A possible mechanism of formation of faculae

We propose a new heating mechanism of faculae. We think that the formation of faculae is a result of the Joule dissipation of the Hall current generated by the interaction of the convection field of granules in an active region and the inter-granular magnetic field. For a region to generate effectively Hall current, its characteristic length must be such that the magnetic Reynolds number is less than 1. The equation of energy balance in the facula region is

$\text{16σT}{}^3{}_{\mathrm{p}}\text{(T}{}_{\mathrm{l}}\text{? T}{}_{\mathrm{p}}\text{)}{}^{\mathrm{n}}\text{HP}{}_{\mathrm{s}}{}^{\mathrm{a}}\text{H}{}^{\mathrm{?}}\text{=}\text{Qn}{}_{\mathrm{s}}\text{m}{}_{\mathrm{i}}\text{u}{}_{\mathrm{x}}{}^2\text{(τ}{}^2{}_{\mathrm{in}}\text{ω}{}_{\mathrm{i}}\text{)}$

.For five observational models of faculae, we calculated the corresponding velocity fields, and the results are in basic agreement with the observed fields. The present mechanism explains the dependence of the facula brightness on the magnetic and velocity fields, the apparent distribution of the faculae on the solar disk and suggest a possible interpretation of the five structures of faculae.     

Es-q layer at huancayo during the March 1970 geomagnetic storm

The paper describes a comparison of vertical electron drift in the F-region (V_z) measured by VHP incoherent scatter radar at Jicamarca with the corresponding variations of geomagnetic horizontal field (H) and the maximum frequency reflected from The E_s layer (E_s) at Huancayo during the geomagnetic storm period 7–9 March, 1970. The V_z is generally upward during the daytime at the equator, but during 7–9, March, 1970, V_z was negative for brief periods associated with negative bays in H. These periods of abnormally low or of downward V_z correspond closely with the period of complete disappearance of the q type of E_s layer. The magnetic bays associated with the intensification of ring current do not affect the equatorial E_s- q and it is only the negative bays in H at the equator due to the ionospheric current flowing westward, that cause sudden disappearance of E_s? q. It is suggested that the q type of E_s is due to cross-field instability created in the electrojet region due to interaction of northward magnetic field and vertical upward Hall polarization electric field when the plasma density gradient is upward. The sudden disappearances of E_s? q are due to the reversal of the horizontal electric field in the equatorial ionosphere and thereby due to the reversal of the equatorial electrojet currents. These reversals of electric field may be due to the imposition on the normal S_q field of another westward electric field.     

The moon’s permanent magnetic field: a cratered-shell model

As part of our study of the larger-scale remanent magnetic field of the Moon, we have examined the effects of cratering in an otherwise spherically symmetrical shell magnetized by a concentric dipolar magnetic fieldH _o to an intensity of magnetizationc H _o, wherec is a constant. In our initial model, we assume that the material excavated from the craters is distributed with random orientation and thus does not contribute to the remanent dipole momentM _g . We further assume that the mare fill does not contribute significantly toM _g . We choose the magnetizing dipole momentM _o and the constantc such that the magnitude of the productcH _o ≃ 3 × 10⁻⁴Г at the outer surface of the shell in the equatorial plane of the dipole. This value of the intensity of remanent magnetization was chosen to be within the range 10⁻⁷−10⁻³Г’; the intensities of thermo-remanent magnetization exhibited by Apollo samples. Finally, we use the locations and diameters of the 10 largest craters on the Moon and the depth-to-diameter ratios of Pike’s formulation to model approximately the excavation of the magnetized shell. The remanent dipole momentM _g was calculated for each of three orthogonal orientations of the magnetizing dipoleM _o. The three magnitudes ofM _g fall in the range 4 × 10¹⁸−1 × 10¹⁹Г cm³ which is close to the upper limit of 10¹⁹Г cm³ estimated forM _g from the field measurements obtained with the Apollo subsatellites. Further, the distribution of the craters is such as to produce a significant transverse component ofM _g with acute angles between the spin axis andM _g in the range 51°–77°.     

Deuteron whistler and trans-equatorial propagation of the ion cyclotron whistler

New ion cyclotron whistlers which have the asymptotic frequency of one half the local proton gyrofrequency, $\text{G}{}_{\mathrm{p}}\text{2}$, and the minimum (or equatorial) proton gyrofrequency, G_pm, along the geomagnetic field line passing through the satellite have been found in the low-latitude topside ionosphere from the spectrum analysis of ISIS VLF electric field data received at Kashima, Japan. Ion cyclotron whistlers with asymptotic frequency of G_pm or $\text{G}{}_{\mathrm{pm}}\text{2}$ are observed only in the region of $\text{B}{}_{\mathrm{m}}\text{>}\text{B}\text{2}$ or rarely $\text{B}{}_{\mathrm{m}}\text{>}\text{B}\text{4}$, where B is the local magnetic field and B_m is the mini magnetic field along the geomagnetic field line passing through the satellite.The particles with one half the proton gyrofrequency may be the deuteron or alpha particle. Theoretical spectrograms of the electron whistlers (R-mode) and the ion cyclotron whistlers (L-mode) propagating along the geomagnetic field lines are computed for the appropriate distributions of the electron density and the ionic composition, and compared with the observed spectrograms.The result shows that the ion cyclotron whistler with the asymptotic frequency of $\text{G}{}_{\mathrm{p}}\text{2}$ is the deuteron whistler, and that the ion cyclotron whistlers with the asymptotic frequency of G_pm or $\text{G}{}_{\mathrm{pm}}\text{2}$ are caused by the trans-equatorial propagation of the proton or deuteron whistler from the other hemisphere.     

Neutron star sequences and the starquake glitch model for the Crab and the Vela pulsars

We construct for the first time, the sequences of stable neutron star (NS) models capable of explaining simultaneously, the glitch healing parameters, Q, of both the pulsars, the Crab (Q≥0.7) and the Vela (Q≤0.2), on the basis of starquake mechanism of glitch generation, whereas the conventional NS models cannot give such consistent explanation. Furthermore, our models also yield an upper bound on NS masses similar to those obtained in the literature for a variety of modern equations of state (EOSs) compatible with causality and dynamical stability. If the lower limit of the observational constraint of (i) Q≥0.7 for the Crab pulsar and (ii) the recent value of the moment of inertia for the Crab pulsar (evaluated on the basis of time-dependent acceleration model of the Crab Nebula), I _Crab,45≥1.93 (where I ₄₅=I/10⁴⁵ g cm²), both are imposed together on our models, the models yield the value of matching density, E _b =9.584×10¹⁴ g cm⁻³ at the core-envelope boundary. This value of matching density yields a model-independent upper bound on neutron star masses, M _max≤2.22M _⊙, and the strong lower bounds on surface redshift z _R ≃0.6232 and mass M≃2.11M _⊙ for the Crab (Q≃0.7) and the strong upper bound on surface redshift z _R ≃0.2016, mass M≃0.982M _⊙ and the moment of inertia I _Vela,45≃0.587 for the Vela (Q≃0.2) pulsar. However, for the observational constraint of the ‘central’ weighted mean value Q≈0.72, and I _Crab,45>1.93, for the Crab pulsar, the minimum surface redshift and mass of the Crab pulsar are slightly increased to the values z _R ≃0.655 and M≃2.149M _⊙ respectively, whereas corresponding to the ‘central’ weighted mean value Q≈0.12 for the Vela pulsar, the maximum surface redshift, mass and the moment of inertia for the Vela pulsar are slightly decreased to the values z _R ≃0.1645, M≃0.828M _⊙ and I _Vela,45≃0.459 respectively. These results set an upper and lower bound on the energy of a gravitationally redshifted electron-positron annihilation line in the range of about 0.309–0.315 MeV from the Crab and in the range of about 0.425–0.439 MeV from the Vela pulsar.     

Comment on the group velocity of surface waves on the plasma sheet boundary

In the recent estimation by Maltsev and Lyatsky (1984) of the group velocity of surface waves on the inner boundary of the plasma sheet, the effect of the curvature of the field lines of the ambient magnetic field of the Earth on the spectrum has been assessed. The authors have not accounted for the fact, however, that the group velocity of the compressional surface magnetohydrodynamic waves itself is nonzero transverse to the magnetic field—a characteristic which has been omitted in the spectrum of Chen and Hasegawa (1974), being used by Maltsev and Lyatsky.This characteristic of compressional surface MHD waves is inherent for the spectrum $\text{ω = (}\text{k}{}_6\text{k}\text{)V}{}_{\mathrm{A}}\text{(k}{}^2{}_6\text{+ 2k}{}^2{}_{\mathrm{\perp }}\text{)}{}^{\mathrm{<mtext>1</mtext><mtext>2</mtext>}}$, obtained by Nenovski (1978) in the cold plasma limit V_A ? V_S(V_A is Alfvén velocity, and V_S, sound velocity). A comment has been made on the restrictions, proceeding from the approximation, used by Maltsev and Lyatsky. The estimation of the velocities for movements of auroral riometer absorption bays have been reviewed.     

Pulsational instability of yellow hypergiants

Instability of population I (X = 0.7, Z = 0.02) massive stars against radial oscillations during the post-main-sequence gravitational contraction of the helium core is investigated. Initial stellar masses are in the range 65M _⊙ ≤ M _ZAMS ≤ 90M _⊙. In hydrodynamic computations of self-exciting stellar oscillations we assumed that energy transfer in the envelope of the pulsating star is due to radiative heat conduction and convection. The convective heat transfer was treated in the framework of the theory of time-dependent turbulent convection. During evolutionary expansion of outer layers after hydrogen exhaustion in the stellar core the star is shown to be unstable against radial oscillations while its effective temperature is T _eff > 6700 K for M _ZAMS = 65M _⊙ and T _eff > 7200 K for M _ZAMS = 90M _⊙. Pulsational instability is due to the κ-mechanism in helium ionization zones and at lower effective temperature oscillations decay because of significantly increasing convection. The upper limit of the period of radial pulsations on this stage of evolution does not exceed ≈200 day. Radial oscillations of the hypergiant resume during evolutionary contraction of outer layers when the effective temperature is T _eff > 7300 K for M _ZAMS = 65M _⊙ and T _eff > 7600 K for M _ZAMS = 90M _⊙. Initially radial oscillations are due to instability of the first overtone and transition to fundamental mode pulsations takes place at higher effective temperatures (T _eff > 7700 K for M _ZAMS = 65M _⊙ and T _eff > 8200 K for M _ZAMS = 90M _⊙). The upper limit of the period of radial oscillations of evolving blueward yellow hypergiants does not exceed ≈130 day. Thus, yellow hypergiants are stable against radial stellar pulsations during the major part of their evolutionary stage.     

Stability of motion in the Sitnikov 3-body problem

We study the stability of motion in the 3-body Sitnikov problem, with the two equal mass primaries (m ₁ = m ₂ = 0.5) rotating in the x, y plane and vary the mass of the third particle, 0 ≤ m ₃ < 10⁻³, placed initially on the z-axis. We begin by finding for the restricted problem (with m ₃ = 0) an apparently infinite sequence of stability intervals on the z-axis, whose width grows and tends to a fixed non-zero value, as we move away from z = 0. We then estimate the extent of “islands” of bounded motion in x, y, z space about these intervals and show that it also increases as |z| grows. Turning to the so-called extended Sitnikov problem, where the third particle moves only along the z-axis, we find that, as m ₃ increases, the domain of allowed motion grows significantly and chaotic regions in phase space appear through a series of saddle-node bifurcations. Finally, we concentrate on the general 3-body problem and demonstrate that, for very small masses, m ₃ ≈ 10⁻⁶, the “islands” of bounded motion about the z-axis stability intervals are larger than the ones for m ₃ = 0. Furthermore, as m ₃ increases, it is the regions of bounded motion closest to z = 0 that disappear first, while the ones further away “disperse” at larger m ₃ values, thus providing further evidence of an increasing stability of the motion away from the plane of the two primaries, as observed in the m ₃ = 0 case.     

A Study on the CO Isotopic Lines of the Star Forming Region AFGL 5157

By the mapping observations simultaneously at the ¹²CO (J=1-0), ¹³CO (J=1-0), and C¹⁸O (J=1-0) lines on the area of 24’×24’ (12 pc×12 pc) of the star forming region AFGL 5157, we have obtained the distribution and averaged physical parameters for the respective ¹³CO and C¹⁸O cores of this molecu- lar cloud. At the edge of the molecular cloud, the isotopic abundance ratio is X [(¹³CO)/(C¹⁸O)] ≈ 10, close to the ratio of a giant molecular cloud. The viral masses of the ¹³CO and C¹⁸O cores are less than the masses of the molecu-lar cloud cores, so the molecular cloud cores are gravitationally unstable, and the C¹⁸O molecular cloud core is more easy to collapse. The column density distributions of the C¹⁸O molecular cloud core in the northeast and southwest directions are, respectively, 1.1 × 10²³× z^−0.43 and 4.6 × 10²⁵× z^−0.58, where z is the distance from the center of the molecular cloud core. The high velocity molecular out?ow has been con?rmed from our ¹²CO spectra, the mass loss rate of the out?ow has been estimated, and the mass-velocity relation of the out?ow is ?tted by a power-law function of m ∝ v^−1.8. The star formation rate of the ¹³CO molecular cloud core is as high as 23%, probably, under the in?uence of     

Emission line shapes produced by dissociative excitation of atmospheric gases

Special line shapes are derived fro the λ 1356 Å (${}^5\text{S}{}^0\text{-}{}^3\text{P}$) transition of atomic oxygen from metastable (${}^5\text{S}{}^0\text{-}{}^3\text{P}$) time-of-flight spectra produced by electron impact dissociative excitation of O₂, CO₂, CO, and NO, and they are compared with the broadened λ 1304 A resonance line shapes deduced by Poland and Lawrence (1973) from atomic oxygen absorption studies. The non-thermal line shapes for both airglow emission features are shown to have an effective width comparable to a 60,000 K thermal doppler line shape for an electron impact energy of 100eV. The variation of the effective line width with electron-impact energy from threshold to 300 eV is given. Since the effective line width of the resonance radiation produced by dissociative excitation is very large compared with the doppler absorption widths of the ambient O atoms at normal exospheric temperatures, the anomalously broadened resonance lines will propagate through a planetary atmosphere as though they were optically thin. Thus, electron-impact dissociation of CO and CO₂ will contribute to the observed optically thin component of the λ 1304 Å emission in the upper atmospheres of Venus and Mars. However, the process cannot account for more than 10% of the observed optically thin emission because of the small magnitude of the excitation cross-section and the comparatively high-energy threshold for the process. The possibility that the source of the kinetically energetic $\text{O}\text{(}{}^3\text{S}$) atoms is the dissociative recombination of vibrationally excited CO₂⁺ ions is discussed.     

Magnetic fields and quasi-periodic oscillations of black hole radiation

Various relations are found between the key parameters of black holes and active galactic nuclei. Some have a statistical property, others follow from the theoretical consideration of the evolution of these objects. In this paper we use a recently discovered empirical relation between the characteristic frequency of quasi-periodic oscillations of radiation ν _br of black holes, their masses and matter accretion rates to determine the magnetic field strength B _H at the black hole event horizon. Since the characteristic frequency can be determined from observations, the use of a new relation for the estimations of magnetic field B _H can yield more definite results, since we are decreasing the number of the unknown or poorly-determined parameters of objects (it especially concerns the accretion rate Ṁ). The typical values which we have found are B _H ≃ 10⁸G for the stellar mass black holes, and B _H ≃ 10⁴G for the supermassive black holes. Besides, we demonstrate that if the linear polarization of an object is caused by the radiation of a magnetized accretion disk, then the degree of observable polarization is p ∼ ν _br ^−1/2.     

The determination and analysis of the orbit of NIMBUS 1 ROCKET, 1964-52B: Part 2: Variations in orbital eccentricity

The influence of aerodynamic drag and the geopotential on the motion of the satellite 1964-52B is considered. A model of the atmosphere is adopted that allows for oblateness, and in which the density behaviour approximates to the observed diurnal variation. A differential equation governing the variation of the eccentricity, e, combining the effects of air drag with those of the Earth's gravitational field is given. This is solved numerically using as initial conditions 310 computed orbits of 1964-52B.The observed values of eccentricity are modified by the removal of perturbations due to luni-solar attraction, solid Earth and ocean tides, solar radiation pressure and low-order long-periodic tesseral harmonic perturbations. The method of removal of these effects is given in some detail. The behaviour of the orbital eccentricity predicted by the numerical solution is compared with the modified observed eccentricity to obtain values of atmospheric parameters at heights between 310 and 430 km. The daytime maximum of air density is found to be at 14.5 hours local time. Analysis of the eccentricity near 15th order resonance with the geopotential yielded values of four lumped geopotential harmonics of order 15, namely: $\text{10}{}^9\text{C}{}^{\mathrm{1,0}}{}_{15}\text{= ?78.8 ± 7.0}$, $\text{10}{}^9\text{S}{}^{\mathrm{1,0}}{}_{15}\text{= ?69.4 ± 5.3}$, $\text{10}{}^9\text{C}{}^{\mathrm{?1,2}}{}_{15}\text{= ?41.6 ± 3.5}$$\text{10}{}^9\text{S}{}^{\mathrm{?1,2}}{}_{15}\text{= ?26.1 ± 8.9}$, at inclination 98.68°.     

Models of the formation of the solar nebula

Models of the collapse of a protostellar cloud and the formation of the solar nebula reveal that the size of the nebula produced will be the larger of $\text{R}{}_{\mathrm{<mtext>CF</mtext>}}\text{≡ J}{}^2\text{/k}{}^2\text{GM}{}^3\text{and}\text{R}{}_{\mathrm{<mtext>V</mtext>}}\text{≡ (GMv/2c}{}_{\mathrm{<mtext>c</mtext>}}{}^3\text{)}\text{1}\text{2}$ (where J, M, and c_s are the total angular momentum, total mass, and sound speed of the protosetellar material; G is the gravitational constant; k is a number of order unity; and v is the effective viscosity in the nebula). From this result it can be deduced that low-mass nebulas are produced if P ≡ (R_V/R_CF)² ? 1; “massive” nebulas result if P ? 1. Gravitational instabilities are expected to be important for the evolution of P ? 1 nebulas. The value of J distinguishes most current models of the solar nebula, since P ∝ J^?4. Analytic expressions for the surface density, nebular mass flux, and photospheric temperature distributions during the formation stage are presented for some simple models that illustrate the general properties of growing protostellar disks. It does not yet seem possible to rule out either P ? 1 or P < 1 for the solar nebula, but observed or possible heterogeneities in composition and angular-momentum orientation favor P < 1 models.     

Hidden Mass in the Asteroid Belt

The total mass of the asteroid belt is estimated from an analysis of the motions of the major planets by processing high precision measurements of ranging to the landers Viking-1, Viking-2, and Pathfinder (1976-1997). Modeling of the perturbing accelerations of the major planets accounts for individual contributions of 300 minor planets; the total contribution of all remaining small asteroids is modeled as an acceleration caused by a solid ring in the ecliptic plane. Mass M_ring of the ring and its radius R are considered as solve-for parameters. Masses of the 300 perturbing asteroids have been derived from their published radii based mainly on measured fluxes of radiation, making use of the corresponding densities. This set of asteroids is grouped into three classes in accordance with physical properties and then corrections to the mean density for each class are estimated in the process of treating the observations. In this way an improved system of masses of the perturbing asteroids has been derived.The estimate M_ring≈(5±1)×10⁻¹⁰M_⊙ is obtained (M_⊙ is the solar mass) whose value is about one mass of Ceres. For the mean radius of the ring we have R≈2.80 AU with 3% uncertainty. Then the total mass M_belt of the main asteroid belt (including the 300 asteroids mentioned above) may be derived: M_belt≈(18±2)×10⁻¹⁰M_⊙. The value M_belt includes masses of the asteroids which are already discovered, and the total mass of a large number of small asteroids—most of which cannot be observed from the Earth. The second component M_ring is the hidden mass in the asteroid belt as evaluated from its dynamical impact onto the motion of the major planets.Two parameters of a theoretical distribution of the number of asteroids over their masses are evaluated by fitting to the improved set of masses of the 300 asteroids (assuming that there is no observational selection effect in this set). This distribution is extrapolated to the whole interval of asteroid masses and as a result the independent estimate M_belt≈18×10⁻¹⁰M_⊙ is obtained which is in excellent agreement with the dynamical finding given above.These results make it possible to predict the total number of minor planets in any unit interval of absolute magnitude H. Such predictions are compared with the observed distribution; the comparison shows that at present only about 10% of the asteroids with absolute magnitude H<14 have been discovered (according to the derived distribution, about 130,000 such asteroids are expected to exist).     

Instability of LBV stars against radial oscillations

Hydrodynamic calculations of nonlinear radial oscillations of LBV stars with effective temperatures 1.5 × 10⁴ K ⩽ T _eff ⩽ 3 × 10⁴ K and luminosities 1.2 × 10⁶ L _⊙ ⩽ L ⩽ 1.9 × 10⁶ L _⊙ have been performed. Models for the evolutionary sequences of Population I stars (X = 0.7, Z = 0.02) with initial masses 70M _⊙ ⩽ M _ZAMS ⩽ 90M _⊙ at the initial helium burning stage have been used as the initial conditions. The radial oscillations develop on a dynamical time scale and are nonlinear traveling waves propagating from the core boundary to the stellar surface. The amplitude of the velocity variations for the outer layers is several hundred km s⁻¹, while the bolometric magnitude variations are within ΔM _bol ⩽ 0_· ^m 2. The onset of oscillations is not related to the κ-mechanism and is attributable to the instability of a self-gravitating envelope gas whose adiabatic index is close to its critical value of Γ₁ = 4/3 due to the dominant contribution of radiation in the internal energy and pressure. The interval of magnitude variation periods (6 days ≤ II ≤ 31 days) encompasses all currently available estimates of the microvariability periods for LBV stars, suggesting that this type of nonstationarity is pulsational in origin.     

First comprehensive study of W-type W UMa eclipsing binary V1036 Her

CCD photometry of the eclipsing binary system V1036 Her was performed using Johnson V filter in Dr. Mojtahedi Observatory of the University of Birjand during July-August 2017 and July 2018. Moreover, the spectroscopy of the system was carried out by TRES during April 2018. A mass ratio of $3.07\left(27\right)$is obtained and an initial effective temperature of 5500 was suggested. For the first time, the relative and absolute parameters of the system are determined by analyzing the light curve and radial velocity data. The results indicate that V1036 Her is a W-subtype W UMa system with a degree of overcontact of 22%. The analysis of the period change shows that the period of the system changes with the rate of $\dot{P}=2.23\left(4\right)\times {10}^{-7}\text{day}/\text{year}$. With the assumption of the system mass conservation, a mass transfer rate from the primary to the secondary component of ${\dot{m}}_1=1.00\left(3\right)\times {10}^{-7}{M}_{\odot }/\mathit{year}$ is probable. Additionally, a periodic behavior with a period of about 10 years is observed in the O-C curve, which predicts the possibility of a third body with a minimum mass of $0.14\left(1\right)M_{\odot }$.     

Temperature of interstellar warm ionized medium

This investigation on the temperature of the interstellar warm ionized medium (WIM) is characterized by the number and energy balance of the constituents of the WIM complex plasma viz. H plasma (electrons/ions/neutral atoms) and graphite dust, having a size distribution, characterized by the MRN (Mathis, Rumpl and Nordsieck) power law. Ionization of neutral atoms, electron–ion recombination, photoemission of electrons from and accretion on the dust and cooling through electron collisional excitation, followed by radiative decay of atoms has been included in the analysis. An appropriate expression for the rate of emission and mean energy of photoelectrons emitted from the surface of positively charged dust particles has been used which takes into account the dependence of absorption efficiency on wavelength of the radiation, radius of the particle and spectral irradiance distribution. The results of the parametric analysis have been displayed graphically. It is seen that the consensus values of temperature, surface potential on the dust particles and electron/ion/neutral atom densities, characteristic of interstellar warm ionized medium can be explained on the basis of plausible combinations of the dust particle density n _d and the parameter f _ex α _ex , where f _ex is the fraction of the energy of the neutral gas atoms which gets irradiated, α _ex n _e n _n is the number of the neutral atoms, which get excited per unit volume per unit time and n _e (n _n ) correspond to the density of electrons (neutral atoms).     

Numerical studies of the transport of solar protons in interplanetary space

Numerical solutions of the Fokker-Planck equation governing the transport of solar protons are obtained using the Crank-Nicholson technique with the diffusion coefficient represented by K_r=K₀r^b where r is radial distance from the Sun and b can take on positive or negative values. As b ranges from +1 to ?3, the time to the observation of peak flux decreases by a factor of 5 for 1 MeV protons when $\text{V}\text{K}{}_0\text{=}\text{3 AU}{}^{\mathrm{b?1}}$ where V is the solar wind speed. The time to peak flux is found to be very insensitive to assumptions concerning the solar and outer scattering boundary conditions and the presence of exponential time decay in the flux does not depend on the existence of an outer boundary. At $\text{V}\text{K}{}_0\text{? 15 AU}{}^{\mathrm{b?1}}$, 1 MeV particles come from the Sun by an almost entirely convective process and suffer large adiabatic deceleration at b?0 but for b=+1, large Fermi acceleration is possible at all reasonable $\text{V}\text{K}{}_0$ values. Implications of this result for the calculation and measurement of particle diffusion coefficients is discussed. At $\text{b?0}$, the pure diffusion approximation to transport overestimates by a factor 2 or more the time to peak flux but as b becomes more negative, the additional effects of convection and energy loss become less important.     

Twenty-Seven-Day Variation of Galactic Cosmic Rays

Neutron monitor data observed at Climax (CL) and Huancayo/Haleakala (HU/HAL) have been used to calculate the amplitude A of the 27-day variation of galactic cosmic rays (CRs). The median primary rigidity of response, R _m, for these detectors encompasses the range 18 ≤R _m≤46 GV and the threshold rigidity R ₀ covers the range 2.97≤R ₀≤12.9 GV. The daily average values of CR counts have been harmonically analyzed for each Bartels solar rotation (SR) during the period 1953 – 2001. The amplitude of the 27-day CR variation is cross-correlated to solar activity as measured by the sunspot number R, the interplanetary magnetic field (IMF) strength B, the z-component B _z of the IMF vector, and the tilt angle ψ of the heliospheric current sheet (HCS). It is anticorrelated to the solar coronal hole area (CHA) index as well as to the solar wind speed V. The wind speed V leads the amplitude by 24 SRs. The amplitude of the 27-day CR variation is better correlated to each of the these parameters during positive solar polarity (A>0) than during negative solar polarity (A<0) periods. The CR modulation differs during A>0 from that during A<0 owing to the contribution of the z-component of the IMF. It differs during A ₁>0 (1971 – 1980) from that during A ₂>0 (1992 – 2001) owing to solar wind speed.