Issue 
J. Space Weather Space Clim.
Volume 10, 2020



Article Number  30  
Number of page(s)  24  
DOI  https://doi.org/10.1051/swsc/2020033  
Published online  16 July 2020 
Research Article
Semiannual, annual and Universal Time variations in the magnetosphere and in geomagnetic activity: 2. Response to solar wind power input and relationships with solar wind dynamic pressure and magnetospheric flux transport
^{1}
Department of Meteorology, University of Reading, RG6 6BB Reading, UK
^{2}
Institute of Space and Atmospheric Studies, University of Saskatchewan, Saskatoon, SK S7N 5A5, Canada
^{3}
School of Physics and Astronomy, University of Southampton, SO17 1BJ Southampton, UK
^{*} Corresponding author: m.lockwood@reading.ac.uk
Received:
13
September
2019
Accepted:
22
June
2020
This is the second in a series of papers that investigate the semiannual, annual and Universal Time (UT) variations in the magnetosphere. We present a varied collection of empirical results that can be used to constrain theories and modelling of these variations. An initial study of two years’ data on transpolar voltage shows that there is a semiannual variation in magnetospheric flux circulation; however, it is not as large in amplitude as that in geomagnetic activity, consistent with the latter showing a nonlinear (quadratic) variation with transpolar voltage. We find that during the persistent minimum of the UT variation in geomagnetic activity, between about 2 and 10 UT, there is also a persistent decrease in observed transpolar voltage, which may be, in part, caused by a decrease in reconnection voltage in the nightside crosstail current sheet. We study the response of geomagnetic activity to estimated power input into the magnetosphere using interplanetary data from 1995 onwards, an interval for which the data are relatively free of data gaps. We find no consistent variation in the response delay with timeofyear F and, using the optimum lag, we show that the patterns of variation in Fyear spectrograms are very similar for geomagnetic activity and power input into the magnetosphere, both for average values and for the occurrence of large events. The Russell–McPherron (R–M) mechanism is shown to be the central driver of this behaviour. However, the (R–M) effect on power input into the magnetosphere is small and there is a nonlinear amplification of the semiannual variation in the geomagnetic response, such that a very small asymmetry in power input into the magnetosphere P_{α} between the “favourable” and “unfavourable” polarities of the IMF B_{Y} component generates a greatly amplified geomagnetic response. The analysis strongly indicates that this amplification is associated with solar wind dynamic pressure and its role in squeezing the nearEarth tail and so modulating the storage and release of energy extracted from the solar wind. In this paper, we show that the equinoctial pattern is found in the residuals of fits of P_{α} to the am index and that the amplitude of these equinoctial patterns in the am fit residuals increases linearly with solar wind dynamic pressure. Similarly, the UT variation in am is also found in these fit residuals and also increases in amplitude with solar wind dynamic pressure.
Key words: geomagnetic activity / semiannual variation / solar wind magnetosphere coupling
© M. Lockwood et al., Published by EDP Sciences 2020
This is an Open Access article distributed under the terms of the Creative Commons Attribution License (https://creativecommons.org/licenses/by/4.0), which permits unrestricted use, distribution, and reproduction in any medium, provided the original work is properly cited.
1 Introduction
Paper 1 in this series (Lockwood et al., 2020, hereafter “Paper 1”) reviewed the semiannual variation in geomagnetic activity, as seen in a number of indices, and highlighted a number of puzzles that require explanation. We study these in the present paper empirically using the power input into the magnetosphere, P_{α}, deduced from interplanetary measurements using the formula that was originally derived theoretically by Vasyliunas et al. (1982) from dimensional analysis. This formula is given by equation (2) of Paper 1 and the derivation has also recently been given, and expanded upon, by Lockwood (2019). Analysis by Finch & Lockwood (2007) showed that P_{α} performed consistently better than a basket of other widelyused solarwind/magnetosphere coupling functions on all averaging timescales tested (between 1 day and 1 year). An important consideration for all coupling functions was studied by Lockwood et al. (2019a), namely the effect of data gaps. These have often been ignored on the grounds that their effects average out, but this was shown not to be the case by Lockwood et al. who introduced synthetic data gaps into nearcontinuous data. This analysis revealed that data gaps add noise to the correlation studies and lead to erroneous fits. In particular, there is a problem with “overfitting”, whereby the use of too many free fit parameters can generate seemingly good fits to the training data by fitting to the noise, leading to incorrect fits that have reduced, little or even, in extreme cases, zero predictive power when applied to data other than the training dataset. Overfitting is a problem that is well recognized in disciplines such as climate science and population studies but has not often been considered in space physics. A key point about P_{α} is that it has just one free fit parameter, the coupling exponent α, and this minimises the risk of overfitting. At the same time, P_{α} makes allowance for solar wind speed, V_{SW}, number density, N_{SW}, mean ion mass, m_{SW}, and Interplanetary Magnetic Field (IMF) strength, B, as well as an orientation factor, A_{θ}. The test devised by Vasyliunas et al. (1982) shows that the optimum form of A_{θ} is sin^{4}(θ/2), where θ is the IMF clock angle in the GSM frame (see the supporting information file to Lockwood et al., 2019b and Sect. 1.1). These factors could all be given their own fit exponent but this greatly increases the risk of overfitting; to avoid this P_{α} uses the theory to make the exponent of each variable a fixed function of the one free fit parameter, α.
As also discussed in Paper 1, P_{α} is based on the dominant energy flux in the solar wind, namely the bulkflow kinetic energy flux of the particles, whereas some other coupling functions incorrectly use the Poynting flux in the solar wind, which is very small in comparison. The energy input to the magnetosphere across the magnetopause is in the form of Poynting flux, but most of that is generated from the bulkflow solar wind kinetic energy by currents that flow in the bow shock and magnetosheath (Cowley, 1991; Lockwood, 2004, 2019; Pulkkinen et al., 2016). Lockwood (2019) has recently added the minority solar wind Poynting flux to the Vasyliunas et al. (1982) formulation for P_{α} at the expense of adding a second free fit parameter. The improvements to the correlations with geomagnetic activity are sometimes statistically significant but only on short timescales and they are always extremely small. We do not here include this additional term in solar wind Poynting flux because it increases the risk of overfitting.
One particular puzzle about the semiannual variation in geomagnetic activity is the origin of the equinoctial pattern of geomagnetic response with timeofyear, F, and Universal Time, UT (de La Sayette & Berthelier, 1996; Cliver et al., 2000; Chambodut et al., 2013). This has been a matter of debate, with a large number of mechanisms proposed as to how the implied association with Earth’s dipole tilt originates and it is not clear which, if any, of the proposed mechanisms is active (see review in Sect. 1 of Paper 1).
A related, unexplained feature, also highlighted by Paper 1, is that the semiannual variation in geomagnetic activity is undoubtedly caused by the modulation of power input into the magnetosphere P_{α} by the Russell–McPherron (R–M) effect (Russell & McPherron, 1973) but is amplified, being of larger fractional amplitude in most geomagnetic indices than in P_{α}. The need for some such amplification was noted by Russell and McPherron in their original paper but the mechanism responsible for it has never been identified.
There is a third major puzzle concerning large storms. The R–M mechanism is based on the idea that the IMF lies close to its average orientation, parallel to the solar equatorial plane (the XY plane of the Geocentric Solar Ecliptic, GSEQ, reference frame) and that the component in the +Y or −Y direction of GSEQ is converted into “geoeffective” southward field in the Geocentric Solar Magnetospheric (GSM) reference frame by the tilt of Earth’s magnetic moment in the YZ plane. This tilt sets the consequent rotation of the GSM frame with respect to the GSEQ frame about their common Xaxis. The GSEQ and GSM reference frames are organised with respect to the solar rotation axis and the Earth’s magnetic dipole axis, respectively, and formal definition and discussion of both is given in Paper 1. The R–M effect predicts that, through the dipole tilt effects on the rotation between the GSEQ and GSM frames, the IMF field with [B_{Y}]_{GSEQ} < 0 gives southward field in the GSM frame ([B_{Z}]_{GSM} < 0) around the March equinox, whereas [B_{Y}]_{GSEQ} > 0 generates southward field in GSM around the September equinox. Hence, as noted by Zhao & Zong (2012), splitting the data into these two polarities of [B_{Y}]_{GSEQ} allows us to directly identify the contribution of the R–M effect by which of the two equinoxes is enhanced by which polarity of [B_{Y}]_{GSEQ}. We here refer to the [B_{Y}]_{GSEQ} polarity which, at a given equinox, increases/decreases the southward [B_{Z}]_{GSM} (and hence increases/decreases geomagnetic activity) as the “favoured”/“unflavoured” polarity. Paper 1 uses this dependence on the IMF [B_{Y}]_{GSEQ} polarity to reveal that the R–M effect is at the heart of the equinox peaks of several parameters, including: mean power input into the magnetosphere, P_{α}; the mean of the am geomagnetic index; and of the occurrence of both P_{α} and am values in the top 5% of their overall distributions, f[P_{α} > q(0.95)] and f[am > q(0.95)], respectively. Interestingly, this means that the R–M effect is at the heart of even the equinoctiallike timeofday (UT) – timeofyear (F) patterns seen for geomagnetic data, which, as discussed in Paper 1, are distinct from the pattern that is predicted for the R–M effect. Hence, although the equinoctial patterns indicate that another mechanism is at work, the core cause of their semiannual variation is undoubtedly the R–M effect.
The occurrence of large storms generates another puzzle in relation to the R–M effect because it has long been recognized that the largest storms are produced by interplanetary disturbances during which the IMF, at some stage at least, points strongly southward in the GSM frame (e.g., Taylor et al., 1994; Webb et al., 2000; Echer et al., 2005; Xie et al., 2006; Kilpua, et al., 2017; Li & Yao, 2020). Crooker et al. (1992) proposed a mechanism to give the observed increased occurrence of large storms at the equinoxes that is consistent with the R–M effect by arguing that in the CME sheath the IMF remains in the equatorial plane of the GSEQ frame but its component in the ±Y direction is enhanced by the compression due to the event. This enhanced equatorial [B_{Y}]_{GSEQ} is then converted into enhanced southward field in the GSM frame by the R–M effect. However, Lockwood et al. (2016) and Paper 1 show unambiguously that large southward field in GSM is almost always associated with large southward field in GSEQ, and not the nearzero southward field that Russell & McPherron (1973) and Crooker et al. (1992) invoked in applying the R–M mechanism. Appendix B of Paper 1 shows that when there is a dominant southward component of the IMF in the GSEQ frame, far from enhancing the geoeffective southward component in GSM, the R–M effect actually reduces it. Paper 1 therefore highlights a paradox in that although the R–M effect is clearly at work to enhance average geomagnetic disturbance levels and the occurrence of small and moderate geomagnetic disturbances, large disturbances are driven by outofequatorial southward field in GSEQ (usually deflected out of the solar equatorial plane by a CME or CIR) for which the R–M effect actually tends to reduce the southward component in GSM and hence the geoeffectiveness of the event. Hence the observed equinoctial peaks in the occurrence of large storms are not attributable to the R–M effect, at least not directly, in the same way that the peaks for smaller disturbances and average conditions are.
This point is further emphasised here by Figure 1, which is an update of the plot by Lockwood et al. (2016). It compares the northward component of the IMF in the GSEQ frame ([B_{Z}]_{GSEQ}) to that in the GSM frame ([B_{Z}]_{GSM}), both averaged over 3h intervals that 1 h ahead of the 3h intervals in which the geomagnetic activity level is measured (here quantified by the am index). In other words, we have employed an average response lag of the am index to the IMF orientation of dt_{o} = 1 h (see Sect. 3.3 below). The plots show the numbers of samples, N, in bins of width 2 nT in both [B_{Z}]_{GSEQ} and [B_{Z}]_{GSM}. The data cover 1995–2017, inclusive, the interval of nearcontinuous IMF observations. Results are shown for: (a) all data; (b) am exceeding its 90percentile, am > q(90) (i.e., am is in the top 10% of all its measured values); (c) am exceeding its 95percentile, am > q(95); and (d) am exceeding its 99percentile, am > q(99). Figure 1 confirms the result presented in Paper 1, that large [B_{Z}]_{GSEQ} < 0 is actually the dominant driver of large southward field in the GSM frame ([B_{Z}]_{GSM} < 0) and the R–M effect gives relatively minor deviations from the diagonal mauve line that is at [B_{Z}]_{GSEQ} = [B_{Z}]_{GSM}. It can be seen that for larger geomagnetic events, the IMF data are increasingly in the [B_{Z}]_{GSEQ} < 0, [B_{Z}]_{GSM} < 0 sector. Note that there are some exceptions to this with a small number values in the [B_{Z}]_{GSEQ} > 0, [B_{Z}]_{GSM} > 0 sector: these can be attributed to the fact that using a 3h prior IMF orientation averaged over 3 h with a lag dt_{o} of 1 h is not always appropriate (see Sect. 3.3). The pure R–M effect invokes [B_{Z}]_{GSEQ} = 0 and so points would lie along the vertical white line only: the small deviations from the diagonal show that the R–M effect is contributing much less to [B_{Z}]_{GSM} < 0 than does [B_{Z}]_{GSEQ} < 0 in most cases. In Figure 1a (for all data), the R–M effect causes both increases and decreases in [B_{Z}]_{GSM} around [B_{Z}]_{GSEQ} = 0, as expected and depending on the polarity of the IMF [B_{Y}]_{GSEQ}. For small, moderate and some large geomagnetic disturbances (Figs. 1b–1d), respectively we see more points below the mauve diagonal line at small [B_{Z}]_{GSEQ} which demonstrates that the R–M effect is having an effect (by making [B_{Z}]_{GSM} more negative during the favourable polarity of [B_{Y}]_{GSEQ} and so enhancing the am index). However, for large disturbances (Fig. 1d) and strongly negative [B_{Z}]_{GSEQ} there are actually more points above the diagonal mauve line than below it, showing the R–M effect is tending to reduce the southward field in GSM in these cases, as predicted in Appendix B of Paper 1.
Fig. 1 The southward interplanetary magnetic field (IMF) in the Geocentric Solar Magnetospheric (GSM) frame, [B_{Z}]_{GSM} as a function of its value in the Geocentric Solar Equatorial (GSEQ) frame, [B_{Z}]_{GSEQ} from a survey of 24 years’ interplanetary and geomagnetic activity data (1995–2017, inclusive). The number of samples in 0.2 nTby0.2 nT bins, N, as a ratio of the total number, ΣN, is colour contoured as a function of [B_{Z}]_{GSM} and [B_{Z}]_{GSEQ}. The data are averages over 3h intervals for 1995–2017, inclusive, which are the intervals over which each am value is compiled, shifted by an average optimum am response lag dt_{o} = 1 h (see Sect. 3.3). (a) Is for all data whereas (b–d) are for am exceeding, respectively, its 90%, 95% and 99% quantile for the years studied. The diagonal mauve lines are [B_{Z}]_{GSEQ} = [B_{Z}]_{GSM}. 
1.1 The paradox of the R–M effect and of southward field in the GSEQ frame
Figure 2 is a demonstration of the important point made in Appendix B of Paper 1. The graphs are taken from the supporting information to Lockwood et al. (2019b), in which their derivation is explained in greater detail. This plot demonstrates the test of IMF orientation factors, A_{θ}, in solarwind magnetosphere coupling functions that was devised by Vasyliunas et al. (1982) and evaluates the various proposed forms for A_{θ}. The black dots are based on the geomagnetic SML index data (the SuperMAG version of the auroral AL index but compiled from a northern hemisphere network of over 100 stations) and show a linear regression of SML/G against θ where G is the bestfit estimate of the power input into the magnetosphere from solar wind data but without an IMF orientation factor (i.e., G = P_{α}/A_{θ}) and θ is the IMF “clock angle” in the GSM frame, θ = arctan ([B_{Y}]_{GSM}/([B_{Z}]_{GSM}). Note that this definition means that θ is independent of the polarity of [B_{Y}]_{GSM} and that θ varies from zero for purely northward IMF in GSM ([B_{Z}]_{GSM} = B_{YZ}, where B_{YZ} is the magnitude of the field in the YZ plane that is the same in the GSEQ, GSE, and GSM reference frames) to 180° for purely southward IMF in GSM ([B_{Z}]_{GSM} = −B_{YZ}). It can be seen from Figure 2 that the optimum fit to the SML data is for A_{θ} = sin^{4}(θ/2) (the blue line). Almost identical plots to Figure 2 for the AE and am geomagnetic indices have been presented in Figure 4 of Bargatze et al. (1986) and Figure 9 of Lockwood (2019), respectively. The original formulation of the R–M effect by Russell & McPherron (1973) used the average IMF orientation with [B_{Z}]_{GSEQ} = 0, for which the rotation between the GSEQ and GSM frames will cause θ to vary over a year (with fraction of the year F and UT) between 90 − 33.5 = 56.5° and 90 + 33.5 = 123.5° (the range shaded light pink in Fig. 2). In addition, the area shaded dark pink shows the range for the annual effect of the Earth’s orbit without the UT effect. In these ranges of θ, the optimum A_{θ} is increasing in gradient (i.e., d^{2}A_{θ}/dθ^{2} > 0). This is true for all the proposed coupling functions shown except sin^{2} (θ/2) (in orange), which we do not consider as it is not a good fit to the data (for all tested geomagnetic indices, the A_{θ} = sin^{4} (θ/2) formulation yields the best fit). The original R–M paper employed a “halfwave rectified” coupling function such as U(θ) cos(θ) (shown in black) where U(θ) = 0 for [B_{Z}]_{GSM} > 0 and U(θ) = 1 for [B_{Z}]_{GSM} ≤ 0 for which the change in slope (with d^{2}A_{θ}/dθ^{2} > 0) all takes place at θ = 90°. The nonlinearity of the A_{θ}(θ) curve around θ = 90° is vital to the R–M explanation of the semiannual variation. This is because the distribution of θ values in the GSM frame will be approximately symmetric around a mode value of 90°. However, the nonlinearity in A_{θ} means that increases in geomagnetic activity for θ > 90° (caused by the “favourable” polarity of [B_{Y}]_{GSEQ} at a given F) will be greater in amplitude than the decreases caused by θ < 90° (caused by the other, “unfavourable” [B_{Y}]_{GSEQ} polarity). Hence the net effect in this case is a rise in average activity and this provides the R–M explanation of the semiannual variation even though the occurrence of favourable [B_{Y}]_{GSEQ} is the same as that of unfavourable [B_{Y}]_{GSEQ}. The same applies for all coupling function that show d^{2}A_{θ}/dθ^{2} > 0 around θ = 90°.
Fig. 2 The test devised by Vasyliunas et al. (1982) of proposed IMF orientation factors, A_{θ}(θ), where θ is the IMF clock angle in the GSM frame: (orange line) sin^{2}(θ/2); (cyan line) sin^{3}(θ/2); (blue line) sin^{4}(θ/2); (mauve line) sin^{5}(θ/2); and (black line) U(θ)cos(θ), where U(θ) = 0 for [B_{Z}]_{GSM} > 0 and U(θ) = 1 for [B_{Z}]_{GSM} ≤ 0. The black dots are s(SML/G) + c where s and c are the bestfit linear regression coefficients, SML is the SuperMAG westward electrojet index, and G = P_{α}/A_{θ}, where P_{α} is the best fit estimate of the power input into the magnetosphere. The dark pink shaded area is the range over which θ varies due to the R–M effect for [B_{Z}]_{GSEQ} = 0 (i.e., IMF lying in the solar equatorial plane) and the range of variation of the dipole tilt over a full year caused by Earth’s rotational axis tilt. The light pink area is for the full range caused by the annual variation plus the diurnal variation due to the offset of Earth’s rotational and magnetic axes. The light and dark orange regions are the corresponding ranges for [B_{Z}]_{GSEQ} = B (i.e., purely southward field normal to the solar equatorial plane). 
However, contrast this to what happens if [B_{Z}]_{GSEQ} is not zero but rather is strongly negative, as is the case for the solar wind transients that drive most major geomagnetic disturbances (as demonstrated by Fig. 1). At θ near 120° all the A_{θ} coupling functions become approximately linear (d^{2}A_{θ}/dθ^{2} ≈ 0). This means that there is no longer an R–M effect because the effects of the “favourable” increases in θ are cancelled by the effects of the “unfavourable” decreases in θ. At the largest θ, d^{2}A_{θ}/dθ^{2} < 0. Here the dark and light orange regions are the ranges in θ caused by the annual and the annualplusUT variations in dipole tilt, respectively, for [B_{Z}]_{GSEQ} = −B_{YZ} (purely southward field normal to the solar equator). In this case, any effect of the dipole tilt reduces the value of A_{θ} and so has the opposite effect to the classical R–M effect (which relies on [B_{Z}]_{GSEQ} being near zero). This point is made in greater detail in Appendix B to Paper 1.
The key point is that the nonlinearity in A_{θ}, upon which the R–M effect depends, is present only for small [B_{Z}]_{GSEQ}. A large [B_{Z}]_{GSEQ} contribution to [B_{Z}]_{GSM} during large geomagnetic events (as shown in Fig. 1) means that the coupling function is far removed from the part of the response curve with the required nonlinearity. This leads to either a reduced effect, no effect or even an effect in the opposite sense to the R–M effect, as it is presently understood (which applies only when [B_{Z}]_{GSEQ} is small). This means that, because of the increased role of nonzero [B_{Z}]_{GSEQ,} the R–M effect on the semiannual variation should be weaker or absent for the occurrence of larger events, the opposite of the behaviour that is observed (as shown in the original paper by Russell & McPherron, 1973 and, for example, by Paper 1). Hence Figures 1 and 2 present a considerable puzzle in terms of understanding the equinox peaks in the occurrence of large storms that are driven by IMF that deflected strongly southward by transient events such as CMEs and CIRs.
1.2 The timing of the peaks of the semiannual variation
Large southward field in the GSEQ frame during large geomagnetic disturbances has another implication. As discussed in Paper 1, many studies attempt to use the precise timing of the equinox peaks in average geomagnetic activity levels, or in the occurrence of storms, as a potential discriminator of the various mechanisms (Russell & McPherron, 1973; Le Mouël et al., 2004). However, the random nature of events of large negative [B_{Z}]_{GSEQ} hitting the Earth’s magnetosphere means that data from a great many years are needed before a systematic peak can be observed such that the precise timing of the peak can be properly assessed (Russell & McPherron, 1973).
1.3 The role of solar wind dynamic pressure
Another factor that the literature indicates that we should consider is solar wind dynamic pressure (, where m_{SW} is the mean ion mass, V_{SW} the number density and V_{SW} the speed of the solar wind). Caan et al. (1973) showed that the magnetic energy density in the nearEarth tail lobes was increased by both prior intervals of southwardpointing IMF (i.e., substorm growth phases) and by increased p_{SW}. Midlatitude “range” indices (such as am) are strongly modulated by the auroral electrojet of the substorm current wedge and so have a strong correlation with indices such as AE, AL, SME and SML. Hence they are influenced by substorm expansion phases when the stored tail lobe energy is released (Adebesin, 2016; Lockwood et al., 2019a, 2019d, Appendix A of Paper 1). These midlatitude range indices also have a strong dependence on and hence on p_{SW} (Lockwood, 2013; Lockwood et al., 2014). Using the standard deviation of geomagnetic variations, Finch et al. (2008) found that the nightside auroral electrojet was the source both of the equinoctial F–UT pattern and of the dependence. Lockwood (2013) pointed out that this implies that p_{SW} influences the auroral electrojet and the equinoctial F–UT pattern by constraining the nearEarth tail such that on appending open flux to the tail lobes of the magnetosphere (during periods of southward IMF) the lobe field (and hence the stored energy density and magnetic shear across the nearEarth crosstail current sheet) increases by a greater factor if p_{SW} is large. This does not happen further down the tail where the magnetopause boundary becomes aligned with the solar wind flow: here the magnetic field, the energy density and the crosstail current, are all set by the static pressure in interplanetary space with no influence of p_{SW}. At these large negative X_{GSE}, adding open flux just causes the tail to flare in crosssectional area and the field in the lobes and the magnetic shear across (and hence total current in) the crosstail current all remain constant for a given solar wind static pressure.
An important effect of p_{SW} was demonstrated directly by Karlsson et al. (2000) who showed that nearEarth tail energy content was reduced if p_{SW} decreased and that such sudden decreases caused quenching of any substorm expansion that had recently begun. Conversely, increases in p_{SW} have been seen to trigger onsets of full substorm expansion phases (Schieldge & Siscoe, 1970; Kokubun et al., 1977; Yue et al., 2010). Various studies suggest that increased p_{SW} enhances general magnetospheric convection and fieldaligned current systems as well as geomagnetic activity (e.g., Lukianova, 2003; Lee et al., 2004; Palmroth et al., 2004; Boudouridis et al., 2005; Stauning & Troshichev, 2008): this is beyond and separate to the known generation of transient filamentary field aligned currents and travelling convection vortices by the boundary deformation (e.g., Lühr et al., 1996). Some of these observations are interpreted as being the result of enhanced magnetopause reconnection; however, in many cases the response delay appears to be too long for this to be the explanation.
1.4 The F–UT response of the am index and the aσ indices
For studies of the F–UT pattern of geomagnetic activity, the am index (Mayaud, 1980) is by far the best geomagnetic index to employ because it is based on data from longitudinal rings of magnetometers in both hemispheres that are as uniform as possible (and deploys weighting functions to reduce the effects of necessary nonuniformity of the rings): this gives the am index an exceptionally uniform F–UT response pattern, especially at higher levels of geomagnetic activity, as demonstrated by the modelling of the index response by Lockwood et al. (2019d). As in Paper 1, we also make use of the aσ indices (Chambodut et al., 2013) which are based on the same data as the am index but restrict the data used to stations that are within 6h magnetic local time (MLT) sectors around dawn (03–09 h MLT, giving aσ_{dawn}), noon (09–15 h MLT, giving aσ_{noon}), dusk (15–21 h MLT, giving aσ_{dusk}), and midnight (21–24 and 00–03 h MLT, giving aσ_{midnight}).
1.5 The aims of the present paper
The main aim of the present paper is to establish some observed behaviours of the semiannual variation. We start with an initial look in Section 2 at the relationship between annual and semiannual variations of magnetospheric flux transport (convection) and geomagnetic activity. This is an initial and interim study of a limited dataset but still reveals some important behaviours. In Section 3 we investigate the geomagnetic response to power input into the magnetosphere, including an analysis of response delays as a function of activity level and time of year, F. In Section 4, we study annual variations in the geomagnetic response to events of large power input to the magnetosphere and to average power input, both with and without the separation by the polarity of the prevailing IMF [B_{Y}]_{GSEQ} component that enables us to identify the R–M effect. We study variations with the solar wind dynamic pressure p_{SW} in Section 5. However, we do not isolate the effects of solar wind mean ion mass m_{SW} number density, N_{SW}, nor speed, V_{SW}, independently and we note that dependencies of p_{SW} detected could arise from a dependence on different combinations of the parameters, potentially leading to different interpretations. The literature described in Section 1.3 strongly suggests that dynamic pressure, as a physical entity, plays a role in magnetospheric responses by squeezing the nearEarth tail; however, the discussion of such effects will be in a later paper that uses magnetopause and magnetospheric models to discuss the mechanisms in the light of the empirical results presented here.
2 Initial survey of semiannual variations in power input into the magnetosphere and its response in geomagnetic activity and magnetospheric flux circulation
In this section, we make use of a database of 20,430 polar cap traversals made by DMSP satellites during 2001 and 2002, analysed using the procedure of Lockwood et al. (2009). We use only passes that cross the dawnside convection reversal boundary at 02–10 MLT and the duskside convection reversal boundary at 14–22 MLT and most of the data come from the DMSPF13 satellite that was in a suitable orbit. Figure 3 shows the semiannual variations in various parameters in fully simultaneous data: in each data sequence, a data gap in any one parameter is introduced into all other data sequences and each case, the data have been averaged into 365 equalsized bins of timeof year F (i.e., of oneday duration for the nonleap years used). This yields between 48 and 63 samples in each 1day bin, with an average of just under 56. A running, boxcar mean was the taken over 27 bins to cover whole solar rotation intervals and give an even mix on the two polarities of the Ycomponent of the IMF. The dawntodusk component, (E_{DD}, in the +Y GSM direction) of the electric field in interplanetary space (), is shown in Figure 3a. This has been halfwave rectified such that all negative (i.e., dusktodawn) E_{DD} values are put to zero. Given that Figure 3 shows averages over 27day intervals, we should expect steadystate to apply in which case, by Faraday’s law, the electric field would be curlfree and map from interplanetary space into the ionosphere down the open polar cap field lines. The AU index (Fig. 3b) shows only a strong annual variation due to the dayside conductivity enhancement in summer, which is also present in AL (Fig. 3c) but to a much lesser extent and the semiannual variation can be seen. The polar cap flux, W_{PC} (Fig. 3d) and dawndusk transpolar voltage Φ_{PC} (Fig. 3e) are derived from the convection reversal boundaries detected by the DMSP satellites using the procedure described by Lockwood et al. (2009): in both cases the black line is for all passes and the red and blue lines are for northern and southern polar cap passes, respectively. The procedure applies a statisticallyderived correction to the observed difference in potential between the dawnside and duskside convection reversal boundaries to normalise to the value Φ_{PC} for an ideal 06–18 MLT orbit. Note that Φ_{PC} is often referred to as the “crosscap potential” but is, in reality, a potential difference and hence a voltage: we here use the term “transpolar voltage”. Note also that, by Faraday’s law, a voltage is physically identical to a magnetic flux transfer rate. The statistical model of Lockwood et al. (2009) also gives the variable polar cap shape which is used with the observed polar cap diameter to generate estimates of the polar cap flux, W_{PC}. All three lines (for northern and southernhemisphere data separately and all data) are similar in Figure 3d, as indeed they should be given that by Maxwell’s equation , the total open flux in the two hemispheres is always identical. The small difference between the independent values for the two hemispheres gives strong support to the method used to compute W_{PC}. In addition, for these 27day averaging intervals, the total rate of flux transfer across the polar caps (the transpolar voltages) should also be very similar, as they are seen to be in Figure 3e. (On shorter timescales nonsteadystate effects can become apparent and the flux transfer rate across the central dawndusk diameters of the two polar caps does not have to be the same at any one instant). Because of the long averaging timescale, the reconnection efficiency can be taken to be the ratio of the voltages across the polar cap and the whole magnetosphere, η_{15} = Φ_{PC}/(2E_{DD} R*) and in Figure 3f is computed assuming that the crosssectional radius of the magnetosphere is constant at R* = 15 R_{E} (where a mean Earth radius 1 R_{E} = 6370 km). The value of η_{15} derived is relatively constant over the year and does not show a semiannual variation. Figure 3g shows the corresponding values of the am index (linearly interpolated from the 3hourly data to the central time of each polar cap pass). The bottom panel (Fig. 3h) plots the ratio am/Φ_{PC} and shows that the semiannual variation in am is amplified compared to that in magnetospheric convection. Figure 3 shows that the average semiannual variations in E_{DD}, Φ_{PC}, W_{PC}, −AL, and am are all very similar in waveform.
Fig. 3 Simultaneous data from the years 2001 and 2002. In each case the data have been averaged into 365 equalsized bins of timeof year F (one day long for these nonleap years) and then a running mean taken over 27 bins to cover whole solar rotation intervals. From top to bottom: (a) the halfwave rectified dawntodusk electric field, E_{DD}, in interplanetary space; (b) the AU index; (c) the AL index; (d) the polar cap flux, W_{PC}, from the DMSP satellite data; (e) the transpolar voltage, Φ_{PC}, from the DMSP satellite data; (f) the reconnection efficiency assuming that the crosssectional radius of the magnetosphere is 15 R_{E} (where a mean Earth radius 1 R_{E} = 6370 km), η_{15}; (g) the am geomagnetic index; and (h) the ratio of the am index to the transpolar voltage, am/Φ_{PC}. In panels (d) and (e) for W_{PC} and Φ_{PC}, the red and blue lines are for the northern and southern polar caps, respectively. 
Figure 4 demonstrates that the amplification of the semiannual variation in am, relative to Φ_{PC}, is a general feature of the am index by plotting daily means of am as a function of daily means of Φ_{PC}. Note that during each day there are typically 25 polar cap satellite passes and there are 8 am measurements that are linearly interpolated to the central times of the passes. The error bars give the standard errors in these daily means. The mauve line is the bestfit 3rdorder polynomial and the grey area is the 2σ uncertainty band in that fit. Appendix details the ensemble fitting procedure used and gives the polynomials for the best fit and its 2sigma uncertainty limits. Appendix also gives the corresponding fits for Φ_{PC} as a function of am, which could be useful for studies of magnetospheric flux transfer (i.e., convection) that wish to make use of the geomagnetic activity data. The coefficients for the and am^{3} terms in the polynomial fits (for am and Φ_{PC}, respectively) are very small and the fitted variations are very close to being quadratic in form. It can be seen that the average am increases monotonically with average Φ_{PC}, but not linearly and this is consistent with the amplification of the semiannual variation in am relative to that in Φ_{PC}. Note that Figure 5 of Lockwood et al. (2019a) shows that daily means of the power input into the magnetosphere, P_{α}, estimated from interplanetary measurements, have a linear relationship with daily means of the am index. If we assume that steady state applies to averages taken over one day, the magnetospheric power input, P_{α} would be equal to the total daily power deposited in the magnetosphere–ionosphere–thermosphere system. This averaging timescale would cover several substorm cycles so should largely smooth these cycles out, which are the dominant oscillation of the storage/release magnetospheric system; however, we note that large storms often last longer than a day and so we should expect this relation to break down during such events. Hence together, Figure 5 of Lockwood et al. (2019a) and Figure 4 of the current paper strongly suggest that the total power deposited has an approximately dependence on transpolar voltage.
Fig. 4 Scatter plot of daily means of the am index as a function of the corresponding daily mean transpolar voltage, Φ_{PC}, as determined for 2001–2002 by DMSP satellites (principally F13) and normalised for the satellite track to an ideal 06–18 MLT path using the procedure of Lockwood et al. (2009). The error bars are plus and minus one standard error in the means. The mauve line is the bestfit 3rd order polynomial fit and the grey area is bounded by the 2sigma uncertainty level in that fit. The fitting procedure is given in Appendix to this paper, along with polynomial expressions for the best fit and the uncertainty band edges. These can be used to estimate the am level associated with a given Φ_{PC}: Appendix also gives the corresponding 3rd order polynomials that allow computation of the Φ_{PC} value associated with a given am value. 
Figure 5 presents timeofyear/time–ofday (F–UT) plots for transpolar voltage Φ_{PC} and polar cap flux W_{PC} for this dataset. Data are sorted into 24 equalwidth bins in F and 16 equalwidth bins of UT using the time of the centre of the polar cap crossing. The data are then smoothed with a 1–3–1 triangularweighting filter, applied in both the F and UT dimensions. The top panels are for passes of the northern polar cap, the lower panels are for passes of the southern polar cap. Figures 5d and 5h show that the number of samples in each F–UT bin, n_{bin}, are high and relatively constant varying between 26 and 30. The transpolar voltage Φ_{PC} (Figs. 5a and 5e) and, to a lesser extent polar cap flux, W_{PC} (Figs. 5b and 5f) show the semiannual variation but no hint of any equinoctial pattern. However, we cannot draw any firm conclusion from this because neither does the am index for simultaneous data (Figs. 5c and 5g). In contrast, Paper 1 demonstrated that the equinoctial pattern is clearly present in larger am datasets. This will be revisited in a later paper that deploys a much longer database of transpolar voltage observations.
Fig. 5 Timeofyear/time–ofday (F–UT) plots for transpolar voltage and polar cap flux for 2001–2002, inclusive from the DMSP satellites. Data are sorted into 24 equalwidth bins in F and 16 equal width bins of UT using the central time of the polar cap traversal. The top panels (a–d) are for passes of the northern polar cap, the lower panels (e–h) are for passes of the southern polar cap. (a and e) Show the transpolar voltage Φ_{PC}; (b and f) show the polar cap flux, W_{PC}; (c and g) show the simultaneous am index value (linearly interpolated from the threehourly values to the time of the centre of the polar cap crossing) and (e and h) shows the number of data points, n_{bin} in each F–UT bin. 
The F–UT plots for Φ_{PC} may not show an equinoctial variation, but they do show a fascinating UT variation with a major minimum at 0–10 UT at all times of year. Figure 5 shows this is seen in both hemispheres and it is also found in both of the two years when analysed in isolation (not shown). Paper 1 showed that in this UT band there is a persistent, but smaller, drop in geomagnetic indices and, notably, am with its highly uniform F–UT response. There is also an interesting UT variation in the polar cap flux seen in both polar caps, with W_{PC} almost in antiphase with Φ_{PC}. This behaviour is seen at all times of year except early in the year (F ≈ 0.2). Again this is seen in both polar caps and in the data from both years when analysed separately.
For the expandingcontracting polar cap model of ionospheric convection excitation (Cowley & Lockwood, 1992), the transpolar voltage measured by a satellite that passes through the centre of the convection polar cap, in the approximation that the polar cap remains circular in shape, is given by (Lockwood, 1991):(1)where Φ_{D} is the voltage along the reconnection Xline (or Xlines) in the dayside magnetopause where open flux is generated and Φ_{N} is the voltage along the reconnection Xline (or Xlines) in the crosstail current sheet where open flux is destroyed. We can apply equation (1) because the Φ_{PC} data have been normalised to be along the 06–18 MLT line by using the procedure of Lockwood et al. (2009). By Faraday’s law, applied to the openclosed field line boundary, the rate of change of open flux in the polar cap flux (W_{O}) is(2)
Equation (2) is effectively a continuity equation for the open flux, W_{O}. The equality with the rate of change of W_{PC} is only approximate because the polar cap flux determined from the spacecraft data is the flux inside the convection polar cap, which is generally greater than the open flux W_{O}; however the difference is generally small and the rate of change of the difference even smaller. By equation (2), the rise in W_{PC} (dW_{PC}/dt > 0) seen at about 22–05 UT implies either that Φ_{D} has increased or that Φ_{N} has decreased, but the former would, by equation (1), cause a rise in Φ_{PC} whereas the latter would cause a fall. Hence the combined behaviour of a fall in Φ_{PC} and a rise in W_{PC} is qualitatively consistent with a fall in Φ_{N} at around 22–05 UT. Figure 5 shows that the rise in W_{PC} does lead the minimum in Φ_{PC} by 2–3 h. This could potentially be explained by the initial motions in the ionospheric footprints of the reconnection lines and the consequent inductive decoupling of the voltages appearing along the Xlines and their ionospheric footprints, the “merging gaps”. This means that ionospheric flows (and hence voltages) are not established straight away as the open/closed boundary moves (Morley & Lockwood, 2006; Lockwood et al., 2006). However, the rates of change in dW_{PC})/dt correspond to voltage differences of between about 5 and 10 kV only, whereas the drop in Φ_{PC} is of order 25 kV so quantitatively the two are not in such good agreement with a common cause of a drop in Φ_{N}. This may arise from the limitations in the method used to derive W_{PC} or may point to another factor contributing to the observed minimum in Φ_{PC} at 0–10 UT. There is a systematic variation with UT of the orbit path of the main spacecraft employed (DMSPF13) and that may have contributed to the 0–10 UT minimum in Φ_{PC}. In theory, this should have been corrected for by the procedure of Lockwood et al. (2009) that normalises Φ_{PC} measured for the actual orbit path to an ideal 06–18 MLT pass: however, if the correction is too small it could give an overestimation of the depth of the 0–10 UT minimum in Φ_{PC}.
To summarise this section, we have presented an initial, limited study of the semiannual variation in transpolar voltage and flux transfer. We find there is a clear semiannual variation in transpolar voltage that mirrors the waveform of that in geomagnetic activity closely but it is not as large in amplitude. This is consistent with the latter showing a nonlinear (quadratic) variation with transpolar voltage. We find that during the persistent minimum of the UT variation in geomagnetic activity defined in Paper 1 there is a persistent decrease in transpolar voltage which may be, at least in part, consistent with a decrease in reconnection voltage in the crosstail current sheet. Confirmation of these findings using a larger dataset of transpolar voltage measurements will be presented in a later paper in this series, but the great consistency with which we find the features described (in both polar caps and in each of the three years studied here) leads us to conclude they are probably real phenomena.
3 The geomagnetic response to power input to the magnetosphere
3.1 Estimation and distributions of power input
As in Paper 1, we use the power input to the magnetosphere, P_{α}, computed from interplanetary measurements using the Vasyliunas et al. (1982) theoretical formulation. To remove some constants in the formulation we divide by the overall mean for the interval studied which is 1995–2017, P_{o} = <P_{α}>_{all}. The correlations between P_{α}/P_{o} and the am index were studied for the near continuous data after 1995 and are very high, being 0.978 for annual means, 0.932 for means over Carrington rotation intervals (≈27 days), 0.908 for daily means, and 0.842 for 3hourly means (Lockwood, 2019).
Figure 6 shows the cumulative distribution functions (c.d.f.s) of P_{α}/P_{o}. The P_{α}/P_{o} data are divided into 288 bins, 36 equalsized bins in fraction of year F and the same eight 3h Universal Time (UT) bins over which the range am index is evaluated. For each F–UT bin the c.d.f. of P_{α}/P_{o} is coloured according to the mean value for that bin, <P_{α}/P_{o}>_{F,UT}. Figure 6 corresponds to the plot for am presented in Figure 8b of Paper 1, covering the same interval (1995–2017), so these two families of c.d.f.s can be directly compared and are found to be very similar. The origin of the distributions (which are lognormal in form) has been explained by Lockwood et al. (2019b, 2019c) and arises from the effect of averaging timescale on the distribution of the IMF orientation factor sin^{4}(θ/2) factor at high (1 min) time resolution, where θ is the IMF clock angle. It should be stressed here that the 4th power exponent in this sin^{4}(θ/2) factor is not a free fit parameter but arises from the analysis proposed Vasyliunas et al. (1982). As discussed in the introduction to this paper and in Paper 1 (Appendix B), the IMF orientation factor is critical to the analysis of the semiannual variation presented here because it gives rise to the R–M effect. Lockwood et al. (2019b) and Lockwood (2019) have shown that this sin^{4}(θ/2) factor performs better than all suggested alternatives, including the “halfwave rectified” southward field which was the basis of the original R–M theory (Russell & McPherron, 1973). The sin^{4}(θ/2) factor preserves the nonlinearity in the geomagnetic response which results in the R–M effect giving the semiannual variation, but avoids the discontinuity in slope at [B_{Z}]_{GSM} = 0 (θ = π/2). The formulation also allows for a continued, lower rate of magnetopause reconnection opening field lines, even when the IMF is northward, as has been deduced in a number of studies, including observations of ionospheric O^{+} ions escaping the magnetosphere on open field lines (Chandler et al., 1999). However, note that although the sin^{4}(θ/2) factor is used here and other studies have found it to be optimum, other forms for the IMF orientation factor will generate the R–M effect and the results may vary in detail, but not in the basic principle.
Fig. 6 Cumulative probability distributions (c.d.f.s) of the power input to the magnetosphere, P_{α}, as a ratio of its mean value P_{o} = <P_{α}>_{all} for the full interval of nearcontinuous interplanetary data (1995–2017, inclusive). The P_{α}/P_{o} data are divided into 288 bins: 36 equalsized bins in fraction of year F and the same eight 3h Universal Time (UT) bins over which the range am index is evaluated. For each F–UT bin the c.d.f. of P_{α}/P_{o} as a ratio of the mean value for that bin <P_{α}/P_{o}>_{F,UT} was plotted and colored according to the value of that ratio. The black line is for all P_{α}/P_{o} data, the orange line for the F–UT bin giving the largest <P_{α}/P_{o}>_{F,UT} and the green line for the F–UT bin giving the smallest <P_{α}/P_{o}>_{F,UT}. 
We use 3hourly means of the power input P_{α}/P_{o} to compare with the 3hourly values of the am index, even though the latter are not mean values over the 3 h, being based on the range (maximum minus minimum) of variation within the 3h interval, as seen at the various magnetometer stations employed by the index. Because over fixed 3h intervals the mean value of P_{α}/P_{o} is proportional to the integrated value, this is consistent with the storage/release concept of magnetospheric behaviour, whereby integrated energy input into the magnetosphere is stored in the geomagnetic tail and released to give bursts of geomagnetic activity that set the range values detected by each the am magnetometers. There is an issue as to what is the most appropriate response lag between the 3hourly power input means and the geomagnetic response and that is analysed in Section 3.3 below. Confirmation that it is appropriate to compare 3hourly means with threehourly geomagnetic range values (such as am) is provided by Appendix A of Paper 1 which compares the am index data with both the averages of the auroral electrojet AE and −AL indices and with their maxima over the 3 h intervals and finds equally good correlations.
Figure 7 stresses the excellent correlations between am and P_{α}/P_{o} for large averaging timescales, τ, but also shows there is structure in the scatter. Part (a) is for annual means (τ = 1 year). The black squares are for all data, to which the green line is the best linear regression fit. The coloured squares are for the data separated into the 8 UT ranges over which the am index is computed, the squares are coloured according to the centre value of those UT ranges using the scale given at the top of the panel. There is a persistent pattern with a fixed P_{α}/P_{o} value giving the lowest am for the 00–06 UT data and the highest am for the 21–24 UT data. This is consistent with the UT variation discussed in Paper 1 and in Section 2 above. Figure 7b shows the scatter plot for τ = 10 days. The points are here colourcoded by separation in time from the closest equinox of the centre of the 10day averaging interval, Δt_{eq}. There is again a consistent pattern with a fixed P_{α}/P_{o} value giving the lowest am around the solstices (which are at Δt_{eq} = 0.25) and highest am at the equinoxes (at which Δt_{eq} = 0). The green line is the leastsquares regression line for all data and is very similar to that for annual data shown in Figure 7a, and the orange and black lines are the linear regression fits to all data taken in intervals of length 0.25 yr around the solstice and equinoxes. The bestfit coefficients and their 2sigma errors for the ordinary least squares (OLS) linear regression are given in Table 1. The slopes of the fit for the equinox data is significant smaller than for the solstices (difference equal to −3.7 × 10^{−3} ± 0.3 × 10^{−3}) so it takes a lower P_{α}/P_{o} to generate a given am value at the equinox. This demonstrates that there is a significant difference between the solstices and equinoxes in the geomagnetic response to a given power input P_{α}/P_{o}. From the linear regression coefficients for the 6month intervals around the equinoxes and solstices used in Table 1, this amplification of the equinox am value over the solstice am value for a given P_{α}/P_{o} is by a factor of 1.085. However, this figure is obtained by dividing all data in the year into either solstice and equinox and will greatly underestimate the amplitude of the semiannual variations seen in higher timeresolution data. This was investigated in Paper 1, using 36 equalwidth bins of timeofyear, F (each just over 10 days long) and for am the semiannual variation amplification factor was found to be about 2 at this 10day time resolution. The important point about Figure 7b and Table 1 is that it establishes that the ratio am/P_{α} varies with time of year and is statistically higher at the equinoxes than at the solstices. Similarly Figure 7a shows that am/P_{α} is reduced at 00–06 UT. This means that at least some of both the semiannual variation and the UT variation arises not from the solar windmagnetosphere energy input but the response of the magnetosphere–ionosphere–thermosphere system to that energy input.
Fig. 7 The relationship of am and normalized power input to the magnetosphere P_{α}/P_{o} for large averaging timescales τ: (a) τ = 1 year and (b) τ = 10 days. In (a) the black circles are for all data, for which the linear correlation coefficient is r = 0.97 and the best fit OLS (Ordinary LeastSquares) linear regression, shown by the green line, is (P_{α}/P_{o}) = s·am + c. Averages for the 8 UTs of the am index are shown by the squares, colored by the scale shown at the top of the figure: it can be seen that am is persistently a little smaller than the average response at 0–9 UT and persistently a little larger at 15–4 UT. In (b) the points are colored by the separation of timeofyear F from the value at the closest equinox, Δt_{eq}. The linear correlation for all data is r = 0.96 and the best fit linear regression is again the green line. The am response at low Δt_{eq} (around the equinoxes) is persistently greater than at high Δt_{eq} (around the solstices) showing that there is a contribution to the semiannual variation in am is not associated with that in (P_{α}/P_{o}) and hence not directly attributable to the R–M effect on solar windmagnetosphere coupling. The black and orange lines are, respectively, the bestfit linear regressions for around the equinoxes (Δt_{eq} ≤ 0.125) and around the solstices (Δt_{eq} > 0.125). The ordinary least squares (OLS) linear regression coefficients and errors are given in Table 1. 
3.2 Amplification of the semiannual variation and the equinoctial F–UT pattern
Paper 1 used the four aσ indices, compiled by Chambodut et al. (2013) from stations in four 6h magnetic local time (MLT) intervals around 06, 12, 18 and 24 MLT, to show that all display a semiannual variation and an equinoctial F–UT pattern. It is interesting to note how the index compilation has influenced behavior. The aσ indices are based on the range (between maximum and minimum) of the horizontal field component detected in a 3h intervals and are all, like the am index, dominated by substorm expansion phases and sawtooth events; in other words, by the unloading part of the storagerelease magnetospheric response.
In Paper 1 it was shown that the variation of aσ_{midnight}/<aσ_{midnight}> over the year is 0.317 ± 0.040 whereas for P_{α}/P_{o} the corresponding amplitude is 0.136 ± 0.035. Thus the semiannual variation in the aσ_{midnight} index is significantly amplified compared to that in power input to the magnetosphere and the best estimate of the factor is 0.317/0.136 ≈ 2.33. The aσ_{midnight} index shows the strongest semiannual variation, aσ_{dawn} and aσ_{dusk} show similar amplitude variations to am and aσ_{noon} shows the weakest (with no significant amplification over P_{α}/P_{o}). For both aσ_{dawn} and aσ_{dusk} the best estimate of the amplification factor is 2.01, for aσ_{noon} it is 1.08, and for the am index it is 2.08. Hence all the am and aσ indices, apart from aσ_{noon}, show amplification of the semiannual variation in input power into the magnetosphere and this, as for the equinoctial F–UT patterns reported by Chambodut et al. (2013) is almost nonexistent at noon and strongest at midnight. The studies of McPherron et al. (2013) and Chu et al. (2015) using the AL auroral electrojet AL index and a midlatitude nightside “bay” index find that the R–M effect on solar wind/magnetosphere coupling is about 40% of the geomagnetic response which yields an amplification factor of 2.5, in good agreement with the amplification factor found here for aσ_{midnight}.
3.3 Response lag of the am index to power input
In order to study shorter averaging timescales than the approximately 10day intervals used in Figure 7, we evaluate the correlation between P_{α}/P_{o} and the am index as a function of F, UT and activity level, allowing for an optimum response lag of am, dt_{o}. To do this, we make 1min values of P_{α}/P_{o} using the procedure of Lockwood et al. (2019a). This employs 1min interplanetary observations for 1995–2017 (inclusive), downloaded from the Omni database which is compiled and maintained by the Space Physics Data Facility at NASA’s Goddard Space Flight Center. These Omni interplanetary data have been lagged so that the timings apply to the conditions arriving at the nose of Earth’s bow shock. Oneminute mean ion mass data is taken from the highest resolution available (which is often hourly) using Piecewise Cubic Hermite Interpolating Polynomial (PCHIP) interpolation. This interpolation procedure was tested using available 64s observations of the solar wind mean ion mass from the Advanced Composition Explorer (ACE) which were averaged into hourly values to which the same procedure was then applied. It was found the distribution of differences was Gaussian with mode and mean value at essentially zero and a standard deviation that was 2.5% of the overall average of the mean ion mass. We employ the optimum coupling exponent α of 0.44 found by Lockwood et al., (2019b). These 1min P_{α}/P_{o} values are then combined into hourly means, rejecting values where the predicted error in P_{α}/P_{o} due to data gaps exceeds 5%, using the criterion established by Lockwood et al. (2019a) from the introduction of synthetic data gaps. Note that interpolation, such as used above to obtain 1min mean ion mass values, cannot be used to fill data gaps in the P_{α}/P_{o} sequence because of the much greater variability introduced the IMF orientation factor A_{θ} (Lockwood et al., 2019b, 2019c). These hourly values are then averaged into 3hourly values, keeping data only when all three 1h values are available. This procedure was repeated many times to generate 3hourly values centred on time t_{am}dt, where t_{am} are the times of the centres of the threehourly intervals over which am is computed and the response lag dt was varied between −10 min and +300 min in steps of 1 min.
All 310 of the series of threehourly P_{α}/P_{o} data obtained this way were then subdivided into the 288 F–UT bins used in Figure 3. The total number of am data points in the interval (1 January 1995–1 September 2019) is 75,008 and dividing into 288 F–UT bins gives an average of 260.4 F–UT samples per bin. To make the lag correlations we use one day of data (i.e., eight values around each F–UT sample at times −9 h, −6 h, −3 h, 0, +3 h, +6 h, +9 h and +12 h relative to the UT of the bin) giving an average of 2083.3 observations per bin for all the am data. The data were then further subdivided into nine quantile ranges of geomagnetic activity, as quantified by am. We use the notation that, for example, the 50% quantile of am (i.e., the median) is q(0.5). The quantile ranges used are: q(0) < am < q(1) (i.e., all data, the red points in Fig. 8 for which there are 2083.3 samples in each correlation analysis); five nonoverlapping bands each containing 20% of the data (i.e., an average of 416.7 am samples in each correlation analysis), q(0) < am < q(0.2) (orange points), q(0.2) < am < q(0.4) (pink points), q(0.4) < am < q(0.6) (light green points), q(0.6) < am < q(0.8) (cyan points), and q(0.8) < am < q(1) (dark green points). We also study and the largest 10% of am values q(0.9) < am < q(1) (blue points, each based on 208.3 samples on average), the largest 5% q(0.95) < am < q(1) (the mauve points, each based on 104.2 samples on average) and q(0.99) < am < q(1) (the black points, each based on just 20.8 samples on average). In fact some data subdivisions had fewer data points to correlate than this because of data gaps in the P_{α} data series. In each of these 288 × 9 = 2592 cases, a lag correlogram was generated and the peak correlation, r_{p}, and the lag giving it, dt_{p}, were determined. Figure 8 is a scatter plot of r_{p} against log_{10}(dt_{p}). Points are colour coded by the quantile range using the key given (as described above, the red points are for all data and the black point for the largest 1%). Points are plotted only where the number of data points exceeded 16 and the pvalue of the null hypothesis (that the correlation was zero) was less than 0.05. As expected, the lowest correlations are for the loweractivity quantiles, the higher quantiles generally giving values that exceed the value of 0.842 that is obtained for all 3hourly means. The values of dt_{p} are generally between 15 min and 2 h, although we note peak correlation is sometimes found at larger dt_{p} for the largest 1% of am values – these are long duration storms where the activity level remains high for a day or longer. Note also that these values for the largest am are the least reliable because they are based on the smallest number of samples, although the pvalues of points used always meet the 95% significance threshold. The Omni data have been lagged to the nose of the bow shock and so dt_{p} includes allowance for propagation time across the magnetosheath to the magnetopause, the time for open flux to be appended to the tail lobe and the growth phase accumulation of sufficient open flux in substorm growth phases to drive the onset of expansion phase activity and then a further delay as am rises to its peak response.
Fig. 8 Peak correlation r_{p} between am and 3hourly means of P_{α}/P_{o} as a function of the logarithm of the lag dt_{p} giving that peak correlation, for data from 1995–2017, inclusive. Oneminute P_{α}/P_{o} data were averaged into hourly and then 3hourly intervals, centred on the midpoints of the am data intervals, minus a response lag dt that was varied between −10 min and 5 h in steps of one minute. At each dt the correlation between P_{α}(t − dt)/P_{o} and am(t) is evaluated for each of the data subsets studied. The data are divided and colourcoded into 9 quantile ranges of am given by the legend. We use the notation that 20% of all the data have am values lower than q(0.2). The red dots are for all data, q(0) < am ≤ q(1). Five of the quantile ranges each contain 20% of the data: q(0) < am ≤ q(0.2), q(0.2) < am ≤ q(0.4), q(0.4) < am ≤ q(0.6), q(0.6) < am ≤ q(0.8) and q(0.8) < am ≤ q(1). In addition, the plot studies the top 10%, 5% and 1% of all am values, respectively q(0.9) < am ≤ q(1), q(0.95) < am ≤ q(1), and q(0.99) < am ≤ q(1). Correlations are shown for F–UT–m bins that have more than 16 samples and for which the correlation significance exceeds 95%. 
Figure 9 plots the distributions of 288 dt_{p} values (one for each of the F–UT bins) for various quantile ranges of am. The distributions are taken in 1min ranges and then smoothed with a 10min running mean. The upper panel shows the distribution for all values (q(0) ≤ am ≤ q(1), the red points in Fig. 8) and the lower panel shows the distributions for the eight other quantile ranges, using the same colour scheme as for the points in Figure 8. It can be seen in the lower panel that for lower am the lag is somewhat longer (peaking around 75 min). For am values above q(0.8) there are hints of two peaks to the distribution, one around 65 min, the other near 40 min and as am is further increased, the relative magnitude of these two peaks changes appears to change. It is not certain that there really are two peaks to these distributions, but the net effect of a decrease in the mode value at higher quantile levels is clear. For the largest 1% of am values there is a very broad distribution with a single peak near 45 min. The top (red) curve shows that the optimum single lag is dt_{o} = 60 min. As expected there are almost no values at negative lags and there are very few at lags below 10 min which would be quasiinstantaneous responses. These cases may be chance occurrences or could be the effects of a second pulse of reconnection, as modelled for ionospheric convective flow by Morley & Lockwood (2006). We searched for a consistent pattern in F and UT in this distribution of dt_{p} lags but found none (the lack of variation with UT in particular being expected because values for a whole day were selected around each UT).
Fig. 9 Distribution of response lags for the 288 F–UT bins studied for the response of the am index to the power input to the magnetosphere P_{α}/P_{o} for 1995–2017, inclusive. (Top) The red line shows the distribution of the lags giving peak correlation dt_{p} between P_{α}/P_{o} and am for all the data. (Bottom) The same analysis but the am data have been further subdivided into the 8 other quantile ranges used in Figure 8. The distributions are taken in 1min ranges and then smoothed with a 10min running mean. 
In summary, in this section we have studied the response of geomagnetic activity, as quantified by the am index, to power input into the magnetosphere. We find no consistent variation in the response delay with timeofyear, F, and the use of oneday intervals around nominal UTs (to keep sample numbers high and correlations significant) precluded the detection of any UT variation. The mode of the distribution consistently moves to slightly lower lags with increased activity levels and the width of the distribution increases. However, we note that the lower number of samples for the highest am values means that these results are the least reliable. Overall, a lag of 1 h is optimum for the whole dataset (dt_{o} = 60 min), but there is considerable spread with the 2sigma points being near 10 and 100 min. For all studies in the remainder of this paper that compare the am index with any interplanetary parameters (for example power input into the magnetosphere, P_{α}/P_{o}, solar wind dynamic pressure, p_{SW}, or the IMF [B_{Y}]_{GSEQ} component) we make allowance for this mean lag by taking the interplanetary data dt_{o} = 60 min before the time of the am value (which is the midpoint of the 3h interval over which the range and hence the am value is determined).
4 Yeartoyear fluctuations in the semiannual variations in power input to the magnetosphere and the geomagnetic response
As discussed in Paper 1 and in the introduction to the present paper, the semiannual variation generates a paradox because the R–M effect, which is undoubtedly active, enhances the average magnetic activity levels and the occurrence of small and moderate storms. However, because the largest storms are driven by southward pointing field in the GSEQ frame, the R–M effect should actually, on average, decrease the occurrence of large storms at the equinoxes, the opposite of what is observed. Hence, in this section we investigate the occurrence of large geomagnetic disturbances. Panels (d)–(i) of Figure 10 are “Fyear spectrogram” plots, and allow us to study the yeartoyear fluctuations in the timeofyear variation in the am index and in the aσ indices and compare to the corresponding behaviour of P_{α}/P_{o}. Each panel gives the normalized mean value in the 36 equalsized bins of timeofyear F as a function of year and F: hence, for example, the pixels in Figure 10e are coloured according to the value of <am>_{F}/<am>_{1yr}. We normalize by dividing by the annual means, <am>_{1yr}, so that the timeofyear variations can always be seen, even at sunspot minimum. The top row of plots shows the solar cycle variations by plotting the variations in the annual means as a ratio of the overall mean value: Panel (a) is for power input into the magnetosphere, <P_{α}>_{1yr}/P_{o}; (b) is for the am index, <am>_{1yr}/<am>_{all} (and also shows in green the number of substorm onsets N_{o} derived from the SML index, <N_{o}>_{1yr}/<N_{o}>_{all} as used in Paper 1); (c) shows the annual means for the four aσ indices, <aσ>_{1yr}/<aσ>_{all} . The solarcycle variations of annual means shown in parts (a)–(c) of Figure 10 are all very similar indeed. The Fyear patterns in parts (d)–(i) are also very similar indeed, in terms of both average levels and individual large events (storms), which show up as orange and yellow pixels. The colour scale used in all panels is the same and one can see the amplitude in the am and aσ response patterns is larger than for P_{α}/P_{o}, particularly during the larger events. Because the results for am and the four aσ indices are so similar, we hereafter show only the results for am.
Fig. 10 Top panel: plots of annual means of an index or variable x, <x>_{1yr}, for the years 1995–2017, inclusive, where x is: (a) the power input into the magnetosphere, P_{α}; (b) the am index (in black) and the number of substorm onsets derived from the SML index, N_{o} (in green); (c) the aσ indices for dawn (orange), noon (cyan), dusk (mauve) and midnight (blue). Middle and lower panels: yearF spectrogram plots of the means in 36 equalsized bins of F, divided by the mean for that year, <x>_{F}/<x>_{1yr}, where x is: (d) P_{α}; (e) am; (f) aσ_{midnight}; (g) aσ_{dawn}, (h) aσ_{noon}, and (i) aσ_{dusk}. The data have been smoothed by applying a 3point running mean to the time series of means in the bins of width (1/36) yr. 
As in Paper 1, in Figure 11 we subdivide the data for P_{α}/P_{o} and am into times when the [B_{Y}]_{GSEQ} was positive (middle column) or negative (righthand column), in order to identify the role of the R–M effect. We allow for the optimum lag dt_{o} = 60 min between the IMF conditions at the nose of the bow shock and the am response that was derived on Section 3.3. The lefthand column repeats the variations for all data for comparison. The top row (Figs. 11a–11c) shows the variations in annual means of P_{α}/P_{o} (in black) and am (in mauve). These solar cycle variations are all very similar, although the amplitude is slightly smaller for [B_{Y}]_{GSEQ} < 0 than for [B_{Y}]_{GSEQ} > 0 for the years covered by this study. The middle row shows <P_{α}/P_{o}>_{F}/<P_{α}/P_{o}>_{1yr} and the bottom row <am>_{F}/<am>_{1yr}. The R–M effect is clearly evident with persistent increases at the September and March equinoxes for [B_{Y}]_{GSEQ} > 0 and [B_{Y}]_{GSEQ} < 0, respectively. The annual variations for the two [B_{Y}]_{GSEQ} polarities separately are, again, slightly greater for am (Figs. 11h and 11i) than for P_{α}/P_{o} (Figs. 11e and 11f).
Fig. 11 Identification of the RussellMcPherron effect using the polarity of the IMF [B_{Y}]_{GSEQ} component. The lefthand column is for all data, the middle column for [B_{Y}]_{GSEQ} > 0 and the righthand column for [B_{Y}]_{GSEQ} < 0. The top row shows annual means of am and P_{α} divided by their overall means for the whole interval: respectively, <am>_{1yr}/<am>_{all} (mauve lines) and <P_{α}>_{1yr}/P_{o} (black lines). The middle row shows the variations of P_{α}/P_{o} as a function of fractionofyear, F and year, <P_{α}>_{F}/<P_{α}>_{1yr} and the bottom row the same for am, <am>_{F}/<am>_{1yr}. As for Figure 10, data are for 1995–2017 and have been averaged into 36equal width bins of F in each year, and a 3point running mean applied to the time series to smooth the data. 
For both P_{α}/P_{o} and am, Figure 11 demonstrates a clear R–M effect in mean values, detected by the dependence of which is the favoured equinox for a given IMF [B_{Y}]_{GSEQ}. However, the picture is not so clear in the occurrence of large storms, particularly in am. Taking a threshold of <am>_{F}/<am>_{1yr} = 1.7 (yellow pixels), we can see in Figure 11h that for [B_{Y}]_{GSEQ} > 0, 13 events occurred at the favoured March equinox, 5 at the unfavoured September equinox, and 4 around the solstices. In Figure 11i we see that for [B_{Y}]_{GSEQ} < 0 there were 10 events at the favoured September equinox, 3 at the unfavoured March equinox, and 6 around the solstices. Hence the tendency for events to cluster around the favoured equinox is present as predicted for the R–M effect, but Paper 1 (and Appendix B in particular) has explained why we might expect the R–M effect to, if anything, be reducing the magnitude of the largest storms (and hence the number of events over a large threshold) at the favoured equinox.
One additional point to note about the unseparated means (Figs. 11d and 11g) is that we are averaging over roughly 10day intervals and, given that for a 2sector, 3sector and 4sector structures in the heliospheric field, we remain in one sector for, on average, 13.5 days, 9 days and 6.75 days, this means that the averages are usually taken for one dominant polarity of IMF [B_{Y}]_{GSEQ}, and there is, in most cases, relatively minor cancellation of effects that depend on the polarity of [B_{Y}]_{GSEQ}. This is no longer true if we average over all years (instead of one year at a time as in Fig. 11) which will mean that for every F bin we are averaging roughly equal numbers of data with [B_{Y}]_{GSEQ} > 0 and [B_{Y}]_{GSEQ} < 0 with very similar occurrence distributions of [B_{Y}]_{GSEQ} for the two polarities. The results are radically and revealingly different, as shown in Figure 12. The lefthand plots in Figure 12 are for am and the righthand plots are for P_{α}/P_{o}. The top row shows the variations with F of the means for all years for all data (thick black line), for data with [B_{Y}]_{GSEQ} > 0 (thin line joining open circles) and for data with [B_{Y}]_{GSEQ} < 0 (thin line joining solid triangles). The results now show a significant difference in the behaviour for am and for P_{α}/P_{o}. For P_{α}/P_{o} (Fig. 12b) both the [B_{Y}]_{GSEQ} > 0 and [B_{Y}]_{GSEQ} < 0 cases have almost sinewave forms that are in antiphase. In fact, the positive deflection (at the equinox for which the [B_{Y}]_{GSEQ} polarity is favoured) is only very slightly larger than the negative deflection (at the equinox for which the [B_{Y}]_{GSEQ} polarity is unfavoured) and so the average for all data has only a very weak semiannual variation, because the nearperfect asymmetry between the variations for the two polarities means that they almost cancel and the net semiannual variation is small. Contrast this with the corresponding variations for am (Fig. 12a) for which the positive deflections at the favoured equinox in the singlepolarity [B_{Y}]_{GSEQ} curves are considerably larger than the negative deflections at the unfavoured equinox. As a result, the variation for all data (thick black line) has a much more marked semiannual variation. It is this much larger difference between the results at any one equinox between the “favourable” and “unfavourable” polarities of IMF [B_{Y}]_{GSEQ} that causes the amplification of the semiannual variation in the am geomagnetic activity index, compared to that in P_{α}/P_{o}.
Fig. 12 Analysis of the differences between semiannual variations in the am index (lefthand panels) and in power input to the magnetosphere, P_{α}/P_{o} (righthand panels). In all panels a thick line denotes the variation for all data, a thin line connecting open circles is for IMF [B_{Y}]_{GSEQ} > 0 and a thin line connecting filled triangles is for [B_{Y}]_{GSEQ} < 0. The lines in parts (a and b) are mean values whereas (c–f) show the average number of days per year N when the daily means of am (panels c and e) or of P_{α}/P_{o} (panels d and f) exceed their 90% quantile (blue lines), 95% quantile (orange lines) and 99% quantile (mauve lines). Parts (g) and (h) plot N/<N>, the number of days per year N normalized by their average value. Note that the variations of N in (e) and (f) for the separate IMF [B_{Y}]_{GSEQ} polarities are plotted on a scale that is different to that for all data in (c) and (d). The 90% quantile levels, q(0.9), are 2.031 for am/<am>_{all} and 2.081 for P_{α}/P_{o}; the 95% quantile levels, q(0.95), are 2.556 for am/<am>_{all} and 2.588 for P_{α}/P_{o}; the 99% quantile levels, q(0.99), are 4.193 for am/<am>_{all} and 4.400 for P_{α}/P_{o}. The close similarity between these pairs quantile values arises from the similarity of the normalised c.d.f.s shown in Figure 6 of the present paper and Figure 8b of Paper 1. 
This demonstrates that although the R–M effect is working, it is not working in quite the way that is commonly thought, which is the way that was envisaged by Russell & McPherron (1973) in their seminal paper. If the semiannual variation were working solely through the modulation of solar windmagnetosphere coupling by the nonlinearity of the response of power input to the magnetosphere to IMF orientation (as envisaged in the original paper), we would expect am and P_{α}/P_{o} to show the same behaviour in Figures 12a and 12b. This is clearly not the case. Figure 12 shows that for average P_{α}/P_{o}, the R–M effect is effective, but the resulting semiannual variation is much smaller than that in average am. The am index reflects the increase in P_{α}/P_{o} for the favourable polarity of [B_{Y}]_{GSEQ} but does not reflect the almost equal magnitude decrease in P_{α}/P_{o} for the unfavourable polarity of [B_{Y}]_{GSEQ}. Hence it is not the nonlinearity of the solar windmagnetosphere energy coupling that gives us most of the R–M effect, rather it is the nonlinearity of the geomagnetic response to that energy input. Hence separating the two polarities of the IMF [B_{Y}]_{GSEQ} component reveals that it is the nonlinearity of the am response to P_{α}/P_{o}, rather than just that of the P_{α}/P_{o} response to IMF orientation, that is causing the two subsets for [B_{Y}]_{GSEQ} > 0 and [B_{Y}]_{GSEQ} < 0 to combine to give such a strong semiannual variation in average am.
Figures 12a and 12b deal only with average values and the lower panels of Figure 12 look at the corresponding behaviour in the occurrence of large events. The thick blue, orange and mauve lines in Figures 12c–12f plot the number of events, N, in which a 3hourly am exceeds the 90, 95 and 99 percentile, respectively, i.e., am>q(0.90), am>q(0.95), and am>q(0.99). Figure 12c is for am and Figure 12d is for the corresponding P_{α}/P_{o} for the same data subset, both being for both IMF [B_{Y}]_{GSEQ} polarities. Similar behaviour is seen as for the mean values, with the amplitude of the semiannual variation in am always exceeding the corresponding one in P_{α}/P_{o}. Figures 12e and 12f break the variations in Figures 12c and 12d into their [B_{Y}]_{GSEQ} > 0 and [B_{Y}]_{GSEQ} < 0 subsets, using the same symbols as in parts (a) and (b). It can be seen that the amplification of the semiannual variation of the occurrence of large am events occurs for the same reasons as for the mean values. Lastly Figures 12g and 12h compare the normalised variations of all events above the thresholds by plotting N/<N> for both polarities. By definition, these variations in N/<N> are always about unity and the amplitudes of the semiannual variations in event occurrence can be compared. It can be seen in Figures 12g and 12h that for both am and P_{α}/P_{o}, the fractional amplitude of the semiannual variation in storm occurrence is larger for the larger storms. Comparing each line in Figure 12g with its corresponding line in Figure 12h, we can see the amplification of the semiannual variations in the occurrence of large am events compared to the corresponding variation for P_{α}/P_{o}. However, this also shows the amplification is greatest for the smaller events and decreases with a higher event threshold, such that for the top 1% of events, the amplitude of the semiannual variation in P_{α}/P_{o} is approaching that in am and the amplification is not much greater than unity. This runs counter to the Russell & McPherron (1973) explanation of large storms which was that small variations in P_{α}/P_{o} induced by the R–M effect caused a large increase in the number of great storms through the amplification effect. Figure 12 shows that in fact the amplification factor is smallest for the largest storms.
This variation in the amplification in the am response to P_{α}/P_{o} variations is puzzling. We know it is not a general feature of the overall variation of am with P_{α}/P_{o} as that is linear on all timescales. However, there is scatter about that linear variation and the effect must be hidden in that scatter. We will return to this point in Section 7. Figure 13 clarifies how the behaviour noted in Figure 12 is occurring. The upper plots show the variation of the probability distribution functions (p.d.f.s) of P_{α}/P_{o} with F for (left) IMF [B_{Y}]_{GSEQ} < 0 and (right) [B_{Y}]_{GSEQ} > 0. These are consistent with the means and occurrence frequencies for P_{α}/P_{o} shown in the right hand panels of Figure 12. The bottom panels of Figure 13 shows the am amplification, expressed as (am/<am>_{all})/(P_{α}/P_{o}), as a function of F and for the same P_{α}/P_{o} bins as for the p.d.f.s in the upper panels. What is noticeable is that the amplification of am, relative to P_{α}, is greater at the equinoxes and is particularly effective in raising am at low P_{α}/P_{o}. It is this effect that causes increases in am for the unfavourable [B_{Y}]_{GSEQ} polarity which means that the decreases for the unfavourable [B_{Y}]_{GSEQ} polarity cancel the increases for the favourable [B_{Y}]_{GSEQ} polarity to a relatively small extent only. This cancellation occurs to a much greater the extent for P_{α}/P_{o}, giving a smaller semiannual variation. A key point about Figure 13 is that it shows that the mechanism responsible for this amplification in am is only working at the equinoxes because we do not see any such amplification at the solstices.
Fig. 13 Plots of (top) the probability distribution functions of (P_{α}/P_{o}) as a function of F and (bottom) the am amplification factor, (am/<am>_{all})/(P_{α}/P_{o}), as a function of F and in the same (P_{α}/P_{o}) bins as the p.d.f.s in the top panels. The left hand panels are for IMF [B_{Y}]_{GSEQ} < 0, the right hand panels are for IMF [B_{Y}]_{GSEQ} > 0. 
To summarise this section, we have shown that the pattern of variations in Fyear spectrogram plots in geomagnetic activity and power input into the magnetosphere are very similar in both average values and the occurrence of large events and that the R–M effect is at the heart of both; however, there is a nonlinear amplification of the semiannual variation in the geomagnetic response. We show that there is only a very small asymmetry in the increase in power input into the magnetosphere P_{α}/P_{o} at a given equinox for the “favourable” polarity of [B_{Y}]_{GSEQ} (via the R–M) effect, compared to the decrease for the “unfavourable” polarity of [B_{Y}]_{GSEQ} but that this difference is much greater in the am response and it is this that amplifies the semiannual variation in am. We conclude that there is a second mechanism at work at the equinoxes that amplifies the am response to a given P_{α}/P_{o}.
5 Variations with solar wind dynamic pressure
As discussed in the introduction, there is evidence that geomagnetic activity depends on both the amount of magnetospheric open flux and on the solar wind dynamic pressure, p_{SW}. Given that the dominant mechanism that allows power input into the magnetosphere is the generation of open flux, this implies that geomagnetic activity should be increased by both solar wind power input to the magnetosphere and the solar wind dynamic pressure. This is confirmed by Figure 14 which colourcodes the average am in bins of width 0.15 in both P_{α}/P_{o} (along the x axis) and p_{SW}/<p_{SW}>_{all} (along the y axis). Figure 14 shows that am increases both with increased P_{α} at a fixed p_{SW} and with increased p_{SW} at a fixed P_{α}. (For both p_{SW} and P_{α}/P_{o}, values are taken a time dt_{o} before the am data, where dt_{o} = 60 min. is the optimum response lag, as derived in Sect. 3.3). The grey points in Figure 15a are a scatter plot of 3hourly am data as a function of the 3hourly P_{α}/P_{o} values (again using the optimum response lag dt_{o} of 60 min) and Figure 15b studies the extent to which the scatter around the best regression is explained by p_{SW}. In Figure 15a the mauve line is the linear regression fit to the 3hourly data, given by(3)
Fig. 14 Mean values of am as a function of normalised power input into the magnetosphere, P_{α}/P_{o} (horizontal axis) and normalised solar wind dynamic pressure, p_{SW}/<p_{SW}>_{all} (vertical axis). The mean values of am are evaluated in bins of width 0.15 in both P_{α}/P_{o} and p_{SW}/<p_{SW}>_{all} and pixels colourcontoured according the scale shown. Only bins containing at least 5 samples are considered. 
Fig. 15 Analysis of 3hourly am data. (a) The grey points form a scatter plot of am against 3hourly means of power input into the magnetosphere, P_{α}/P_{o}, generated by averaging over 3 h intervals that are shifted forward in time by the derived best lag of dt_{o} = 60 min (see Fig. 7a) relative to the threehourly intervals in which am is evaluated. The orange points are the mean values averaged in 1% quantile ranges of P_{α}/P_{o}, i.e., q(0) ≤ <P_{α}/P_{o}>_{τ=3h} < q(0.01), q(0.01) ≤ <P_{α}/P_{o}>_{τ=3h} < q(0.02), up to q(0.99) ≤ <P_{α}/P_{o}>_{τ=3h} < q(1). The black error bars are the plus and minus one standard deviation in those means. The mauve is the bestfit OLS linear regression to the 3hourly data. (b) the grey points are a scatter plot of the fit residuals Δam for the fit shown in (a): this is the difference between each threehourly am value and the best fit linear regression value based on the corresponding <P_{α}/P_{o}>_{τ=3h} value (Δam = am − am_{fit}), plotted as a function of the simultaneous threehourly mean of the normalized solar wind dynamic pressure <p_{SW}>_{τ=3h}/<p_{SW}>_{all}. The orange points in (b) are means in 1% quantile ranges of <p_{SW}>_{τ=3h} and error bars are plus and minus one standard deviation in the mean Δam. The mauve line is the best linear regression to the 3hourly values. 
Higherorder polynomial fits were carried out but differences were always negligible. The orange points are mean values in 1 percentile ranges of P_{α}/P_{o} and the errors bars are plus and minus one standard deviation. The correlation coefficient between am and P_{α}/P_{o} for this 3h timescale is r_{c} = 0.867 which means P_{α}/P_{o} is explaining 100 × = 75.2% of the variation in am. The fit residuals for this linear regression, given by(4)were then evaluated and are plotted in Figure 15b as a function of normalized solar wind dynamic pressure p_{SW}/<p_{SW}>_{all}. As expected from Figure 14, the average values of Δam increase with p_{SW}/<p_{SW}>_{all} but there is extremely large scatter, indicating that although p_{SW} is a factor, it is certainly not the only factor influencing these fit residuals. However, Figures 14 and 15 show that the amplification of the am response to a given P_{α}/P_{o}, identified in Figures 12 and 13 as a key component of the semiannual variation, increases with increased solar wind dynamic pressure, p_{SW}.
Figure 16 plots the fit residuals Δam as a function of F and UT. Given that am shows an equinoctial pattern and that P_{α}/P_{o} does not, it is not surprising that the fit residuals Δam form an equinoctial pattern. This can be seen in Figure 16a which is for all data. Figures 16b–16d subdivide the data into the three terciles of solar wind dynamic pressure, p_{SW}. These plots all show the equinoctial pattern and comparing them reveals that this pattern grows in amplitude as p_{SW} increases, revealing that solar wind dynamic pressure is a key factor in the generation of the equinoctial pattern as well as a contributor to the scatter in the relationship between P_{α}/P_{o} and the am index.
Fig. 16 F–UT pattern plots of the fit residuals Δam in Figure 15b for 1995–2017, inclusive. In each case the plots are constructed using hourly means of the solar wind dynamic pressure p_{SW} and the threehourly fit residuals Δam are interpolated to the midpoint of those hourly intervals. All plots are then smoothed with a 1–3–1 triangular weighting filter in both the F and UT dimension. (a) The overall mean value of Δam. (b) The mean Δam for the lower tercile of the simultaneous solar wind dynamic pressure, <Δam>_{F,UT} for q(0) ≤ <p_{SW}>_{τ=1h} < q(0.33). (c) The mean Δam for the middle tercile of the simultaneous solar wind dynamic pressure, i.e., <Δam>_{F,UT} for q(0.33) ≤ <p_{SW}>_{τ=1h} < q(0.67). (d) The mean Δam for the upper tercile of the simultaneous solar wind dynamic pressure, <Δam>_{F,UT} for q(0.67) ≤ <p_{SW}>_{τ=1h} < q(1). Note that in parts (b–d) the same color scale is used to emphasize that the equinoctial pattern, although present for all three ranges of p_{SW}, is of amplitude that increases with p_{SW} and is much larger in amplitude for the largest p_{SW} values. This relationship is further studied by Figure 17. 
Figure 17 studies this relationship in more detail using 20 quantile ranges in p_{SW} that each include 5% of the data. It shows that the amplitude of the equinoctial F–UT pattern, here quantified by the standard deviation of the 288 data points that go into each pattern, σ(<Δam>_{F,UT}), varies almost linearly with p_{SW}. The alternate grey and white bands define the quantile ranges in p_{SW} that were used.
Fig. 17 The amplitude of the equinoctial pattern as a function of normalised solar wind dynamic pressure. The pattern amplitude is quantified by the standard deviation of the mean values of Δam in the 288 F–UT bins used to construct patterns like those shown in Figure 10, σ(<Δam>_{F,UT}). Patterns were constructed for 20 quantile ranges of the hourly mean solar wind dynamic pressure p_{SW} and σ(<Δam>_{F,UT}) for each is plotted as a function of <p_{SW}>_{q}/<p_{SW}>_{all}, where <p_{SW}>_{q} is the mean dynamic pressure in the quantile range and <p_{SW}>_{all} is the mean for all the data (which are for 1995–2017). The grey and white bands define the 20 quantile ranges of p_{SW} employed. There are a total of 192,719 valid hourly means of p_{SW} in the dataset and so each of the 20 quantile ranges contains 9636 samples. 
To summarise this section, we have confirmed the results of previous studies discussed in the introduction (Schieldge and Siscoe, 1970; Caan et al., 1973; Kokubun et al., 1977; Karlsson et al., 2000; Yue et al., 2010) that indicated enhanced solar wind dynamic pressure enhances geomagnetic activity. We have quantified the effect and show that it is responsible for the equinoctial pattern of the geomagnetic response which increases in amplitude with enhanced solar wind dynamic pressure.
6 The timeofyear and Universal Time variation
Figure 18 analyses the contributions to the (a) timeofyear F and (b) UT variations in am. The orange lines are am_{fit} variations, derived from the P_{α}/P_{o} data using the linear regression equation (1). As expected from Figure 12, the semiannual variation with F (averaged over all UT) in am_{fit} is much smaller amplitude than that in am (black line). The mauve, blue and cyan lines show the means of the fit residuals Δam from equation (4) for the upper, middle and lower tercile ranges of the solar wind dynamic pressure, p_{SW}. Figure 18a shows that the semiannual variation is present for the lower and middle terciles of p_{SW} but the largest contribution comes from the highest p_{SW}. Figure 18b is the corresponding plot for the UT variations. In this case, there is no UT variation in P_{α}/P_{o} and hence am_{fit}. The UT variation, like the F variation, is seen in all of the average variations of Δam but again is strongest for the larger tercile of p_{SW}. This indicates that p_{SW} also plays a role in generating the UT variation.
Fig. 18 Analysis of the contributions to the (a) semiannual and (b) UT variations of the am index observed in the interval 1995–2017, inclusive. Mean values are shown as a function of (a) timeofyear F and (b) UT of: (black lines) the observed am data; (orange lines) am_{fit}, the bestfit of power input to the magnetosphere, P_{α}/P_{o} to am; (mauve lines) the fit residual Δam for the upper tercile of the solar wind dynamic pressure, q(0.67)<p_{SW} ≤ q(1); (blue lines) the fit residual Δam for the middle tercile of the solar wind dynamic pressure, q(0.33) < p_{SW} ≤ q(0.67); and (cyan lines) the fit residual Δam for the lower tercile of the solar wind dynamic pressure, q(0) < p_{SW} ≤ q(0.33). 
This section has provided further tests of theoretical and model interpretations of the variations with dynamic pressure. It shows that not only does the amplification of the semiannual variation increase with increased solar wind dynamic pressure, so do the amplitudes of the equinoctial pattern and of the UT variation.
7 Discussion
We have presented a purely empirical analysis of the response of the am geomagnetic index (and its 4 MLT sector subindices, aσ_{dawn}, aσ_{noon}, aσ_{dusk}, and aσ_{midnight}) to interplanetary conditions arriving at the nose of Earth’s bow shock as a function of time of year, F, Universal Time, UT, and geomagnetic activity level, using oneminute interplanetary data averaged into 3h windows. We have made use of the analysis of the effect of data gaps by Lockwood et al. (2019a) to ensure that the error they cause to estimates of power input into the magnetosphere is below 5%. In order to compare the interplanetary data with the geomagnetic data we need to know the best response lag to employ to ensure we are comparing the most relevant interplanetary data. In Section 3.3, we used averages over same interval durations as used to compile the am geomagnetic data (3hourly) but the averaging intervals were shifted in steps of one minute in order to determine the best response lag (taken to be the lag giving peak correlation). We show that there is considerable variability in response times between about 10 min and about 100 min and there is some systematic variation in the derived distribution of response lag with activity level  the mode of the overall distribution being close to 60 min. Hence using a single value for the response lag is an approximation; however, in this context it must be remembered that using an ideal individual lag for each case would also be an approximation because the driving power input may be relevant over an interval that could be longer, or shorter, than the 3 h used. However, using 3h interplanetary data and neglecting such a response lag completely (by using simultaneous data) would result in, on average, one third of each interplanetary sample not being the most relevant data and using an overall average lag is hence better than using simultaneous data.
We have presented a number of empirical analyses that will be used to constrain interpretation in a later paper in this series which will employ both empirical and MHD models of the magnetosphere.
The results presented in the current paper confirm that the R–M effect is active and the key part of the semiannual variation in geomagnetic activity. In the midnight sector we find that 42.9% of the semiannual variation in geomagnetic activity can be explained as the semiannual variation in the power input into the magnetosphere, P_{α}, which is almost entirely due to the RussellMcPherron effect. This agrees well with the results of McPherron et al. (2013) who employed the AL auroral electrojet and a simpler interplanetary coupling function (the halfwave rectified dawntodusk interplanetary electric field E_{DD} = V_{SW}B_{S}) and the linear prediction filter formalism which enabled them to allow for systematic variations in the relationship between E_{DD} and the AL index (for example with time of year because AL is a northern hemisphere index only and so has a strong seasonal response). At the same time, we find that in 3hourly means the linear correlation coefficient between the am index and P_{α} is 0.867, meaning that 75.2% of the variation of am is explained by P_{α}. Furthermore, in this paper and in Paper 1, by sorting the data by [B_{Y}]_{GSEQ} we have shown that the R–M effect dominates the am response in both average levels and the occurrence of large events. The patterns in the Fyear spectrogram plots presented in Figure 10 for P_{α,} am and the four aσ indices are very similar, it is just the amplitude of the pattern varies with the amplification factor for the geomagnetic index in question. Furthermore, sorting by the polarity of the IMF [B_{Y}]_{GSEQ} component confirms the patterns originate from the R–M effect. Hence we conclude that it is not that the R–M effect is partially responsible for the semiannual variation, rather it is almost wholly responsible. However, Figures 12 and 13 show us how the amplification of a very small semiannual variation in P_{α} to give a much larger semiannual variation in am occurs. The R–M effect only causes a slight asymmetry in how much average P_{α} (or the occurrence of large P_{α} events) is enhanced for the “favourable” polarity of the IMF [B_{Y}]_{GSEQ} at a given equinox compared to how much it is decreased by the “nonfavourable” polarity of the IMF [B_{Y}]_{GSEQ}. However, because of a nonlinearity in the am response this drives a larger asymmetry in geomagnetic activity. We note that in their original paper, Russell & McPherron (1973) did comment on the need for this amplification; however, the discussion was conflated with that on the occurrence of large storms which, as discussed in the introduction, raises the involvement of large, negative [B_{Z}]_{GSEQ} which, when it becomes dominant over [B_{Y}]_{GSEQ} (i.e., clock angles in GSEQ become greater than about 120°), make the R–M have the opposite effect to that for the traditional R–M effect as proposed by Russell & McPherron (1973).
However, there is still a puzzle to resolve here because when we look at a scatter plot of all am data as a function of P_{α} (as for the 3hourly means in Fig. 15a) we see only a very slight nonlinearity but not one large enough to generate the contrasting the behaviours of am and P_{α} seen in Figure 12.
Figures 14 and 15b both demonstrate that am depends on both P_{α} and the solar wind dynamic pressure, p_{SW}. These two factors have common elements but differ in the terms that depend on the IMF and its orientation. For the bestfit coupling exponent used here of α = 0.44, as derived for all averaging timescales by Lockwood et al. (2019b), P_{α} varies as where is the mean solar wind ion mass, N_{SW} is the number density, V_{SW} is the solar wind speed, B the IMF magnitude and θ is the IMF clock angle in the GSM frame. On the other hand, and so the two factors are different but also have terms in common and so there will be some systematic dependence of one on the other. Figure 12b shows that the fit residuals of the best fit of P_{α}/P_{o} to am depends on p_{SW} implying there is an additional dependence of am on p_{SW}. Interestingly the fit residuals (Fig. 16), show the equinoctial F–UT pattern (Cliver et al., 2000) which increases in amplitude linearly with increased p_{SW} (Fig. 17). Given that the equinoctial pattern is present in am but absent in P_{α} this strongly implies that p_{SW} introduces and additional variation to geomagnetic activity that includes the equinoctial F–UT pattern. Naturally, because these fit residuals show an equinoctial F–UT pattern, they will also show an additional semiannual variation that is not accounted for by P_{α} on its own. Figure 18a confirms this and that the amplitude of the semiannual variation increases with and p_{SW}. An interesting finding is presented in Figure 18b that shows that these fit residuals also show the UT variation which also increases in amplitude as p_{SW} increases.
Figure 19 investigates further the complex interrelationship between p_{SW} and P_{α} and its implications. The combined am, p_{SW} and P_{α}/P_{o} data are averaged into 36equal width bins of F and divided into the two polarities of IMF [B_{Y}]_{GSEQ} (only to show there is no significant difference between the two) and 5 quantile ranges of p_{SW}. The averages of am are plotted as a function of the corresponding averages of P_{α}/P_{o} for each of the 5 quantile ranges of p_{SW}: the [B_{Y}]_{GSEQ} polarities are distinguished by the symbols used, and the p_{SW} quantile ranges by the colors used. The black lines are 3rdorder polynomial fits to the data and are drawn only between the largest and smallest mean value of P_{α}/P_{o} for the quantile range of p_{SW} in question: (note that, because of the common factors in p_{SW} and P_{α}, that range varies with p_{SW}). These polynomial fit lines are replotted in Figure 19c, using the same colour scheme as for the data points in part (a), and superposed on a scatter plot of averages for all data in time bins of 1/36 yr (the grey points). It can be seen that at each constant p_{SW} value, am has nonlinear dependence on P_{α}/P_{o} which is not apparent in the overall scatter plot. Figures 19a and 19c show the variation of am with P_{α}/P_{o} in 5 equal sized quantile ranges of p_{SW}. Figures 19b and 19d show the corresponding variation of am with p_{SW} in 5 equal sized quantile ranges of P_{α}/P_{o}: they show that am also depends nonlinearly on p_{SW} at a fixed P_{α}/P_{o}.
Fig. 19 Analysis of the effects on am of (left column) power input to the magnetosphere P_{α}/P_{o} at a given solar wind dynamic pressure p_{SW} and of (right column) the effects of p_{SW} at a given P_{α}/P_{o}. Values of am/<am>_{all} are plotted as a function of corresponding normalised values of (left) P_{α} and (right) p_{SW}. In (a) data are averaged into 5 quantile ranges of the normalised solar wind dynamic pressure p_{SW}/<p_{SW}>_{all}, 36 bins of timeofyear F, and by the polarity of the IMF Y component in GSEQ, [B_{Y}]_{GSEQ}: the means of am/<am>_{all} are plotted against the corresponding means of P_{α}/P_{o}. In (b) data are averaged into 5 quantiles of normalised power input to the magnetosphere P_{α}/P_{o}, 36 bins of timeofyear F, and by the [B_{Y}]_{GSEQ} polarity and the means of am/<am>_{all} are plotted against the corresponding means of p_{SW}/<p_{SW}>_{all}. The black lines in the upper panels are 3rdorder polynomial fits in each case, fitted to data for all [B_{Y}]_{GSEQ} drawn over the range of the averaged data for the quantile range in question. The points are coloured according to the quantile range as given by the key. The lines are also plotted (this time using the same colours as the points in the upper panels) in the corresponding lower panel. The grey points show a scatter plot of all 10day average values. 
The key point about Figure 19 is that when we look at a constant solar wind dynamic pressure p_{SW} we do see a highly nonlinear increase in geomagnetic activity (am) with the power input into the magnetosphere P_{α}, of the kind needed to explain the amplification of the geomagnetic response (and hence of the semiannual variation) to a relatively small R–M effect in P_{α}. This is lost in the overall scatter plots of P_{α} against am (which appear linear with some scatter) because of the common factors in p_{SW} and P_{α}. However, Figure 19 strongly suggest that p_{SW} has a distinct and different physical effect from P_{α} and this is supported by the findings presented in the current paper that the amplification of the am response, the equinoctial pattern and the UT variation arise from the effect of p_{SW} rather than the effect of P_{α}. The concept of a distinct and different influence of p_{SW} on the nightside magnetosphere and the substorm phenomenon has been reported several times in the literature. Increases in p_{SW} have been reported to trigger onsets of full substorm expansion phases by, for example, Schieldge & Siscoe (1970), Kokubun et al. (1977) and Yue et al. (2010) and Caan et al. (1973) showed that, statistically, the magnetic energy density in the nearEarth tail lobes (that powers substorm expansion phases) was increased by both by prior intervals of southwardpointing IMF and solar wind dynamic pressure, p_{SW}. An effect of p_{SW} was directly demonstrated directly by Karlsson et al. (2000) who showed that nearEarth tail energy content was reduced if p_{SW} decreased and that sudden decreases caused quenching of any substorm expansion that had recently begun. Lockwood, 2013; Lockwood et al., 2014) and Finch et al. (2008) have noted that indices that are strongly influenced by the substorm current wedge have a strong dependence on and hence p_{SW}. Finch et al. (2008) noted that only magnetometer stations closest to the nightside auroral electrojet showed both the equinoctial F–UT pattern and the dependence. Lockwood (2013) argues that this implies that p_{SW} influences the auroral electrojet and the equinoctial F–UT pattern by constraining the nearEarth tail such that the appending of open flux to the tail lobes of the magnetosphere increases the field and hence the stored energy density and magnetic shear across the crosstail current sheet. This does not happen further down the tail where the magnetopause boundary becomes aligned with the solar wind flow and the magnetic pressure in the lobes, and hence the field, is set by the balance with static pressure in interplanetary space. Hence, at these large negative X_{GSE}, adding open flux just causes the tail to flare in crosssectional area and the field in the lobes and the magnetic shear across, and hence total current in, the crosstail current all remain constant. Thus there is considerable evidence for a separate role of solar wind dynamic pressure p_{SW} in modulating the storage and release of the energy extracted from the solar wind, P_{α} in the nearEarth tail of the magnetosphere.
8 Conclusions
This paper has presented a collection of empirical results concerning the semiannual variation in geomagnetic activity. Interpretation will largely be left to a subsequent paper in this series which will make use of empirical and numerical MHD models of the magnetosphere. In this section we bring together the summary conclusions of each section of this paper.
We have presented an initial study of 2 years’ data on the semiannual variation in transpolar voltage and flux transfer. We find there is a semiannual variation in transpolar voltage but it is not as large in amplitude as that in geomagnetic activity which is consistent with the latter showing a nonlinear (quadratic) variation with transpolar voltage. We find that during the persistent minimum of the UT variation in geomagnetic activity defined in Paper 1 there is a persistent decrease in transpolar voltage which is potentially consistent with a decrease in reconnection voltage in the crosstail current sheet. Confirmation of these findings using a larger dataset of transpolar voltage measurements will be presented in a later paper in this series, but the great consistency with which we find the features described (in both polar caps and in each of the two years studied here) leads us to believe that they are real effects.
We have studied the response of geomagnetic activity to power input to the magnetosphere, estimated using interplanetary data from 1995 onwards which is relatively free of data gaps. This is important because previous work has demonstrated that large errors are introduced into studies of solar wind magnetosphere coupling by data gaps (Lockwood et al., 2019a). We find no consistent variation in the response delay with timeofyear, F. There is a 2 sigma spread in the lag values of approximately 10–100 min and some systematic change in the distribution with activity level but, overall, a lag of 60 min is optimum for the whole dataset. Using this lag we have shown that the pattern of variations in Fyear spectrogram plots in geomagnetic activity and power input into the magnetosphere are very similar, both for average values and the occurrence of large events, and that the R–M effect is at the heart of both. However the effect on power input into the magnetosphere is very small and there is a nonlinear amplification of the semiannual variation in the geomagnetic response such that a very small asymmetry in power input into the magnetosphere P_{α} between the “favourable” and “unfavourable” polarities of [B_{Y}]_{GSEQ} generates a greatly amplified geomagnetic semiannual response.
The origin of this amplification will be discussed in full in a later paper. However, as discussed in the last section, the analysis presented here indicates strongly that it is associated with solar wind dynamic pressure and its role in squeezing the near Earth tail and so modulating the storage and release of energy extracted from the solar wind in the nearEarth tail where the tail radius is still flaring with the −X_{GSE} coordinate and so dynamic pressure is a factor. In this paper we have shown that the equinoctial pattern is found in the residuals of fits of P_{α} to the am index and that the amplitude of these equinoctial patterns in the am fit residuals increases linearly with solar wind dynamic pressure p_{SW}. Similarly, the UT variation in am is found in these fit residuals and also increases in amplitude with solar wind dynamic pressure. This strongly suggests the solar wind pressure is a distinct influence of geomagnetic activity. In a later paper we will use empirical and numerical MHD models of the magnetosphere to study the role of solar wind dynamic pressure as a function of dipole tilt and see if the models can produce the influence of solar wind dynamic pressure on the equinoctial F–UT pattern, on the UT variation, and on the amplification of the variation of power input that is inferred here. One finding that needs explanation is why the amplification of the am response is greater at the equinoxes, as shown by Figure 13. With understanding gained about these effects we aim to revist the paradox of large geomagnetic storms and the fact that their occurrence peaks at the equinoxes, yet they are often driven by large southward field in the GSEQ frame, which should make them less common at the equinoxes.
Acknowledgments
This work is supported by a number of grants. CEW, MJO, CJS and ML at the University of Reading are supported by STFC consolidated grant number ST/M000885/1 and CEW is also supported by STFC grant ST/R000921/1 and NERC grant NE/P017274/1. The work of ML, LAB and MJO at University of Reading is also supported by the SWIGS NERC Directed Highlight Topic Grant number NE/P016928/1/. The work of JCC at the University of Southampton is supported by the UK Natural Environment Research Council (NERC) grant number NE/L007177/1 and by Science and Technology Facilities Council (STFC) Ernest Rutherford grant ST/L002809/1 and Consolidated grant ST/R000719/1. Funding for KAW at University of Saskatchewan was provided by the Canadian Foundation for Innovation (CFI), the Province of Saskatchewan, and a Discovery Grant from the Natural Sciences and Engineering Research Council (NSERC) of Canada. Initial work by KAW for this paper was carried out at University of Reading on sabbatical leave from University of Saskatchewan. The authors are also grateful to the staff of the Space Physics Data Facility, NASA/Goddard Space Flight Center, who prepared and made available the OMNI2 dataset used: these interplanetary data were downloaded from http://omniweb.gsfc.nasa.gov/ow.html. We are also grateful to the staff of L’École et Observatoire des Sciences de la Terre (EOST), a joint of the University of Strasbourg and the French National Center for Scientific Research (CNRS) and the International Service of Geomagnetic Indices (ISGI) for making the am and aσ indices available from http://isgi.unistra.fr/data_download.php. The editor thanks two anonymous reviewers for their assistance in evaluating this paper.
Appendix
Relationships between the am index and the transpolar voltage
Figure 4 of the main text presents a scatter plot of daily means of the transpolar voltage Φ_{PC} against the means of the am index for the same day. The Φ_{PC} values are derived from observations in 2001–2002 by the DMSP satellites and are normalised to an ideal satellite path along the 0618 MLT polar cap diameter using the procedure of Lockwood et al. (2009). The error bars in the plot are plus and minus one standard error in the means. The mauve line is the bestfit 3rd order polynomial fit and the grey area is bounded by the 2sigma uncertainties in that fit.
The procedure employed to make the fit was as follows. The data points were fitted 10,000 times using a 3rdorder polynomial fit. For each fit, every data point was shifted in both am and Φ_{PC}, by an error drawn at random from a Gaussian distribution with mean value equal to the observed mean value for the point in question and of standard deviation equal to the observed standard deviation around that mean. The c.d.f. of the ensemble of fitted am values at each value Φ_{PC} was then computed and the bestfit for that Φ_{PC} then taken to be the median value of this distribution. The maximum and minimum uncertainties were taken to be the 5% and 95% percentiles of these distributions (i.e. the upper and lower 2sigma points of the ensemble).
The best fit (ensemble median) is(A.1)where Φ_{PC} is in kV and am is in nT. The upper 2sigma level is(A.2)and the lower 2sigma level is(A.3)
In addition, in Figure A.1 we present the same data but plotted with am as the abscissa (along the horizontal axis) and Φ_{PC} as the ordinate (along the vertical axis). The fit procedure described above was repeated but this time the c.d.f. of Φ_{PC} was taken at each fitted am value and the best fit, and its 2sigma uncertainties, derived in the same way. This gives the corresponding polynomial fits to allow us to convert am into the corresponding Φ_{PC} value. The best fit is(A.4)the upper 2sigma level is(A.5)and the lower 2sigma level is(A.6)
Fig. A.1 Scatter plot of daily means of the transpolar voltage Φ_{PC}, as determined for 2001–2002 by the DMSP satellites and normalised for the satellite track to an ideal 06–18 MLT path using the procedure of Lockwood et al. (2009), as a function of the means of the am index for the same day. The error bars are plus and minus one standard error in the means. The mauve line is the bestfit 3rd order polynomial fit of Φ_{PC} and the grey area is bounded by the 2sigma uncertainty level in the fit. 
References
 Adebesin BO. 2016. Investigation into the linear relationship between the AE, Dst and ap indices during different magnetic and solar activity conditions. Acta Geod Geophys 51(2): 315–331. https://doi.org/10.1007/s4032801501282. [CrossRef] [Google Scholar]
 Bargatze LFB, McPherron RL, Baker DN. 1986. Solar windmagnetosphere energy input functions. In: Solar windmagnetosphere coupling, Kamide Y., Slavin J.A. (Eds.), Terra Scientific Publishing Company, Tokyo, Japan, and D. Reidel Publishing Company, Dordrecht, Holland, pp. 101–109. https://doi.org/10.1007/9789400947221_7. [Google Scholar]
 Boudouridis A, Zesta E, Lyons LR, Anderson PC, Lummerzheim D. 2005. Enhanced solar wind geoeffectiveness after a sudden increase in dynamic pressure during southward IMF orientation. J Geophys Res 110: A05214. https://doi.org/10.1029/2004JA010704. [NASA ADS] [CrossRef] [Google Scholar]
 Caan MN, McPherron RL, Russell CT. 1973. Solar wind and substormrelated changes in the lobes of the geomagnetic tail. J Geophys Res 78(34): 8087–8096. https://doi.org/10.1029/ja078i034p08087. [CrossRef] [Google Scholar]
 Chambodut A, Marchaudon A, Menvielle M, ElLemdani F, Lathuillere C. 2013. The Kderived MLT sector geomagnetic indices. Geophys Res Lett 40: 4808–4812. https://doi.org/10.1002/grl.50947. [CrossRef] [Google Scholar]
 Chandler MO, Fuselier SA, Lockwood M, Moore TE. 1999. Evidence of component magnetic merging equatorward of the cusp. J Geophys Res 104: 22623–22648. https://doi.org/10.1029/1999JA900175. [CrossRef] [Google Scholar]
 Chu X, McPherron RL, Hsu TS, Angelopoulos V. 2015. Solar cycle dependence of substorm occurrence and duration: Implications for onset. J Geophys Res Space Phys 120: 2808–2818. https://doi.org/10.1002/2015JA021104. [CrossRef] [Google Scholar]
 Cliver EW, Kamide Y, Ling AG. 2000. Mountains versus valleys: Semiannual variation of geomagnetic activity. J Geophys Res 105: 2413–2424. https://doi.org/10.1029/1999JA900439. [NASA ADS] [CrossRef] [Google Scholar]
 Cowley SWH. 1991. Acceleration and heating of space plasmas  Basic concepts. Ann Geophys 9: 176–187. [Google Scholar]
 Cowley SWH, Lockwood M. 1992. Excitation and decay of solarwind driven flows in the magnetosphereionosphere system. Ann Geophys 10: 103–115. [Google Scholar]
 Crooker NU, Cliver EW, Tsurutani BT. 1992. The semiannual variation of great geomagnetic storms and the postshock RussellMcpherron effect preceding coronal mass ejecta. Geophys Res Lett 19(5): 429–432. https://doi.org/10.1029/92GL00377. [CrossRef] [Google Scholar]
 de La Sayette P, Berthelier A. 1996. The am annualdiurnal variations 1959–1988: A 30year evaluation. J Geophys Res 101(A5): 10653–10663. https://doi.org/10.1029/96JA00165. [CrossRef] [Google Scholar]
 Echer E, Alves MV, Gonzalez WD. 2005. A statistical study of magnetic cloud parameters and geoeffectiveness. J Atmos Terr Phys 67(10): 839–852. https://doi.org/10.1016/j.jastp.2005.02.010. [CrossRef] [Google Scholar]
 Finch ID, Lockwood M. 2007. Solar windmagnetosphere coupling functions on timescales of 1 day to 1 year. Ann Geophys 25: 495–506. https://doi.org/10.5194/angeo254952007. [CrossRef] [Google Scholar]
 Finch ID, Lockwood M, Rouillard AP. 2008. The effects of solar wind magnetosphere coupling recorded at different geomagnetic latitudes: Separation of directlydriven and storage/release systems. Geophys Res Lett 35: L21105. https://doi.org/10.1029/2008GL035399. [CrossRef] [Google Scholar]
 Karlsson SBP, Opgenoorth HJ, Eglitis P, Kauristie K, Syrjäsuo M, Pulkkinen TI, Lockwood M, Nakamura R, Reeves G, Romanov S. 2000. Solar wind control of magnetospheric energy content: substorm quenching and multiple onsets. J Geophys Res 105: 5335–5356. https://doi.org/10.1029/1999JA900297. [CrossRef] [Google Scholar]
 Kilpua EKJ, Balogh A, von Steiger R, Liu YD. 2017. Geoeffective properties of solar transients and stream interaction regions. Space Sci Rev 212(3/4): 1271–1314. https://doi.org/10.1007/s1121401704113. [Google Scholar]
 Kokubun S, McPherron RL, Russell CT. 1977. Triggering of substorms by solar wind discontinuities. J Geophys Res 82(1): 74–86. https://doi.org/10.1029/ja082i001p00074. [CrossRef] [Google Scholar]
 Lee DY, Lyons LR, Yumoto K. 2004. Sawtooth oscillations directly driven by solar wind dynamic pressure enhancements. J Geophys Res 109: A04202. https://doi.org/10.1029/2003JA010246. [Google Scholar]
 Le Mouël JL, Blanter E, Chulliat A, Shnirman M. 2004. On the semiannual and annual variations of geomagnetic activity and components. Ann Geophys 22: 3583–3588. https://doi.org/10.5194/angeo2235832004. [CrossRef] [Google Scholar]
 Li D, Yao S. 2020. Stronger southward magnetic field and geoeffectiveness of ICMEs containing prominence materials measured from 1998 to 2011. Astrophys J 891(1): 79. https://doi.org/10.3847/15384357/ab7197. [CrossRef] [Google Scholar]
 Lockwood M. 1991. On flow reversal boundaries and crosscap potential in average models of high latitude convection. Planet Space Sci 39: 397–409. https://doi.org/10.1016/00320633(91)90002R. [CrossRef] [Google Scholar]
 Lockwood M. 2004. Solar outputs, their variations and their effects of Earth. In: “The Sun, Solar Analogs and the Climate”, Proc. Saas‐Fee Advanced Course, 34 by J.D. Haigh, M. Lockwood and M.S. Giampapa, Rüedi I., Güdel M., Schmutz W. (Eds.), Springer, pp. 107–304. [Google Scholar]
 Lockwood M. 2013. Reconstruction and prediction of variations in the open solar magnetic flux and interplanetary conditions. Liv Rev Sol Phys 10: 4. https://doi.org/10.12942/lrsp20134. [Google Scholar]
 Lockwood M. 2019. Does adding solar wind Poynting flux improve the optimum solar wind – magnetosphere coupling function? J Geophys Res Space Phys 124(7): 5498–5515. https://doi.org/10.1029/2019JA026639. [CrossRef] [Google Scholar]
 Lockwood M, Lanchester BS, Morley SK, Throp K, Milan SE, Lester M, Frey HU. 2006. Modeling the observed proton aurora and ionospheric convection responses to changes in the IMF clock angle: 2. Persistence of ionospheric convection. J Geophys Res 111: A02306. https://doi.org/10.1029/2003JA010307. [Google Scholar]
 Lockwood M, Hairston MR, Finch ID, Rouillard AP. 2009. Transpolar voltage and polar cap flux during the substorm cycle and steady convection events. J Geophys Res 114: A01210. https://doi.org/10.1029/2008JA013697. [Google Scholar]
 Lockwood M, Nevanlinna H, Barnard L, Owens MJ, Harrison RG, Rouillard AP, Scott CJ. 2014. Reconstruction of geomagnetic activity and nearEarth interplanetary conditions over the Past 167 years: 4. NearEarth solar wind speed, IMF, and open solar flux. Ann Geophys 32: 383–399. https://doi.org/10.5194/angeo323832014. [Google Scholar]
 Lockwood M, Owens MJ, Barnard LA, Bentley S, Scott CJ, Watt CE. 2016. On the origins and timescales of geoeffective IMF. Space Weather 14: 406–432. https://doi.org/10.1002/2016SW001375. [CrossRef] [Google Scholar]
 Lockwood M, Bentley S, Owens MJ, Barnard LA, Scott CJ, Watt CE, Allanson O. 2019a. The development of a space climatology: 1. Solarwind magnetosphere coupling as a function of timescale and the effect of data gaps. Space Weather 17: 133–156. https://doi.org/10.1029/2018SW001856. [Google Scholar]
 Lockwood M, Bentley S, Owens MJ, Barnard LA, Scott CJ, Watt CE, Allanson O, Freeman MP. 2019b. The development of a space climatology: 2. The distribution of power input into the magnetosphere on a 3hourly timescale. Space Weather 17: 157–179. https://doi.org/10.1029/2018SW002016. [CrossRef] [Google Scholar]
 Lockwood M, Bentley S, Owens MJ, Barnard LA, Scott CJ, Watt CE, Allanson O, Freeman MP. 2019c. The development of a space climatology: 3. The evolution of distributions of space weather parameters with timescale. Space Weather 17: 180–209. https://doi.org/10.1029/2018SW002017. [CrossRef] [Google Scholar]
 Lockwood M, Chambodut A, Finch ID, Barnard LA, Owens MJ, Haines C. 2019d. Timeofday/timeofyear response functions of planetary geomagnetic indices. J Space Weather Space Clim 9: A20. https://doi.org/10.1051/swsc/2019017. [Google Scholar]
 Lockwood M, Owens MJ, Barnard LA, Haines C, Scott CJ, McWilliams KA, Coxon JC. 2020. Semiannual, annual and Universal Time variations in the magnetosphere and in geomagnetic activity: 1. Geomagnetic data. J Space Weather Space Clim 10: 23. https://doi.org/10.1051/swsc/2020023. [Google Scholar]
 Lühr H, Lockwood M, Sandholt PA, Hansen TL, Moretto T. 1996. Multiinstrument groundbased observations of a travelling convection vortex event. Ann Geophys 14: 162–181. https://doi.org/10.1007/s005859960162z. [CrossRef] [Google Scholar]
 Lukianova R. 2003. Magnetospheric response to sudden changes in solar wind dynamic pressure inferred from polar cap index. J Geophys Res 108(A12): 1428. https://doi.org/10.1029/2002JA009790. [CrossRef] [Google Scholar]
 Mayaud PN. 1980. Derivation, meaning and use of geomagnetic indices. In: Geophysical monograph, 22. American Geophysical Union, Washington, DC. https://doi.org/10.1029/GM022. [Google Scholar]
 McPherron RL, Baker DN, Pulkkinen TI, Hsu TS, Kissinger J, Chu X. 2013. Changes in solar wind–magnetosphere coupling with solar cycle, season, and time relative to stream interfaces. J Atmos SolTerr Phys 99: 1–13. https://doi.org/10.1016/j.jastp.2012.09.003. [CrossRef] [Google Scholar]
 Morley SK, Lockwood M. 2006. A numerical model of the ionospheric signatures of timevarying magnetic reconnection: 3. Quasiinstantaneous convection responses in the CowleyLockwood paradigm. Ann Geophys 24: 961–972. https://doi.org/10.5194/angeo249612006. [CrossRef] [Google Scholar]
 Palmroth M, Pulkkinen TI, Janhunen P, McComas DJ, Smith CW, Koskinen HEJ. 2004. Role of solar wind dynamic pressure in driving ionospheric Joule heating. J Geophys Res 109: A11302. https://doi.org/10.1029/2004JA010529. [CrossRef] [Google Scholar]
 Pulkkinen TI, Dimmock AP, Lakka A, Osmane A, Kilpua E, Myllys M, Tanskanen EI, Viljanen A. 2016. Magnetosheath control of solar wind–magnetosphere coupling efficiency. J Geophys Res Space Phys 121: 8728–8739. https://doi.org/10.1002/2016JA023011. [CrossRef] [Google Scholar]
 Russell CT, McPherron RL. 1973. Semiannual variation of geomagnetic activity. J Geophys Res 78: 82–108. https://doi.org/10.1029/JA078i001p00092. [Google Scholar]
 Schieldge JP, Siscoe GL. 1970. A correlation of the occurrence of simultaneous sudden magnetospheric compressions and geomagnetic bay onsets with selected geophysical indices. J Atmos Terr Phys 32(11): 1819–1830. https://doi.org/10.1016/00219169(70)90139x. [CrossRef] [Google Scholar]
 Stauning P, Troshichev OA. 2008. Polar cap convection and PC index during sudden changes in solar wind dynamic pressure. J Geophys Res 113: A08227. https://doi.org/10.1029/2007JA012783. [CrossRef] [Google Scholar]
 Taylor JR, Lester M, Yeoman TK. 1994. A superposed epoch analysis of geomagnetic storms. Ann Geophys 12(7): 612–624. https://doi.org/10.1007/s0058599406124. [CrossRef] [Google Scholar]
 Vasyliunas VM, Kan JR, Siscoe GL, Akasofu SI. 1982. Scaling relations governing magnetospheric energy transfer. Planet Space Sci 30: 359–365. https://doi.org/10.1016/00320633(82)900411. [CrossRef] [Google Scholar]
 Webb DF, Cliver EW, Crooker NU, St. Cyr OC, Thompson BJ. 2000. Relationship of halo coronal mass ejections, magnetic clouds, and magnetic storms. J Geophys Res Space Phys 105(A4): 7491–7508. https://doi.org/10.1029/1999JA000275. [NASA ADS] [CrossRef] [Google Scholar]
 Xie H, Gopalswamy N, Manoharan PK, Lara A, Yashiro S, Lepri S. 2006. Longlived geomagnetic storms and coronal mass ejections. J Geophys Res Space Phys 111(A1): A01103. https://doi.org/10.1029/2005JA011287. [Google Scholar]
 Yue C, Zong QG, Zhang H, Wang YF, Yuan CJ, Pu ZY, Fu SY, Lui ATY, Yang B, Wang CR. 2010. Geomagnetic activity triggered by interplanetary shocks. J Geophys Res 115: A00I05. https://doi.org/10.1029/2010JA015356. [Google Scholar]
 Zhao H, Zong QG. 2012. Seasonal and diurnal variation of geomagnetic activity: RussellMcPherron effect during different IMF polarity and/or extreme solar wind conditions. J Geophys Res 117: A11222. https://doi.org/10.1029/2012JA017845. [Google Scholar]
Cite this article as: Lockwood M, McWilliams KA, Owens MJ, Barnard LA, Watt CE, et al. 2020. Semiannual, annual and Universal Time variations in the magnetosphere and in geomagnetic activity: 2. Response to solar wind power input and relationships with solar wind dynamic pressure and magnetospheric flux transport. J. Space Weather Space Clim. 10, 30.
All Tables
All Figures
Fig. 1 The southward interplanetary magnetic field (IMF) in the Geocentric Solar Magnetospheric (GSM) frame, [B_{Z}]_{GSM} as a function of its value in the Geocentric Solar Equatorial (GSEQ) frame, [B_{Z}]_{GSEQ} from a survey of 24 years’ interplanetary and geomagnetic activity data (1995–2017, inclusive). The number of samples in 0.2 nTby0.2 nT bins, N, as a ratio of the total number, ΣN, is colour contoured as a function of [B_{Z}]_{GSM} and [B_{Z}]_{GSEQ}. The data are averages over 3h intervals for 1995–2017, inclusive, which are the intervals over which each am value is compiled, shifted by an average optimum am response lag dt_{o} = 1 h (see Sect. 3.3). (a) Is for all data whereas (b–d) are for am exceeding, respectively, its 90%, 95% and 99% quantile for the years studied. The diagonal mauve lines are [B_{Z}]_{GSEQ} = [B_{Z}]_{GSM}. 

In the text 
Fig. 2 The test devised by Vasyliunas et al. (1982) of proposed IMF orientation factors, A_{θ}(θ), where θ is the IMF clock angle in the GSM frame: (orange line) sin^{2}(θ/2); (cyan line) sin^{3}(θ/2); (blue line) sin^{4}(θ/2); (mauve line) sin^{5}(θ/2); and (black line) U(θ)cos(θ), where U(θ) = 0 for [B_{Z}]_{GSM} > 0 and U(θ) = 1 for [B_{Z}]_{GSM} ≤ 0. The black dots are s(SML/G) + c where s and c are the bestfit linear regression coefficients, SML is the SuperMAG westward electrojet index, and G = P_{α}/A_{θ}, where P_{α} is the best fit estimate of the power input into the magnetosphere. The dark pink shaded area is the range over which θ varies due to the R–M effect for [B_{Z}]_{GSEQ} = 0 (i.e., IMF lying in the solar equatorial plane) and the range of variation of the dipole tilt over a full year caused by Earth’s rotational axis tilt. The light pink area is for the full range caused by the annual variation plus the diurnal variation due to the offset of Earth’s rotational and magnetic axes. The light and dark orange regions are the corresponding ranges for [B_{Z}]_{GSEQ} = B (i.e., purely southward field normal to the solar equatorial plane). 

In the text 
Fig. 3 Simultaneous data from the years 2001 and 2002. In each case the data have been averaged into 365 equalsized bins of timeof year F (one day long for these nonleap years) and then a running mean taken over 27 bins to cover whole solar rotation intervals. From top to bottom: (a) the halfwave rectified dawntodusk electric field, E_{DD}, in interplanetary space; (b) the AU index; (c) the AL index; (d) the polar cap flux, W_{PC}, from the DMSP satellite data; (e) the transpolar voltage, Φ_{PC}, from the DMSP satellite data; (f) the reconnection efficiency assuming that the crosssectional radius of the magnetosphere is 15 R_{E} (where a mean Earth radius 1 R_{E} = 6370 km), η_{15}; (g) the am geomagnetic index; and (h) the ratio of the am index to the transpolar voltage, am/Φ_{PC}. In panels (d) and (e) for W_{PC} and Φ_{PC}, the red and blue lines are for the northern and southern polar caps, respectively. 

In the text 
Fig. 4 Scatter plot of daily means of the am index as a function of the corresponding daily mean transpolar voltage, Φ_{PC}, as determined for 2001–2002 by DMSP satellites (principally F13) and normalised for the satellite track to an ideal 06–18 MLT path using the procedure of Lockwood et al. (2009). The error bars are plus and minus one standard error in the means. The mauve line is the bestfit 3rd order polynomial fit and the grey area is bounded by the 2sigma uncertainty level in that fit. The fitting procedure is given in Appendix to this paper, along with polynomial expressions for the best fit and the uncertainty band edges. These can be used to estimate the am level associated with a given Φ_{PC}: Appendix also gives the corresponding 3rd order polynomials that allow computation of the Φ_{PC} value associated with a given am value. 

In the text 
Fig. 5 Timeofyear/time–ofday (F–UT) plots for transpolar voltage and polar cap flux for 2001–2002, inclusive from the DMSP satellites. Data are sorted into 24 equalwidth bins in F and 16 equal width bins of UT using the central time of the polar cap traversal. The top panels (a–d) are for passes of the northern polar cap, the lower panels (e–h) are for passes of the southern polar cap. (a and e) Show the transpolar voltage Φ_{PC}; (b and f) show the polar cap flux, W_{PC}; (c and g) show the simultaneous am index value (linearly interpolated from the threehourly values to the time of the centre of the polar cap crossing) and (e and h) shows the number of data points, n_{bin} in each F–UT bin. 

In the text 
Fig. 6 Cumulative probability distributions (c.d.f.s) of the power input to the magnetosphere, P_{α}, as a ratio of its mean value P_{o} = <P_{α}>_{all} for the full interval of nearcontinuous interplanetary data (1995–2017, inclusive). The P_{α}/P_{o} data are divided into 288 bins: 36 equalsized bins in fraction of year F and the same eight 3h Universal Time (UT) bins over which the range am index is evaluated. For each F–UT bin the c.d.f. of P_{α}/P_{o} as a ratio of the mean value for that bin <P_{α}/P_{o}>_{F,UT} was plotted and colored according to the value of that ratio. The black line is for all P_{α}/P_{o} data, the orange line for the F–UT bin giving the largest <P_{α}/P_{o}>_{F,UT} and the green line for the F–UT bin giving the smallest <P_{α}/P_{o}>_{F,UT}. 

In the text 
Fig. 7 The relationship of am and normalized power input to the magnetosphere P_{α}/P_{o} for large averaging timescales τ: (a) τ = 1 year and (b) τ = 10 days. In (a) the black circles are for all data, for which the linear correlation coefficient is r = 0.97 and the best fit OLS (Ordinary LeastSquares) linear regression, shown by the green line, is (P_{α}/P_{o}) = s·am + c. Averages for the 8 UTs of the am index are shown by the squares, colored by the scale shown at the top of the figure: it can be seen that am is persistently a little smaller than the average response at 0–9 UT and persistently a little larger at 15–4 UT. In (b) the points are colored by the separation of timeofyear F from the value at the closest equinox, Δt_{eq}. The linear correlation for all data is r = 0.96 and the best fit linear regression is again the green line. The am response at low Δt_{eq} (around the equinoxes) is persistently greater than at high Δt_{eq} (around the solstices) showing that there is a contribution to the semiannual variation in am is not associated with that in (P_{α}/P_{o}) and hence not directly attributable to the R–M effect on solar windmagnetosphere coupling. The black and orange lines are, respectively, the bestfit linear regressions for around the equinoxes (Δt_{eq} ≤ 0.125) and around the solstices (Δt_{eq} > 0.125). The ordinary least squares (OLS) linear regression coefficients and errors are given in Table 1. 

In the text 
Fig. 8 Peak correlation r_{p} between am and 3hourly means of P_{α}/P_{o} as a function of the logarithm of the lag dt_{p} giving that peak correlation, for data from 1995–2017, inclusive. Oneminute P_{α}/P_{o} data were averaged into hourly and then 3hourly intervals, centred on the midpoints of the am data intervals, minus a response lag dt that was varied between −10 min and 5 h in steps of one minute. At each dt the correlation between P_{α}(t − dt)/P_{o} and am(t) is evaluated for each of the data subsets studied. The data are divided and colourcoded into 9 quantile ranges of am given by the legend. We use the notation that 20% of all the data have am values lower than q(0.2). The red dots are for all data, q(0) < am ≤ q(1). Five of the quantile ranges each contain 20% of the data: q(0) < am ≤ q(0.2), q(0.2) < am ≤ q(0.4), q(0.4) < am ≤ q(0.6), q(0.6) < am ≤ q(0.8) and q(0.8) < am ≤ q(1). In addition, the plot studies the top 10%, 5% and 1% of all am values, respectively q(0.9) < am ≤ q(1), q(0.95) < am ≤ q(1), and q(0.99) < am ≤ q(1). Correlations are shown for F–UT–m bins that have more than 16 samples and for which the correlation significance exceeds 95%. 

In the text 
Fig. 9 Distribution of response lags for the 288 F–UT bins studied for the response of the am index to the power input to the magnetosphere P_{α}/P_{o} for 1995–2017, inclusive. (Top) The red line shows the distribution of the lags giving peak correlation dt_{p} between P_{α}/P_{o} and am for all the data. (Bottom) The same analysis but the am data have been further subdivided into the 8 other quantile ranges used in Figure 8. The distributions are taken in 1min ranges and then smoothed with a 10min running mean. 

In the text 
Fig. 10 Top panel: plots of annual means of an index or variable x, <x>_{1yr}, for the years 1995–2017, inclusive, where x is: (a) the power input into the magnetosphere, P_{α}; (b) the am index (in black) and the number of substorm onsets derived from the SML index, N_{o} (in green); (c) the aσ indices for dawn (orange), noon (cyan), dusk (mauve) and midnight (blue). Middle and lower panels: yearF spectrogram plots of the means in 36 equalsized bins of F, divided by the mean for that year, <x>_{F}/<x>_{1yr}, where x is: (d) P_{α}; (e) am; (f) aσ_{midnight}; (g) aσ_{dawn}, (h) aσ_{noon}, and (i) aσ_{dusk}. The data have been smoothed by applying a 3point running mean to the time series of means in the bins of width (1/36) yr. 

In the text 
Fig. 11 Identification of the RussellMcPherron effect using the polarity of the IMF [B_{Y}]_{GSEQ} component. The lefthand column is for all data, the middle column for [B_{Y}]_{GSEQ} > 0 and the righthand column for [B_{Y}]_{GSEQ} < 0. The top row shows annual means of am and P_{α} divided by their overall means for the whole interval: respectively, <am>_{1yr}/<am>_{all} (mauve lines) and <P_{α}>_{1yr}/P_{o} (black lines). The middle row shows the variations of P_{α}/P_{o} as a function of fractionofyear, F and year, <P_{α}>_{F}/<P_{α}>_{1yr} and the bottom row the same for am, <am>_{F}/<am>_{1yr}. As for Figure 10, data are for 1995–2017 and have been averaged into 36equal width bins of F in each year, and a 3point running mean applied to the time series to smooth the data. 

In the text 
Fig. 12 Analysis of the differences between semiannual variations in the am index (lefthand panels) and in power input to the magnetosphere, P_{α}/P_{o} (righthand panels). In all panels a thick line denotes the variation for all data, a thin line connecting open circles is for IMF [B_{Y}]_{GSEQ} > 0 and a thin line connecting filled triangles is for [B_{Y}]_{GSEQ} < 0. The lines in parts (a and b) are mean values whereas (c–f) show the average number of days per year N when the daily means of am (panels c and e) or of P_{α}/P_{o} (panels d and f) exceed their 90% quantile (blue lines), 95% quantile (orange lines) and 99% quantile (mauve lines). Parts (g) and (h) plot N/<N>, the number of days per year N normalized by their average value. Note that the variations of N in (e) and (f) for the separate IMF [B_{Y}]_{GSEQ} polarities are plotted on a scale that is different to that for all data in (c) and (d). The 90% quantile levels, q(0.9), are 2.031 for am/<am>_{all} and 2.081 for P_{α}/P_{o}; the 95% quantile levels, q(0.95), are 2.556 for am/<am>_{all} and 2.588 for P_{α}/P_{o}; the 99% quantile levels, q(0.99), are 4.193 for am/<am>_{all} and 4.400 for P_{α}/P_{o}. The close similarity between these pairs quantile values arises from the similarity of the normalised c.d.f.s shown in Figure 6 of the present paper and Figure 8b of Paper 1. 

In the text 
Fig. 13 Plots of (top) the probability distribution functions of (P_{α}/P_{o}) as a function of F and (bottom) the am amplification factor, (am/<am>_{all})/(P_{α}/P_{o}), as a function of F and in the same (P_{α}/P_{o}) bins as the p.d.f.s in the top panels. The left hand panels are for IMF [B_{Y}]_{GSEQ} < 0, the right hand panels are for IMF [B_{Y}]_{GSEQ} > 0. 

In the text 
Fig. 14 Mean values of am as a function of normalised power input into the magnetosphere, P_{α}/P_{o} (horizontal axis) and normalised solar wind dynamic pressure, p_{SW}/<p_{SW}>_{all} (vertical axis). The mean values of am are evaluated in bins of width 0.15 in both P_{α}/P_{o} and p_{SW}/<p_{SW}>_{all} and pixels colourcontoured according the scale shown. Only bins containing at least 5 samples are considered. 

In the text 
Fig. 15 Analysis of 3hourly am data. (a) The grey points form a scatter plot of am against 3hourly means of power input into the magnetosphere, P_{α}/P_{o}, generated by averaging over 3 h intervals that are shifted forward in time by the derived best lag of dt_{o} = 60 min (see Fig. 7a) relative to the threehourly intervals in which am is evaluated. The orange points are the mean values averaged in 1% quantile ranges of P_{α}/P_{o}, i.e., q(0) ≤ <P_{α}/P_{o}>_{τ=3h} < q(0.01), q(0.01) ≤ <P_{α}/P_{o}>_{τ=3h} < q(0.02), up to q(0.99) ≤ <P_{α}/P_{o}>_{τ=3h} < q(1). The black error bars are the plus and minus one standard deviation in those means. The mauve is the bestfit OLS linear regression to the 3hourly data. (b) the grey points are a scatter plot of the fit residuals Δam for the fit shown in (a): this is the difference between each threehourly am value and the best fit linear regression value based on the corresponding <P_{α}/P_{o}>_{τ=3h} value (Δam = am − am_{fit}), plotted as a function of the simultaneous threehourly mean of the normalized solar wind dynamic pressure <p_{SW}>_{τ=3h}/<p_{SW}>_{all}. The orange points in (b) are means in 1% quantile ranges of <p_{SW}>_{τ=3h} and error bars are plus and minus one standard deviation in the mean Δam. The mauve line is the best linear regression to the 3hourly values. 

In the text 
Fig. 16 F–UT pattern plots of the fit residuals Δam in Figure 15b for 1995–2017, inclusive. In each case the plots are constructed using hourly means of the solar wind dynamic pressure p_{SW} and the threehourly fit residuals Δam are interpolated to the midpoint of those hourly intervals. All plots are then smoothed with a 1–3–1 triangular weighting filter in both the F and UT dimension. (a) The overall mean value of Δam. (b) The mean Δam for the lower tercile of the simultaneous solar wind dynamic pressure, <Δam>_{F,UT} for q(0) ≤ <p_{SW}>_{τ=1h} < q(0.33). (c) The mean Δam for the middle tercile of the simultaneous solar wind dynamic pressure, i.e., <Δam>_{F,UT} for q(0.33) ≤ <p_{SW}>_{τ=1h} < q(0.67). (d) The mean Δam for the upper tercile of the simultaneous solar wind dynamic pressure, <Δam>_{F,UT} for q(0.67) ≤ <p_{SW}>_{τ=1h} < q(1). Note that in parts (b–d) the same color scale is used to emphasize that the equinoctial pattern, although present for all three ranges of p_{SW}, is of amplitude that increases with p_{SW} and is much larger in amplitude for the largest p_{SW} values. This relationship is further studied by Figure 17. 

In the text 
Fig. 17 The amplitude of the equinoctial pattern as a function of normalised solar wind dynamic pressure. The pattern amplitude is quantified by the standard deviation of the mean values of Δam in the 288 F–UT bins used to construct patterns like those shown in Figure 10, σ(<Δam>_{F,UT}). Patterns were constructed for 20 quantile ranges of the hourly mean solar wind dynamic pressure p_{SW} and σ(<Δam>_{F,UT}) for each is plotted as a function of <p_{SW}>_{q}/<p_{SW}>_{all}, where <p_{SW}>_{q} is the mean dynamic pressure in the quantile range and <p_{SW}>_{all} is the mean for all the data (which are for 1995–2017). The grey and white bands define the 20 quantile ranges of p_{SW} employed. There are a total of 192,719 valid hourly means of p_{SW} in the dataset and so each of the 20 quantile ranges contains 9636 samples. 

In the text 
Fig. 18 Analysis of the contributions to the (a) semiannual and (b) UT variations of the am index observed in the interval 1995–2017, inclusive. Mean values are shown as a function of (a) timeofyear F and (b) UT of: (black lines) the observed am data; (orange lines) am_{fit}, the bestfit of power input to the magnetosphere, P_{α}/P_{o} to am; (mauve lines) the fit residual Δam for the upper tercile of the solar wind dynamic pressure, q(0.67)<p_{SW} ≤ q(1); (blue lines) the fit residual Δam for the middle tercile of the solar wind dynamic pressure, q(0.33) < p_{SW} ≤ q(0.67); and (cyan lines) the fit residual Δam for the lower tercile of the solar wind dynamic pressure, q(0) < p_{SW} ≤ q(0.33). 

In the text 
Fig. 19 Analysis of the effects on am of (left column) power input to the magnetosphere P_{α}/P_{o} at a given solar wind dynamic pressure p_{SW} and of (right column) the effects of p_{SW} at a given P_{α}/P_{o}. Values of am/<am>_{all} are plotted as a function of corresponding normalised values of (left) P_{α} and (right) p_{SW}. In (a) data are averaged into 5 quantile ranges of the normalised solar wind dynamic pressure p_{SW}/<p_{SW}>_{all}, 36 bins of timeofyear F, and by the polarity of the IMF Y component in GSEQ, [B_{Y}]_{GSEQ}: the means of am/<am>_{all} are plotted against the corresponding means of P_{α}/P_{o}. In (b) data are averaged into 5 quantiles of normalised power input to the magnetosphere P_{α}/P_{o}, 36 bins of timeofyear F, and by the [B_{Y}]_{GSEQ} polarity and the means of am/<am>_{all} are plotted against the corresponding means of p_{SW}/<p_{SW}>_{all}. The black lines in the upper panels are 3rdorder polynomial fits in each case, fitted to data for all [B_{Y}]_{GSEQ} drawn over the range of the averaged data for the quantile range in question. The points are coloured according to the quantile range as given by the key. The lines are also plotted (this time using the same colours as the points in the upper panels) in the corresponding lower panel. The grey points show a scatter plot of all 10day average values. 

In the text 
Fig. A.1 Scatter plot of daily means of the transpolar voltage Φ_{PC}, as determined for 2001–2002 by the DMSP satellites and normalised for the satellite track to an ideal 06–18 MLT path using the procedure of Lockwood et al. (2009), as a function of the means of the am index for the same day. The error bars are plus and minus one standard error in the means. The mauve line is the bestfit 3rd order polynomial fit of Φ_{PC} and the grey area is bounded by the 2sigma uncertainty level in the fit. 

In the text 
Current usage metrics show cumulative count of Article Views (fulltext article views including HTML views, PDF and ePub downloads, according to the available data) and Abstracts Views on Vision4Press platform.
Data correspond to usage on the plateform after 2015. The current usage metrics is available 4896 hours after online publication and is updated daily on week days.
Initial download of the metrics may take a while.