveal and peripheral (20° off fovea) light sensitivity, that creates the discrepancy. As a result, a star midway in visual brightness between a magnitude of 0.0 and 1.0 is actually magnitude 0.46, not 0.50; a star midway between magnitudes 0.0 and 2.0 is magnitude 0.85, not 1.0; a star midway in brightness between magnitudes 0.0 and 5.0 is magnitude 1.65 (!), not 2.50. In addition (assuming a naked eye limit magnitude of 7), a strong threshold effect appears at around magnitude 5 where the magnitude and brightness curves diverge. Near the visual threshold, small changes in magnitude correspond to increasingly larger changes in visual brightness.

The key points for the visual interpretation of the magnitude scale are that (1) "brightness" in the magnitude scale is strictly a metric of log illuminance or log radiant flux (I), not of perceived brightness (J); (2) the numerical midpoint between two values of the magnitude scale identifies a star that is fainter than the perceptual midpoint between the corresponding two visual brightnesses, and this discrepancy increases across larger magnitude differences; and (3) magnitude differences indicate much larger brightness differences in parafoveal or averted vision — about 2 magnitudes above the visual or aperture limit magnitude.

Since magnitude is a measure of radiant flux, not perceived brightness, the magnitude system can be applied to flux across any part of the electromagnetic spectrum, including bandwidths that cannot be perceived by the eye. Most astronomical spectral radiances are measured within defined sections of the spectrum. The historically most important for visual astronomy were visual magnitude, measured primarily on the flux between 504 to 592 nm ("green" to "yellow orange" wavelengths of light), and photographic magnitude, centered at around 425 nm ("blue violet") and originally measured as the diameter of the star image in UV sensitive photographic emulsions. Color indices such as the B–V color index are simply the flux of an object measured with a V (visual) filter subtracted from the flux measured through a B (blue) filter. In the Johnson-Morgan system of filters, the benchmark illuminance was defined so that a "white" type A0 star had 0.0 magnitude in all filters: again, the historical standard was Vega.

(2) Geometry of Magnitude. The logarithmic basis of the magnitude scale means that any equal magnitude difference defines the same flux ratio. Thus, a 5th magnitude star produces only 1/100 of the illuminance of a 0 magnitude star such as Vega, and the same illuminance ratio applies between stars of magnitudes –5 and 0 or 5 and 10.

The fundamental relationships among magnitude intervals, relative flux ratios, the aperture diameter required to grasp equal illuminance from fainter magnitudes, and the parsec distance of a star of 5.0 absolute magnitude (close to the Sun's 4.83 absolute magnitude) are listed in the table below down across the stellar magnitude range of a 150 mm telescope.

basic ratios of the magnitude system







*Assuming a near solar absolute magnitude = 5.0.


illuminance ratios of stellar magnitudes –1 to 7, roughly the naked eye range, illustrated as differences in circular area

The values for brightness ratio and proportion define the magnitude system. However, several physical insights can be derived from comparisons made with the aperture diameters and physical distances.

The aperture diameters represent multiples of aperture in any measurement unit (e.g., centimeters or inches) that produce the tabulated differences in magnitude; these can be used to equate magnitudes across aperture:

(1) The dark adapted eye pupil has an aperture diameter of approximately 0.6 cm, and the nominal naked eye limit magnitude is 6.5. A 60 cm telescope is 100 times larger, and this aperture ratio (100:1) is equated in the table with a difference of 10 magnitudes. So the limit magnitude of a 60 cm telescope will be 10+6.5, or 16.5.

(2) A 12th magnitude star is observed through a 10 inch telescope, which is is roughly 40 times the 0.25" diameter of dark adapted eye. In the table an aperture ratio of 40 is associated with a magnitude difference of 8. This implies the 12th magnitude star will appear about as bright as a 4th magnitude star observed with the naked eye. (Light absorption and scattering in the telescope, and the power function response of the eye, make this only an approximate calculation.)

(3) If we equate the dark adapted pupil aperture of 0.6 cm with 0 magnitude, then we find it collects as much light from Vega as does a 100 cm aperture from a 11th magnitude star.

(4) The owner of a 4 inch telescope wants to know the gain in the limit magnitude obtained by switching to a 16 inch telescope, an aperture four times as large. In the table, an aperture ratio of 4.0 yields an increase of 3 magnitudes. The owner of a 6" telescope would like to increase his limit magnitude by 2 magnitudes. The aperture ratio for two magnitudes in the table is 2.5. Therefore he would need a 6"·2.5 = 15" aperture. (See also the limit magnitude graph, below.)

The distance metrics are governed by the inverse square law, which states that moving a light source N times farther away will reduce its brightness by 1/N2.

(1) If a star is at a 1 parsec distance, then moving it to 10 parsecs will reduce its flux by 1/102 = 1/100th; consulting the column of flux ratios indicates that this will increase its magnitude by 5.

(2) Given two stars of equal absolute magnitude, one 7 times farther than the other, the farther star will be 1/72 = ~1/50 as bright. In the table the flux ratio 1/50 falls between magnitudes 4 and 5 but is much closer to 4, so we can estimate that the farther star will appear to be about 4.2 magnitudes fainter. (Exact calculation with formula [2] below yields a magnitude difference of 4.27.)

(3) The Sun has an absolute magnitude of ~4.8 (at 10 parsecs), and a 14 cm refractor has a limit magnitude of 13.3, a difference 8.5 magnitudes. This corresponds to a distance ratio of about 50, which means stars as bright as the Sun cannot be detected in the 14 cm refractor if they are more than 50·10 = 500 parsecs away. The absolute magnitude of the Sun is 6.5-4.8 = 1.7 magnitudes higher the naked eye limit, which is between the distance ratios 1.58 and 2.51; assuming that the corresponding distance ratio is about 2.2, the Sun would not be visible to the naked eye if it were 10·2.2 or more than 22 parsecs distant. (Exact calculation with the distance modulus yields 501 and 21.9 parsecs, respectively.)

(4) If Vega (7.7 parsecs distant) were 25 times farther away (at 192 parsecs) then it would appear to be a 7th magnitude star.

These examples illustrate the many different ways that the relationships between magnitude, flux, aperture and distance can be exploited to find one value from the others, and how the brightness ratios can be used directly with a minimum of calculation.

(3) Magnitude and Illuminance. There are several equivalent standard formulas to calculate the magnitude difference between two stars from the ratio of their radiant fluxes. Given a magnitude difference m'm (where m is the magnitude of the brighter star, m' > m), to find the illuminance ratio b/b' (where b is the illuminance of the brighter star, b > b'):

[3a] b/b' = 2.512(m'm), or

[3b] b/b' = 100.4(m'm), or

[3c] log(b/b') = 0.4(m'm)

Given an illuminance ratio, to find the equivalent magnitude difference:

[4] m'm = 2.5·log(b/b')

Note that 0.4 is the inverse of 2.5. Note also that illuminances are defined a ratio and magnitudes are defined as a difference because division is done in an exponential scale (the magnitude scale) by subtraction of exponents (the magnitudes).

The illuminance ratio across all naked eye stars (from Sirius at –1.46 to the 6.5 magnitude limit) is then:

b/b' = 2.5126.5––1.46 = 2.5127.96 = 1528

However averted vision is required to achieve the low end of this dynamic range. If we put the foveal threshold (stars visible with direct fixation) at around magnitude 4.5 under conditions of typical sky brightness, then the dynamic range of dark adapted foveal vision is at minimum 2.512(4.5––1.46) = 2.5125.96 = 242 (see the table, above). This two magnitude difference suggests that averted vision is about 6.3 times more sensitive to light than foveal vision.

In double stars imaged as a single light source, the magnitude of the combined source is not the sum of the separate magnitudes. Instead, the incremental brightness (x) is calculated from m'm, and this is subtracted from the magnitude of the brighter star (m):

[5] x = log(1+(1/2.512m'–m))/0.4, then M = mx

Plugging in some basic values, we find that the combined magnitude of two stars of nearly equal magnitude (m'm = ~0) increases by log(1+1/1)/0.4 = 0.75 magnitudes, and this is subtracted from the magnitude of the brighter star. The magnitude increase is x = log(1+(1/2.512))/0.4 = 0.36 magnitudes when the magnitude difference is 1.0, x = 0.16 when the magnitude difference is 2.0, and x = 0.07 when the difference is 3.0. (The human eye is usually insensitive to magnitude changes less than ~0.2.)

To anchor these relationships as measures of illuminance or brightness, the light gathering power of the telescope can be expressed in three ways: (1) as the increase in light gathered over the light that enters an unaided, dark adapted eye; (2) as the limit magnitude of stars visible with averted vision, both visual parameters; and (3) as the illuminance of the telescopic image on a film or CCD sensor, a photographic parameter.

(4) Aperture & Light Grasp. The light gathering power of the telescope is a function of its aperture area (Ao): as the aperture increases, more light enters the system and is concentrated into the usable image. For a refracting telescope this area is usually calculated as:

Ao = Do2

Because the secondary mirror in a reflecting telescope blocks a certain amount of light from reaching the primary mirror, its area (Ds2) must be subtracted from the area of the objective in order to get the effective aperture (Deff):

[6] Deff = sqrt(Do2Ds2)

Thus, a 250 mm (10") telescope with a 50 mm (2") diagonal diameter has an effective aperture of 9.6" or 245 mm.

The relative light grasp of any two apertures is proportional to the square of their areas, and the increase in light grasp with aperture scales as the square of the aperture ratio:

[7] ΔAo = (D1/D2)2

Thus, compared to a 250 mm (10") objective, a 400 mm (16") objective gathers (400/250)2 = 1.62 = 2.56 times more light.

(5) Telescopic Gain Over Naked Eye Brightness. The first criterion requires an estimate of the dark adapted pupil aperture. To measure: obtain a complete Allen wrench set. At night, once your eyes are fully dark adapted (roughly 1/2 hour after the last bright light exposure), close one eye and pass an allen wrench back and forth in front of the other eye. Hold the long shaft of the wrench vertically with the short end pointed in the direction you are looking. If the wrench fully occludes a bright star, causing it to "blink", then the wrench is larger than your pupil opening. Try different wrenches until you discover the largest wrench that dims but does not completely block the star image. Then the width of that wrench is the diameter of your dark adapted pupil (δ).

Then the measure of a telescope's light grasp in comparison to your naked eye is the ratio of the two areas:

[8] Lo = (Deff/δ)2t

where t is an estimate of the system's light transmission, which is 0.98 for each surface of aluminum high transmission catoptrics, 0.90 for aluminum standard catoptrics, and approximately 2% per centimeter of fully multicoated refractor or eyepiece glass. My dark adapted pupil is about 6 mm wide, so a 10" reflecting telescope with a 2" diagonal will gather at most (245/6)2 x 0.98 x 0.98 = ~1600 times more light than my dark adapted eye, not accounting for light absorption by the eyepiece.

Magnification is a critical qualification in estimates of the visual brightness of a telescopic image: the objective gathers light, but image magnification spreads a fixed density of light over a square visual area. Thus the ratio of telescopic to visual brightness becomes:

[9] Lo = (Deff2/Mt2δ2)t

(6) Telescope Limit Magnitude. Because visual magnitude differences are a direct measure of luminance ratios (expressed in logarithms), the telescope/eye ratio implies a 10" telescope limit magnitude that is 1656 times fainter than the visual limit magnitude.

Using (as a first approximation) the standard brightness ratio formula [3c] (where b' and m' refer to the fainter brightness and magnitude, respectively), assuming a visual limit magnitude of 6.5, and rearranging to solve for m' yields:

m'250 = log(2452/62)/0.4 + 6.5 = 14.56

for a Deff = 245 mm (10") telescope limit magnitude and a dark adapted pupil aperture of 6 mm. This standard formula is more often cited in the following equivalent form:

[10] m'o = mne – 5log(δ) + 5log(Deff)

This form shows that the three critical variables are effective aperture (Deff) and pupil aperture (δ), both in millimeters, and naked eye magnitude limit mne (a function of zenith angle, atmospheric seeing, transparency, diffusion and light pollution, which are not measured directly). Assuming a pupil aperture δ = 6 mm and a "dark sky" naked eye limit magnitude mne = 6.5, the formula simplifies to:

[11] m'o = 2.61 + 5log(Deff)

and for a Deff = 245 mm (10") telescope this yields a value m'o = 14.56 identical to the value calculated using formula [3c]. (To calculate using different values for pupil aperture or naked eye limit magnitude, use formula [10] to derive the constant to replace 2.61.)

Attempts have been made to refine this straightforward calculation by including more variables — magnification, zenith angle, observer experience, optical quirks of the eye — under the statistically dubious assumption that a greater number of inputs will increase accuracy.

The graph (right) shows estimated limit magnitude, as calculated by Bradley Schaefer (1990), across 6" to 16" apertures and 5 levels of magnification, under the condition that the faintest stars are viewed with averted vision. The predictions assume δ = 6.7 mm, a naked eye limit magnitude mne = 6.0, a telescope transmission of 80% and atmospheric transmission of 70%. The predicted values correspond reasonably well to the actual limit magnitudes in 314 observing reports from over four dozen observers. The main effects shown in the graph are that limit magnitude increases with aperture and with increasing magnification (from 30x to 400x), but that the incremental increases decline significantly as aperture or magnification increases.

Schaefer claims to include the effects of human dark adapted visual sensitivity, atmospheric absorption, observer age corrected pupil size, telescope transmission, magnification, scotopic eye transmission, star and sky color temperatures and observer experience. But the largest additional impact is from the sky brightness (naked eye limit magnitude, NELM). The blue curves assume a visual limit magnitude of 6.0; the single green curve shows the effect of reducing the visual limit to 4.0 at a magnification of 100x. Note that this increase in the limit magnitude of a 10" objective between a suburban and dark sky site (a, Δm = 1.25) is roughly twice as large as the increase between a 10" and 16" objective (b, Δm = 0.66)! This illustrates the significant benefit of a medium aperture but compact and portable field telescope for observers with access to dark sky sites. Note too that increasing magnification from 100x to 400x increases the limit magnitude of a 250 mm aperture to slightly more than the limit magnitude of a 400 mm aperture at 100x (b). This illustrates the large impact of magnification on limit magnitude.

Schaefer's complex model predicts m'o = 14.15 for a 10" aperture at 60x magnification. One can ask whether the difference of 0.09 magnitude between this model and the estimate of 14.06 given by the simple formula [10] (using the same values for δ, mne and Do) is worth the added complexity. Subsequent and more intricate proposals by Nils Olof Carlin and Chris Lord do not improve the validity of the methods by using different variables or assumptions, because the error in the actual observed magnitude across observers, due to factors omitted or impractical to measure, remains on the order of ±0.5 magnitude — the difference between a 6" and 8" telescope.

Schaefer's model underestimates limit magnitudes with a ±0.75 average magnitude error, and he comments that "the limit of accuracy for my model is much larger than I would have hoped or expected." That is because his expectations were unrealistic. Limit magnitude is difficult to predict because many stellar, atmospheric, instrumental and observer variables can affect it, most of them inconvenient to measure with the accuracy necessary for reliable prediction. Worse, estimates of visual acuity are an exercise in predictive psychophysics, which has always proven difficult, even with carefully controlled stimuli in vision laboratory experiments, due to the very large differences in visual capabilities across individual observers.

My suggestion is to use formula [10] to obtain a baseline estimate, then revise this value with the judgment of experience to correspond to changes in magnification or observing conditions. Schaefer actually demonstrates this approach at the end of his research paper, where he makes several ad hoc adjustments to his model in order to "predict" Steven O'Meara's remarkable visual acuity.

(6) Photographic illuminance. The photographic illuminance of the telescopic image increases with aperture and decreases with relative aperture, making it proportional to:

[12] Io = (Do/No)2

Eyepiece magnification is not a factor, because the sensor is utilizing the prime focus of the telescope.

Relative & Numerical Aperture

The relative aperture (No) of a telescope is the focal length of the objective (ƒo) divided by the diameter of its entrance pupil or clear aperture (Do):

[1] No = ƒo/Do .

This ratio of projection distance to aperture is frequently termed the f-number, what photographers refer to as the focal ratio. A larger relative aperture is photographically slower, either because the focal length is longer or the aperture has been made smaller.

This is distinct from the numerical aperture (NA), which is measured as the field angle of the field stop (θ) at the system focal point:

[2] NA = no sin(θ) .

Numerical aperture is simply Snell's law applied to the last refracting surface of the optical system. Relative and numerical aperture coincide when viewing objects at optical infinity in optics surrounded by air: then N = 1/(2no sin(θ)).

Relative aperture is a scale free dimension, meaning the numerical aperture of a telescope does not tell us how large it is. In contrast, numerical aperture is related to the maximum field radius (ρ) produced by a telescope objective:

[3] ρ = 2arctan(1/2No) radians

which is equivalent to the angle subtended by the aperture as viewed from the focal point.

Relative aperture is sometimes referenced in the context of optical aberrations. Thus, coma is considered to be a significant aberration in Newtonian reflectors shorter than ƒ/5, and the spherical aberration in a spherical mirror or lens becomes negligible (by the Rayleigh criterion) when

[4] N ≥ 3.4(Dcm)1/3 or ≥ 4.6(Din)1/3

Both depth of focus (defocus tolerance) and depth of field increase with relative aperture.

In photometric contexts, relative aperture expresses the proportion of unit illumination (aperture area) per unit magnification (focal length) produced by the telescope objective. However, these units are dimensionless, and comparatively useful only by holding one element constant. For example, as the relative aperture gets larger for a constant aperture, the quantity of light projected onto a unit area of the image plane decreases because the image is magnified. Equivalently, as the relative aperture gets larger for a constant focal length, the quantity of light projected into the unit image area decreases because a smaller aperture admits less light.

Relative aperture determines the ratio between eyepiece focal length and exit pupil diameter, independent of aperture or focal length. Thus, an ƒ/10 telescope will produce an exit pupil that is 1/10th the focal length of the eyepiece, regardless of the dimensions of the objective or the image magnification.

Because magnification is proportional to focal length, but light grasp is proportional to the square of aperture, the relative image illuminance produced by different relative apertures (different focal lengths for the same aperture, or different apertures for the same focal length) scales as the ratio of the relative apertures squared (chart, above).

The geometry of relative aperture is luminance (Lo), calculated as the product of image area and aperture solid angle, or as aperture solid angle when image area and luminous flux are held constant. Then the formula for image luminance reduces to:

[5] Lo = E/N2

where E is the illuminance at the aperture in lux. This is just the inverse squared distance effect of decreasing illuminance but measuring distance in aperture units. The table below shows the aperture angular diameter for each relative aperture, with its area in steradians and illuminance in proportion to the luminance at ƒ/2. The diagram (right) illustrates relative aperture as distance from a small circular aperture in a dark room: as the image surface is placed farther from the aperture, the light falling on its surface decreases. Note that ƒ/2 is near the lowest feasible aperture, ƒ/4 is useful for richest field telescopes, and ƒ/8 is the lower value for traditional Newtonian reflectors, achromat refractors, and Schmidt Cassegrains. (The case ƒ/250 is relevant to the discussion of exit pupil.)

No2βaperture area
relative illuminance



greater distance from a small circular window in a darkened room
reduces the illumination on a surface in the same way that
increasing relative aperture reduces image brightness

Thus an ƒ/2 system has roughly six times the image illuminance of an ƒ/5 system ((5/2)2 = 2.52 = 6.25), an ƒ/7 system has half the illuminance, an ƒ/10 system one fourth, and so on.

The standard system of photographic ƒ stops doubles the area of the diaphragm stop (effective aperture) at each step, which means the ƒ stops increase the image illuminance by powers of the square root of 2 (1.414): 1, 1.414, 2, 2.83, 4, 5.65, 8, 11.3 and 16.


The resolution of an optical system is measured as the smallest linear width that the system can image at the focal surface, which is equivalent to the smallest angular width that can be observed in the image space. In astronomical systems this limit is determined by three things: the quality of the optics of the system, the turbulence or distortion introduced by the atmosphere, and the quantum wave properties of light. Only the last property can be calculated directly as the reciprocal of aperture, 1/Do.

The Diffraction Artifact. Calculation of the aperture resolution limit depends on a quantum (physical) rather than geometrical analysis of light.

If the light wavefronts from a distant star are traveling in direction x parallel to the optical axis of a telescope, then the width of each wavefront is measured along a dimension y perpendicular to the direction of light. Because the star is so far away, the scale of the y dimension is effectively too large to measure.

The quantum indeterminacy of the location and momentum of any photon in the wavefront is distributed across the entire dimension y, and both the energy and location of the photon can have a precise value, which means it can carry information about the emitting source of the photons.

In a telescope, the wavefront location for each photon is fixed at the focal surface of the objective. But because the wavefront has passed through the aperture stop, the indeterminacy distributed across the entire y dimension is reduced to the aperture width D. This introduces a quantum uncertainty or "error" in the specific location of each photon in the image, which must be no smaller than:

U = hν

where h is Planck's constant, 6.626 x 10-34, and ν is the frequency (energy) of the light, and where the speed of light c is their product, c = νλ. Casting these quantum relationships in a form that defines the indeterminacy of a circular aperture of diameter D yields:

[1] UD = hc / hνD = hνλ / hνD = λ / D

If we specify the energy of light by wavelength rather than frequency, for example as the commonly used λ = 550 nm (approximately the eye's peak photopic sensitivity), then the angular width of the indeterminacy in a 10" (254 mm) telescope becomes:

UD = 0.00055 / 254 = 2.165 x 10-6 radians

In the light from a far distant star, the result of the cumulative uncertainty from all photons focused to a single image point does not have a simple "fuzzy" image structure around the point, due to the destruction or reinforcement of light waves as different parts of the wavefront are superimposed at the same location on the image plane. Instead it appears as a central concentration of light surrounded by one or more rings, the center and rings separated by concentric dark intervals.

The central concentration of this diffraction artifact is called the Airy disk (after the Cambridge professor George Biddell Airy who first described it mathematically in 1834), surrounded by concentric diffraction rings. The angular dimension of these features, in relative units of λ/D, the relative luminance (as a proportion of the peak Airy disk luminance) and the percentage of total light energy enclosed by each dark ring are given in the table (below).

the "point" source diffraction artifact of a circular aperture
feature dimension (k) in λ/D radianspeak
percent of

radius of

diameter of

Airy disk FWHM0.5141.0280.50050.0%
first dark ring(1.220)(2.440)083.8%
first bright ring1.6353.2700.0175.

the Airy disk and diffraction rings of a "point" light source in a circular aperture

the visible diffraction artifact differs from its mathematical specification in several details (see text below); turbulence, magnification, aberrations, collimation and focus all have a significant distorting effect on the diffraction appearance

second dark ring(2.233)(4.466)091.0%
second bright ring2.6795.3580.0042.
third dark ring(3.238)(6.476)093.8%
third bright ring3.6997.3980.0016.
fourth dark ring(4.241)(8.482)095.3%

Angular Resolution Criteria. This diffraction artifact — the Airy disk and encircling rings —is the smallest possible physical image of a "point" light source. As a result, the diffraction artifacts created by two "incoherent" or physically unrelated light sources (such as two stars) will appear to merge on the image plane when the angular separation between them is smaller than the angular diameter of the artifacts.

What is the minimum angular separation between two point sources necessary to resolve their overlapping diffraction artifacts? This is the angular resolution limit (Ro) for a given aperture. It is determined by the wavelength of light taken as the reference, the unit of the diffraction artifact used as the minimum resolvable detail of the image, and the aperture diameter:

[2]  Ro = k·λ/Do·206265 arcseconds

where λ and Do are measured in millimeters, and the constant 206265 converts the tangent (radians) into arcseconds of angular width.

The factor k is used to define a perceptual (visual or sensor relevant) resolution criterion in units of the fundamental λ/Do resolution metric. When k = 1.0 we have the λ/D resolution limit or Abbe resolution limit, first applied to microscope optics by the German physicist and optical designer Ernst Abbe in 1873. If k is greater or less than 1.0 we have one of the many perceptual standards that have been proposed as resolution criteria. The three most often encountered in astronomy are:

angular criterion (year)criterion formulabasis for the criterion

[3a] Rayleigh criterion (1879)RR = 1.220·λ/Do·206265 arcsecondsline resolution of a slit spectroscope
[3b] Dawes criterion (1832)RD = 1.021·λ/Do·206265 arcsecondstelescopic observation of double stars
[3c] λ/D (Abbe) limit (1873)Ro = 1.000·λ/Do·206265 arcsecondsoptical theory
[3d] Sparrow criterion (1916)RS = 0.950·λ/Do·206265 arcsecondslaboratory visual experiments

Adopting a reference wavelength of λ = 0.00055 mm, the more convenient formulas to calculate the resolution criteria of an aperture are:

Do in millimeters

Do in inches

[4a] Rayleigh (k = 1.22)RR = 138.4"/DoRR = 5.449"/Do
[4b] Dawes (k = 1.02)RD = 115.7"/DoRD = 4.560"/Do
[4c] λ/D or Abbe (k = 1.00)Ro = 113.4"/DoRo = 4.466"/Do
[4d] Sparrow (k = 0.95)RS = 107.8"/DoRS = 4.243"/Do

Note the equivalence of the λ/D limit and Dawes criterion: for example, the minimum interval resolved in a 12 inch (305 mm) aperture at the λ/Do limit is Ro = 113.4"/305 = 0.372"; with the Dawes criterion, it is 0.379". The diagram (below) illustrates the visual implication of various resolution intervals with FWHM Airy disks.

visual aspect of different resolution criteria (k) in relation to the FWHM Airy disk

Angular Telescope Magnification. Once the angular size of the smallest resolvable interval is calculated for a specific aperture, the practical question is: how much magnification is necessary for the observer to resolve it visually?

Assuming that both the eye and telescope are diffraction limited, a common rule of thumb is that the required magnification will equal the aperture diameter in millimeters: a 305 mm aperture will require about 305x magnification.

A more accurate answer requires a measure or estimate the observer's angular visual resolution limit (Rv). For normal vision this is usually given as 120 arcseconds. Then the necessary magnification is:

[5a] Mt = Rv/Ro

[5b] ƒe = ƒo / (Rv/Ro)

Given the standard angular visual resolution limit Rv = 120", the required magnification in a 305 mm aperture for the λ/D limit would be Mt = 120"/0.372" = 323x. My visual limit is about Rv = 100", which requires only about 270x. This large range in the estimated necessary magnification is typical of the differences in visual acuity across "normal" individual observers.

Linear Resolution Criteria. Resolution can also be expressed as a physical interval (in millimeters) on the image plane, where the image is much too small to be utilized by naked eye. The function of the eyepiece is to provide the dioptric magnification necessary to make these tiny image details visible. This requires comparison with the observer's linear visual resolution limit (Rv) or, without the eyepiece in astrophotography, the width of individual pixels in a CCD receptor.

Unlike the angular resolution in image space, which scales only with aperture Do, the linear or physical width of the diffraction artifact on the image plane (So) scales only with relative aperture No, independent of aperture and focal length:

[6] So = k·λ·No millimeters

where k is the scaling factor of a resolution limit based on a diffraction feature (see formulas 3a-3d, above). Given an ƒ/8 telescope and λ = 0.00055 mm, the linear diameter of the FWHM Airy disk is 2.03 x 0.00055 x 8 = 0.009 mm — about the size of a single red blood cell. The eyepiece enlarges this small interval so that it is visible to the eye.

For reference, the table below shows the linear metric (at No = 1) and the matching naked eye linear criterion (at No = 250) corresponding to the four angular resolution criteria (for λ = 0.00055 mm).


angular metric R
(Ro = R/Do)

linear metric S
(So = S*No)

matching naked
eye criterion

[7a] Rayleigh criterionRR = 138.4"S1 = 0.000671 mmS250 = 0.16775 mm
[7b] Dawes criterionRD = 115.7"S1 = 0.000562 mmS250 = 0.14039 mm
[7c] λ/D (Abbe) limitRo = 113.4"S1 = 0.000550 mmS250 = 0.13750 mm
[7d] Sparrow criterionRS = 107.8"S1 = 0.000523 mmS250 = 0.13063 mm

The linear metric must be multiplied by No (or the matching naked eye criterion by the quotient of No/250) to obtain the linear resolution of a specific relative aperture. For example, in an ƒ/10 telescope with the λ/D limit, the linear resolution limit is either So = (10/250)*0.1375 = 0.04*0.1375 = 0.0055 mm, or So = 0.00055*10 = 0.0055 mm.

Linear Eyepiece Magnification. How much eyepiece magnification is necessary? As before, that depends on the observer's visual acuity, defined by a linear visual resolution limit (Sv). This is measured as line pairs per millimeter (lp/mm) or as the spacing between parallel white and black lines viewed at the near point (250 mm) with the naked eye. The typical value is 6 to 8 lp/mm or an average line spacing of about 0.1429 mm.

In the example of an ƒ/8 objective given above, we'd require an eyepiece magnification of 0.1429/0.009 = 15.9x to barely detect the FWHM Airy disc; the formula for dioptric magnification indicates this is roughly a 250/15.9 = 16 mm eyepiece.

The formula to calculate directly the eyepiece focal length (ƒe) at which the diffraction limited image details will be just visible is:

[8] ƒe = 250·(So/Sv) millimeters

Thus the eyepiece focal length that will resolve the smallest image detail at the λ/D limit in an ƒ/8 telescope (So = 0.0044 mm), given a visual resolution of Sv = 0.1429 mm, is

ƒe = 250·0.0044/0.1429 = 7.7 mm

which is an eyepiece magnification Me = 250/7.7 = 32.5x. (See also the comparative table here.)

Line Resolution Criterion. In an afocal system with circular aperture, the unweighted quantum formula is used to compute the minimum resolvable interval between two lines is:

[9] νo = λ/Ds·206265 arcseconds

As already explained, this is equivalent to the spacing between dark rings of the point diffraction artifact or the dark bands produced by a two slit aperture mask, where Ds is the distance between the centers of the two slits.

Choice of Resolution Criterion. Because these are practical resolution criteria, and implement different conceptions of what resolution means, none of them should be interpreted as anything more than guidelines to aperture performance. As a general caution, Williams & Coletta (1987) comment:

"The concept of an optical resolution limit has proven difficult to formulate; it is not clear that the resolving power of an optical instrument is a particularly meaningful figure of merit. Even the cutoff frequency of a diffraction limited system can be surpassed under some circumstances. The concept of visual resolution suffers from a similar lack of generality, and the theoretical limits of performance are set by the task as much as by the underlying visual architecture. The sampling theorem [that predicts resolution based on Nyquist foveal cone spacing] correctly specifies the highest frequency possible for image reconstruction without aliasing. However, it does not necessarily prevent an observer from extracting enough critical features of a supra-Nyquist grating to be confident that he sees it." ("Cone spacing and the visual resolution limit," 1987, p.1521)

We might expect, then, that published resolution limits are inconsistent with each other and have limited generality: and this is what we find.

John William Strutt, Lord Rayleigh (pronounced RAY-lee) was originally concerned with the linear diffraction artifacts produced in spectral lines by the rectangular slit of a spectroscope, and he observed in 1879 that "a double [spectral] line cannot be fairly resolved unless its components subtend an angle exceeding that subtended by the wave-length of light at a distance equal to the horizontal aperture" (in other words, the factor k must be greater than 1.0). He proposed as a criterion that "the central line in the image of one [spectral line diffraction artifact] coincides with the first zero of brightness in the image of the other." This is given as k = 1.22 or Ro = 1.22·λ/Do radians in the table above. In telescopic images it places the central peak intensity of each Airy disk over the minimum of the first dark ring of the other (diagram, right), where the merged minimum brightness between the two stars is roughly 70% of the peak brightness of each Airy disk. Note that the Rayleigh criterion is commonly used in astronomy despite the fact that Rayleigh later (1880) determined by experimental tests that resolution with circular apertures was practicable at k = 1.09.

Jean-Bernard-Léon Foucault (c.1876), observing a scale of evenly spaced dark lines, found a resolution limit of about 4/3" with a 100 mm telescope, which is Ro =133"/Do and k = 1.17 at λ = 0.00055 mm.

William Dawes established his criterion in around 1832 after he "examined with a great variety of apertures a vast number of double stars, whose distances seemed to be well-determined," then calculated that a one inch (25.4 mm) aperture would just resolve a matched pair of 6th magnitude stars separated by 4.56 arcseconds (k = 1.02). He then extrapolated the ratio 4.56"/Dinches to apertures as large as 30 inches (762 mm), with the bald assertion that "the separability of all magnitudes is nearly the same, provided the state of the air is such as to bear well the increase of power."

C.M. Sparrow (1916) proposed that two point images can be resolved when the joint intensity function (the luminance sum of the overlapping diffraction artifacts) falls below a straight line connecting the two intensity peaks (k = 0.95). In two stars of equal magnitude, this means the intensity function between the two peaks of a point source pair is not perfectly flat but very slightly concave, and the visible circumferences of the Airy disks appear to touch (diagram, above). Note that no dark interval is perceived between the two point images, which does not qualify as resolution in some uses of the word. However this is considered a more realistic resolution limit for both photographic optics and human visual acuity.

Other criteria exist: the Schuster criterion requires the full width Airy disks to be completely separated, which is twice the Rayleigh criterion; the Houston criterion uses the FWHM in the same way. These are rarely used in astronomy.

Because these criteria reflect assumptions that are questionable in some circumstances and arbitrary in others, there is clarity in the simple choice of Ro = λ/Do (that is, k = 1.00). This λ/D resolution limit directly implements the 1/D rule first deduced by Fraunhofer (1823) and verified by the quantum theory of light. Empirically, it is equivalent to the average separation between two adjacent dark interference bands in the diffraction artifact of a circular aperture (the smallest diffraction detail, average k = 1.01), and it is the criterion often used for the resolution of line pairs (see formula [9], above). It is insignificantly different from both the empirical Dawes resolution criterion (k = 1.02) and the FWHM diameter of the Airy disk (k = 1.03) as shown in the diagram (above). Whether this λ/Do interval can be resolved by the observer's eye depends on the image magnification, the seeing, the optical quality of the telescope, and the observer's visual acuity. At an exit pupil of 1.0 (where magnification is equal to the aperture diameter in millimeters) it requires a visual resolution limit of Rv = 113.4 arcseconds, which is within the range of normal vision.

However, magnification and visual acuity are not the only factors that affect a practical resolution criterion:

• Sharp eyed and experienced observers are able to detect a "rodlike" or "elliptical" elongation of the star image when two stars are separated by as little as k = 0.37 to 0.50 (a resolution threshold of Rv = ~0.45), although this detection criterion seems to rapidly deteriorate in apertures greater than about 90 cm because large apertures are limited by the effects of atmospheric turbulence.

• Sidgwick (Amateur Astronomer's Handbook, pp.45-7) notes that matched 9th magnitude stars (fainter than the 4th to 6th magnitude range in which resolution is typically evaluated) require a separation of from k = 1.37 to 1.90, which implies a low contrast or low luminance criterion of roughly Rv = 135 to 190 arcseconds.

• All the resolution criteria place the Airy disk of a binary component within the diffraction rings of the primary; these obscure the fainter star if the primary star is bright (m < 5), produces glare or scatter in the image, or if the companion is faint m > ~7.5 and is more than 2 magnitudes fainter than the primary.

• The more stringent Dawes resolution criterion applies only to two stars of equal brightness viewed through a circular telescope aperture with no central obstruction. In addition, the Sparrow criterion becomes larger as the magnitude difference between the two stars increases. The Rayleigh criterion is applied without regard to the magnitude or difference in magnitude between two stars, the presence of a central obstruction, or differences in aperture.

• Reflectors, because of the central obstruction, push some light into the diffraction rings, slightly reducing the brightness and diameter of the Airy disk. This can provide improved resolution for very close pairs of similar magnitude, but can obscure a faint companion located within the diffraction rings of the brighter star.

• The reference wavelength of light is usually chosen to match the peak foveal response, which is not well defined. Different measures of the photopic sensitivity function place the average peak response between 545 and 555 nm, and individual differences in color vision enlarge this over a rather broad range between approximately 530 to 580 nm. The choice of 550 nm is purely conventional.

Clearly any theoretical resolution limit is at best a general guideline which must be applied carefully depending on the specific visual task.

Appearance of the Diffraction Artifact. Given a fixed aperture and focal ratio, optical theory stipulates that — regardless of the magnitude of the "point source" — the diffraction artifact will always have the same linear and angular diameter and the same proportional illuminances between the Airy disk and any of the rings around it.

Vasco Ronchi (1961) noted the important difference between the ethereal image or unseen physical characteristics of the actual star, the calculated image or theoretical representation of this image in an optical instrument, and the detected image that actually appears to the astronomer's eye. Remarkably, optical theory in the form of the calculated image, a "mere mathematical construction," does not describe the detailed visual appearance of a stellar detected image.

First, as Sidgwick explains (Amateur Astronomer's Handbook, p.39):

"At normally encountered ƒ/ratios (say, ƒ/5 to ƒ/20) no intensity gradient across the disc is perceptible, the border between the disc and the first minimum appears nearly sharp, and the rings are brighter than theory would indicate, the first ring being not much fainter than the disc itself. ... The visible extent of the disc, like the number of rings visible, varies for a given instrument with the brightness of the source, although the discs are in fact the same size, irrespective of brightness."

In the 17th and 18th centuries the variable diameter of the Airy disks led some astronomers to believe that they were observing the surfaces of stars of different diameters or distances. However early 19th century astronomers such as William Herschel concluded this was an error, and coined the term spurious disc to denote it. In 1828 John Herschel declared flatly that "these are not the real bodies of stars, which are infinitely too remote to be ever visible with any magnifiers we can apply, but spurious, or unreal images, resulting from optical causes." (Note that his intent was to disparage the disk as the image of a physical object, not to disparage the visible disk in contrast to its mathematical description, as some amateur astronomers use the term today.) He continued:

"The apparent size of the disc is different for different stars, being uniformly larger the brighter the star. This cannot be a mere illusion of judgment; because when two unequally bright stars are seen at once, as in the case of a close double star, so as to be directly compared, the inequality of their spurious diameters is striking; nor can it be owing to any real difference in the stars, as the intervention of a cloud, which reduces their brightness, reduces also their apparent discs till they become mere points." (Encyclopedia Metropolitana, 1828)

The paradox of the varying disk diameter and number of rings was explained by Airy in 1835 as due to the luminance threshold of the eye. This is illustrated by plotting the diffraction artifact using a log scale for the luminance (diagram, right), because the relative brightness of lights viewed against a dark background is a logarithmic function:

Δb = b2/b1 = 100.4(m2-m1)

The dotted horizontal lines in the diagram show the relative effect of a visual threshold that extends down to either ~1/10th to ~1/1000th the peak Airy disk brightness. If a magnitude 2 star is bright enough to present three rings, this explanation requires the third ring of a mag. 2 star to appear as bright as a magnitude 9.5 star (diagram below, A).

Surprisingly, a threshold explanation for the changing angular size of the disk does not logically explain the appearance. For example, the Rayleigh resolution criterion results in a clearly visible dark gap between the Airy disks of two equal magnitude stars, which the diagram of the joint intensity function (above) shows is a brightness of about 73% of the peak disk brightness. But the brightest (first) ring of the diffraction artifact is only 1% of the peak brightness (table, above), and therefore should always appear to be "dark" (invisible) for any matched pair that appear to be two separate disks at the Rayleigh angular separation.

In addition, a faint star in a large aperture reflector can still display the first diffraction ring even when the Airy disk is noticeably reduced in size, and the diagram (above) shows that this is also inconsistent with a threshold explanation — the threshold that would produce a reduced disk will be far above the peak brightness of any ring, so the ring would be invisible.

Finally, I observe the ring to contract around this smaller disk, maintaining the same dark gap width but becoming slightly thinner (diagram right, B). The ring certainly does not keep the same diameter regardless of the star magnitude, as theory requires: this would produce in very faint stars a visibly enlarged first gap around a much reduced Airy disk (diagram right, C). This contraction of the first ring around the reduced Airy disk also cannot be explained by the threshold model.


Magnification is the increase in the angular size (apparent width) of an object in comparison to the naked eye size of the object. In telescope optics, magnification results from the combined effect of the telescope objective and eyepiece. But the magnification is obtained in opposite ways: by means of projection magnification in the objective (a function of ƒo), and by means of dioptric magnification in the eyepiece (a function of 1/ƒe). In both cases, image dimensions formed by the naked eye are the benchmark to calculate magnification.

The Near Point. The eye has a variable focal length: it can focus on objects at different distances within the fixed distance from the cornea and lens to the retina. An object appears larger — is effectively magnified — as it is moved nearer to the eye, and this also increases the distance behind the eye at which the object image comes to a focus. To compensate, the eye lens can increase its optical power or accommodate to shorten the back focal length and bring the image of near objects into focus on the retina.

What object distance provides the benchmark for the calculation of magnification? It is the distance that provides the maximum image size on the retina, the minimum distance at which an object will still appear comfortably and completely in focus. This is the point of maximum accommodation, where the lens has contracted to produce its greatest optical power (shortest focal length).

For the average human observer this near point is 250 mm (25 cm) in front of the principal point of the eye, defined by the eye's aperture stop (the iris).

Dioptric (Eyepiece) Magnification. An object moved closer to the eye than the near point will appear larger, but also out of focus because the lens cannot increase its optical power any further. Focusing the image of this closer object is made possible by adding the refractive or dioptric power of a magnifier lens or eyepiece.

Dioptric power is measured in diopters, the reciprocal of focal length (1/ƒ) measured in meters. The inverse form indicates that a shorter focal length means greater refractive power.

Magnification can be defined as the ratio between the half angular height β of the object at its magnified (apparently closer) distance to the eye divided by its half angular height α when viewed by the naked eye at the near point distance of 25 cm, and using the tangent of the angles for precision (diagram, above):

[1] Me = tan(β)/tan(α)

The distance size rule states that the ratio between the tangents of the two visual angles is the inverse of the ratio between the two distances. Thus, an object with an angular height α = 4° that is magnified to an angular height β = 20° will appear magnified by Me tan(20)/tan(4) = 5.2 times, and therefore will appear enlarged as if placed at a distance 250/5.2 = 48 mm (4.8 cm) from the eye.

This principle leads naturally to the calculation of the eyepiece magnification as the focal length of the eyepiece standardized on the near point distance of 250 mm (diagram below, right):

[2] Me = 250/ƒe

If the ocular is used as an actual magnifier, with the object held at the front focal plane of the eyepiece, then its magnification is 250/ƒe+1.

Projection (Objective) Magnification. The projection magnification of a telescope objective is fundamentally different from eyepiece magnification: it is determined by focal length (ƒo), not by the convergence angle of optical power.

The comparison here is to the magnification produced by a pinhole camera, which is produced only by the distance between the pinhole and the surface receiving the image. A pinhole has no optical power (no focal length), so to determine projection magnification we use the angular separation α of two rays from an imaged object.

In optics, a chief ray is the light path on which a lens or mirror has no optical effect, because it leaves an optical element at the same angle that it entered. This ray projects a greater field height y, measured perpendicular to the optical axis, as focal length increases. Therefore the magnification produced by the telescope objective is a linear function of its focal length (diagram above, left).

What telescope focal length will produce an image of a distant object (such as the Moon) that, when the image is viewed at the near point, will appear exactly the same size as the naked eye view of the Moon in the sky (x)? The answer: when ƒo = 250 mm. So the objective magnification is equal to:

[3] Mo = ƒo/250

For example, the objective in a 250 mm ƒ/8 telescope has a magnification of (2032)/250 = 8.13 times.

The Moon is 3475 km in diameter at a distance of 384,400 km, but in terms of optical calculations the linear width of its image at the near point is 2.26 mm. Therefore its image diameter at the focal point of an ƒo = 2032 system is 8.13 x 2.26 = 18.4 mm.

Telescope Magnification. A telescopic image is actually the compound of two different kinds of magnification: projection magnification that increases with increasing objective focal length, and dioptric magnification that increases with decreasing eyepiece focal length.

Because one magnification has a multiplying effect on the other, the magnification of the combined telescope and eyepiece system is the product of the separate telescope and eyepiece magnifications. The value 250 cancels out, leaving the ratio of the two focal lengths:

[4] Mt = Mo·Me = ƒo/250 x 250/ƒe = ƒo/ƒe

Or, starting with formula [1], and defining the tangents as α = h/ƒo and β = h/ƒe, we have:

Mt = β/α = (h/ƒe)/(h/ƒo) = (h/ƒe)·(ƒo/h) = ƒo/ƒe

Thus, in a 10" (254 mm) ƒ/8 telescope with a 20 mm eyepiece, the magnification is:

Mt = (254 x 8)/250 x 250/20 = 8.13 x 12.5 = 101.6

The diagram (below) summarizes these relationships. The magnification of the objective Mo depends on the distance that the real image is projected beyond the objective near point. The magnification of the eyepiece Me depends on the distance that the virtual image is located in front of the eye near point (green arrows in diagram). The optical effects are reciprocal: objective magnification, as in 17th century instruments, made telescopes larger (with larger relative apertures); high power, wide field dioptric magnification, the main innovation of 20th century eyepieces, allowed telescopes to be more compact (with smaller relative apertures).

Note that the field stop linear radius h is transformed into the virtual field stop limiting the apparent field of view (the eyepiece projection angle β) by the eyepiece focal length; it can be calculated from those parameters without actually measuring the physical edge of the field stop in the eyepiece. Also, as a feature of the eyepiece, the field stop does not define the maximum image width that can be produced by the objective, only the maximum image width that can be viewed through the eyepiece. The largest feasible image diameter is usually the interior diameter of the largest eyepiece that can be used with the objective, which in commercial 2" eyepiece barrels is about 1.8 inches or 46 mm.

When calculating magnification, keep in mind that manufacturer supplied data on telescope focal length, eyepiece focal length and eyepiece field stop diameter is often inexact, and therefore magnifications calculated from them should be considered approximate. Dioptric magnification greater than about 10x (ƒe < 25 mm) and free of serious aberrations is not possible with a single lens; a compound lens is required.

Magnification & Image Brightness. Magnification affects the image illuminance: as magnification increases, the image becomes dimmer. This is because magnification causes a smaller part of the sky (the true field of view) to be passed to the eye, and therefore a smaller quantity of light is being used to fill the apparent field of the eyepiece. The magnification required to produce an exit pupil of a certain size, and the effect of this on image brightness (illuminance), can be calculated starting with the exit pupil equality de = Do/Mt = δ and assuming an eye pupil nominal diameter of 6 mm (table, below).

image luminance relative to the naked eye

for apertures at specified exit pupils of 6 to 0.1 mm and equivalent magnification


(naked eye = 1)

magnification at column aperture (mm)
100 150 200 250 300 350 400 500


Thus, an exit pupil of 2.0 transmits only (2/6)2 = 1/9 or 0.1111 the light of an exit pupil of 6.0, which is produced by a magnification of 100x in a 200 mm objective or 175x in a 350 mm objective.

The telescopic image of an extended object (the Moon) cannot be brighter than the naked eye image of the same celestial area, due to "stopping down" of larger beams by the eye pupil. The telescopic image of a star brightens with aperture, because the larger quantity of light gathered by the larger aperture is concentrated into a star image that, at exit pupils above ~2.0, effectively remains a "point" source of light.

Magnification Selection

Various guidelines have been proposed to decide the optimal or most useful astronomical magnification. In this nook of the astronomical literature there are three strategies or approaches, based on exit pupil, target angular width, or image quality.

The first two approaches require the observer's visual resolution limit (Rv), the minimum angular width that can be resolved by the naked eye. Estimates of this limit for a normal eye range from 60 up to 300 arcseconds, depending on the type of resolution required. Typically 120 arcseconds is the value used for visual double star astronomy, although the limit becomes larger for faint, low contrast (deep sky) image detail. (Your personal resolution value is easy to measure with an aperture mask and a zoom eyepiece or a series of eyepieces of different focal lengths.) 

Three Magnification Criteria. Aperture (Do) determines the dimensions of the diffraction effects that limit the angular resolution of the telescopic image. In some situations it is useful to select magnification in relation to the visualization of diffraction artifacts or resolution limits. This is the exit pupil approach.

The effects of a smaller exit pupil include: (1) greater visibility of the diffraction artifacts in the image, including the resolution limit of the aperture; (2) greater visibility of the effects on diffraction limited resolution of objective aberrations such as lateral color, coma or astigmatism, (3) a reduction in the effect of aberrations created by the observer's eye, due to concentration of the image beam at the center of the lens and cornea, and (4) the shadow projection of distracting "floaters" in the eye or of dust particles on the eyepiece eye lens.

If Rv is known (or assumed), then the optimal magnification can be determined by the angular size of the diffraction resolution criterion (Ro) of the aperture as:

Mopt = k·Rv·(Do/113) = k·Rv/Ro

where Do is measured in millimeters, 113 signifies the λ/D resolution criterion, and k is a "comfort factor" that depends on how much larger than this "just barely" visual limit you want the feature to appear. For example, the Rayleigh resolution criterion uses k=1.22, and in a 140 mm ƒ/7 telescope this is 1.22·113/140 = 1 arcsecond, which (assuming Rv = 120) requires a magnification of 120x obtained with an 8 mm eyepiece.

This magnification is sufficient to make all the image detail just visible to the observer's eye. As an analogy, it is equivalent to holding a printed image just close enough to the eye that the halftone dots in the image are visible, since the dots are the smallest detail the image contains. Any lower magnification will disguise the smallest details, and any higher magnification (for example, k = 2) will make the diffraction details larger and easier to see.

Because magnification by definition makes the image of the object you observe larger, it can be recommended by considering the angular size ρ of the smallest feature you are attempting to see, calculated as:

Mopt = k·Rv/ρ

Here k = 2 to 3 is more common. For example, a typical crater on the Moon is 10 kilometers wide, or about 5 arcseconds when viewed from Earth, which in a 140 mm ƒ/7 telescope requires a comfortable magnification (assuming k = 3) of about 3·120/5 = 70x obtained with a 14 mm eyepiece. This object magnification approach guarantees that the object you are looking for will be visually small but nevertheless clearly visible.

The difference between diffraction magnification and object magnification is illustrated (below) as the separation in a matched binary about 2.0" apart (the object interval) viewed through three different telescopes, which produce Airy discs of different diameters (the diffraction interval). Simply by switching eyepieces between the largest and smallest telescopes, different diffraction magnifications appear at similar object magnifications (M ~250x), or similar diffraction magnifications (de ~1.0) appear at different object magnifications.

However, magnification influences many other aspects of the image besides the visibility of diffraction artifacts or the apparent size of features in object space. These include: (1) a darkened sky background that increases the visibility of faint stars or the contrast in extended faint objects such as nebulae, (2) a reduction in objective optical aberrations like coma, distortion or field curvature that increase with field height on the image plane, (3) a fainter and less distinct image, (4) amplification of the visual effects of atmospheric turbulence, focus inaccuracy and instrument vibrations, and (5) greater difficulty centering and manually tracking a target object. A suggested magnification based on any of these purely visual criteria has adopted the image quality approach.

Image quality is a vague criterion, because it is judgmental rather than calculated, and because image quality depends on many things, for example the optical quality of your objective, the angular width of the object or feature you want to see, collimation, light pollution, and so on. The recommendation to limit your magnification to "what the seeing will allow" is one example of the image quality approach: there is no consideration of the exit pupil or object apparent size involved in this suggestion.

Magnification Recommendations. Whether it is useful to see the resolution limit of the aperture or the smallest feature of an astronomical object depends on whether the image structure or the object structure is the observational interest. Whether you are able to use the magnification necessary to see the artifact or feature you are looking for depends on the quality of the image that results. This in turn depends on your equipment.

As a practical reality, the available configurations of amateur astronomical equipment are limited. The available range of apertures is approximately 80 mm to 600 mm; relative apertures range from ƒ/4 to ƒ/20, eyepiece focal lengths vary from about 5 mm to 40 mm. By combining these extremes we can estimate the equipment limit values for exit pupil (0.25 to 10), magnification (8x to 2400x) and resolution (1.41" to 0.19") that suggest practicable opportunities. Suggested magnifications are choices within those ranges supported by good visual reasons.

Magnification that is defined as a multiple of aperture (in inches or millimeters) is actually a recommendation for the diameter of the exit pupil (de), because the exit pupil can be defined as aperture divided by magnification (Do/M; see below), so magnification is just aperture multiplied by the reciprocal of the exit pupil, Mt = 1/de·Do. Magnification that is defined per se, for example as 200x, is specifically an object magnification recommendation, and is usually given in relation to a specific astronomical target (e.g., the Cassini division in Saturn's rings, the disk of Venus, a 3" separation in a binary star, the Hyginus Rille on the Moon). Image quality recommendations can use either the exit pupil or object magnification, or the quality of the seeing and light pollution.

The astronomical literature commonly adopts the exit pupil approach, both because magnification varies inversely with exit pupil and because absolute magnification values (such as 200x) can potentially be anywhere within the range of magnifications possible with a single telescope and eyepiece selection, meaning they can have widely different image quality implications. The exit pupil serves to anchor all those variations along a dimension of image quality that uses diffraction detail as a magnification boundary.

The five common exit pupil recommendations, and the justifications offered for them, are:

• Mt = Do/δ, ƒe = δNNormal magnification – This defines the magnification at which the telescope exit pupil equals the eye pupil aperture of the observer. The rationale for this limit is that any lower magnification produces an exit pupil that is larger than the eye pupil, so that not all the light can enter the observer's eye: this effectively stops down the objective (reduces the illumination of the image) with no visible change in the image resolution. (In a 180 mm ƒ/10 refractor or newtonian reflector with a 60° eyepiece field, the normal magnification yields a true field of view of 2°.)

This is a dubious rule. Few observers have actually measured their dark adapted pupil aperture (the customary value of 7 or even 8 mm is typically too large ... mine is 5.9 mm), and observing bright objects such as the Moon can contract the pupil to 3 mm or less, wasting even more aperture. Low magnification has benefits unrelated to light grasp: the true field of view is larger, pointing and tracking become easier and less critical, and the effects of poor seeing are minimized. Commercially available eyepiece focal lengths rarely exceed ƒe = 40 mm, so telescopes at relative apertures above ƒ/6.7 (most refractors and Newtonians, and all Cassegrain format reflectors) will not reach the prescribed "normal magnification" anyway. Dobsonian users are most often concerned (needlessly) with this lower limit. A more important lower bound is:

• Mt = 0.25Do to 0.33Do, ƒe = 3N to 4N – The so called Minimum useful magnification – A magnification between 0.25Do and 0.5Do is sometimes recommended for deep sky observing when a turbulent atmosphere precludes detailed observation of planets or the Moon. However atmospheric turbulence is sufficiently masked at exit pupils roughly half the "minimum useful" range, so turbulence is rarely the reason to choose a magnification this low.

Instead, the critical fact is that the optical aberrations of the cornea and lens are minimized as the exit pupil is reduced, while the size of the diffraction artifact is reduced as the exit pupil is increased. The crossover between these two curves, which occurs in most normal eyes at an exit pupil of around 3.5 (Mt = 0.3Do), produces the smallest diameter stellar image. This is therefore the brightest stellar image that can be produced with the aperture. Because the magnification has also reduced the sky brightness to less than half its naked eye luminance, the relative brilliance achieved through spot size minimization and increased contrast is at its peak. Faint extended objects are also judged to appear at optimal contrast with an exit pupil of 3.0, or ƒe = 3N. The only countervailing issues are when the relative aperture of the telescope is so low that it introduces distracting aberrations in wide field eyepieces at low powers, or when the angular diameter of the target (star cluster, extended nebula, etc.) is very large — that is to say, really large (a 180 mm ƒ/10 refractor or newtonian reflector with a 60° AFOV eyepiece at the minimum useful magnification yields a true field of view of 1.0° to 1.3°.)

• Mt = 0.5Do, ƒe = 2.0NMinimum optimal magnification – This is the lowest magnification that will render the largest diffraction limited image detail (the Airy disk) just visible, assuming a very optimistic estimate for the observer's visual resolution (i.e., ~1 arcminute or 60 arcseconds, which gives M = 60/Ro). Note that the observer will not be able to split two matched double stars separated by the telescope's resolution limit. When the Airy disk becomes visible, even small atmospheric turbulence becomes noticeable as well, so this is often the highest magnification useful under conditions of mediocre or poor seeing. (In a 180 mm ƒ/10 refractor or newtonian reflector with a 60° eyepiece field, the minimum optimal magnification yields a true field of view of 40 arcminutes.)

• Mt = 1.0Do, ƒe = 1.0NMaximum optimal magnification – This is usually the magnification that renders the smallest diffraction limited image detail (the aperture resolution limit, or the dark interval between the Airy disk and first diffraction ring) just visible, given a conservative criterion for the observer's visual resolution (i.e., ~2 arcminutes or 120 arcseconds, which gives M = 120"/Ro). It is also easiest to remember, since the magnification is just the telescope aperture (in millimeters) and the necessary eyepiece focal length is just the telescope relative aperture. Note that this is equivalent to "Whittaker's rule", usually stated as Mt = 25Do when aperture is in inches, because 25 is just the number of millimeters in an inch. (In a 180 mm ƒ/10 refractor or newtonian reflector with a 60° eyepiece field, the maximum optimal magnification yields a true field of view of 20 arcminutes.)

Here is where resolution or contrast criteria dominate the choice of magnification. At the low end the magnification is "the largest that the seeing will allow," and at the high end it is bounded by the appearance of fresnel diffraction around star images. This is using diffraction resolution as an image quality criterion. Both effects usually occur within the range 0.5DM ≤ 1.0D.

It is useful to try several eyepieces within the range of eyepiece focal lengths prescribed by the minimum and maximum "optimal magnification" factors. This is because variations in magnification cause sky brightness to decrease and make very faint stars more visible, and cause changes in apparent contrast in faint extended objects, making subtle contrast areas larger and more visible. Optimal seeing will allow magnification that shows the diffraction artifacts clearly. Mediocre seeing, atmospheric dispersion or light pollution will push magnification the other way.

• Mt = 2.0Do to 4.0Do, ƒe = 0.5N to 0.25NMaximum useful magnification – This is the magnification that causes the smallest (diffraction limited) image detail to be just visible given the most pessimistic estimate for the observer's visual resolution (~5 arcminutes, which gives M = 300/Ro). This rule is highly dependent on observing conditions and the observing target. Planetary and lunar observers will generally prefer a lower magnification, within "the maximum that the seeing will allow" given the low contrast of planetary detail. In contrast, because the Airy disks of bright stars tolerate high magnification very well, even under poor seeing, double star observers will use magnification beyond the high end of this range. (In a 180 mm ƒ/10 refractor or newtonian reflector with a 60° eyepiece field, the maximum useful magnification yields a true field of view of 10 to 6.6 arcminutes.)

William Herschel was quite explicit about using magnifications over 3000x, which created an exit pupil of 0.05 millimeter in his 165 mm reflector. The question arises whether exit pupils that small become the delimiting pupil in the optical train. The answer seems to be no. Optical studies show an exit pupil below 2 mm down to 0.5 mm does not filter high frequency wavelengths as a foil pinhole aperture would, and therefore they can transmit the objective image very well. The maximum useful magnification has been exceeded when the image is too dark, too blurry (due to the quality of the optics or atmospheric turbulence), or too difficult to view (due to "floaters" in the observer's eye, small eye relief, etc.).

The image quality rule — use the magnification that produces the most useful image — summarizes the way that the different image quality and diffraction resolution criteria combine, and leaves the question of object magnification out of consideration. Atmospheric turbulence, the magnitude of stars or the surface brightness and surface contrast of extended objects, the sky brightness, the type and optical quality of the instrument, the quality of the observer's eyes and the specific visual task or pastime — all these contribute to the choice of the optimal and most useful magnification.

The corollary is: change the eyepiece often. No matter what the visual task or observing target, observing the same target with different eyepieces will verify the best choice, and can reveal unexpected details in the peripheral field or in the target. This desire to see alternative views is one sign that the image effects of magnification are really understood.

Exit Pupil

Physical Attributes. The exit pupil is fundamentally the naked eye image of the telescope entrance pupil (aperture stop) as it appears through the eyepiece (image, right). If the telescope is turned toward the daytime sky, the exit pupil as a virtual image appears as a small, bright disk hovering just above the eye lens of the eyepiece. Like the virtual image in a mirror, it appears to move as you change your vantage.

This virtual image is formed by the convergence of all real image rays focused by the eyepiece. At this point the cross section of the projected, afocal real image of the telescope is physically smallest (as shown here). In this fixed position, centered on and perpendicular to the optical axis, it is traditionally called the Ramsden disc in the United Kingdom, in honor of the 18th century scientific instrument maker who first described it.

All real image rays means that a physical baffle with the same internal diameter as the exit pupil could be placed around it without obstructing any of the light exiting the eyepiece. The observer's correctly positioned iris mimics a baffle or stop when the telescope is used for visual astronomy; in that context the intersection of the exit pupil with the optical axis is called the eye point.

Finally, the exit pupil is the point where the telescope afocal image is most concentrated and therefore brightest; it is used as the benchmark luminance in the calculation of changes in image brightness caused by magnification.

These three prosaic physical attributes of the exit pupil — virtual image, smallest real image, brightest projected image — belie the almost fetishistic importance that some astronomers attach to it. In the evaluation of an optical system, as Smith, Ceragioli & Berry (2012) recommend doing it, "all that matters is the diameter of that exit pupil!" That claim invites a detailed examination.

Definitions. The most primitive method to specify the exit pupil is to measure its diameter (in millimeters) with a reticule magnifier. This often precedes the use of the exit pupil to calculate the telescope magnification as the ratio between objective aperture (Do) and exit pupil diameter (de) — the afocal beam compression of the system:

[1] Mt = Do/de

In practice, measuring the size of the exit pupil is inconvenient (especially when it is small), so it is simpler and usually more accurate to calculate the telescope magnification as the ratio of objective focal length (ƒo) to eyepiece focal length (ƒe). Then the exit pupil diameter can be calculated as the ratio between the objective aperture and the telescope magnification:

[2] de = Do/Mt

The logic of this definition is rarely explained. It is based on the fact that the exit pupil is the naked eye image of the entrance pupil (the objective mirror or lens). If we replaced the objective/eyepiece combination with a single objective lens, its object magnification would be equal to its focal length (ƒo). As explained above, this single objective will have a magnification Mo = ƒo/250. Substituting, we have

de = Do/Mo = Do/(ƒo/250) = 250·(Do/ƒo) = 250/No

So the exit pupil calculation [2] based on the ratio of aperture to magnification is identical to the apparent size of the entrance pupil in a single objective system whose relative aperture (No) produces the same magnification. And this single objective must have a very large relative aperture to achieve the necessary magnification — for example, No = 250 when de = 1.0 mm and Mt = Do.

Because the eyepiece has been omitted, this single objective would not form a classic Ramsden disc of real image rays. We can, however, make the comparison in this way. If we place a 1 mm exit pupil at the near point, 250 mm from the eye (the distance of an "actual size" magnification of the object being viewed), its angular width will be 1/250 radian (0.23° or 13.75 arcminutes), and this matches the angular diameter of the single objective entrance pupil when viewed from the N = 250 focal distance. The exit pupil is identical to the image of the aperture.

The significance of relative aperture in the system magnification of formula [2] becomes obvious if we substitute the common formula for magnification mentioned above (Mt = ƒo/ƒe), which yields a second definition of the exit pupil:

[3] de = Do/(ƒo/ƒe) = ƒe·(Do/ƒo) = ƒe/No

Again, start with a single N = 250 objective: this will have an aperture image equivalent to an exit pupil de = 1.0. The entrance pupil, viewed at the focal distance, will subtend an angular width of 1/250 radians = 0.23° or 13.75 arcminutes. But at smaller relative apertures the focal point will be closer to the entrance pupil and therefore the aperture image will appear much larger: for example, in an ƒ/10 objective, the entrance pupil will appear to be 5.73° wide.

Eyepiece magnification, because its effect on image magnification is the same as extending the objective focal length, is equivalent to reducing the entrance pupil angular size by division. As explained above, eyepiece magnification is calculated as Me = 250/ƒe. So an ƒe = 10 mm eyepiece (which according to formula [3] yields an exit pupil of 1.0 in the ƒ/10 system) provides a magnification of 250/10 = 25x, and adopting the 5.73° entrance pupil image in an ƒ/10 objective, we have 5.73°/25 = 0.23° or 13.75 arcminutes. As before, the telescope exit pupil viewed from the near point is identical to an objective entrance pupil of equivalent magnification (focal length) viewed from its focal point.

Object vs. Diffraction Magnification. This second formula indicates that there are only two practical ways to change the diameter of the exit pupil with a telescope of given aperture: change the eyepiece focal length ƒe, or change the objective relative aperture No, either by using a Barlow lens or a focal reducer.

There is however a third way to reduce the exit pupil: by means of an aperture stop (an opening that is smaller than the full aperture). This demonstrates that we can change the size of the exit pupil without changing the object magnification of the system. A reduced entrance pupil does not affect the focal length of the objective or the eyepiece, so the telescope magnification — the apparent size of any physical object we observe with the telescope — is unchanged. Instead, a reduced entrance pupil increases the linear resolution of the system, in effect magnifying the physical diameter of a star diffraction artifact at the focal surface, and it does so simply by changing the relative aperture (No) of the objective.

The diagram (below) summarizes these relationships between eyepiece and magnification, and between relative aperture and aperture, in the definition of the exit pupil. Notice, in order to get from formula [3] to both values in the ratio of formula [2], we must scale both elements of formula [3] by dividing them into the physical dimension of objective focal length (ƒo). In effect, formula [3] is a dimensionless or scale free version of formula [2].

For clarity, I refer to these two forms of magnification by the nonstandard terms object magnification and diffraction magnification, respectively.

The resolution of a system is commonly used to benchmark, if not limit, the high end of magnification. But because we have two definitions of magnification, we require two matching definitions of minimum diffraction width: (1) as an angular width (visual extent in image space) expressed in arcseconds, which corresponds to object magnification dependent on aperture; and (2) as a linear width (physical extent on the image plane) expressed in millimeters, which corresponds to the diffraction magnification dependent on the relative aperture:

[4a] angular width: Ro = 206265·λ/Do arcseconds

[4b] linear width: So = λ·No millimeters

where 206265 (the number of arcseconds in a radian) projects the metric into a visual space of angular dimensions and the relative aperture No projects the metric onto an image plane of physical dimensions. The metric defined by diffraction bands of "yellow green" light (λ = 0.00055) yields Ro = 113.4/Do arcseconds of angular width and So = 0.00055·No mm of linear width.

To determine whether these angular or linear intervals will be visible to the human eye, they must be compared to a naked eye resolution limit defined with the same measurement geometry:

(1) The angular limit (Rv) is the minimum angular separation in cycles per degree or equivalent arcseconds that can be resolved by an observer with normal vision. A commonly used value is: Rv = 120 arcseconds.

(2) The linear limit (Sv) is (2) the minimum linear spacing in a square wave grating viewed at the near point distance (250 mm), and stated as line pairs per millimeter or the reciprocal single pair width in millimeters, that can be resolved by an observer with normal vision. A commonly used value is: Sv = 0.1454 millimeters.

Now we can determine the exit pupil at which diffraction limited detail can be visually resolved in the image of telescopic objects as the ratio of the matching telescopic and visual resolution limits:

[5] de = Ro/Rv = So/Sv

and when Ro = Rv then this "exit pupil of visual resolution" will be de = 1.0 by definition.

The reference values of Rv and Sv are population averages and, depending on the condition of your eyes, your individual visual acuity may be larger or smaller than the normative values. For example, my naked eye resolution is approximately Rv = 100 arcseconds, so diffraction limited detail (the interference bands) is visible to me at de = 1.13. The fact that Rv is typically close to Ro and diffraction limited detail is visible to a normal observer at an exit pupil of ~1.0, indicates that normal human visual acuity is approximately diffraction limited.

Formula [5] allows us to translate the exit pupil specification into the equivalent object magnification (Mt) or longest eyepiece focal length (ƒe) sufficient to resolve diffraction limited detail:

[6] Mt = Do(Rv/Ro) or Do(Sv/So)

[7] ƒe = No(Ro/Rv) or No(So/Sv)

Note that the ratio of optical to visual resolution is inverted between the two formulas — one definition is the reciprocal of the other.

"Wasted" Light & "Empty" Magnification. The exit pupil is a convenient metric for selecting an eyepiece magnification suitable for the parameters of any specific telescope system. It's easy math to multiply focal ratio by the desired exit pupil value to get the necessary eyepiece focal length.

However, the exit pupil has also been used to define the factitious twin "problems" of wasted light (or wasted aperture) and empty magnification. These are worth debunking.

The astronomer is often advised to prevent "wasted" light by avoiding magnification that creates an exit pupil larger than the dark adapted eye pupil, where de > δ. The argument is that an eye pupil smaller than the exit pupil will block some light from entering the eye, effectively stopping down the objective aperture.

The premise for a "wasteful" judgment is demonstrably feeble. With most commercially available astronomical equipment, it is unusual for a practicable exit pupil to exceed the eye pupil aperture. With an ƒ/6 telescope and a pupil aperture of 6 mm (my measured dark adapted pupil size), an eyepiece with focal length greater than 36 mm would be necessary. Even in an ƒ/4 telescope, only eyepieces above 24 mm would create the situation.

So why would one want to use a "wasteful" 30 mm eyepiece in an ƒ/4 telescope? Because a longer focal length eyepiece offers a wider, lower power view of an extended object, and provides maximum eye relief for any type of eyepiece design. Low magnification is used to increase the true field of view and viewing comfort at the cost of "wasted light" in the same way that high magnification is used to increase visual resolution at the cost of reduced image luminance and true field size. In fact, the most serious drawback when deδ is that image contrast is maximally reduced by the brightness of the background sky, making faint stars or extended objects more difficult to see.

At the other extreme, "empty" magnification is said to be any enlargement greater than sufficient to make the smallest diffraction detail visible to the eye. The logic is that you are no longer resolving detail in the celestial object of interest (e.g., the surface details of a planet or the Moon) but resolving the diffraction that actually obscures detail, like the pixelation of a digital image.

This limiting exit pupil can be calculated, for any visual resolution criterion Rv, as:

[8] de = 138.4"/Rv mm (Rayleigh criterion)

[9] de = 113.4"/Rv mm (λ/D or Abbe limit)

Thus, adopting the commonly cited visual resolution threshold of 120" and the Rayleigh resolution criterion predicts that the dark rings around the Airy disk will be resolved when the exit pupil is de = 138.4/120 = 1.15, and the Airy disk will be detectable when the exit pupil is twice as large (de = 2.3). Conventional wisdom asserts that an exit pupil half as large (de < 0.6 mm) only provides "empty magnification" because the eyepiece magnification is then more than twice the amount required to resolve the diffraction limited image detail.

This abstract reasoning ignores any situation where the diffraction artifact is actually what you want to observe. Because the Airy disk is quite robust to magnification when the seeing is mediocre, exit pupils below 0.2 are routinely valuable when resolving close, bright double stars and counteracting the image motion of turbulence. (William Herschel employed exit pupils as small as 0.03 when studying double stars with his "most excellent" ƒ/13 reflector.) Very small exit pupils are also necessary when conducting a star test of optical quality or to evaluate the seeing. There is no such thing as "empty magnification" — only magnification that is not optimal for the visual task at hand under actual observing conditions.

The linear diameter of the Ramsden disc does have important visual consequences. A small exit pupil (less than 1.0 mm) projects light like a pinhole, which can cause vivid shadowing of suspended internal debris or "floaters" on the retina. At the same time the tiny Ramsden disc must be centered in the plane of the iris aperture to produce a view of the entire eyepiece field, which becomes a kinesthetic challenge in very wide angle eyepieces or awkward viewing postures unless the observer has a chair or ladder for support. On the other hand, a smaller exit pupil is only refracted by the central portion of the eye lens and cornea, minimizing the effects of astigmatism. For that reason, many older viewers prefer smaller exit pupils.

More significant problems with the exit pupil occur when observing during daylight with a binocular or spotting telescope. The high luminance of daylight images contracts the pupil aperture to 2 mm or less, and this makes alignment of the eye pupil and exit pupil(s) more difficult. It also makes any spherical aberration of the exit pupil more noticeable.

Image Brightness. In addition to estimates of magnification in relation to diffraction limited resolution, the exit pupil is also used as a benchmark for relative image brightness. Here eye pupil diameter δ (also measured in millimeters) is required for the calculations.

The first [2] definition of the exit pupil represents the same proportion that defines the telescope image illuminance in relation to the retinal illuminance — the ratio in the brightness of the telescope image to the brightness of the naked eye image of the same area of sky:

[10] It = (Do/Mtδ)2, therefore It = (de/ δ)2

As explained above, the size of the exit pupil is related to the system magnification, so that as magnification increases the true field of view gets smaller and the exit pupil, through which all the light exiting the telescope must pass, also gets smaller. As a result, variations in the exit pupil diameter affect the exiting illuminance of the telescope image in the same way that the opening in a camera diaphragm determines the amount of light that can enter a camera. As the exit pupil gets larger, the image luminance increases (graph, left).

The "shutter" and magnification effects combine to brighten or darken the image in inverse square proportion, and the difference in the illuminance of the telescope image equals the ratio of two exit pupils squared. The peak image brightness occurs when de = δ. At that point, the telescopic image is theoretically just as bright as the naked eye image of the same area of sky, omitting transmission loss in the instrument (which in practice is usually significant).

Referring again to the concept of the exit pupil as the image of a single objective entrance pupil viewed from the objective focal distance, an exit pupil of 6.0 indicates the image brightness of an ƒ/(250/6) = ƒ/42 optical system with an entrance pupil apparent diameter of 83 arcminutes (1.38°).

Because any larger exit pupil diameter will exceed the nominal eye pupil aperture, de > δ, the fixed eye pupil becomes the limiting factor as the exit pupil is increased. The area sampled by the fixed eye pupil size is a smaller part of the total exit pupil area, effectively stopping down the objective aperture at the output end. Image brightness again declines.

Apparent Field of View

The AFOV is measured as the projection angle of the radius (β) or diameter (2β) of the field stop (diagram, below). The tangent of this angle is equal to the radius of the field stop opening divided by the eyepiece back focal length (ƒBFL).

The physical diameter of the eyepiece barrel and location of the field stop are dictated by the optics of the eyepiece (its abaxial aberrations and design AFOV) and the ƒ ratio of the objective (see here). The field stop excludes abaxial light from the objective that would produce aberrations in the eyepiece image or an unacceptably dark image, and blocks stray light reflected from the interior of the telescope tube. The barrel diameter limits the maximum internal diameter of the field stop, to ~27 mm in a 1.25" eyepiece barrel and ~46 mm in a standard 2" eyepiece barrel. In all well made eyepieces, the field stop will define the apparent field of view and will be crisply in focus.

The apparent field of view of an eyepiece is normally specified by the manufacturer, along with the eyepiece focal length and design type. However, the nominal AFOV of an eyepiece is often inflated due to the effects of positive distortion, especially in "wide field" eyepieces (AFOV > 60°).

When the manufacturer's apparent field specification appears inaccurate, or is not known, then the AFOV can be measured in three ways:

(1) Projective Measurement. Support the eyepiece so that it is lying on its side with the eye lens about 2 feet from a smooth wall or flat surface. (The seam between pages of an open hardcover book, or between two books of the same thickness laid spine to spine, is convenient.) Support a high powered flashlight in a similar way so that it shines into the field lens (barrel end of the eyepiece) from a distance of a few feet. (The eyepiece and flashlight should be lying perfectly level, the optical axis of the eyepiece and flashlight must be perpendicular to the wall, and the flashlight must be "collimated" or centered on and parallel to the eyepiece optical axis, which can be arranged by centering the shadow of the flashlight bulb inside the projected circle of light.) The eyepiece will project the light from the flashlight onto the wall as a circle of light larger than the flashlight beam. Project the image onto a small white card held close to the eye lens, and move the card back and forth until the projected circle appears smallest: this is the location of the exit pupil. Measure the greatest width of the circle (D) and the projection distance (P) from the eyepiece exit pupil to the wall. Then:

AFOV = 2arctan(D/2P).

For example, if the imaged diameter is 82.3 cm, and the wall to exit pupil distance is 73.3 cm, then:

AFOV = 2arctan(D/2P) = 2 x arctan(82.3/(2 x 73.3)) = 2 x 29.3° = 58.6°.

(2) Visual Measurement. Tape to a convenient wall a long strip of paper or tape measure on which small increments (inches or centimeters) are clearly marked and visible from a distance (Dm) of about 100 cm. Hold the eyepiece over one eye so that the full field is clearly visible and superimposed over the measurement tape, visible with the other eye. Manipulate the eyepiece position until the tape appears to measure the full diameter of the eyepiece field (De). Then:

AFOV = 2arctan(0.5De/Dm).

For example, viewing a metric measuring tape from 100 cm, you measure an apparent width of the eyepiece field as 93 cm. Then:

AFOV = 2arctan(0.5 x 93/100) = 2arctan(0.47) = 2 x 25° = 50°

(3) Physical Measurement. Use calipers to measure the interior diameter of the field stop, assuming it exists and can be recognized. If in doubt, look through the eyepiece and use the tip of a toothpick or pencil to locate the field stop: this is either the end of the eyepiece barrel or the interior edge of the lock nut holding the field lens. (Note that many "super wide field" eyepieces locate the field stop inside the eyepiece where it cannot be measured.) Then the apparent field is simply the visual angle α that the field stop radius would subtend at 250 mm, divided by the apparent field radius β. The distance size equation then makes this equal to:

AFOV = 2arctan(0.5Dfs/ƒe).

For example, if the field stop diameter is 23 mm and the eyepiece focal length is 20 mm, then:

AFOV = 2arctan(0.5Dfs/ƒe) = 2arctan(11.5/20) = 2arctan(0.58) = 2 x 29.9° = 59.8°.

Lateral View of AFOV. Often a wider field of view requires looking to one side of the field, which causes vignetting by the edge of the exit pupil. The field stop is usually apparent in the peripheral field of centered vision, but directing the gaze to the edge of field requires the observer to move his head in the opposite direction to avoid obstruction by the edge of the exit pupil.

For an eye of normal focal length (24 mm), the minimum diameter exit pupil that allows direct lateral viewing of the field edge from a fixed eye position is:

EPρ = 12·tan(AFOV/2-arctan(δ/24)).

where δ is the diameter of the viewer's dark adapted pupil. For example, given an eyepiece with AFOV = 70° and an eye pupil δ = 6 mm, we have:

EPρ = 12·tan(70°/2-arctan(δ/24)) = 12·tan(35°-14°) = 4.6 mm.

With a 10 mm eyepiece, this limit is reached with an ƒ/2.2 or larger focal ratio objective, or for any object magnification greater than 0.22Do, which is below the minimum useful magnification established by image quality criteria. Solving for a maximum 6 mm eye pupil, a fully visible field of view requires an exit pupil no smaller than 5.8 mm with an 80° apparent field of view, 3.4 mm with a 60° AFOV, 2.3 mm with a 50° AFOV, and 1.8 mm with a 45° AFOV.

True Field of View

The exit window admits to the eye an image that has been magnified by both the telescope and eyepiece. The effective size of the exit window, if the image area it defines were viewed with the naked eye, would appear much smaller and much farther away (diagram, above). The area of sky that would be visible through that opening at that distance is equivalent to the true field of view (TFOV) of the objective/eyepiece combination.

Standard Formulas. The TFOV is equal to twice the maximum field angle (α) focused by the objective on the image plane that is admitted to the eyepiece by the physical radius of the field stop (h'). Since α is a linear function of the focal length (ƒo), the TFOV is:

TFOVo = 2αmax = 2h'max/ƒo.

By substitution, the TFOV is calculated as the apparent field of view (AFOV) divided by the image magnification (Mt):

TFOV = 60·AFOV/Mt arcminutes.

Thus, if the apparent field of view of an eyepiece is 50°, and the magnification of the telescope and eyepiece combination is 200x, then the true field of view is:

TFOV = 50°/200 = 0.25°; or 0.25° x 60 = 15 arcminutes

In the diagram (above), an eyepiece with a 50° field of view offers the same visual area as a circular opening 23 cm in diameter held perpendicular to the eye at a 25 cm distance. Then the true field of view is equivalent to the area of sky seen through the same circular opening placed at a distance of 50 meters (164 feet).

Maximum Posible TFOV. If the field stop diameter and objective focal length are known, then the maximum possible TFOV can be calculated as:

TFOV = 57.3·(Dfs/ƒo) degrees.

This formula can also be used to calculate the maximum true field of view possible with an eyepiece barrel or a visual back of a certain diameter. For example, in a 254 mm (10") ƒ/8 telescope and a 1.25" (27 mm) eyepiece barrel, the maximum true field of view is:

TFOV = 27/2032 x 57.3° = 0.76° x 60 = 46 arcminutes.

Using a 2" (46 mm) eyepiece barrel, the maximum true field of view is:

TFOV = 46/2032 x 57.3° = 1.30° x 60 = 78 arcminutes.

Star Drift Measurement. If the field stop diameter and apparent field of view are not known, the true field of view can be measured by means of star transit times. This requires at least three measurements of the time (in seconds) it takes a bright star, placed outside the eyepiece field of view, to appear at one side of the eyepiece field, cross the center of the field, and disappear at the opposite side.

Care must be taken to hold the head in a fixed position at the eye point, so that the entire circumference of the field edge is visible, and to pass the star through the center of the eyepiece field.

Then the true field is the average of the three timings multiplied by the cosine of the star's declination (stars closer to the celestial pole will require a longer time to transit the eyepiece field):

TFOV = 0.25cos(Declstart.

The factor 0.25 is necessary to convert from seconds in time to arcminutes of angular width. For example, using Regulus (alp LEO) as the transit star (Declination = +11°58'), three eyepiece transit times are 134, 131 and 138 seconds. Then:

TFOV = 0.25 x cos(11.97°) x (134+131+138)/3 = 0.25 x 0.978 x 134.3 = 32.8 arcminutes.

Given the TFOV (in arcminutes) and system magnification, the AFOV (in degrees) is:

AFOV = TFOV·Mt / 60.

The TFOV determined by star transit measurements will typically not correspond exactly to the TFOV calculated from the manufacturer supplied eyepiece apparent field of view and the magnification calculated from the manufacturer supplied eyepiece focal length. This is usually because the eyepiece AFOV and ƒe are nominal, inferred from the computer optical design rather than the physical construction of the eyepiece and because star transit times includes some systematic measurement error due to curvature of field; taking transit times in a commercial SCT, whose focal length changes with mirror focusing, also can create small discrepancies.

These errors are typically not large. The chart (right) compares the TFOV of over 60 different eyepieces, as calculated from the nominal eyepiece AFOV and ƒe and as measured by star transits. The calculated AFOV overestimates the measured TFOV by about 4% on average; the overestimation is not correlated with AFOV but is largest (up to 22% greater) at the shortest eyepiece focal lengths (highest magnifications) and in certain eyepiece brands or designs (Orthoscopics).

Finally, what is the telescope configuration that maximizes or minimizes the true field of view?

The four charts (left) show the true fields that result from all possible combinations of aperture (from D = 80 mm to 500 mm), relative aperture (from N = 4 to 20), objective focal length (from ƒo = 320 mm to 10,000 mm) and object magnification resulting from eyepiece focal lengths from ƒe = 3.5 mm to 40 mm with a moderately wide (70°) apparent field of view. The charts assume that a ~50 mm eyepiece barrel can be used with all apertures.

The vertical spread in the charts for aperture, relative aperture and focal length are produced by variation in the other two factors and eyepiece focal length; the single curve for magnification results because the combination of the other system factors produces a unique value. The vertical spread — variation within any single factor caused by the other factors — is largest in relative aperture, nearly as large in aperture, and only somewhat reduced in objective focal length. Only magnification has a strict relation to field of view, as we'd conclude from its role as the denominator in the calculation of TFOV from AFOV.

However, it's the degree of slope or curve in the plots across their range of values that signifies how important each factor is to TFOV, and clearly the most potent feature is magnification, followed by the major component of magnification, objective focal length (where values are spread vertically by eyepiece focal length). Aperture, by itself, has a much reduced effect and relative aperture (or "fast" optical design) has the weakest independent effect of any factor.

Richest Field Telescopes

A richest field telescope (RFT) is the combination of telescope objective and eyepiece design that provides the brightest image within the widest practicable field of view — or, in the original definition, that can show "more stars in the average Milky Way view than any smaller or larger telescope can show."

This design criterion is dependent on three parameters: exit pupil, objective aperture and eyepiece apparent field of view (AFOV). Separate from those design features, relative aperture (focal ratio) and eyepiece focal length should be chosen to minimize optical aberrations or vignetting:

• Exit pupil (de) determines the diameter of the compressed light beam that is passed to the eye by the aperture. To ensure a maximally illuminated image, the common recommendation is an exit pupil that matches the diameter of the observer's dark adapted eye pupil (nominally 7 mm, but in older observers more often between 5 mm to 6 mm). However, as explained above, the brightest stellar images are actually produced at an exit pupil of around 3.5 mm, so exit pupils down this far are also acceptable.

• Aperture (Do) determines the limit magnitude of the system and therefore the total number of stars (or features of dim extended objects) that will be visible above the observer's dark adapted luminance threshold. Because exit pupil is fixed, and Mt = Do/de, the choice of aperture also defines the object magnification (Mt).

• By the selection of eyepiece apparent field of view, the object magnification determines the true field of view. A larger aperture increases the limit magnitude but reduces the true field of view, and this can be partially compensated by choosing eyepieces with the widest practicable field of view.

• Finally, because aperture and magnification are fixed, the relative aperture (No) only affects the eyepiece focal length (ƒe) as ƒe = No/de. If the exit pupil is fixed at 4 mm, then the largest relative aperture is around N = 10, because the focal lengths in commercially available wide field or super wide field eyepiece series are usually no longer than 40 mm.

Although not included in the original RFT concept, optical aberrations can become a practical constraint on RFT design and may suggest starting the design specification with an exit pupil smaller than the eye pupil. Eyepiece focal lengths should not be so long as to introduce coma, astigmatism or spherical aberration of the exit pupil at low powers or in super wide field optical designs. If (to avoid aberrations) eyepieces are used with focal lengths in the 20 mm to 25 mm range, then we'd require telescope relative apertures of between 3.5 to 4. On the other hand, the objective must not introduce excessive coma or chromatic aberration, which suggests a relative aperture of 5 or 6, which requires eyepieces in the 30 mm to 36 mm range. These complications are normally ignored in an RFT specification, although balancing the optical quality of eyepiece and objective usually suggests a relative aperture of 4 to 6 for eyepieces in the 16 mm to 25 mm range.

Given these considerations, what is the optimal combination of aperture and magnification (exit pupil) for the objective, and of the relative aperture and eyepiece focal length? To answer that question some measure of the system performance is required.

S.L. Walkden, in the original RFT paper (1936), proposed using an estimate of the number of stars per square degree of sky visible with a given aperture. This star count density is calculated as the magnitude limit of the aperture divided by the squared magnification that is delivered by the exit pupil at that aperture. Given a specific aperture Do and exit pupil (de) measured in millimeters, the objective magnitude limit mD, aperture star count per unit area of sky (KD) and magnification adjusted star count RFTD can be estimated as follows:

mD = 2.8+5*log(Do)

KD = 10(–0.0003*mD3+0.0019mD2+0.484mD–3.82)

Mt = Do/de


With this metric one finds that the star count peaks at around D = 70 mm (diagram above, blue curve). This is also Walkden's original conclusion, and illustrates his intent to find a unique advantage for small apertures in an era of giant telescopes. However the range of high values is approximately 40mm to 120mm, which suggests the optimal RFT telescope is actually a large aperture, low power binocular — especially as a binocular view is usually judged to be about ~0.4 magnitude brighter than a monocular view at the same aperture and magnification.

Other authors, for example Glenn Shaw (Sky & Telescope March, 1980), use empirical star counts in different parts of the sky, and this method gives widely ranging results, including apertures above 200 mm. This demonstrates that a star count metric by itself not determinative. Star count is a heuristic that leads to the concept of an RFT, but does not usefully define the parameters of an RFT.

However, scaling the RFT aperture on a rudimentary luminance weighted star count formula

RFTD = mD·KD/Mt2

produces a gently peaking curve at 150 mm with a broad maximum between 90 mm to 260 mm (~3.5" to ~10"; diagram above, yellow curve); these apertures deliver a large true field of view between 6.2° to 2.2° with modern super wide angle (80° AFOV) eyepieces. Note that the size of the exit pupil does not affect the shape or placement of these two curves.

This finding is consistent with Walkden's admission that a 7" aperture might be more suitable as an RFT, because "with the larger telescopes, the proportion of strikingly bright looking stars increases rather rapidly and the proportion of faint looking ones decreases rather rapidly; so that except in actual number of stars to be counted, the fields of the larger telescopes generally look richer, and decidedly more splendid."

These sentiments match the recommendation by Robert Ayers that the optimal RFT refractor has an aperture of 80 mm to 200 mm with relative apertures between N = 4 to 6. There is a trivial change in the luminance weighted star counts across this aperture range, which verifies Shaw's observation that he could see no difference between 6" and 8" RFT scopes. And as a historical note, E.E. Barnard used a 250 mm, ƒ/5 Petzval or "portrait" lens for his high resolution, wide field and splendidly rich photographs of the Milky Way.

Given that the longest focal length offered in most commercial eyepiece lines is 40 mm (which gives a ~7 mm exit pupil at N = 6) and that optical aberrations become severe at relative apertures smaller than N = 4 or (at the same relative aperture) in larger apertures, we can reorder the terms and restate the abstract RFT specification as follows:

an RFT is a large aperture binocular or modest size telescope used for visual astronomy; it must have a sufficiently small aperture (Do <260 mm) and short focal length to deliver a low magnification and aberration free image through a large (~4 mm) exit pupil and low power (ƒe = 20 mm to 30 mm) wide angle or super wide angle (AFOV >70°) eyepiece.

Note that this specification replaces a bracket of relative apertures with a range of eyepiece focal lengths, as the two are equivalent ways to describe an RFT.

Ayers offers cogent arguments in favor of choosing a refractor over a reflector for an RFT although, as already pointed out, large aperture, low magnification binoculars (D = 60 to 100 mm) are highly suitable for very wide field views. If one is interested to use the fastest feasible ƒ ratio, a large aperture (~250 mm) Newtonian reflector with coma corrector is an attractive alternative — disregarding the vignetting effect of the central obstruction, which is rarely noticeable to the visual astronomer.

Effect of Central Obstruction

Resolution is affected by wave diffraction effects, but so too is the potential of the objective to transmit contrast information at different spatial frequencies. Transmission is specifically affected by a central obstruction, such as the secondary mirror in a Newtonian or Cassegrain design, which creates an annular rather than circular aperture geometry and diffraction at four edges on the diameter rather than just two. This obstruction ratio is defined as:

obstruction ratio (η) = Ds / Do

where Ds is the diameter of the secondary obstruction, the mounting for the secondary mirror.

The standard method to evaluate optical performance across a range of spatial frequencies is the modulation transfer function or MTF. The MTF shows the proportion of luminance contrast that is retained in the image formed by an optical system as the spatial frequency of the luminance variations is increased. Geometrically, the MTF calculation is equivalent to the area of overlap between two aperture superimposed at a displacement equal to the spatial frequency.

The stimulus used is either a pattern of crisp black and white lines (to evaluate the transmission of sharp edges) or a sine wave pattern of modulated grays (to evaluate the transmission of tonal gradients). Both methods converge on similar conclusions.

Modulation transfer is defined as the reduction in image contrast between the range of the luminance in the image (LI) as a proportion of the range of luminance in the target or stimulus (LS):

contrast ratio = (LImaxLImin) / (LSmaxLSmin)

Because diffraction produces a tiny amount of smearing in the image, contrast is imperceptibly reduced at low spatial frequencies (wide spacing) but becomes severe at high spatial frequencies (very narrow spacing). So the MTF is scaled to zoom in on the high spatial frequencies, defined as the inverse of the angular spacing in multiples of the minimum theoretically resolvable angular spacing (λ/Dmm). For example, the Rayleigh criterion of resolution is defined as 1.22λ/D, so its spatial frequency in an MTF is 1/1.22, or 0.82. (The MTF is sometimes graphed in terms of line pairs per millimeter (lp/mm) for a specific telescope relative aperture, but this produces the same focus on high frequency spacing and the compression of low frequency spacing at the far left of the graph.)

The green line in the generic MTF diagram (right) shows the curve for an optically perfect and unobstructed circular aperture transmitting a sine wave pattern with a spacing frequency up to the resolution limit — which in a 250 mm ƒ/10 telescope is a stimulus angular separation of 0.45 arcseconds or a spacing on the image plane of 0.0055 mm or 182 line pairs per millimeter.

The yellow line shows the effect on the contrast ratio of a 25% obstruction, typically the upper limit in an ƒ/8 Newtonian reflector, and a 50% obstruction, the upper limit in a Cassegrain design optimized for astrophotography (Ritchey Chrétien or Schmidt Cassegrain). Both patterns can be summarized (reading the graph left to right) as (1) a negligible (less than 10%) contrast reduction in spatial frequencies from the full image diameter down to about 10 times the Rayleigh limit, (2) a significant drop in resolution down to about 2 or 3 times the Rayleigh limit, (3) an increase in resolution that can exceed an unobstructed aperture by a small amount (due to enhanced contrast caused by overlapping wave reinforcement), and (4) resolution matching the theoretical minimum at the Rayleigh limit.

Using the 50% obstruction curve as an example, the greatest relative loss of contrast occurs between 2 and 3 times the minimum resolution, or an angular width of about 1.5 arcseconds. This means the most severe image degradation will occur in lunar features that are about 3 kilometers wide, but will produce a slight crispening in features half that size. It will most affect features that are about 3% of the width of Jupiter's disk at opposition or about 6% the width of Mars's disk at opposition. Contrast on the Galilean satellites of Jupiter and on the disks of Mercury, Uranus and Neptune would be strongly obscured or obliterated. Double stars separated by about 1.3 arcseconds would be most compromised, but very close pairs would be easier to separate.

Whether and how much a central obstruction matters depends somewhat on the size of the aperture, but more on the specific task the telescope is used to perform. Astrophotographers can be tolerant of a large central obstruction because they typically work on an image scale of 1 pixel = 5x the resolution limit, and image contrast can be enhanced with image stacking and digital manipulation.

For visual astronomers the problem is more complex. Inherently high contrast targets such as the moon and double stars will display relatively minor contrast reduction, and the effect will depend on magnification, because human contrast sensitivity peaks at around 7 cycles per degree (around 10 arcminutes) and declines as frequency increases. At the other extreme, low contrast and dim targets such as deep sky objects will be minimally affected because the fovea cannot function at their low luminances. This leaves planetary observation and close double stars widely different in magnitude as the targets where the size of the central obstruction can be a critical issue, one reason that Newtonian or refractor telescopes are often preferred by planetary and double star observers.

Optics of the Eye

The eye is an equal component with objective and eyepiece in the performance of an optical system. The critical performance dimensions for astronomy are dark adaptation, hue discrimination, and resolution as a function of stimulus brightness and stimulus/background contrast. However I begin with a topic that is customarily minimized or ignored in discussions of visual performance: individual differences.

Individual Differences

Human vision is generally characterized in three ways: by describing the physical structure and theoretical optics of the eye; by calculating the optical performance of "ideal" or mathematical eyes; and by measuring individual performance in a variety of visual tasks.

Visual performance is measured using tasks of detection (the stimulus is declared to be visible or not visible), matching (two stimuli are declared to be or are adjusted by the observer until they appear to be the same), discrimination (similar stimuli can be identified as different in some way) or resolution (two separate stimuli can be perceived as separate rather than as a single stimulus).

In all cases, statistical theory is required to summarize performance across individuals. Stimulus thresholds in detection or discrimination are defined as the stimulus value that produces a 50% or 95% probability that an average person will respond with the necessary detection or discrimination. Matching stimuli are defined by the average and standard error of stimulus values judged to be the same, or the average just noticeable difference (JND) that produces detection that they are different. These measurements are then used to create models of average, typical or "normal" human performance.

All descriptions of average performance deviate from actual performance in one important respect: they fail to indicate the extent to which individual performance will vary from the average. These individual differences appear in all visual tasks, however they are measured. Peter Kaiser and Robert Boynton describe the situation this way: "Comparisons between different observers, whether in the same or a different experiment, present a discouraging picture. Although observers agree on certain major trends, individual differences are best described as enormous."

Standard psychometric methods disguise these typically "enormous" individual differences in perceptual thresholds of detection, matching, discrimination and resolution. These vary with the type of task used to measure visual acuity, on the specific way in which the visual stimulus is created and presented, and most importantly, on the perceptual and cognitive attributes of the observer, including past experience with the task — it is possible to learn or practice your way to better visual acuity.

Two examples will illustrate the scale of the problems. On the subject of star color, the first diagram (below left) shows the location on a standard hue circle of the range (variation) in the single wavelengths chosen as a "pure" or best shade of red, yellow, green and blue light by subjects with normal color vision; the inset shows the number of individual subjects (circles) who chose a specific wavelength in the "green" matching task. On the subject of calculated telescope resolution criteria, the second diagram (below right) shows the variation in the judged brightness of wavelengths across the visible spectrum; the "peak" wavelength, which is used to define aperture resolution criteria, varies across a 40 nm range from about 530 to 570 nm.

In the color matching task, there seems to be good agreement among subjects in the locus of a "pure" yellow wavelength — in fact, this constitutes the Nagal anomaloscope test for red/green color blindness. But in the green and blue color matching the performance can best be described as scattered: the green values chosen by different subjects range from a "green blue" below 500 nm to a "green yellow" near 560 nm.

In terms of photopic luminance, individual peak sensitivity spans a range of at least 40 nm (from 535 to 575 nm). The commonly used value of 550 nm is near the peak of the 1964 CIE V* photopic luminosity function (orange curve) or a similar function, which are actually not defined by tests of brightness matching but by a weighted average of the L and M cone response curves measured with color mixture matching.

If you are curious about your visual performance as a visual astronomer, it will be very helpful to discuss testing for color normal vision, resolution and contrast sensitivity with your ophthalmologist. It is easy to measure pupil aperture yourself, when your eye is fully dark adapted. This information will provide an empirical basis for your assessment of telescopic optical performance.

Illuminance & Luminance

For the purposes of describing a light stimulus, the key distinctions are between illuminance and luminance.

Illuminance (E) is the quantity of light incident on a unit area of surface or passing through a unit area of aperture, defined as:

E = LI/d2

where LI is the luminous intensity of the light (the quantity of light radiated into a solid angle of 1 steradian, in lumens) and d is the distance in from the light source. When distance is in meters and the unit area is a 1 meter square surface or 113 cm diameter circular aperture, illuminance is measured in lux.

Luminance is the quantity of light emitted by a surface onto a surface, defined as:

L = E/s

where s is the solid angle of the light source (in steradians or square radians) as observed from the receiving surface or aperture. When both the light source surface and the receptor surface are standardized as 1 meter square, luminance is measured in candelas per square meter.

In optical systems, the objective focal length ƒo and aperture area D2 and reciprocal relative aperture govern the relationship between light grasp and image luminance:

Lo = D2E/ƒo2 = E/N2

As explained above, changes in relative aperture produce squared changes in image luminance.

The key distinctions between illuminance and luminance are:

• Illuminance is completely invisible to the eye; all light perception is a form of imaged luminance, perceived either as an emitting light source or a reflecting material. All brightness and lightness perceptions are luminance perceptions.

• Illuminance decreases as the light source is at a larger distance from the receiving surface. Luminance decreases as the same quantity of light is emitted by a physically larger source. Illuminance and luminance both decrease as the same amount of light is incident on a larger surface area.

• Illuminance provides no information about the size, distance or intensity of the light source (the same quantity of illuminance can come from a nearby dim light or a distant bright light). Luminance is a specific measure of the visual geometry of the illuminance from a light source in relation to a light receptor; it depends both on the angular area (surface area and distance) of the light source and on the area of the light receptor (e.g., the diameter of the pupil), but is invariant with distance because a light source that decreases in illuminance with distance decreases in the same inverse square proportion in angular area.

• Illuminance decreases with increasing distance between the light source and receiving surface. Luminance remains constant regardless of the distance between the emitting light and receptor surface — the angular area of the source decreases with distance in the same inverse square proportion that the illuminance from the source decreases.

• The power or light concentrating capability of an optical system is proportional to D2/ƒ2 = N2; a 2" aperture astrograph operating at ƒ/4 delivers more concentrated light to a photoreceptor than an ƒ/10 instrument of 20" aperture. Regardless of aperture, the luminance of an optical image is equal to the luminance of the source.

Illuminance and luminance are measurement units in photometry, which is limited to electromagnetic radiation that produces a response in the eye, usually considered to be wavelengths between 750 to 380 nm. Irradiance and radiance are the equivalent measurement units in radiometry, which includes all electromagnetic energy from radio waves through xrays. Thus, a star that radiates very powerfully but almost entirely in either the infrared or ultraviolet will appear visually faint, because little of the energy from the star can stimulate the eye.

Luminance Adaptation

Vision can be characterized by three adaptation states, each associated with a range of environmental incident light (illuminance) levels:

• Scotopic. Only the rods respond to light, generally at illuminance levels below 0.1 lux. Several visual criteria define this adaptation, chief among them is the loss of color sensation, although memory color may intrude to tint familiar objects near the scotopic threshold.

• Photopic. The cones predominate in the response to light, which nominally occurs at illuminance levels above 10 lux or more. In fact, the rods are still fully functioning at illuminance levels even 10 times higher, up to light levels typical of noon sunlight. Photopic adaptation characterizes the state in which color perception can be entirely described in terms of the cone sensitivities alone, and any increase in illuminance will not improve visual acuity (resolution of small angular intervals).

• Mesopic. The transition adaptation between scotopic and photopic vision. Color perception is present, but color saturation and contrast are muted, and hues are somewhat biased by the rods.

Light adaptation is the response of the eye to increased luminance levels. It is generally prompt, occurring within a fraction of a second at the pupil and within less than a minute in the retina (via photopigment bleaching) and brain.

Dark adaptation is the response of the eye to decreased luminance levels. It occurs within 2 minutes or more, depending on the light exposure prior to darkness.

The astronomer is more interested in the dynamics of dark adaptation and the corresponding sensitivity of the eye to faint light in darkness. The common lore is that dark adaptation is best preserved by using a dim red light source in darkness, and this is borne out by dozens of research studies. There is no colored light source that is comparable to extreme red light, and most are equivalent to a white light. The main alternatives are red or white: white should be preferred only when it is necessary to perceive differences in color, for example in equipment electrical wiring or printed star charts.

Studies of the time required to recover from exposure to red or white light consistently show an advantage for red light, which becomes disproportionately greater if the light exposure is brighter than about 65 lumens (the intensity of a desk lamp). All light sources induce a greater recovery time (return to dark adaptation) as they become brighter, so regardless of lighting choice the first principle is to work in as dim a light environment as practicable.

It is sometimes claimed that full dark adaptation (recovery of the maximum sensitivity to low light levels) requires hours, days or even weeks to occur. This is entirely false. The photopigment rhodopsin is regenerated in a light shielded eye at the rate of 0.25% per second of the residual (unregenerated) quantity of bleached photopigment. This means, if 100% of the photopigment is bleached by exposure to an extremely bright light, that it would take 77 minutes for 99.999% of the bleached photopigment to be regenerated. This is effectively complete darkness adaptation. Performance measurements of visual sensitivity or acuity (diagrams, above, from Luria & Kobus, 1984) clearly indicate that dark adaptation after exposure to a relatively dim light source (not brighter than a desk lamp) occurs within about 10 minutes for white light and 4 minutes for red.

Again, a dim light level is as critical as a red light environment. Hecht & Hsia (1945) found that adaptation from a white light environment took 15 minutes (!) longer than adaptation from a red light environment if both lights illuminated at 3500 lux, but the white adaptation took only 2 minutes longer if both lights were at 30 lux. Hulbert (1951) measured the difference to adapt from 1070 lux at 14 minutes, and from 11 lux at 5 minutes. As most commercial incandescent "red" light bulbs actually emit a significant amount of "white" (broadband) light, the importance of the dimmest practicable light levels is obvious.

However, the astronomer does not reach the potential limit of dark adaptation, because the ambient "darkness" is typically much brighter. The physiological dark adapted threshold to light is around 0.000001 lux (reflected from a white surface), but the ambient illuminance from a "dark site" starry sky is about 0.001 lux and the light from a full moon is around 0.1 lux, five orders of magnitude greater than the absolute threshold. Suburban or periurban light pollution can also increase the sky brightness of a starry, moonless sky and raise the threshold that must be reached, reducing the recovery time needed to adapt to it.

Some research reports disparage as insignificant the difference in time required to recover from white versus red light exposure. The criterion used for this judgment is usually the time required to adapt only once from a subdued red or white light environment to complete darkness. The astronomer, alternately working at the eyepiece and with charts, guides, computer displays and notetaking or drawing materials, may make the transition from an illuminated to a dark visual environment several dozen times in a single evening. If the difference in adaptation between white or red light regimes amounts to only two minutes, this would result in a cumulative difference in adaptation times between red and white lighting of a half hour to an hour or more. This repetitive adaptation means that the most desirable lighting regime for the visual astronomer is the dimmest practicable red light, obtained solely from LED light sources, sufficient to perform the necessary visual tasks.

Light Grasp

The light grasp of the eye as a function of aperture is quite modest. At a dark adapted pupil diameter of about 6 mm, the pupil area allows only 3.5% the light grasp of a 150 mm (6") telescope. However the eye is capable, under optimal conditions, of perceiving a light source emitting only a few photons per second. Thus, what it lacks in aperture it makes up in retinal sensitivity.


The minimum angular interval or linear dimension that can be resolved by a normal human eye under optimal viewing conditions varies with the type of resolution required.

The angular criterion is usually given as Rv = 60 arcseconds or 0.00029 radians for the resolution of two point light sources, and 120 arcseconds or 0.00058 radians for the resolution of parallel lines. A study by M.A. Danjon (1928) gave 74 arcseconds resolution for high (90%) contrast detail and 127 arcseconds for low (2.5%) contrast detail, which equals linear resolution criteria between 0.0897 to 0.1539 mm. The value 120 arcseconds seems to be commonly adopted, but resolution can clearly be better than that. My own resolution limit for the high contrast interference bands in a λ/D diffraction artifact is around 100 arcseconds.

The diagram (right) describes the eye's optical resolution in the abstract, as the spot diameter of a point light source on the retina as a function of pupil diameter, based on calculations with the Navarro ideal or mathematical "eye".

The Navarro eye allows the spot diameter to be separated into the contributions from the size of the diffraction artifact (the Airy disk, which gets smaller with larger pupil aperture) and from optical aberrations in the cornea and lens (geometric spreading, which increases with an eye pupil larger than ~3 mm). At the pupil sizes characteristic of most visual astronomy, the spot diameter in a normal eye is generally well above 25 μm (0.025 mm), which means the resolution of dim extended objects with an emmetropic, dark adapted eye (~6 mm pupil diameter) is roughly four times worse than the photopic resolution, in the range of 240 to 480 arcseconds (0.0012 to 0.0023 radians), and may be twice that when averted vision is used. (Note that a pupil diameter of around 3.5 mm would still apply to most lunar and some planetary observing in larger apertures, and that resolution values will generally be lower for older and visually impaired observers.)

The common method for estimating the eye's peak visual acuity is to calculate the average spacing between two cones in the fovea as the diameter of a single cone, and to assume that visual resolution is equal to this average spacing (the gap between two point stimuli is as wide as one cone). The average cone density is estimated to be around 190,000 per mm2 in the foveola (diameter on the retina of about 0.35 mm), and about 100,000 per mm2 in the fovea (diameter of about 1.85 mm). This works out to an average cone diameter of about 0.0025 mm in the foveola and 0.0034 mm in the fovea, given the characteristic hexagonal "tiling" of cone aperture.

With the eye's average internal focal length of 17 mm, these widths define an angular separation of between 30 to 41 arcseconds. However, the minimum spot diameter of the eye under optimal conditions (diagram, above) is about 7 μm or 0.007 mm, which is roughly the diameter of three cones. Thus the cone spacing method for predicting visual acuity overstates the eye's detection threshold under optimal conditions by at least two times — and this is consistent with the measured performance acuity of the normal eye under bright light, which is between 1 and 2 arcminutes across the fovea. Thus, the limitation on the eye's visual acuity is primarily optical, not sensory, and optical performance strongly depends on the pupil aperture. Calculations based on retinal tiling are also imprecise because neural networks perform aliasing of the image that can actually exceed the theoretical resolution limit of the eye's optics.

Finally, the task dependence of visual acuity is reflected in the different discrimination thresholds that characterize different monocular visual tasks:

• Shape recognition - Black on white letter shapes can be resolved (e.g., rotational orientations of the shape "E" can be identified) when elements of the letter are at least 5 arcminutes wide.

• Grating resolution - Black on white line gratings (the lines and the spaces between them are of equal width) can be discriminated from a uniform gray background when the spacing between black lines is at least 2 arcminutes.

• Point recognition - Two black points on a white background can be resolved when the space between them is at least 1 arcminute; white points on a black background can be discriminated at about 2 arcminutes (due to spreading of the image).

• Vernier (Nonius) acuity - The misalignment between two black lines placed end to end or the vertical misalignment of two dots can be perceived when the lateral displacement is at least 10 arcseconds. This form of acuity is the principle behind the heliometer and the reason why it was the tool that first measured stellar parallax (Bessel and 61 Cygni).

Contrast Sensitivity

Contrast or modulation range

The troland is a measure of retinal illuminance equal to the illumination from a surface emitting 1 lux per square meter into a pupil area of one square millimeter; it increases both with an increase in source brightness and an increase in pupil aperture. For comparison, 1000 trolands is approximately the brightness of a "white" paper viewed under bright indoor lighting, and 0.1 trolands is approximately the brightness of white paper viewed under moonlight.

This degradation of visual acuity with increasing pupil aperture is clearly indicated in the modulation transfer functions of normal eyes at various pupil apertures, measured by photographing the spot size actually imaged on the retina of subjects (diagram, right). For low contrast targets such as nebulae or faint planetary detail, the limit imposed by a large pupil

Most "deep sky" astronomical targets present retinal illuminances of less than 0.1 trolands, which implies a visual acuity limited at around 2 to 5 cycles per degree or Rv = 360 to 900 arcseconds.

Field of View

Because the eye's normal point is the standard for visual magnification, calculation shows that the eye's true field of view is equal to its apparent field of view.

As a simple test of your true field of view: stand and look straight ahead with both eyes, then extend your arms on either side, fingers facing forward. Wiggle your fingers while moving both your arms forward and backward until you find the point where the movement of the fingers of both hands is just visible on both sides of your visual field. The angle between your arms will be about 120°.

However, the fact that your fingers are barely visible at 120° means that your useable field is much smaller. How small? At the minimum, the fovea is only able to resolve the central ~1° of your visual field. Perceptual research also shows that viewers will say a single object fills their visual field when it subtends about a 20° visual angle — a generalization first made by Leonardo da Vinci. For example, the industry recommended wide screen TV viewing distances produce a visual angle to the wide screen diagonal of between 20° to 40°, with the average around 25° — an image width that is said to produce "an immersive feel". By that standard, the 40° apparent field of traditional eyepieces would be fully immersive.

Currently available commercial eyepieces offer apparent fields of view between 40° to 100°. The choice of one end of that range over the other seems to tip at an apparent field of around 60°, and can depend on several considerations: esthetic preference (the "majesty factor" as opposed to "looking through a drinking straw"), visual task (sweeping for comets as opposed to resolving double stars), the telescope mounting and balance sensitivity (super wide field eyepieces can be very heavy, especially at long eyepiece focal lengths; super wide field eyepieces allow extended "drift" viewing of objects with hand guided Dobsonian altazimuthal mountings), objective relative aperture (off axis eyepiece aberrations are more noticeable at smaller focal ratios) and price (super wide field eyepieces are generally much more expensive than traditional designs).

Off axis image quality, in telescopes with focal ratios down to about ƒ/6, can be superior in some super wide angle eyepieces than in some traditional eyepieces, and image quality at the center of the field (on axis) is, for almost all oculars operating at all focal ratios, optically perfect. Only at objective focal ratios below ƒ/5 do optical considerations become important.

On the other hand, super wide field eyepieces require the observer to "look around" in the wide available field in order to see clearly, and this movement of the eye and head can cause vignetting or aberrations to appear. (Bias disclosure: I prefer eyepieces with apparent fields no greater than 70°, mainly to avoid these effects.)

In general, preference for wide or "narrow" apparent field eyepieces is not an issue of optical quality or the "immersive" visual area useful to the steady gaze of the human eye, but depends on the astronomer's personal preferences in visual esthetics, ergonomics and visual tasks, the convenience of use with a specific type of telescope ... and price.

Focus Diopters

Diopters are used to measure the distance that a lens must be moved forward or backward along the optical axis to produce a change in the image focus equal to one diopter of the eye. This focus diopter is found as:

ƒe2/1000 mm.

Thus the movement required to adjust the focus of a 20 mm eyepiece by one diopter is (202)/1000, or ±0.4 mm.

Focus diopters may be positive or negative. At zero focus diopter, an eyepiece projects collimated light rays of light parallel to the optical axis, which can then focused by a relaxed normal eye as if the light came from a far distant source. Therefore the optimal focus diopter for a normal eye is zero or very slightly negative.

Further Reading

Astronomical Optics, Part 1: Basic Optics - an overview of basic optics.

Astronomical Optics, Part 2: Telescope & Eyepiece Combined - the design parameters of astronomical telescopes and eyepieces, separately and combined.

Astronomical Optics, Part 4: Optical Aberrations - an in depth review of optical aberrations in astronomical optics.

Astronomical Optics, Part 5: Eyepiece Designs - an illustrated overview of historically important eyepiece designs.

Astronomical Optics, Part 6: Evaluating Eyepieces - methods to test eyepieces, and results from my collection.

Amateur Astronomer's Handbook by J.B. Sidgwick - excellent basic chapters in optics, light gathering, resolution, magnification and more.

Telescopic Limiting Magnitudes by Bradley Schaefer - an attempt to predict telescopic limit magnitudes using visual data, star color, age of observer, sky brightness and other factors.

Human Eye from Handbook of Optical Systems, Vol. 4: Survey of Optical Instruments by Herbert Gross, Fritz Blechinger & Bertram Achtner (eds.). (Berlin, DR: Wiley-VCH, 2008).

SECTION 17. VISION from Bioastronautics Data Book NASA SP-3006 edited by Paul Webb. (Washington, DC: National Aeronautics & Space Administration, 1964).

Visual Acuity - Summary of the various methods used to test the human visus.

Average optical performance of the human eye as a function of age in a normal population by A Guirao, C González, M Redondo, E Geraghty, S Norrby and P Artal. Investigative Ophthalmology & Visual Science Jan. 1999, pp. 203-213.

The N.A.A. Telescope Calculator - handy and accurate calculator page for most optical parameters of a telescope/eyepiece combination.

Eyepiece Focal Length Measurement - Jim Easterbrook explains how to measure eyepiece focal lengths in grating projections.

The Collimation - an excellent optical discussion by Thierry Legault

Notes on LEDs - technical overview of the observatory light source most desirable for astronomers.

The Relative Effectiveness of Red and White Light for Subsequent Dark Adaptation (1984) - a useful literature review of the different effects of red and white lights on dark adaptation.

What Is a MTF Curve? - Basics of the MTF.


Last revised 12/05/15 • ©2015 Bruce MacEvoy