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In chemistry and physics, the iron group refers to elements that are in some way related to iron. These elements are relatively abundant both on Earth and elsewhere in the universe. The term is ambiguous in different contexts, and almost obsolete in chemistry.
Contents
Periodic table
The iron group in the periodic table referred to the elements iron, cobalt and nickel, that is the first row of group VIII (or VIIIB under the old numbering system.) In modern numbering, the iron group appears as three columns numbered group 8, 9 and 10. These metals, and the platinum group immediately below them, were set aside from the other elements as they show obvious similarities among themselves in their chemistry, but are not obviously related to any of the other groups.
The similarities in chemistry along what is now known as the first row of the transition metals were noted by Adolph Strecker in 1859. Newlands' "octaves" (1865) were harshly criticized for separating iron from cobalt and nickel. Mendeleev stressed that groups of "chemically analogous elements" could have similar atomic weights as well as atomic weights which increase by equal increments, both in his original 1869 paper and his 1889 Faraday Lecture.
Analytical chemistry
In the traditional methods of qualitative inorganic analysis, the iron group consists of those cations which
The main cations in the iron group are iron itself (Fe2+ and Fe3+), aluminium (Al3+) and chromium (Cr3+). If manganese is present in the sample, a small amount of hydrated manganese dioxide is often precipitated with the iron group hydroxides. Less common cations which are precipitated with the iron group include beryllium, titanium, zirconium, vanadium, uranium, thorium and cerium.
Astrophysics
The iron group in astrophysics is the group of elements from chromium to nickel which are substantially more abundant in the universe than those that come after them – or immediately before them – in order of atomic number. The study of the abundances of iron group elements relative to other elements in stars and supernovae allows the refinement of models of stellar evolution.
The explanation for this relative abundance can be found in the process of nucleosynthesis in certain stars, specifically those of about 8–11 Solar masses. At the end of their lives, once other fuels have been exhausted, such stars can under a brief phase of "silicon burning". This involves the sequential addition of helium nuclei 4
2He
(an "alpha process") to the heavier elements present in the star, starting from 28
14Si
:
All of these nuclear reactions are exothermic, that is they release energy: the energy that is released partially offsets the gravitational contraction of the star. However, the series ends at 56
28Ni
, as the next reaction in the series,
is endothermic. With no further source of energy to support itself, the core of the star collapses on itself while the outer regions are blown off in a Type II supernova.
Nickel-56 is unstable with respect to beta decay, and the final stable product of silicon burning is 56
26Fe
.
Nuclide stability
It is often incorrectly assumed that iron-56 is the most stable of all the nuclides. This is not quite true: 62
28Ni
and 58
26Fe
have slightly higher binding energies per nucleon – that is, they are slightly more stable as nuclides – as can be seen from the table on the right. However, there are no rapid nucleosynthetic routes to these nuclides. There are several stable nuclides of elements from chromium to nickel around the top of the stability curve, accounting for their relative abundance in the universe. The nuclides which are not on the direct alpha-process pathway are formed by the so-called S-process, the capture of slow neutrons within the star.