### 4.2 The timing model

Ideally, in order to model the rotational behaviour of the neutron star, we require TOAs
measured by an inertial observer. An observatory located on Earth experiences accelerations with
respect to the neutron star due to the Earth’s rotation and orbital motion around the Sun and is
therefore not in an inertial frame. To a very good approximation, the solar system centre-of-mass
(barycentre) can be regarded as an inertial frame. It is now standard practice [130] to transform the
observed TOAs to this frame using a planetary ephemeris such as the JPL DE405 [299]. The
transformation is summarised as the difference between barycentric () and observed () TOAs:
where is the position of the observatory with respect to the barycentre, is a unit vector in the
direction of the pulsar at a distance and is the speed of light. The first term on the right hand side
of Equation (9) is the light travel time from the observatory to the solar system barycentre. Incoming
pulses from all but the nearest pulsars can be approximated by plane wavefronts. The second term,
which represents the delay due to spherical wavefronts, yields the parallax and hence . This
has so far only been measured for five nearby millisecond pulsars [276, 320, 174, 287]. The
term represents the Einstein and Shapiro corrections due to general relativistic time
delays in the solar system [18]. Since measurements can be carried out at different observing
frequencies with different dispersive delays, TOAs are generally referred to the equivalent time
that would be observed at infinite frequency. This transformation is the term (see also
Equation (1)).
Following the accumulation of a number of TOAs, a surprisingly simple model is usually
sufficient to account for the TOAs during the time span of the observations and to predict the
arrival times of subsequent pulses. The model is a Taylor expansion of the rotational frequency
about a model value at some reference epoch . The model pulse phase is

where is the barycentric time and is the pulse phase at . Based on this simple
model, and using initial estimates of the position, dispersion measure and pulse period, a “timing
residual” is calculated for each TOA as the difference between the observed and predicted pulse
phases.
A set of timing residuals for the nearby pulsar B1133+16 spanning almost 10 years is shown in
Figure 22. Ideally, the residuals should have a zero mean and be free from any systematic trends (see
Panel a of Figure 22). To reach this point, however, the model needs to be refined in a bootstrap fashion.
Early sets of residuals will exhibit a number of trends indicating a systematic error in one or more of the
model parameters, or a parameter not incorporated into the model.

From Equation (10), an error in the assumed results in a linear slope with time. A parabolic trend
results from an error in (see Panel b of Figure 22). Additional effects will arise if the
assumed position of the pulsar (the unit vector in Equation (9)) used in the barycentric time
calculation is incorrect. A position error results in an annual sinusoid (see Panel c of Figure 22). A
proper motion produces an annual sinusoid of linearly increasing magnitude (see Panel d of
Figure 22).
After a number of iterations, and with the benefit of a modicum of experience, it is possible to identify
and account for each of these various effects to produce a “timing solution” which is phase coherent over the
whole data span. The resulting model parameters provide spin and astrometric information with a precision
which improves as the length of the data span increases. Timing observations of the original millisecond
pulsar B1937+21, spanning almost 9 years (exactly 165,711,423,279 rotations!), measure a period of
[155, 152] defined at midnight UT on December 5 1988!
Astrometric measurements are no less impressive, with position errors of being presently
possible [333].